diff --git "a/batch_s000049.csv" "b/batch_s000049.csv"
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+source,target
+ We set the nominal values of ο=3/2 and 7/2 for the thin and. thick shell cases. respectively.," We set the nominal values of $g=3/2$ and $7/2$ for the thin and thick shell cases, respectively."
+ Thus. (he only [ree parameter of the reverse shock emission is Dp.," Thus, the only free parameter of the reverse shock emission is $\Gamma_0$."
+ For the forward shockwe use the time evolution of the svnchrotron spectrum in the appropriate regime (i.e. spherical [SAL [or /—ἐς Sari.Piran&Naravan1998.. and an expanding jet [or />/;: Sari.Piran&Llalpern 1999)): here /;. the jet break time. is the epoch at which T—8; land 0; is the half opening angle ofthe jet.," For the forward shockwe use the time evolution of the synchrotron spectrum in the appropriate regime (i.e. spherical ISM for $t<t_j$: \citealt{spn98}, and an expanding jet for $t>t_j$: \citealt{sph99}) ); here $t_j$, the jet break time, is the epoch at which $\Gamma\sim\theta_j^{-1}$ ,and $\theta_j$ is the half opening angle ofthe jet."
+ To account [or possible extinction within the host galaxy. tet. we use the parametric extinction curvesof Cardelli. and Fitzpatrick&Massa (1988).. alongwith the interpolation caleulated by Reiehart(2001 )..," To account for possible extinction within the host galaxy, $A_V^{\rm host}$ , we use the parametric extinction curvesof \citet{ccm89} and \citet{fm88}, , alongwith the interpolation calculated by \citet{rei01}. ."
+Observed flat rotational curves of many galaxies have been sibjeet of long-term controversy.,Observed flat rotational curves of many galaxies have been subject of long-term controversy.
+ The observational [act that the azimuthal velocity of gas aud stars in the galactic plane is constaut over a large range of the distauces [rom the ceutre of a galaxy Las yjelded two main explanations., The observational fact that the azimuthal velocity of gas and stars in the galactic plane is constant over a large range of the distances from the centre of a galaxy has yielded two main explanations.
+ Iu an attempt tosave the assertion that the Newtoniau gravitatioual tleory holds over the cosmological distauces. oue such theory assumes the presence of non-baryonie 1jassive dark halo surrounding a spiral disk.," In an attempt to save the assertion that the Newtonian gravitational theory holds over the cosmological distances, one such theory assumes the presence of non-baryonic massive dark halo surrounding a spiral disk."
+ In this scenario. gravitational acceleration GAL«(r)/r which balances the centrifugal acceleration. rode/V7(r)/r. is. assumed to vary as yfL/r.," In this scenario, gravitational acceleration $GM_<(r)/r^2$ which balances the centrifugal acceleration $V^2(r)/r$, is assumed to vary as $1/r$."
+ This. means that tle inass enclosed withit a certain radius r2 Ade(r). scales as xr.," This means that the mass enclosed within a certain radius $r$ , $M_<(r)$, scales as $\propto r$."
+ However. this is not what is observed at large radii of tLe Galaxy.," However, this is not what is observed at large radii of the Galaxy."
+ The secoud possible explanation of the flat rotational curves is 1iat the Newtoulan grwvity does 100 apply on Cosinological scales and further imodificatious are due (Milgrom 1950)., The second possible explanation of the flat rotational curves is that the Newtonian gravity does not apply on cosmological scales and further modifications are due (Milgrom 1983).
+ Histo‘ically the atter explanation was not favoured due to abseuce of the general reativistic extension ofthe theory., Historically the latter explanation was not favoured due to absence of the general relativistic extension of the theory.
+" However. this drawback was alleviated by the formulation of the gejeralisation of Einstein"" general 'elativity based on a pseudo-Bieimauniau metric tensor aud a skew-symuinetric rauk tl ο. ielcl. called inetric-skew-teusor gravity (MSTC)."," However, this drawback was alleviated by the formulation of the generalisation of Einstein's general relativity based on a pseudo-Riemannian metric tensor and a skew-symmetric rank three tensor field, called metric-skew-tensor gravity (MSTG)."
+ The latter leads to a moclilied acceleration law hat cau explain the flat rotation curves of galaxies aud cluster lensiug without postulating exotic dark matter (Mollat 2005)., The latter leads to a modified acceleration law that can explain the flat rotation curves of galaxies and cluster lensing without postulating exotic dark matter (Moffat 2005).
+ Recently. Brownstein MolIat. (2006) have shown that MSTC cau," Recently, Brownstein Moffat (2006) have shown that MSTG can"
+This is because theholes.,This is because the.
+ We do not consider the possibility that stars are made in the nucleus of the galaxies (as would be required by AL) and then redistributed throughout the galaxy. in. stellar-kinematical mergers because (1) the dense stellar cores so produced: will not be disrupted. by mergers (or secular evolution). (2) we would. require about LOO mergers. per L elliptical and the observed. merger rate (c.g. Carlberg et al.," We do not consider the possibility that stars are made in the nucleus of the galaxies (as would be required by ) and then redistributed throughout the galaxy in stellar-kinematical mergers because (1) the dense stellar cores so produced will not be disrupted by mergers (or secular evolution), (2) we would require about 100 mergers per $^{*}$ elliptical and the observed merger rate (e.g. Carlberg et al."
+ 2000) is not that high. and (3) the number density concordance described above is then Lost. due to the Large number of mergers.," 2000) is not that high, and (3) the number density concordance described above is then lost due to the large number of mergers."
+" Van der Marel 1999. finds the black hole mass (in solar units) log),Mia&183|logy,Lv."," Van der Marel 1999 finds the black hole mass (in solar units) $\log_{10} M_{\rm BH} \approx -1.83 +
+\log_{10} L_V$."
+ Pherefore for a Schechter (1976). L elliptical galaxy (Ady= 21.5) at redshift z=0. the central black hole has mass 5.1.LO°AL.., Therefore for a Schechter (1976) $L^{*}$ elliptical galaxy $M_V = -21.5$ ) at redshift $z=0$ the central black hole has mass $5.1 \times 10^{8} {\rm M}_{\odot}$.
+" For the same elliptical galaxy. the stellar core mass. is Aly,=1.60M. assuming. that the core is. approximately. isothermal at small raclii Ge<r, where ο. is the core raclius) that the galaxy has a de Vaucouleurs +t profile at. [large radi. that the ratio of +, to the de Vaucouleurs. elfective racdus ris 0.033 (derived from the sample of Lauer 1985). and a value of (AL/Ly)=3.86 for elliptical galaxies (see the discussion in Section 2.1.3 of Fukugita et al."," For the same elliptical galaxy, the stellar core mass is $M_{*c} = 1.6 \times 10^{9} {\rm M}_{\odot}$, assuming that the core is approximately isothermal at small radii $r<r_c$ where $r_c$ is the core radius), that the galaxy has a de Vaucouleurs $r^{1/4}$ profile at large radii, that the ratio of $r_c$ to the de Vaucouleurs effective radius $r_e$ is 0.033 (derived from the sample of Lauer 1985), and a value of $\left(M/L_V \right) = 3.86$ for elliptical galaxies (see the discussion in Section 2.1.3 of Fukugita et al."
+ LOOS this assumes the stellar initial mass function. (EME) of Gould. Baheall Flynn 1996. whieh turns over at about 0.5 AL.. similar to the LME of Ixroupa. Tout Gilmore (1993) that we adopt below).," 1998 – this assumes the stellar initial mass function (IMF) of Gould, Bahcall Flynn 1996, which turns over at about 0.5 $_{\odot}$, similar to the IMF of Kroupa, Tout Gilmore (1993) that we adopt below)."
+" ""Therefore. for tvpical luminous elliptical galaxies Given the scatter in the correlations presented. by. van der Alarel and Lauer. and the fact that many cores deviate significantly from isothermalitv. the most extreme galaxies may deviate from this ratio by as much an order. of maenitucle."," Therefore, for typical luminous elliptical galaxies Given the scatter in the correlations presented by van der Marel and Lauer, and the fact that many cores deviate significantly from isothermality, the most extreme galaxies may deviate from this ratio by as much an order of magnitude."
+ For accretion onto a black hole. where 5 is the ellicicney of accretion anc AZ is the mass accretion rate.," For accretion onto a black hole, where $\eta$ is the efficiency of accretion and $\dot{M}$ is the mass accretion rate."
+ For dust-enshroudecd star. formation where almost all the πας is absorbed and reraciated: at far-infrarecl wavelengths. the star formation rate (in AL. 1) is (Rowan-Robinson et al.," For dust-enshrouded star formation where almost all the flux is absorbed and reradiated at far-infrared wavelengths, the star formation rate (in $_{\odot}$ $^{-1}$ ) is (Rowan-Robinson et al."
+" 1997) where £, (in L.) is the total luminosity generated. by the voung stars.", 1997) where $L_{*}$ (in $_{\odot}$ ) is the total luminosity generated by the young stars.
+ This relation adopts a value of ὁ=0.45 appropriate to the stellar EME of Ixroupa et al. (, This relation adopts a value of $\phi = 0.45$ appropriate to the stellar IMF of Kroupa et al. (
+1993). and à value ofc=1 (see Section 3 of Rowan-Robinson et al.,"1993), and a value of $\epsilon = 1$ (see Section 3 of Rowan-Robinson et al."
+ 1997 for a definition of those parameters and a discussion of the derivation of this equation and other approaches)., 1997 for a definition of those parameters and a discussion of the derivation of this equation and other approaches).
+ The LAL chosen here has been shown to be consistent with most current observations of both star-forming regions and. the solar neighbourhood (Gilmore Llowell 1998)., The IMF chosen here has been shown to be consistent with most current observations of both star-forming regions and the solar neighbourhood (Gilmore Howell 1998).
+ Combining the two previous equations vields Accretion is a lot more ellicient than star-formation at eenerating. luminosity⋠⋠ per unit. mass: a 12 L. galaxy has a star formation rate of about SO M. vr+ but a2 12 L. ⊏↥⊔⋜↧⊳∖⋜⊔⋅↓↥⋜↧⊳∖⋜⋯⋯⇍≼∼↓⋅⋖⊾↿↓∪⊔↓⋅⋜⋯⊾∪⇂⋜↧∣⋡⋯∐↓↳∖↓⋅∙∖⇁↓⋅⇂∪↓⋅↗∣∶∪⋅↓ . ⊥⋅ (e.g. Chokshi Turner 1992)., Combining the two previous equations yields Accretion is a lot more efficient than star-formation at generating luminosity per unit mass: a $^{12}$ $_{\odot}$ galaxy has a star formation rate of about 80 $_{\odot}$ $^{-1}$ but a 2 $\times$ $^{12}$ $_{\odot}$ quasar has an accretion rate of about 1 $_{\odot}$ $^{-1}$ for $\eta = 0.1$ (e.g. Chokshi Turner 1992).
+ Pherclore averaged over time in the SCUBA sources: This ratio is very big. meaning that the total power eencrated in the SCUBA sources comes almost. entirely [rom accretion if is correct.," Therefore averaged over time in the SCUBA sources: This ratio is very big, meaning that the total power generated in the SCUBA sources comes almost entirely from accretion if is correct."
+ Note that. throughout this caleulation. we have been working backwards. deriving the properties of the progenitor from the properties of the remnants under the assumption that we can match particular progenitors to particular remnants: this is quite a different caleulation from one in which we start out with a cloud of gas and attempt to track its progress.," Note that throughout this calculation, we have been working backwards, deriving the properties of the progenitor from the properties of the remnants under the assumption that we can match particular progenitors to particular remnants; this is quite a different calculation from one in which we start out with a cloud of gas and attempt to track its progress."
+ The fact that the power sources are so heavily accretion-dominated means that most SCUBA sources must be dust-enshrouded active galactic nucle: CAGNs) if is correct., The fact that the power sources are so heavily accretion-dominated means that most SCUBA sources must be dust-enshrouded active galactic nuclei (AGNs) if is correct.
+ This is not surprising. since the SCUBA sources contribute at least 10 per cent of the bolometric Luminosity density of the Universe. and this could not be generated from the stars in gl cores alone. which comprise less than 1. per cent. by mass of the stars in the Universe. and. these are the only stars that the SCUBA sources are allowed to make if is correct.," This is not surprising, since the SCUBA sources contribute at least 10 per cent of the bolometric luminosity density of the Universe, and this could not be generated from the stars in gE cores alone, which comprise less than 1 per cent by mass of the stars in the Universe, and these are the only stars that the SCUBA sources are allowed to make if is correct."
+ There exist. other reasons [or suggesting a connection between the SCUBA sources and dust-enshrouded ACNs., There exist other reasons for suggesting a connection between the SCUBA sources and dust-enshrouded AGNs.
+ Apt Phe local supermassive DII density (Magorrian et al., 4pt The local supermassive BH density (Magorrian et al.
+ 1998. van der Marel 1999) is high and its production generates a bolometric background of (Trentham Blain. 2000) if they racliate at one-tenth of the Eddington luminosity and if the luminosity density of obscureedl ACGNs follows that measured by Bovle Terlevich (1998) for optical quasars.," 1998, van der Marel 1999) is high and its production generates a bolometric background of (Trentham Blain, 2000) if they radiate at one-tenth of the Eddington luminosity and if the luminosity density of obscured AGNs follows that measured by Boyle Terlevich (1998) for optical quasars."
+ This equation follows [rom a consideration of the energy. released by accretion processes per comoving volume element. given a final MDO density some part pacNabers Of which was generated. by dust-enshrouded accretion.," This equation follows from a consideration of the energy released by accretion processes per comoving volume element given a final MDO density some part $\rho_{\rm AGN,dusty}$ of which was generated by dust-enshrouded accretion."
+ Ehe fiducial pcs«deis is the estimate of Salucci et al. (," The fiducial ${\rho_{\rm AGN,dusty}}$ is the estimate of Salucci et al. ("
+1999) for obscurecl quasars.,1999) for obscured quasars.
+ The SCUBA sources in Fig., The SCUBA sources in Fig.
+ L generate about 7 nWm“sr which is close to this value.," 1 generate about 7 ${\rm nW}{\rm m}^{-2}{\rm sr}^{-1}$ ,which is close to this value."
+ 4ρι Phere is growing evidence that the hard. (30 keV) X-ray background. originates from absorbed: quasars which reracdiate energy absorbed at optical. ultraviolet. and soft. X-ray at. far-infrared and submillimetre wavelengths.," 4pt There is growing evidence that the hard (30 keV) X-ray background originates from absorbed quasars which reradiate energy absorbed at optical, ultraviolet, and soft X-ray at far-infrared and submillimetre wavelengths."
+Feibclman 2000).,Feibelman 2000).
+ Very much like in the case of Ix. 1-47. hese PNe have consicerably high. electron. temperatures (average 1). from optical. ISO and IUIS spectra for the three of them are 17200 Ix. up to 18300. IN).," Very much like in the case of K 4-47, these PNe have considerably high electron temperatures (average $T_e$, from optical, ISO and IUE spectra for the three of them are 17200 K up to 18300 K)."
+ Therefore. we have he somewhat contradictory result that while Ix 4-47. seems o be an extreme bipolar PN in terms of its morphology and He anc N abundances. its O. Ne and S abundances rather resemble PNe in the Galactic halo. as its relatively arge height on the Galactic plane might also suggest. (see Section 1). ancl contrary to bipolar PNe which are highly concentrated toward the Galactic plane (Corradi Schwarz 1995).," Therefore, we have the somewhat contradictory result that while K 4-47 seems to be an extreme bipolar PN in terms of its morphology and He and N abundances, its O, Ne and S abundances rather resemble PNe in the Galactic halo, as its relatively large height on the Galactic plane might also suggest (see Section 1), and contrary to bipolar PNe which are highly concentrated toward the Galactic plane (Corradi Schwarz 1995)."
+ To further investigate these issues. in the next two Sections we will model the sspectrum using. first. a shock excitation code. and. then. à pure photoionization moclel.," To further investigate these issues, in the next two Sections we will model the spectrum using, first, a shock excitation code, and, then, a pure photoionization model."
+ Shock models have successfully reproduced. the emission line ratios. line profiles ancl. recently. velocity maps in LLL objects (Beck et al.," Shock models have successfully reproduced the emission line ratios, line profiles and, recently, velocity maps in HH objects (Beck et al."
+ 2004)., 2004).
+ To check the possible inlluence of shock excitation on the sspectrum we have explored. ai variety of steady γα] models. for which we use the photoionization-shock code ALAPPINGS Le (Dopita. Binette Tuohy 1984: Dinette. Dopita Tuohy 1985).," To check the possible influence of shock excitation on the spectrum we have explored a variety of steady plane-parallel models, for which we use the photoionization-shock code MAPPINGS Ic (Dopita, Binette Tuohy 1984; Binette, Dopita Tuohy 1985)."
+ One important characteristic of shock mocels is that the oreclictecl spectra strongly. depends on the preshock tioon., One important characteristic of shock models is that the predicted spectra strongly depends on the preshock on.
+ We have compiled a set of fully pre-ionized shock models (i.c. the incident gas has been ionized) and a set of ocal equlibrium preionization for the gas entering the shock (as ceseribed by Shull Melxee 1979. Llartigan et al.," We have compiled a set of fully pre-ionized shock models (i.e. the incident gas has been ionized) and a set of local equlibrium preionization for the gas entering the shock (as described by Shull McKee 1979, Hartigan et al."
+ LOST)., 1987).
+" telative abundances typical of ""Evpe E PNe were assumed. and two values for the pre-shock density were considered (100. 1000 )."," Relative abundances typical of Type I PNe were assumed, and two values for the pre-shock density were considered (100, 1000 $^{-3}$ )."
+ The input shock velocities were varied rom 95 to HO Km sο. ic. the range producing large rratios (Llartigan A)et al.," The input shock velocities were varied from 95 to 140 km $^{-1}$, i.e. the range producing large ratios (Hartigan et al."
+ 1987). as observed.," 1987), as observed."
+ Note. however. hat the measured expansion velocity of the iis smaller than ((Corracli et al.," Note, however, that the measured expansion velocity of the is smaller than (Corradi et al."
+ 2000). pointing to lower shock velocities.," 2000), pointing to lower shock velocities."
+ None of the predicted spectra could reproduce the large observed ratios of intermediate to high excitation emission ines with respect a Η Balmer line (ic./1147..5007A//1L2... Ne 111] aand He 4686/X//10L3)).," None of the predicted spectra could reproduce the large observed ratios of intermediate to high excitation emission lines with respect a H Balmer line (i.e., [Ne ] and He )."
+ Moreover. the shocked. spectra xedieted by our models imply aand larger than the observed.| values.," Moreover, the shocked spectra predicted by our models imply and larger than the observed values."
+ A Iarge discrepancy was also found. with the observed cemission line ratio. which is more than 2 times/ larger than he model values.," A large discrepancy was also found with the observed emission line ratio, which is more than 2 times larger than the model values."
+ And. finally. iis very high. = 300. because the intensity of iis largely uncderstimated by these mocdels.," And, finally, is very high, $\approx$ 300, because the intensity of is largely understimated by these models."
+ We conclude that the plane-parallel shock models are not able to reproduce the spectrum of the core of Ix 4-47., We conclude that the plane-parallel shock models are not able to reproduce the spectrum of the core of K 4-47.
+ We have used the phototonization code CLOUDY 95.06 (Forlancl et al., We have used the photoionization code CLOUDY 95.06 (Ferland et al.
+ 1998)., 1998).
+ CLOUDY needs as input the information on the shape and intensity of the radiation from the ionizing source. the chemical composition and geometry of the nebula. as well as its censity and size.," CLOUDY needs as input the information on the shape and intensity of the radiation from the ionizing source, the chemical composition and geometry of the nebula, as well as its density and size."
+ As mentioned in Section 1. the distance of Ix 4-47. and thus its size and the luminosity of the central star. are poorly known.," As mentioned in Section 1, the distance of K 4-47, and thus its size and the luminosity of the central star, are poorly known."
+ We adopt the distance of 5.9 kpc computed by TFajitsu Tamura (1998). although extending he calculations to the full range of distances from 3 kpc to 7 kpe proposed. by Corradi et al. (," We adopt the distance of 5.9 kpc computed by Tajitsu Tamura (1998), although extending the calculations to the full range of distances from 3 kpc to 7 kpc proposed by Corradi et al. ("
+2000).,2000).
+ A spherical geometry. ancl a filling factor of 1.0 have en assumed.," A spherical geometry, and a filling factor of 1.0 have been assumed."
+ Using the Ho and images from Corradi ct al. (, Using the $\alpha$ and images from Corradi et al. (
+2000) a core diameter. of 1.0 aresee was estimated contour extension corrected or the finite resolution: Tvlenda et al.,2000) a core diameter of 1.9 arcsec was estimated contour extension corrected for the finite resolution; Tylenda et al.
+ 2003)., 2003).
+ The density was initially kept to constant ancl equal to its empirical value (1900. em. see. Table. 1).," The density was initially kept to constant and equal to its empirical value (1900 $^{-3}$, see Table 1)."
+ However. as noted in he Introduction. Aacquist Kwok (1990). discovered. a wight and very compact radio core at the centre of Ix. 4- implying a high density of 72000 cm. (assuming an optically thin nebula at 5 CGllz. the distance of 5.9 kpe and ollowing CGoucdis 1982).," However, as noted in the Introduction, Aaquist Kwok (1990) discovered a bright and very compact radio core at the centre of K 4-47, implying a high density of 72000 $^{-3}$ (assuming an optically thin nebula at 5 GHz, the distance of 5.9 kpc and following Goudis 1982)."
+ For this reason. models with much ugher density values were also explored.," For this reason, models with much higher density values were also explored."
+ Dust grains have been included. since they have an important effect. particularly on the temperature structure of PNe.," Dust grains have been included, since they have an important effect, particularly on the temperature structure of PNe."
+ This effect depends on the tvpe of grain (graphite and/or silicate). the grain abundances. and. the grain-size distribution. being more relevant in the inner regions of the nebula (see Dopita Sutherland 2000).," This effect depends on the type of grain (graphite and/or silicate), the grain abundances, and the grain-size distribution, being more relevant in the inner regions of the nebula (see Dopita Sutherland 2000)."
+ Both graphite anc silicate grains were considered., Both graphite and silicate grains were considered.
+ As we do not know the grain size distribution and grain abundances for Ix 4-47 and more in general for PNe. we adotped ISM size distribution. with ISAL gas-phase depletions due to grains. following van Lloof et al. (," As we do not know the grain size distribution and grain abundances for K 4-47 and more in general for PNe, we adotped ISM size distribution, with ISM gas-phase depletions due to grains, following van Hoof et al. ("
+2004).,2004).
+ We adopted a blackbocs spectrum for the central star. with celerived from theLL3.. aand ecemission line ratios. using a modified. Zanstra method as done for instance by Alikolajewska ct al. (," We adopted a blackbody spectrum for the central star, with derived from the, and emission line ratios, using a modified Zanstra method as done for instance by ajewska et al. ("
+1999).,1999).
+" Dhis eives Z,.rr—11300000 Ix. (from 4686AX//1L3)). and 0000 Ix. (πο 4686A//5876)).", This gives 000 K (from ) and 000 K (from ).
+ AA Lower limit to the central star luminosity comes from the LIGAS spectral energy distribution (Lajitsu Tamura 1998). which gives Lz 16 D? L. (where D is the distance in kpe) vielding 550 L. at 5.9 kpc.," A lower limit to the central star luminosity comes from the IRAS spectral energy distribution (Tajitsu Tamura 1998), which gives $\ge$ 16 $^2$ $L_{\odot}$ (where D is the distance in kpc) yielding 550 $L_{\odot}$ at 5.9 kpc."
+" As with the elemental abundances. we first run models with the empirical. abundances given in ‘Table 2. and subsequently with ""average"" values for either normal (Ivpe-1D. or Type-L PNe. following Ixingsburgh Barlow (1994)."," As with the elemental abundances, we first run models with the empirical abundances given in Table 2, and subsequently with “average” values for either normal (Type-II), or Type-I PNe, following Kingsburgh Barlow (1994)."
+ CLOUDY models with the empirical abundances cdo not match the observed. line. Duxes., CLOUDY models with the empirical abundances do not match the observed line fluxes.
+. Both helium and (6548 omission lines are largely overestimated. while oxygen. sulphur and neon lines are uncerestimated. mainlv because the very low Of. Νο anc 41 input abundances are not compensated. by ai high electron temperature in the model as empiricallv determined.," Both helium and ) emission lines are largely overestimated, while oxygen, sulphur and neon lines are underestimated, mainly because the very low O/H, Ne/H and S/H input abundances are not compensated by a high electron temperature in the model as empirically determined."
+ Tests with Pype-LE PNe abundances were also not. successful because the predicted intensities are very low. which is expected as the main difference between “Pype-Lb and Ἔνροη abundances is that the latter are depleted. in. nitrogen by ao factor of —4 (dingsburgh Barlow 1994).," Tests with Type-II PNe abundances were also not successful because the predicted intensities are very low, which is expected as the main difference between Type-I and Type-II abundances is that the latter are depleted in nitrogen by a factor of $\sim$ 4 (Kingsburgh Barlow 1994)."
+ Assuming, Assuming
+at 23 make this estimate far too optimistic. but all of the models that match the observations require rj to be no largere than 500 kpe (which is likely still too optimistic:1 sec JaguiOla}). and are thus compatible with constraints placed on the metal distribution by travel-time arguments.,"at $z\gg 3$ make this estimate far too optimistic, but all of the models that match the observations require $\renrich$ to be no larger than $500$ kpc (which is likely still too optimistic; see ), and are thus compatible with constraints placed on the metal distribution by travel-time arguments."
+ Optical depth distributions are caleulated. by firing 10 randomly. chosen lines of sight through the simulation volume and caleulating absorption spectra for both Ε1 and CIN. following the procedure outlined in e.g. Appendix At of(10998)., Optical depth distributions are calculated by firing $10^3$ randomly chosen lines of sight through the simulation volume and calculating absorption spectra for both $\hi$ and $\civ$ following the procedure outlined in e.g. Appendix A4 of.
+. The mean HI optical depth in our simulations al z=3 is map=0.388. which is consistent with observations2008).," The mean $\hi$ optical depth in our simulations at $z=3$ is $\tau_{\rm eff}=0.388$, which is consistent with observations."
+. In order to match observations with ILLIS on the Weck telescope. we convolve our spectra with a Gaussian line-spreacd function with a full-width-at-half-maximun of 6.6 km/s. and resample our spectra onto 1 kms pixels.," In order to match observations with HIRES on the Keck telescope, we convolve our spectra with a Gaussian line-spread function with a full-width-at-half-maximum of 6.6 km/s, and resample our spectra onto 1 km/s pixels."
+ We do not add. noise to our spectra. but we have verified. that the addition of Gaussian. noise with a signal-to-noise ratio of greater than 25 docs not significantly alect anv of our results.," We do not add noise to our spectra, but we have verified that the addition of Gaussian noise with a signal-to-noise ratio of greater than 25 does not significantly affect any of our results."
+ We generate absorption spectra for two transitions: ILI (1215.67A)) and the €IN. doublet(1548.20A.. 1550.78A)).," We generate absorption spectra for two transitions: $\hi$ ) and the $\civ$ doublet, )."
+ In order to compare to the observed TOITH. distribution. we then bin pixels in 7Hi and calculate the median. zciv. corresponding to the redshifts of the pixels in each ILI bin.," In order to compare to the observed $\tauciv-\tauhi$ distribution, we then bin pixels in $\tauhi$ and calculate the median $\tauciv$ corresponding to the redshifts of the pixels in each $\hi$ bin."
+" For each run. in addition to measuring the row7Hi relation. we calculate the fraction of the total mass (f) and volume (fy) that has been enriched from πο m; and h; are the SPLL particle mass and smoothing kernel. respectively, anc the sums in the numerator of cach fraction extend only over particles with non-zero metallicity."," For each run, in addition to measuring the $\tauciv-\tauhi$ relation, we calculate the fraction of the total mass $\fmass$ ) and volume $\fvol$ ) that has been enriched from Here, $m_i$ and $h_i$ are the SPH particle mass and smoothing kernel, respectively, and the sums in the numerator of each fraction extend only over particles with non-zero metallicity."
+ We have verified that using ορ; instead of A? gives nearly identical volume filling fractions. as expected.," We have verified that using $m_i/\rho_i$ instead of $h_i^3$ gives nearly identical volume filling fractions, as expected."
+ For the solar we use the metal mass fraction Z.=0.0127., For the solar we use the metal mass fraction $\zsun=0.0127$.
+ Fig., Fig.
+ 1. shows a thin (1 comoving Mpc// thick) slice of the density field at z=3 through the centre of the simulation., \ref{fig:density} shows a thin (1 comoving $h$ thick) slice of the density field at $z=3$ through the centre of the simulation.
+ Each panel shows a dillerent. combination of rz and m., Each panel shows a different combination of $\renrich$ and $\menrich$.
+ Gas with zero metallicity is shown in grev-scale. while metal-enriched. gas is coloured red.," Gas with zero metallicity is shown in grey-scale, while metal-enriched gas is coloured red."
+" In. addition. the values for fy and fi, are indicated for cach model."," In addition, the values for $\fvol$ and $\fmass$ are indicated for each model."
+ Panels showing metal distributions that are consistent with the observations (see below) are outlined in green., Panels showing metal distributions that are consistent with the observations (see below) are outlined in green.
+ The fractions of the mass and volume that are enriched do not track each other in a simple wav., The fractions of the mass and volume that are enriched do not track each other in a simple way.
+ The ratio fi/fy is always Z1 because metals are only placed around collapsed structures and thus preferentially in overdense regions.," The ratio $\fmass/\fvol$ is always $\ge
+1$ because metals are only placed around collapsed structures and thus preferentially in overdense regions."
+" For m,=üLOYAL. changingH r; from. kkpe to kkpe changes the ratio fii/fy [rom 22.7 to 1.3. as a larger value of rz; at fixed. mz; allows metals to disperse further out of the high-density. peaks."," For $\menrich=10^{10}\,\msun$, changing $\renrich$ from kpc to kpc changes the ratio $\fmass/\fvol$ from 22.7 to 1.3, as a larger value of $\renrich$ at fixed $\menrich$ allows metals to disperse further out of the high-density peaks."
+" Similarly. changing η, at fixed r,=250 kpe. we find that the ratio fu/fv rises from. 2.8 for m,=10""M. to 7.5 for m4=104M. because more massive haloes are preferentially located in higher density environments."," Similarly, changing $\menrich$ at fixed $\renrich=250$ kpc, we find that the ratio $\fmass/\fvol$ rises from 2.8 for $\menrich=10^{9}\,\msun$ to 7.5 for $\menrich=10^{11}\,\msun$, because more massive haloes are preferentially located in higher density environments."
+ Given that we understand how changing the pattern of metal enrichment. alters the relative mass and volume filling factors. we now ask how this impacts the rouTHa relation.," Given that we understand how changing the pattern of metal enrichment alters the relative mass and volume filling factors, we now ask how this impacts the $\tauciv-\tauhi$ relation."
+ The curves in Fig., The curves in Fig.
+ 20 show the relation between rci and τη in the svnthetic absorption spectra., \ref{fig:od} show the relation between $\tauciv$ and $\tauhi$ in the synthetic absorption spectra.
+ Each panel corresponds to a dillerent mz., Each panel corresponds to a different $\menrich$.
+" The solid lines in panels6.. ancl (log,(0;/M.)=9.10.11. respectively) show the predicted rou7Hi relation from the simulation with various imposed metallicity distributions."," The solid lines in panels, and $\log_{10}(\menrich/\msun)={9,10,11}$, respectively) show the predicted $\tauciv-\tauhi$ relation from the simulation with various imposed metallicity distributions."
+" The cot-dashed— lines— in. panels— ancl (log,(n,/M.)28.9. respectively) show the predicted roaτην relations for the small volume. high-resolution simulation.LO0L2N212."," The dot-dashed lines in panels and $\log_{10}(\menrich/\msun)={8,9}$, respectively) show the predicted $\tauciv-\tauhi$ relations for the small volume, high-resolution simulation,."
+. In every panel we compare. our simulated: predictions to the observed. optical depth: pixel statistics of(2003)... for the redshift range 2479<2x4.033. as published in(2005). shown as the vellow points with Lo error bars.," In every panel we compare our simulated predictions to the observed optical depth pixel statistics of, for the redshift range $2.479 \le z \le 4.033$, as published in, shown as the yellow points with $1\sigma$ error bars."
+ The data come from six quasar spectra. 3388. Q1425|604. 1158. ΟΡΟ|230. 2269. and Q1055|461. that were taken with either the Weck/LURIES or the VLTE/UVES.," The data come from six quasar spectra, 388, Q1425+604, 158, Q1422+230, 269, and Q1055+461, that were taken with either the Keck/HIRES or the VLT/UVES."
+ A full description of the sample is given in(2003)., A full description of the sample is given in.
+. In each panel Z was chosen such that LOW10 ry=500 kpe curve exactly matches the observations at 1 highest value of zHi and the metallicity required for this normalisation is given in cach panel., In each panel $\zenrich$ was chosen such that the $\renrich=500$ kpc curve exactly matches the observations at the highest value of $\tauhi$ and the metallicity required for this normalisation is given in each panel.
+ In the present work we are not aiming to reproduce the garape of the τονTHi relation in detail and. indeed. would not necessarily expect our simple models to be capable of oing this.," In the present work we are not aiming to reproduce the shape of the $\tauciv-\tauhi$ relation in detail and, indeed, would not necessarily expect our simple models to be capable of doing this."
+ Rather. we require that the models. predict a median τονTH relation that is consistent with. or larger. aan observed and that the rouyτηι relation is not steeper aan observed.," Rather, we require that the models predict a median $\tauciv-\tauhi$ relation that is consistent with, or larger, than observed and that the $\tauciv-\tauhi$ relation is not steeper than observed."
+" We elfectively combine the two conditions bv scaling Z in cach run to match the normalisation of we observed: τοντΗι relation at log),TH1=2.5. the largest optical depth probed. by the observations. and then requiring the model to predict median τον at lowτη tha are consistent with or greater than observed."," We effectively combine the two conditions by scaling $\zenrich$ in each run to match the normalisation of the observed $\tauciv-\tauhi$ relation at $\log_{10}\tauhi=2.5$, the largest optical depth probed by the observations, and then requiring the model to predict median $\tauciv$ at low$\tauhi$ that are consistent with or greater than observed."
+ We stress that overprediction of τον is not a problem a small τη because our simple models make the assumption that the metallicity inside each enriched bubble is constan with radius., We stress that overprediction of $\tauciv$ is not a problem at small $\tauhi$ because our simple models make the assumption that the metallicity inside each enriched bubble is constant with radius.
+ Phe overestimate at low τον could therefore be solved by imposing a metallicity that decreases with radius., The overestimate at low $\tauciv$ could therefore be solved by imposing a metallicity that decreases with radius.
+ On the other hand. we will assume that underprediction of roa at low rHi signals the failure of the model because a metallicity that increases with radius is likely unphysical.," On the other hand, we will assume that underprediction of $\tauciv$ at low $\tauhi$ signals the failure of the model because a metallicity that increases with radius is likely unphysical."
+ Such an unphysical metallicity gradient would. also have been required. if we hack chosen to scale the metallicity to match the low 7H points. because the unsuccessful. constant metallicity models predict much steeper τονΤΗ relations than observed (see Fig. 2)).," Such an unphysical metallicity gradient would also have been required if we had chosen to scale the metallicity to match the low $\tauhi$ points, because the unsuccessful, constant metallicity models predict much steeper $\tauciv-\tauhi$ relations than observed (see Fig. \ref{fig:od}) )."
+ In fact. such an approach would not be possible for most of the models that we rule out. because they typically prediet the median rou to be zero at low ΤΗ 1.," In fact, such an approach would not be possible for most of the models that we rule out, because they typically predict the median $\tauciv$ to be zero at low $\tauhi$ ."
+ 1n principle. we could produce significantly better fits to," In principle, we could produce significantly better fits to"
+ , 
+"space FE, ?σι,can be identified with⋅⋅ the⋅ parabolic Hardy⋅ space HP*(Eg? 1).∣↽⊓1 ⋅<p<x.having theThe following square function characterizatiou stated informally as: This identification between the above two spaces is the following: Another useful space throughout our analysis is the parabolic inhomogeneous Besov space.","The space $\dot{F}^{0,a}_{p,2}$ can be identified with the parabolic Hardy space $H^{p,a}(\R^{n+1})$, $1\leq p<\infty$, having the following square function characterization stated informally as: This identification between the above two spaces is the following: Another useful space throughout our analysis is the parabolic inhomogeneous Besov space."
+ The main difference in defining this space is the choice of the parabolic dyadic partition of unity that is now altered., The main difference in defining this space is the choice of the parabolic dyadic partition of unity that is now altered.
+" Indeed. we take (ej)joo satisfying: Again.it is clear that vieοeyle)=1. Paleybut now for all z€IE""land in exactly the same way as above. we cau rewrite the Littlewood: decomposition with We then arrive to the following definition: For a detailed study of anisotropic Lizorkin-Triebel auc Besov spaces. we refer the reader to Triebel [21]."," Indeed, we take $(\psi_{j})_{j\geq 0}$ satisfying: Again,it is clear that $\sum_{j\geq 0} \psi_{j}(z) = 1$ , but now for all $z\in \R^{n+1}$, and in exactly the same way as above, we can rewrite the Littlewood-Paley decomposition with We then arrive to the following definition: For a detailed study of anisotropic Lizorkin-Triebel and Besov spaces, we refer the reader to Triebel \cite{Tri1}."
+. We present two embedding resultsfrom Johnsen and Sickel [11].. aud Stécekert [19]..," We present two embedding resultsfrom Johnsen and Sickel \cite{Joh_Sic07}, and Stöcckert \cite{Stockert82}. ."
+mechanism for the production of elliptical galaxies then they must be able to produce the correct scaling relations as well as the evolution of these scaling relations over time.,mechanism for the production of elliptical galaxies then they must be able to produce the correct scaling relations as well as the evolution of these scaling relations over time.
+" ""Theoretical.hydrodynamical studies have shown that simulations of gas-rich galaxy mergers are capable of reproducing the observed scaling relations of elliptical galaxies if the correct. progenitor properties are used (???7?).."," Theoretical studies have shown that simulations of gas-rich galaxy mergers are capable of reproducing the observed scaling relations of elliptical galaxies if the correct progenitor properties are used \citep{Dekel06, RobertsonFP,
+ HopkinsFP,Bournaud:2011a}."
+ This is a step toward verifving the production. of scaling laws through mergers., This is a step toward verifying the production of scaling laws through mergers.
+ However. current computing power only allows the simulation of relatively small numbers of mergers. and the space of possible merger initial conditions and progenitor properties is quite large.," However, current computing power only allows the simulation of relatively small numbers of mergers, and the space of possible merger initial conditions and progenitor properties is quite large."
+ Furthermore. these simulations are not placed. within a cosmological context. making it more cdillieult to explore in detail the origin and evolution of scaling relations.," Furthermore, these simulations are not placed within a cosmological context, making it more difficult to explore in detail the origin and evolution of scaling relations."
+ Currently. the primary theoretical tool for. studving he evolution of statistical samples of galaxies over cosmological timescales is semi-analvtic modeling. (SAM) (7???7?TTTTTTT?3) ," Currently, the primary theoretical tool for studying the evolution of statistical samples of galaxies over cosmological timescales is semi-analytic modeling (SAM) \citep{kwg1993, Cole94, SP1999, Cole00,
+ Galics03, Croton06, DeLucia06, Bower06, S08, Fontanot09, Benson10,
+ BensonBower10, Cook10, Guo10}."
+These models. combine dark matter ido merger trees with analytic recipes for populating the ialos with galaxies., These models combine dark matter halo merger trees with analytic recipes for populating the halos with galaxies.
+ Llowever. semi-analvtic models (SAAIs) do not currently incorporate realistic formulas for predicting he properties of the remnants of galaxy mergers including he effects of dissipation.," However, semi-analytic models (SAMs) do not currently incorporate realistic formulas for predicting the properties of the remnants of galaxy mergers including the effects of dissipation."
+ The agreement between. current SAMSs ancl observed. earlv-tvpe. sealing relations is) not impressive., The agreement between current SAMs and observed early-type scaling relations is not impressive.
+ In. particular. the observed: size-mass relation of earlv-tvpes is steeper than that of their potential type progenitors (7).. and the scatter in the observed. size- relation for earlv-twpes is remarkably small (??)..," In particular, the observed size-mass relation of early-types is steeper than that of their potential late-type progenitors \citep{Shen03}, and the scatter in the observed size-mass relation for early-types is remarkably small \citep{Shen03,
+ Nair10}."
+ The disspationless merger models currently. implemented within SAAS have thus far been unable to reproduce these features (o...77)..," The disspationless merger models currently implemented within SAMs have thus far been unable to reproduce these features \citep[e.g.,][]{Shankar10a,Guo10}."
+ 7. recently. developed a physicallv-motivated: analytic model for predicting the stellar hall-mass radii and central. velocity dispersions of merger remnants (?).., \citet{remnants} recently developed a physically-motivated analytic model for predicting the stellar half-mass radii and central velocity dispersions of merger remnants \citep{remnants}.
+ The parameters in this new merger model were calibrated using a suite of galaxy merecr simulations (see Section 2))., The parameters in this new merger model were calibrated using a suite of galaxy merger simulations (see Section \ref{sec:methods2}) ).
+ Llere we implement a simplified. version of this model using post-processing of merger outputs from the SAMSs developed by 2.. based on he Millennium simulation (?).. and ?..," Here we implement a simplified version of this model using post-processing of merger outputs from the SAMs developed by \citet{Croton06}, based on the Millennium simulation \citep{Springel05}, and \citet{S08}."
+ This results in a »opulation of tens of thousands of merger. remnants. over a lage range of redshifts (0«z 3) complete with oedieted values of size. stellar mass. and velocity cüspersion.," This results in a population of tens of thousands of merger remnants over a large range of redshifts $0<z<3$ ) complete with predicted values of size, stellar mass, and velocity dispersion."
+ Comparison of the mocelecd population of ellipticals with he observed scaling relations provides an important. test of the merger hypothesis as well as physical insight. into he origin and evolution of these relations via merging., Comparison of the modeled population of ellipticals with the observed scaling relations provides an important test of the merger hypothesis as well as physical insight into the origin and evolution of these relations via merging.
+ Future work will implement our merger model self-consistentlv. within the semi-analvtic machinery rather than by post-processing., Future work will implement our merger model self-consistently within the semi-analytic machinery rather than by post-processing.
+ Section 2.1 describes our analytic merger mocel for calculating the properties of stellar spheroids from galaxy mergers including energy. losses from dissipation., Section 2.1 describes our analytic merger model for calculating the properties of stellar spheroids from galaxy mergers including energy losses from dissipation.
+ Section 2.2 explains how we implement our merger model using outputs from the ? and ? SAAIs.," Section 2.2 explains how we implement our merger model using outputs from the \citet{Croton06} and \citet{S08}
+ SAMs."
+ Section 3 systematically explores the elfects of the merger model., Section 3 systematically explores the effects of the merger model.
+ In order to turn the rather shallow size-mass relation of disk galaxies into the steeper size-mass relation of observed. earlv-tvpe. galaxies. we find that idt is essential to include both. dissipation and the decreasing eas content of more massive progenitor disk galaxies.," In order to turn the rather shallow size-mass relation of disk galaxies into the steeper size-mass relation of observed early-type galaxies, we find that it is essential to include both dissipation and the decreasing gas content of more massive progenitor disk galaxies."
+ In Section 4 we summarize the observational results that we compare with our model outputs. focusing especially on Sloan Digital Sky Survey (SDSS) data. [or nearby galaxy size vs. mass ancl other scaling relations. and on cata from several surveys (7). for the evolution of these relations to higher redshifts.," In Section 4 we summarize the observational results that we compare with our model outputs, focusing especially on Sloan Digital Sky Survey (SDSS) data for nearby galaxy size vs. mass and other scaling relations, and on data from several surveys \citep{Trujillo06} for the evolution of these relations to higher redshifts."
+ In section 5 we use the predicted properties of progenitor disk galaxies from the two semi-analvtie models to predict the size-mass and other relations for spheroids formed in major gas-rich mergers and compare them. with the observations out to redshift =3., In section 5 we use the predicted properties of progenitor disk galaxies from the two semi-analytic models to predict the size-mass and other relations for spheroids formed in major gas-rich mergers and compare them with the observations out to redshift $z=3$.
+ Section 6 summarizes these results and. discusses their implications and. some follow-on studies that are in progress., Section 6 summarizes these results and discusses their implications and some follow-on studies that are in progress.
+ Finally. an Appendix. describes an. improvement in the treatment of the central clark matter in the analytic merger model of 7. that we used in ? and in the present paper.," Finally, an Appendix describes an improvement in the treatment of the central dark matter in the analytic merger model of \citet{remnants} that we used in \citet{Covington08} and in the present paper."
+ We use a combination of modeling approaches to construct a theoretical framework for predicting the evolution of carly-tvpe sealing relations over cosmological time., We use a combination of modeling approaches to construct a theoretical framework for predicting the evolution of early-type scaling relations over cosmological time.
+ In. previous work. a large suite of hwdrodynamical galaxy merger simulations were developed (222).," In previous work, a large suite of hydrodynamical galaxy merger simulations were developed \citep{thesis,
+ Cox05, Cox08}."
+ These simulations were performed using the N-body/SPLI code GADGET (77). and include hyverodyvnamies. star formation. and stellar feecback.," These simulations were performed using the N-body/SPH code GADGET \citep{SpGad,SH03}, and include hydrodynamics, star formation, and stellar feedback."
+ The simulation suite contains mergers with a wide variety of progenitor properties. mass ratios. and merger orbits.," The simulation suite contains mergers with a wide variety of progenitor properties, mass ratios, and merger orbits."
+ Variations in progenitor properties include a range of stellar misses. gas fractions. dark matter halo concentrations. bulge fractions. barvonic fractions. and gas clisk sizes.," Variations in progenitor properties include a range of stellar masses, gas fractions, dark matter halo concentrations, bulge fractions, baryonic fractions, and gas disk sizes."
+ 1n subsequent work. Covington ct al. (," In subsequent work, Covington et al. ("
+2008. hereafter COS) constructed à. physicallv-motivated: analytic galaxy merger model capable of predicting the half-mass. raclii. stellar masses. and velocity dispersions of galaxy merger remnants given the properties of the progenitor ealaxies and the initial orbits of the mergers.,"2008, hereafter C08) constructed a physically-motivated analytic galaxy merger model capable of predicting the half-mass radii, stellar masses, and velocity dispersions of galaxy merger remnants given the properties of the progenitor galaxies and the initial orbits of the mergers."
+ This model was calibrated using the galaxy merger suite described. above., This model was calibrated using the galaxy merger suite described above.
+ Unlike previous similar models (??).. this mocel includes the ellects of star formation and energy loss due to dissipation.," Unlike previous similar models \citep{Cole00, Galics03}, this model includes the effects of star formation and energy loss due to dissipation."
+ llere we combine a simplified. version of this new galaxy merger model with merger rates and. progenitor properties predicted: by two semi-analvtic models (SXMSs) in order to explore the ereation and evolution of the scaling relations of earlv-tvpe galaxies., Here we combine a simplified version of this new galaxy merger model with merger rates and progenitor properties predicted by two semi-analytic models (SAMs) in order to explore the creation and evolution of the scaling relations of early-type galaxies.
+"using them comes from our numerical method, which calculates configurations based on a given distortion (axis ratio).","using them comes from our numerical method, which calculates configurations based on a given distortion (axis ratio)."
+" Since magnetic distortions are typically very small, this means that even the axis ratio for the least non-spherical star we can specify (constrained by the grid resolution) corresponds to a very strong magnetic field."," Since magnetic distortions are typically very small, this means that even the axis ratio for the least non-spherical star we can specify (constrained by the grid resolution) corresponds to a very strong magnetic field."
+" This is not necessarily a problem for normal MHD, but superconductivity will be broken at these field strengths, which exceed the second critical field of Πω~10!9 G. We are able to produce these models because the destruction of superconductivity is not built into them; the equations may be solved for any field strength."," This is not necessarily a problem for normal MHD, but superconductivity will be broken at these field strengths, which exceed the second critical field of $H_{c2}\approx 10^{16}$ G. We are able to produce these models because the destruction of superconductivity is not built into them; the equations may be solved for any field strength."
+" Having done so, however, we need to check that these models are consistent with our expectations for NSs with superconducting protons."," Having done so, however, we need to check that these models are consistent with our expectations for NSs with superconducting protons."
+" If they are, we may extrapolate our results back to more realistic models, where B«Πω."," If they are, we may extrapolate our results back to more realistic models, where $\Bav<\bar{H}_{c2}$."
+" We perform this sanity check next, in figure 13.."," We perform this sanity check next, in figure \ref{toroidal_ellips}."
+" We plot the scaling of ellipticity with average field strength B for NS models with toroidal magnetic fields, for normal (left) and type-II superconducting protons (right)."," We plot the scaling of ellipticity with average field strength $\Bav$ for NS models with toroidal magnetic fields, for normal (left) and type-II superconducting protons (right)."
+" Note that all ellipticities are negative, since these configurations are prolate."," Note that all ellipticities are negative, since these configurations are prolate."
+" In the left-hand plot we find that the measured numerical values agree very well with the expected scaling eox B?, with small deviations when B>10!"" G. We now turn to the right-hand plot of figure 13.."," In the left-hand plot we find that the measured numerical values agree very well with the expected scaling $\epsilon\propto B^2$ , with small deviations when $\Bav\gtrsim
+10^{17}$ G. We now turn to the right-hand plot of figure \ref{toroidal_ellips}."
+" This time we plot the ellipticity against B, not B?."," This time we plot the ellipticity against $\Bav$, not $\Bav^2$."
+" First we look at configurations with ¢?-superconductivity, for a central critical field value H«1(0)=1015 G (the points on the line marked (a)) and also for H.1(0)=2x1016 G (the points on line (c))."," First we look at configurations with $\zeta^2$ -superconductivity, for a central critical field value $H_{c1}(0)=10^{16}$ G (the points on the line marked (a)) and also for $H_{c1}(0)=2\times 10^{16}$ G (the points on line (c))."
+" In both cases we see that the points lie virtually on straight lines, showing that εοςB."," In both cases we see that the points lie virtually on straight lines, showing that $\epsilon\propto\Bav$."
+" In addition, line (c) has twice the gradient of line (a)."," In addition, line (c) has twice the gradient of line (a)."
+" This confirms that despite the high field strengths we are obliged to use (see the discussion above), we find the correct scaling of the ellipticity: eoxH«1B."," This confirms that despite the high field strengths we are obliged to use (see the discussion above), we find the correct scaling of the ellipticity: $\epsilon\propto H_{c1}\Bav$."
+ This gives us more confidence about our results and means we can safely extrapolate to more typical NS field strengths., This gives us more confidence about our results and means we can safely extrapolate to more typical NS field strengths.
+ We also present ellipticities for ¢3-superconductivity (points along line (b))., We also present ellipticities for $\zeta^3$ -superconductivity (points along line (b)).
+ The level of distortion inthis case is very similar to that in the C? case., The level of distortion inthis case is very similar to that in the $\zeta^2$ case.
+nmeaniugs as the TLOWWSOLL Cross section. electrou iuass and the speed of light.,"meanings as the Thomson cross section, electron mass and the speed of light."
+ Suce ouly [ree electrous pa‘ticipate in Thomson scattering of the CMB photous. le integra Is €uw olf at the epoch of reiouization fa.," Since only free electrons participate in Thomson scattering of the CMB photons, the integral is cut off at the epoch of reionization $l_{\rm re}$."
+ Sole Iniportant statistics of the SZ elfect inunediately. come to 1iud., Some important statistics of the SZ effect immediately come to mind.
+" The first oO‘ler quantity is y. the mean y pa""anieter average over the whole sky. which measwes the toal theral energy content ofthe univewe."," The first order quantity is $\bar{y}$, the mean $y$ parameter averaged over the whole sky, which measures the total thermal energy content of the universe."
+ Tie angular variation in y can be parameterizec by the two poiit correlation function of le temperatwe ποion ©. or equivaleutly the auguar power sj)ectrun Cy.," The angular variation in $y$ can be parameterized by the two point correlation function of the temperature fluctuation $\Theta$, or equivalently the angular power spectrum $C_l$."
+ For a Claussiall Fallcdoum field. these (wo pa‘ameters would describe the statisics completely.," For a Gaussian random field, these two parameters would describe the statistics completely."
+ Since the 5Z effect is dominated by nolinear structures. non-Craussiaulty may be siguificaut.," Since the SZ effect is dominated by non-linear structures, non-Gaussianity may be significant."
+ So we investigate higher order statistics su] as the ssewness aud kurtosis of the y parameter to quantify the , So we investigate higher order statistics such as the skewness and kurtosis of the $y$ parameter to quantify the non-Gaussianity.
+The SZ elect contalus coutributious [rom all ‘edsüfts and it is challengiug to recover the sineared. redshif informaion., The SZ effect contains contributions from all redshifts and it is challenging to recover the smeared redshift information.
+ We have shown that cross Correlating the SZ ellect with a galaxy photometric recift survey. we can iier the redshilt “ESOvec IGM pressure-galaxy cross correlatjon aud the IGM prSSILEO allo correlation (Zhang&Pe12001)..," We have shown that cross correlating the SZ effect with a galaxy photometric redshift survey, we can infer the redshift resolved IGM pressure-galaxy cross correlation and the IGM pressure auto correlation \citep{Zhang01a}."
+ This method is robust. but does not capture all the iOratio iin the SZ observation.," This method is robust, but does not capture all the information in the SZ observation."
+ Iu this j»aper. we utilize the one point clistributjon function (PDF of he y paraiueter and the distribution of peaks in y to extract more information.," In this paper, we utilize the one point distribution function (PDF) of the $y$ parameter and the distribution of peaks in $y$ to extract more information."
+ Since smoothing is aways present for a real experinent with a finite beam. we calculate the statistics of the y paraineer sinoothed on a given angular scale 0.," Since smoothing is always present for a real experiment with a finite beam, we calculate the statistics of the $y$ parameter smoothed on a given angular scale $\theta$."
+ It we are interestecLin virialized objects. for example clusers aud groups of galaxies. we would expec them to be peaks in the smoothect or filtered πιαps. ViCy>gp).," If we are interested in virialized objects, for example clusters and groups of galaxies, we would expect them to be peaks in the smoothed or filtered $y$ maps. $N(y>y_p)$,"
+ the σιiuulative distribution fuiction (CDF) of peaks with sinoothecl y parameter jeeerMD than certain valle. pp. is the raw observable.," the cumulative distribution function (CDF) of peaks with smoothed $y$ parameter bigger than certain value $y_p$, is the raw observable."
+ The yy €DF is the 5Z analog to a luminosity unction., The $y_p$ CDF is the SZ analog to a luminosity function.
+ I£ we choose a op hat window so that tie observation coue is large ough to include an etire object such as a cluster and is suiall enough hat tydically uo more tlian e such object cau be [οιuid aloug each cone. hen when the cone 15 ceiterecl at the ¢‘enter of each ject. a peak y= ypaj»pears in the sinoothlied nap.," If we choose a top hat window so that the observation cone is large enough to include an entire object such as a cluster and is small enough that typically no more than one such object can be found along each cone, then when the cone is centered at the center of each object, a peak $y\equiv y_p$ appears in the smoothed map."
+ This yp is directly related othe total gas Wass aud temperature of individual object., This $y_p$ is directly related to the total gas mass and temperature of individual object.
+ Asstuning a halo to be isothermal. the to aliass Al ofa halo is related to the gas teuperature T by AL/Als=(F/T4)7.," Assuming a halo to be isothermal, the total mass $M$ of a halo is related to the gas temperature $T$ by $M/M_8=(T/T_8)^{3/2}$."
+ Als=L8x101(054/0.3)/7.IM. is the Inass coutained in a sl!Npe sphere of the universe of mean density today. which is roughly the uiass scale of clusters.," $M_8=1.8\times 10^{14} (\Omega_0/0.3) h^{-1} {\rm M}_{\sun}$ is the mass contained in a $8 {\rm h}^{-1}{\rm Mpc}$ sphere of the universe of mean density today, which is roughly the mass scale of clusters."
+" Zu(z) is the correspoucing temperature of a halo with mass My at reclshilt ο,", $T_8(z)$ is the corresponding temperature of a halo with mass $M_8$ at redshift $z$.
+ We then obtain Here. AQ is the solid angle of the cone. fy is the gas fraction of halos aud d4 is the angular cliameter distance.," We then obtain Here, $\Delta \Omega$ is the solid angle of the cone, $f_g$ is the gas fraction of halos and $d_A$ is the angular diameter distance."
+ For- clusters aud groups. the typicaln augular size. at 2=41 is. about 4/1.," For clusters and groups, the typical angular size at $z=1$ is about $1^{'}$."
+ For- the present cluster number deusity v(T>2keV)~10OfMpc* (Pen1998b).. the average number oL clusters in a cone with augular radius 0~20./ projected to 2 is about one allowing for the evolution of cluster number deusity.," For the present cluster number density $n(T>2 {\rm keV})\sim 
+10^{-5} h^{-3} {\rm Mpc}^3$ \citep{Pen98b}, the average number of clusters in a cone with angular radius $\theta \sim 20^{'}$ projected to $z\sim 2$ is about one allowing for the evolution of cluster number density."
+ So. the size of the smootling scale for the peak analysis should be between these two scales.," So, the size of the smoothing scale for the peak analysis should be between these two scales."
+ In this case. the IN(y2yy) is just the number of halos with gy2 yp.," In this case, the $N(y>y_p)$ is just the number of halos with $y>y_p$ ,"
+(Naze et citeNRM)).,(Naze et \\cite{NRM}) ).
+ Whilst there i$ no obvious X-ray emission from stars A and C. there might be some X-ray emission associated with stars B and D. However. the latter two objects lie in the wings of the X-ray source associated with 1183. and their status as X-ray emitters can therefore not be ascertained with confidence.," Whilst there is no obvious X-ray emission from stars A and C, there might be some X-ray emission associated with stars B and D. However, the latter two objects lie in the wings of the X-ray source associated with 183, and their status as X-ray emitters can therefore not be ascertained with confidence."
+ The lack of X-ray detections for stars A and C is somewhat surprising given their spectral types inferred above., The lack of X-ray detections for stars A and C is somewhat surprising given their spectral types inferred above.
+ One possibility could be a larger interstellar column density. thus larger absorption of the X-ray emission than for other stars in 22 of similar spectral type.," One possibility could be a larger interstellar column density, thus larger absorption of the X-ray emission than for other stars in 2 of similar spectral type."
+ However. this explanation needs to be confirmed using additional high angular resolution optical and X-ray observations.," However, this explanation needs to be confirmed using additional high angular resolution optical and X-ray observations."
+ The various eclipsing binaries that we have studied in this paper yield distance estimates in the range 6.5 to larger than Okkpe., The various eclipsing binaries that we have studied in this paper yield distance estimates in the range 6.5 to larger than kpc.
+ Whilst the uncertainty in these results is. difficult to evaluate. they are in very good agreement with the spectrophotometric distance of (8.0+I4) kkpe inferred in Paper I and clearly consistent with a cluster distance well beyond the kkpe proposed by Ascenso et ((2007)).," Whilst the uncertainty in these results is difficult to evaluate, they are in very good agreement with the spectrophotometric distance of $(8.0 \pm 1.4)$ kpc inferred in Paper I and clearly consistent with a cluster distance well beyond the kpc proposed by Ascenso et \cite{Ascenso}) )."
+ We thus conclude that all available observations of the early-type stellar population of 22 are most consistent with a distance near kkpe in agreement with the distance inferred for 220a., We thus conclude that all available observations of the early-type stellar population of 2 are most consistent with a distance near kpc in agreement with the distance inferred for 20a.
+ This result casts serious doubts on a possible association between the y-ray pulsar JJ1023.0—5746 (Ackermann et citeAckermann)) and the 22 cluster., This result casts serious doubts on a possible association between the $\gamma$ -ray pulsar $-$ 5746 (Ackermann et \\cite{Ackermann}) ) and the 2 cluster.
+ If the distance of the pulsar were confirmed at kkpe. it would then appear more likely that the latter belongs to a foreground stellar population unrelated to the massive stars of 22. but possibly related to the PMS stars discussed by Ascenso et ((2007)). With respect to the multiplicity of early-type stars in 22. we stress that among the eleven O-type stars and one WNha star monitored during our campaign. which were not previously known to be binary systems. only one star 1167) was found to probably be a short-period spectroscopic binary system.," If the distance of the pulsar were confirmed at kpc, it would then appear more likely that the latter belongs to a foreground stellar population unrelated to the massive stars of 2, but possibly related to the PMS stars discussed by Ascenso et \cite{Ascenso}) With respect to the multiplicity of early-type stars in 2, we stress that among the eleven O-type stars and one WNha star monitored during our campaign, which were not previously known to be binary systems, only one star 167) was found to probably be a short-period spectroscopic binary system."
+ The other targets did not display RV variations well above the uncertainties estimated from the RV dispersion of the DIBs., The other targets did not display RV variations well above the uncertainties estimated from the RV dispersion of the DIBs.
+ To assess the significance of our results. we performed Monte Carlo. simulations (Hammersley Handscomb 1964)) of the RVs of a synthetic population of massive binaries.," To assess the significance of our results, we performed Monte Carlo simulations (Hammersley Handscomb \cite{HH}) ) of the RVs of a synthetic population of massive binaries."
+ The parent distribution of the binary parameters of the population of massive binaries is unfortunately unknown. but we can make a series of reasonable assumptions.," The parent distribution of the binary parameters of the population of massive binaries is unfortunately unknown, but we can make a series of reasonable assumptions."
+ Here. we have followed an approach similar to the one of Sana. Gosset Evans (2009)).," Here, we have followed an approach similar to the one of Sana, Gosset Evans \cite{SGE}) )."
+ Following the latter authors. we adopt a bi-uniform period distribution in logP(days) with of the systems in the range 0.3€logP<1.0 and the remainder in the range 1.0«logP.€3.5.," Following the latter authors, we adopt a bi-uniform period distribution in $\log{P(days)}$ with of the systems in the range $0.3 \leq \log{P} \leq 1.0$ and the remainder in the range $1.0 < \log{P} \leq 3.5$."
+ The eccentricity was assumed to be uniformly distributed between 0.0 and 0.9. with all systems with orbital periods shorter than four days assumed to have circular orbits.," The eccentricity was assumed to be uniformly distributed between 0.0 and 0.9, with all systems with orbital periods shorter than four days assumed to have circular orbits."
+ The longitudes of periastron and true anomaly of the first observation were taken to be uniformly distributed in the [0.2] interval and we adopted a uniform distribution of the mass ratio g=m-2/ni between 0.1 and 1.0.," The longitudes of periastron and true anomaly of the first observation were taken to be uniformly distributed in the $[0,2\,\pi]$ interval and we adopted a uniform distribution of the mass ratio $q = m_2/m_1$ between 0.1 and 1.0."
+ The lower limit to the range of mass ratio was chosen because. in practice. it would be very difficult to observe the reflex motion of an O-star for such à low-mass companion.," The lower limit to the range of mass ratio was chosen because, in practice, it would be very difficult to observe the reflex motion of an O-star for such a low-mass companion."
+ Finally. the orbital inclination was taken to be uniform in cos/ between —1.0 and 1.0.," Finally, the orbital inclination was taken to be uniform in $\cos{i}$ between $-1.0$ and $1.0$."
+ For four values of the primary mass (20. 30. 40. and 50M... spanning roughly the range of spectral types for which some observational information on multiplicity is available) we simulated a population of 100000. binary systems.," For four values of the primary mass (20, 30, 40, and $M_{\odot}$, spanning roughly the range of spectral types for which some observational information on multiplicity is available) we simulated a population of 100000 binary systems."
+ For each system. we evaluated the maximum RV difference A RRV = maxAV(¢;)—minRV(¢4;) that would be measured on this system adopting the same temporal sampling as used for our five GIRAFFE-IFUobservations?.," For each system, we evaluated the maximum RV difference $\Delta$ RV = $\max{RV(\phi_i)} - \min{RV(\phi_i)}$ that would be measured on this system adopting the same temporal sampling as used for our five GIRAFFE-IFU."
+. The results are summarized in refMC. and illustrated in refhisto.., The results are summarized in \\ref{MC} and illustrated in \\ref{histo}.
+" If we assumed that a velocity difference of ss! were required to claim the detection of a binarysystem"". we would find that fewer than of the systems containing an VV primary star wwith a mass higher than 30M,.). and having an orbital period of shorter than ddays. would escape detection."," If we assumed that a velocity difference of $^{-1}$ were required to claim the detection of a binary, we would find that fewer than of the systems containing an V primary star with a mass higher than $M_{\odot}$ ), and having an orbital period of shorter than days, would escape detection."
+ Only the lowest inclination systems would indeed remain undetected., Only the lowest inclination systems would indeed remain undetected.
+ For longer orbital periods. the fraction of systems that would be missed by our GIRAFFE-IFU campaign steeply increases: for systems with orbital periods between 50 and ddays. about half of the binaries would escape detection.," For longer orbital periods, the fraction of systems that would be missed by our GIRAFFE-IFU campaign steeply increases: for systems with orbital periods between 50 and days, about half of the binaries would escape detection."
+ However. owing to the assumed preponderance of short orbital-period systems. the total fraction of systems with a period up to ddays that we would miss. would still remain quite low (4 - 5%%)).," However, owing to the assumed preponderance of short orbital-period systems, the total fraction of systems with a period up to days that we would miss, would still remain quite low (4 - )."
+ The same approach was applied to the sampling of the three GIRAFFE-ARGUS observations (see refMC2))., The same approach was applied to the sampling of the three GIRAFFE-ARGUS observations (see \\ref{MC2}) ).
+ As could be expected. this sampling ts less efficient for the detection of binary systems.," As could be expected, this sampling is less efficient for the detection of binary systems."
+ Nonetheless. for primary masses in the range 20 - M.« (corresponding roughly to VV spectral types). our simulations indicate that à binary system with a period shorter than ddays has less than a probability of escaping detection because ARRV 20kkm ss!.," Nonetheless, for primary masses in the range 20 - $M_{\odot}$ (corresponding roughly to V spectral types), our simulations indicate that a binary system with a period shorter than days has less than a probability of escaping detection because $\Delta$ RV $< 20$ $^{-1}$."
+ Therefore. we can safely conclude that the probability that either of the stars 118. 171. 182. 183. 199. or 203 or stars A. B or C is a short-period binary system must be very low.," Therefore, we can safely conclude that the probability that either of the stars 18, 171, 182, 183, 199, or 203 or stars A, B or C is a short-period binary system must be very low."
+ If these stars are indeed binaries. they must have either a very low orbital inclination. a very long orbital period. or both.," If these stars are indeed binaries, they must have either a very low orbital inclination, a very long orbital period, or both."
+ The situation is a bit less clear for 220b for which we do observe a maximum RV difference above our threshold., The situation is a bit less clear for 20b for which we do observe a maximum RV difference above our threshold.
+ However. as stated above. this result must be considered with," However, as stated above, this result must be considered with"
+where the ecometric comoving angular diameter distance terms have been absorbed into ον).,where the geometric comoving angular diameter distance terms have been absorbed into $g(\chi)$.
+ It is therefore now clear. given our earlier discussion. how weal lensing as à cosmological probe is particularly useful in. studies of modified gravity.," It is therefore now clear, given our earlier discussion, how weak lensing as a cosmological probe is particularly useful in studies of modified gravity."
+ For this statistic is sensitive to the erowth of structure via the presence of the linear growth ctor squared (q (a)) in the matter power spectrum I5., For this statistic is sensitive to the growth of structure via the presence of the linear growth factor squared $g^{2}(a)$ ) in the matter power spectrum $P_{\delta}$.
+ lt is sensitive to the expansion history through the ternis in the square brackets and. through the Hubble. drag in he growth of structure and. also. as discussed at the end of the last section. it is sensitive to the relation between he power spectrum of the potentials and density.," It is sensitive to the expansion history through the terms in the square brackets and through the Hubble drag in the growth of structure and also, as discussed at the end of the last section, it is sensitive to the relation between the power spectrum of the potentials and density."
+" In. the equation above the relation from £2,,. to Ps has already oen performed. assuming GR as given routinely in the iterature.", In the equation above the relation from $P_{\phi + \psi}$ to $P_{\delta}$ has already been performed assuming GR as given routinely in the literature.
+ Again. it is worth reiterating that if there is a modification to the Poisson equation. and/or το the anisolropic stress one must augment this powerspectrum with the approapriate prefactor given. for example. as in Equation(17).," Again, it is worth reiterating that if there is a modification to the Poisson equation and/or to the anisotropic stress one must augment this powerspectrum with the approapriate prefactor given, for example, as in Equation."
+. In addition to these sensitivities it is worth adding that because the dellection is given by the potentiaLg which are sourced by mass irrespective of being barvonic or dark. weak lensing does not suller from any unknown bias.," In addition to these sensitivities it is worth adding that because the deflection is given by the potentials, which are sourced by mass irrespective of being baryonic or dark, weak lensing does not suffer from any unknown bias."
+ That is. it probes the entirety of the mass distribution.," That is, it probes the entirety of the mass distribution."
+ While this probe. in principle. is excellent. for our chosen study it is worthwhile noting that the shear signal is a small 174 distortion on the already existing intrinsic ellipticity.," While this probe, in principle, is excellent for our chosen study it is worthwhile noting that the shear signal is a small $1\%$ distortion on the already existing intrinsic ellipticity."
+ This provides a thorough. technical challenge that is being combated with a combination of large galaxy nuniber analyses and refined shear measurement techniques (?.. ?. and ?)).," This provides a thorough technical challenge that is being combated with a combination of large galaxy number analyses and refined shear measurement techniques \citet{Heymans05}, \citet{Massey06} and \citet{Bridle08}) )."
+ Further still. the first. cleteetions of weak lensing are particularly recent. (7... 7... 2? and ?)) and so in this wav lensing is very much a highly promising. vet developing. cosmological probe.," Further still, the first detections of weak lensing are particularly recent \citet{Bacon00}, \citet{Kaiser00}, \citet{Wittman00} and \citet{vanWaerbeke00}) ) and so in this way lensing is very much a highly promising, yet developing, cosmological probe."
+ Despite this there are already a number of papers in the literature that have addressed. the relationship between weak lensing and mocified gravity or dark energy. such as 2.. 2.. 2.. 2. 2. 7 and ?..," Despite this there are already a number of papers in the literature that have addressed the relationship between weak lensing and modified gravity or dark energy, such as \citet{UzanBernardeau01}, \citet{Schimd05}, \citet{DoreMartigMellier07}, \citet{Schimd07}, \citet{AmendolaKunzSapone07}, \citet{Jain07} and \citet{Tsujikawa08}."
+ With these studies and potential modified: gravity attributes in lensing it is imperitive to realise that there does exist. à severe caveat., With these studies and potential modified gravity attributes in lensing it is imperitive to realise that there does exist a severe caveat.
+ This is due to the fact that lensing probes into the non-linear regime., This is due to the fact that lensing probes into the non-linear regime.
+ We are fortunate to be able to use a fitting Function (IZ.g. 2)) for the non- in standard gravity., We are fortunate to be able to use a fitting function (E.g. \citealt{Smith03}) ) for the non-linearities in standard gravity.
+ Unfortunately this is poorly understood in any deviation [from the current. framework and subsequent implementation of the fit would. technically. be invalid.," Unfortunately this is poorly understood in any deviation from the current framework and subsequent implementation of the fit would, technically, be invalid."
+ Vherefore until N-body simulations have been undertaken that could. generalise the fitting function. or accurately quantify deviations from it we must strive to work in the linear regime as much as possible., Therefore until N-body simulations have been undertaken that could generalise the fitting function or accurately quantify deviations from it we must strive to work in the linear regime as much as possible.
+ Although some attempts have been made at quantifving the validity of the present fits (IZ.g. ? and 2)) we keep. for now. a strict and linear only analysis.," Although some attempts have been made at quantifying the validity of the present fits (E.g. \citet{LaszloBean07} and \citet{Oyaizu08}) ) we keep, for now, a strict and linear only analysis."
+ Phere are. in addition. other benefits in avoiding the inclusion of small scales such as the presence of intrinsic ellipticity correlations. shear-shape correlations and the presence of non-Gaussianity in the error.," There are, in addition, other benefits in avoiding the inclusion of small scales such as the presence of intrinsic ellipticity correlations, shear-shape correlations and the presence of non-Gaussianity in the error."
+ We therefore utilise the data provided by FOS based on the CELUELS-wide survey which. due to its range of large angular scales (up to 230 arcminutes) probing the more linear regime. is ideal for work on non-LODAL cosmology such as this.," We therefore utilise the data provided by F08 based on the CFHTLS-wide survey which, due to its range of large angular scales (up to 230 arcminutes) probing the more linear regime, is ideal for work on non-LCDM cosmology such as this."
+ The Canada-Erance-Hawaii Telescope. Legacy Survey (CELUPLS). based on the MEGAPRIAIE/AIEGACAAL instrument. is an ongoing survey with a target of 450 nights extending over 5 vears.," The Canada-France-Hawaii Telescope Legacy Survey (CFHTLS), based on the MEGAPRIME/MEGACAM instrument, is an ongoing survey with a target of 450 nights extending over 5 years."
+ The recent analysis by ? has gone bevond the initial releases and investigations by 7? and ? which themselves were successful in. deriving constraints on the οax degeneracy ancl demonstrating the evolution of the shear signal with redshift., The recent analysis by \citet{Benjamin07} has gone beyond the initial releases and investigations by \citet{Semboloni06} and \citet{Hoekstra06} which themselves were successful in deriving constraints on the $\Omega_{m}- \sigma_{8}$ degeneracy and demonstrating the evolution of the shear signal with redshift.
+ This was achieved in 7/— through a better understanding. of the redshift distribution and having an increased area., This was achieved in \citet{Benjamin07} through a better understanding of the redshift distribution and having an increased area.
+ This. while marking significant progress. is still not the most optimal lensing analvsis for this work.," This, while marking significant progress, is still not the most optimal lensing analysis for this work."
+ This is because they are potentially sensitive to the growth of structures on non-linear scales which. as we emphasised above. is undesirable for a current study of bevond-LEinstcin cosmology and weak lensing.," This is because they are potentially sensitive to the growth of structures on non-linear scales which, as we emphasised above, is undesirable for a current study of beyond-Einstein cosmology and weak lensing."
+ We therefore look to the 3rd vear CELILTLS-wide release CP0003) given by ? (E08)., We therefore look to the 3rd year CFHTLS-wide release (T0003) given by \citet{Fu08} (F08).
+ Although having a smaller Ποιά of view than ? it utilises much larger angular scales (into the linear regime) also avoiding many of the potential systematics mentioned at the end of the last section., Although having a smaller field of view than \citet{Benjamin07} it utilises much larger angular scales (into the linear regime) also avoiding many of the potential systematics mentioned at the end of the last section.
+ Lt is because of this that both works reveal approximately equivalent cosmological constraints ancl little constraining power is lost., It is because of this that both works reveal approximately equivalent cosmological constraints and little constraining power is lost.
+ The current sky coverage of 57deg. approximately 3DAmi’ of the final CELIUELS. target area. is reduced. to 34.2cdee7 alter masking and the removal of various contaminants.," The current sky coverage of $57 \mathrm{deg^{2}}$, approximately $35\%$ of the final CFHTLS target area, is reduced to $34.2 \mathrm{deg^{2}}$ after masking and the removal of various contaminants."
+" Eventually including five bands this / band study stretches to à magnitude of 7,5=24.5 and encapsulating nearly 1.7 million galaxies has an elfective galaxy number density of n—l]33gal/arcmün.", Eventually including five bands this $i'$ band study stretches to a magnitude of $i'_{AB} = 24.5$ and encapsulating nearly 1.7 million galaxies has an effective galaxy number density of $n = 13.3 \mathrm{gal/arcmin^{2}}$.
+ The data (E08) comes in the form of several two point statistics which are relevant to this study., The data (F08) comes in the form of several two point statistics which are relevant to this study.
+ We choose to utilise the LE correlation Function which is shown in Equation and displayed: along with the cosmological best fit in Figure 4.., We choose to utilise the E correlation function which is shown in Equation and displayed along with the cosmological best fit in Figure \ref{fig:paper_E_correlation}.
+" As lor the aperture mass <AL,> and shear top hat variance «[52 two point statistics this is a weighted transform of the convergence power spectrum.", As for the aperture mass $< \! M^{2}_{\mathrm{ap}} \!>$ and shear top hat variance $<|\gamma|^{2}>$ two point statistics this is a weighted transform of the convergence power spectrum.
+ In this case it is given by a zeroth order Bessel function of the first kind Jy., In this case it is given by a zeroth order Bessel function of the first kind $J_{0}$.
+ lt dis in this way that the two point functions vary in their sensitivity to. various aspects of the power spectrum and in turn any svstematios., It is in this way that the two point functions vary in their sensitivity to various aspects of the power spectrum and in turn any systematics.
+£e sullers [from a constant olfset resulting from a mixing of E and. D-mocdes.,$\xi_{E}$ suffers from a constant offset resulting from a mixing of E and B-modes.
+ A finite survey size introduces a maximum angular scale which prevents a complete calculation of the shear correlation function over larger ranges., A finite survey size introduces a maximum angular scale which prevents a complete calculation of the shear correlation function over larger ranges.
+ This is needed. for a separation of I5 and B (?).., This is needed for a separation of E and B \citep{Kilbinger06}.
+ Yo alleviate this we alter the statistic £g to £g|6 including the constant ollset c as an extra. parameter., To alleviate this we alter the statistic $\xi_{E}$ to $\xi_{E} + c'$ including the constant offset $c'$ as an extra parameter.
+ An expression is then obtained for the olfset which represents the best fit ollset (d?de’= 0) for cach parameter choice., An expression is then obtained for the offset which represents the best fit offset $\mathrm{d} \chi^{2}/\mathrm{d} c' = 0$ ) for each parameter choice.
+ This constitutes an analytic mareinalisation over e (2).., This constitutes an analytic marginalisation over $c'$ \citep{Lewis02}.
+ We subsequently find the expression to be, We subsequently find the expression to be
+aud is thus not able to account for two-body relaxation effects in granular media.,and is thus not able to account for two-body relaxation effects in granular media.
+ However. a certain degree of erauularitvis a inherent property of N-body svstemis.," However, a certain degree of granularity is a inherent property of $N$ -body systems."
+ This may then lead to discrepancies. expecially iu the unstable interval of negative specific heat. where ράσο trausitious may be sensitive to s1all scale plivsics.," This may then lead to discrepancies, especially in the unstable interval of negative specific heat, where phase transitions may be sensitive to small scale physics."
+ Whereas thermostatistics is too smooth to account for muicroscopic physics in eranular seltf-eravitatiug media such as the interstellar eas. two-body relaxation is often too strong in N-body system due to computational limitations.," Whereas thermostatistics is too smooth to account for microscopic physics in granular self-gravitating media such as the interstellar gas, two-body relaxation is often too strong in $N$ -body system due to computational limitations."
+. Especially. in. high. force. resolution. simulationsN. the force resolution. is. lareer. than the mass Boghosian resolution., Especially in high force resolution simulations where the force resolution is larger than the mass resolution.
+ We refer to this iu Sect., We refer to this in Sect.
+ 6.1 where we discuss. amoung other thines. the effect of the eramularity on lone-range correlations appearing iu the interval of negativo specific heat.," \ref{corr} where we discuss, among other things, the effect of the granularity on long-range correlations appearing in the interval of negative specific heat."
+ A further discrepancy between nature. analytical models and N-body models may be due to the eutropy used in eravo-thermal statistics (Tara Sakagaii 900111.," A further discrepancy between nature, analytical models and $N$ -body models may be due to the entropy used in gravo-thermal statistics (Taruya Sakagami \cite{Taruya01}) )."
+ Iudeed. the entropy used in analytical models to find equilibrimm states via the maxinuun eutropy principle is the extensive Boltzmaun-Cabbs eutropy. that is im fact not applicable for nou-exteusive sclf-eravitating svsteuis.," Indeed, the entropy used in analytical models to find equilibrium states via the maximum entropy principle is the extensive Boltzmann-Gibbs entropy, that is in fact not applicable for non-extensive self-gravitating systems."
+ Generalized thermostatistics includiug nou-exteusivitv are currently developed (Tsallis L98s8:: Sumivosli 2001:: Latora et al. 2001:, Generalized thermostatistics including non-extensivity are currently developed (Tsallis \cite{Tsallis88}; ; Sumiyoshi \cite{Sumiyoshi01}; Latora et al. \cite{Latora01};
+ Leubner 20013)., Leubner \cite{Leubner01}) ).
+ These formalis sugeest that non-extcusivity changes not all but some of1 the ∖classical thermodynamical∖⋅ry):EM resultsou: (PsallisNNje ΤοΟΟ..1999: 1999)).2⊾d. 1€ which agrees“Wye with∢⊾⋅∖∖↴ ourEM findings.," These formalisms suggest that non-extensivity changes not all but some of the classical thermodynamical results (Tsallis \cite{Tsallis99}; Boghosian \cite{Boghosian99}) ), which agrees with our findings."
+ 1ναιwhere Because currently it is a priori uot known which thermostatistical properties change i non-cxteusive svstenis aud cousisteut theoretical tools are not available. analytical results mustbe considered with caution.," Because currently it is a priori not known which thermostatistical properties change in non-extensive systems and consistent theoretical tools are not available, analytical results mustbe considered with caution."
+tidal forces are time-dependeut and can vary ercatly depending on the orbital phase.,tidal forces are time-dependent and can vary greatly depending on the orbital phase.
+ Here. we show that despite the large tidal forces at periastrou. our boundary conditions are well suited for the imodeliug of such systems.," Here, we show that despite the large tidal forces at periastron, our boundary conditions are well suited for the modeling of such systems."
+ The system we model consists of two maisequence stars with massesof 1.10 and 1.50 AD. with au eccentricity of e=0.15 aud evolved for over four orbits {Port—Ll code units)., The system we model consists of two main-sequence stars with massesof $1.40$ and $1.50$ $_{\odot}$ with an eccentricity of $e=0.15$ and evolved for over four orbits $P_{\textrm{orb}}\simeq 44$ code units).
+ The total uuuber of particles nears ~500.000 aud the location of the boundary is at ~75% of the stars radius. which. as shown in Fieure S. is deep inside the star so that the effects of tidal force are negligible.," The total number of particles nears $\sim500,000$ and the location of the boundary is at $\sim75\%$ of the stars' radius, which, as shown in Figure \ref{fig:encmass}, is deep inside the star so that the effects of tidal force are negligible."
+ Iudeed. Figure 8. shows the radius of the primary euclosiug differcut fractious of the total bound ass (m SPIT particles) as a fiction of time for our binary svsteu.," Indeed, Figure \ref{fig:encmass} shows the radius of the primary enclosing different fractions of the total bound mass (in SPH particles) as a function of time for our binary system."
+ For example. the radii coutaimineg GU to 90% of the total bound mass are shown to not changeX significantly during the whole duration of this simulation.," For example, the radii containing $60\%$ to $90\%$ of the total bound mass are shown to not change significantly during the whole duration of this simulation."
+ Iu fact. only the outer radius of the star. containing over 95% of the bound mass. oscillates ching each orbit.," In fact, only the outer radius of the star, containing over $95\%$ of the bound mass, oscillates during each orbit."
+ Therefore. in this case. the choice of the locaion of the boundary (dotted line} is well justified aud Figure 8. shows jii the use of our method for ecceatric binaries is acequate.," Therefore, in this case, the choice of the location of the boundary (dotted line) is well justified and Figure \ref{fig:encmass}
+ shows that the use of our method for eccentric binaries is adequate."
+ Replacing the core of a star with a ceutral poi ifagnass aud a boundary remains avaid approxinatiou asx long as the boundary is deep enough inside the cuveope of the star., Replacing the core of a star with a central point mass and a boundary remains a valid approximation as long as the boundary is deep enough inside the envelope of the star.
+ We now present the results from the simulation prescuted in 6 .3.., We now present the results from the simulation presented in $\S$ \ref{sect:eccentric}. .
+ In particular. we are interested iu the lass trauster rates observed alone the eccentric orbit.," In particular, we are interested in the mass transfer rates observed along the eccentric orbit."
+function on Nyy for nine of the strongest les. assunius a carbon abunudauce of |C/II|2-2.5 aud a relative abundance pattern similar to that observed in population IT stars.,"function on $N_{\rm HI}$ for nine of the strongest lines, assuming a carbon abundance of [C/H]=-2.5 and a relative abundance pattern similar to that observed in population II stars."
+" A more exhaustive list of lines is presented in [11]. which also discusses the dependence of the LOX on :. O5. aud the normalization of the radiation field J,."," A more exhaustive list of lines is presented in \cite{hel97b}, which also discusses the dependence of the LOX on $z$, $\Omega _b$, and the normalization of the radiation field $J_{\nu}$."
+ This dependence is weak. aud the following predictious are believed to hold rather ecuerally in cosmological models of the forest:," This dependence is weak, and the following predictions are believed to hold rather generally in cosmological models of the forest:"
+et al,et al.
+ as members of the Virgo cluster we assign thei a distance of 17 AIpe., as members of the Virgo cluster we assign them a distance of 17 Mpc.
+ These three galaxies are reported in figure 2b., These three galaxies are reported in figure 2b.
+ For the most luminous (F12212|0919) the upper Πατ of Foo/Fu.2 is compatible with the general treud. the two faint Virgo dwarfs clearly disagree.," For the most luminous (F12242+0919) the upper limit of $\rm F_{60}/F_{0.2}$ is compatible with the general trend, the two faint Virgo dwarfs clearly disagree."
+ Due to their füutuess not much information is available for thoi. F12259|LIT is classified as dE and F12235|0911 dE or Du. A large FIR to UV ratio is not expected for elliptical ealaxies. therefore these objects are probably not dE. We will see in section 6 that even the most FIR bright aud extincted objects known in the Universe follow aud exteud he trend found in figure 2b so the behavior of these two objects is difficult to wuclerstae.," Due to their faintness not much information is available for them, F12259+1141 is classified as dE and F12235+0914 dE or Im. A large FIR to UV ratio is not expected for elliptical galaxies, therefore these objects are probably not dE. We will see in section 6 that even the most FIR bright and extincted objects known in the Universe follow and extend the trend found in figure 2b so the behavior of these two objects is difficult to understand."
+ We can try to estimate an extinction for the objects isted iu able 2., We can try to estimate an extinction for the objects listed in table 2.
+ Ouly two (F12011165]9. F13011|2907) aave been detecος at oth GO aud 100422. For these wo galaxies we have he FIR fux to estimate the UV extinction (a lower lait for F13011]2907) using the oruula {οπιοαι] fit) established in section 3.1.," Only two (F12041+6519, F13041+2907) have been detected at both 60 and $\mu$ m. For these two galaxies we have the FIR flux to estimate the UV extinction (a lower limit for F13041+2907) using the formula (polynomial fit) established in section 3.1."
+ For the eaOs:axies rot detected at LOO µ i we estimate arbitrarily lis flux such as fooπου=0.3 which is imutermediate jetween the values for warni aixl cool dust (Lousdale Uelou 1987)). i£ this value is incompatible witi the upper indt. we adopt the upper lait.," For the galaxies not detected at 100 $\mu$ m we estimate arbitrarily this flux such as $\rm f_{60}/f_{100}=0.3$ which is intermediate between the values for warm and cool dust (Lonsdale Helou \cite{lonsdale}) ), if this value is incompatible with the upper limit, we adopt the upper limit."
+ The extinctions are listed iu table 2., The extinctions are listed in table 2.
+ Adopting the relation of Meurer et al., Adopting the relation of Meurer et al.
+ leads to extinctions larger by 0.1L mae., leads to extinctions larger by 0.4 mag.
+ Three galaxies have a UV extinction huger than 3.5 mae. they are the two objects without any optical identification aud the faintest ealaxy of the table 2 detected in D. The three other cases (two non detectious and the uncertzn one) are less extreme (apyον2.5 mae).," Three galaxies have a UV extinction larger than 3.5 mag, they are the two objects without any optical identification and the faintest galaxy of the table 2 detected in B. The three other cases (two non detections and the uncertain one) are less extreme $\rm a_{UV}>\sim 2.5 ~mag$ )."
+ Note that the upper limits found for these galaxies are compatible with the values fouud for some galaxies of the IRAS/FOCAÀ sample (figures 2)., Note that the upper limits found for these galaxies are compatible with the values found for some galaxies of the IRAS/FOCA sample (figures 2).
+ For example the two most extincted galaxies of our sample. namely M82 aud IC732. have a UV extinction larger than 5 mae and a Foo/Eo» ratio larger than 2 in log unit.," For example the two most extincted galaxies of our sample, namely M82 and IC732, have a UV extinction larger than 5 mag and a $\rm F_{60}/F_{0.2}$ ratio larger than 2 in log unit."
+WII Jised+24: Galactic nebulositv.,WHI J1824+24: Galactic nebulosity.
+ Measurements are given for a ringlike structure. connected to more wisps going olf the edge of the field.," Measurements are given for a ringlike structure, connected to more wisps going off the edge of the field."
+ This is conceivably a PN., This is conceivably a PN.
+ WII J132314-24: Mottled Galactic nebulosity with distant galaxies in the background., WHI J1831+24: Mottled Galactic nebulosity with distant galaxies in the background.
+ COMAW 5-577172: Multiarmed spiral galaxy. behind a lot of Galactic stars.," CGMW 5-5772: Multiarmed spiral galaxy, behind a lot of Galactic stars."
+ WII JiStd4+28: A wisp of Galactic nebulositv., WHI J1844+28: A wisp of Galactic nebulosity.
+ There is more in the field as well as leading out of it: (Bis is the most coherent. compact part.," There is more in the field as well as leading out of it; this is the most coherent, compact part."
+ WII J18562-52: Very faint. (wo wisps of nebulosity forming a part of a circle.," WHI J1856+52: Very faint, two wisps of nebulosity forming a part of a circle."
+ Although it is not much (Gf anv) brighter than flat-fielding residuals. observations at (wo different observing runs eive (he same shape and surface brightness.," Although it is not much (if any) brighter than flat-fielding residuals, observations at two different observing runs give the same shape and surface brightness."
+ WII J1359--45: A few wisps which might outline a larger area of galactic nebulosity., WHI J1859+45: A few wisps which might outline a larger area of galactic nebulosity.
+ WII J19092-50: Galactic nebulosity., WHI J1909+50: Galactic nebulosity.
+ WII J1913+41: The brightest bit of Galactic nebulosity whieh just about fills the field., WHI J1913+41: The brightest bit of Galactic nebulosity which just about fills the field.
+ The main uncertainty in surface brightness comes from not knowing what is skv ancl what is fainter nebulosity., The main uncertainty in surface brightness comes from not knowing what is sky and what is fainter nebulosity.
+ WII J1919--44: Very nice bipolar PN., WHI J1919+44: Very nice bipolar PN.
+ WII J1932+08: Face-on spiral. with a central bar (accentuated by a guiding error in our follow-up image).," WHI J1932+08: Face-on spiral, with a central bar (accentuated by a guiding error in our follow-up image)."
+ WII J19334-55: A roundish piece of nebulosity., WHI J1933+55: A roundish piece of nebulosity.
+ WII J1945+22: Large. [aint nebulositv.," WHI J1945+22: Large, faint nebulosity."
+ Due to the high star density. the surface brightness measurements are even more uncertain than usual.," Due to the high star density, the surface brightness measurements are even more uncertain than usual."
+ WII J20044-64: Swirls of Galactie nebulositv., WHI J2004+64: Swirls of Galactic nebulosity.
+ Little or no Ila., Little or no $\alpha$.
+ WIL J20242-52: A voundish bit of Galactic nebulositv., WHI J2024+52: A roundish bit of Galactic nebulosity.
+ Crowcded field., Crowded field.
+ WII J20314-00: Oval bit of Galactic nebulositv., WHI J2031+00: Oval bit of Galactic nebulosity.
+ ZOAG (G093.12--08.90: A [ace-on. extineted spiral galaxy.," ZOAG G093.12+08.90: A face-on, extincted spiral galaxy."
+ There are wisps of Galactic nebulosity within a few are minutes., There are wisps of Galactic nebulosity within a few arc minutes.
+ IXIXR99-59: A diffuse. oval object. catalogued by Narachentsevοἱal.(1999). as a probable nearby dwarf galaxyv.," KKR99-59: A diffuse, oval object, catalogued by \citet{KKR99} as a probable nearby dwarf galaxy."
+ Is morphology here. together with the [act that it has apparently not been seen in HI by Iluchtmeieretal.(2000a) nor in lla. bv. (2003).. lead us to believe it to be Galactic reflection nebulositv.," Its morphology here, together with the fact that it has apparently not been seen in HI by \citet{HKK00} nor in $\alpha$ by \citet{MKB03}, lead us to believe it to be Galactic reflection nebulosity."
+ WII J2125+44: Bright (and near a bright star)., WHI J2125+44: Bright (and near a bright star).
+ Probably a barred spiral. but. possibly," Probably a barred spiral, but possibly"
+The cosmic microwave background (CMB) was discovered by Penzias and Wilson (1965).,The cosmic microwave background (CMB) was discovered by Penzias and Wilson (1965).
+ Unique information about the earliest phases of the evolution of the Universe can be derived from CMB temperature and polarization maps., Unique information about the earliest phases of the evolution of the Universe can be derived from CMB temperature and polarization maps.
+ Since its discovery. tremendous effort has been made to improve the CMB maps.," Since its discovery, tremendous effort has been made to improve the CMB maps."
+ Significant improvement has been made with the ongoing ASA (WMAP. Benet et al.," Significant improvement has been made with the ongoing NASA (WMAP, Bennet et al."
+ 2003a). and with the Planck mission. launched in Ίαν 2009. it is expected that the sensitivity and angular resolution of the CMB maps will be improved by more than an order of magnitude.," 2003a), and with the Planck mission, launched in May 2009, it is expected that the sensitivity and angular resolution of the CMB maps will be improved by more than an order of magnitude."
+ Unfortunately. the cosmological CMB signal is always mixed with emission from the Milky Way (synchrotron. free-free. and thermal dust emission).," Unfortunately, the cosmological CMB signal is always mixed with emission from the Milky Way (synchrotron, free-free, and thermal dust emission)."
+ To extract the background cosmological information. it is essential to remove the galactic foregrounds without introducing systematic errors.," To extract the background cosmological information, it is essential to remove the galactic foregrounds without introducing systematic errors."
+ Several algorithms have been developed to solve this key issue in CMB research., Several algorithms have been developed to solve this key issue in CMB research.
+ A comprehensive review is given by Delabrouille and Cardoso (2007)., A comprehensive review is given by Delabrouille and Cardoso (2007).
+ Of course. it is most desirable that the method for removing the galactic foregrounds produces both a power spectrum and a CMB map with insignificant systematic errors.," Of course, it is most desirable that the method for removing the galactic foregrounds produces both a power spectrum and a CMB map with insignificant systematic errors."
+ For the Planck mission. 1t is an important requirement. since one of the main scientific goals is to search for non-Gaussian features in the CMB maps.," For the Planck mission, it is an important requirement, since one of the main scientific goals is to search for non-Gaussian features in the CMB maps."
+ A lot of signals of individual sky pixels are averaged in order to derive the power spectrum. therefore. the crucial issue is not so much to minimize the random errors per sky pixel. but to minimize the systematic errors in the CMB map as a whole.," A lot of signals of individual sky pixels are averaged in order to derive the power spectrum, therefore, the crucial issue is not so much to minimize the random errors per sky pixel, but to minimize the systematic errors in the CMB map as a whole."
+ From the FIRAS instrument onboard the COBE satellite. it is known that the CMB spectrum follows a black body spectrum very closely (Mather et al.," From the FIRAS instrument onboard the COBE satellite, it is known that the CMB spectrum follows a black body spectrum very closely (Mather et al."
+ 1999)., 1999).
+ Fortunately. all known non-cosmological signals have very different spectral behaviour from a black body.," Fortunately, all known non-cosmological signals have very different spectral behaviour from a black body."
+ It is thus possible to disentangle the different components of the microwave signals., It is thus possible to disentangle the different components of the microwave signals.
+ The obtained accuracy will. of course. depend on the observational errors and frequency coverage of the data available.," The obtained accuracy will, of course, depend on the observational errors and frequency coverage of the data available."
+ The ESA Planck mission was successfully launched in May 2009. and all systems have been working according to expectations ever since.," The ESA Planck mission was successfully launched in May 2009, and all systems have been working according to expectations ever since."
+ An important part of the preparation of the mission has been evaluation of the available galactic foreground removal algorithms. based on detailed simulations. called the (PSM).," An important part of the preparation of the mission has been evaluation of the available galactic foreground removal algorithms, based on detailed simulations, called the (PSM)."
+ This work was done by Planck Working Group 2. coordinated by J. Delabrouille and G. de Zotti.," This work was done by Planck Working Group 2, coordinated by J. Delabrouille and G. de Zotti."
+ Comparisons of the 8 investigated methods can be found in Leach et ((2008)., Comparisons of the 8 investigated methods can be found in Leach et (2008).
+ Norgaard-ielsen. and Jorgensen (2008. hereafter NNJ) have shown that with observational errors as expected from the Planck satellite. reasonable assumptions about the spectral behaviour of the galactic foregrounds. it is possible to use simple neural networks to extract the CMB temperature signal with negligible systematic errors.," rgaard-Nielsen and rgensen (2008, hereafter NNJ) have shown that with observational errors as expected from the Planck satellite, reasonable assumptions about the spectral behaviour of the galactic foregrounds, it is possible to use simple neural networks to extract the CMB temperature signal with negligible systematic errors."
+ In the analysis of the same PSM data às used by Leach et al.((2008). Norgaard-Nielsen and Hebert (2009. hereafter NNH) have shown that neural networks can also significantly improve the removal of systematic errors in the CMB temperature determination for imaging data.," In the analysis of the same PSM data as used by Leach et (2008), rgaard-Nielsen and Hebert (2009, hereafter NNH) have shown that neural networks can also significantly improve the removal of systematic errors in the CMB temperature determination for imaging data."
+ An analysis of the WMAP 5yr data Is presented here to show the improvement produced by neural networks. also for real observed data.," An analysis of the WMAP 5yr data is presented here to show the improvement produced by neural networks, also for real observed data."
+ It is basically the same method as in NNJ and NNH. so the neural network references can be found there.," It is basically the same method as in NNJ and NNH, so the neural network references can be found there."
+ The frequency maps obtained during the first 5 years of the WMAP mission (K. Ka. Q. V. W. centred at 22GHz. 33GHz. 41GHz. 61GHz. 94GHz. respectively) were taken from the official WMAP website:Anap/eurrent/m_products.," The frequency maps obtained during the first 5 years of the WMAP mission (K, Ka, Q, V, W, centred at 23GHz, 33GHz, 41GHz, 61GHz, 94GHz, respectively) were taken from the official WMAP website:."
+cfit.. The PSM maps were taken from the Planck Working Group 2 Challenge-2 ftp area:2/PSM-maps., The PSM maps were taken from the Planck Working Group 2 Challenge-2 ftp area:.
+v0.. PSM exposures maps (expected hits per sky pixel) are also provided., PSM exposures maps (expected hits per sky pixel) are also provided.
+ In order to derive noise maps for each frequency the algorithm given at the WMAP website hàs been used. assuming that the noise is Gaussianly distributed.," In order to derive noise maps for each frequency the algorithm given at the WMAP website has been used, assuming that the noise is Gaussianly distributed."
+ For each of the WMAP frequencies and each of the components (CMB. synchrotron. free-free. thermal and spinning dust) PSM provides maps without observational errors and no corrections," For each of the WMAP frequencies and each of the components (CMB, synchrotron, free-free, thermal and spinning dust) PSM provides maps without observational errors and no corrections"
+radii and temperatures. as well as My and (V—7) are nicely reproduced.,"radii and temperatures, as well as $M_V$ and $(V-I)$ are nicely reproduced."
+ For 12 Gyr the predicted radii start to be larger than observed. while in the other planes we obtain a good fit.," For 12 Gyr the predicted radii start to be larger than observed, while in the other planes we obtain a good fit."
+ Thus we consider 12 Gyr as an upper limit to the ASAS-04 age., Thus we consider 12 Gyr as an upper limit to the ASAS-04 age.
+ Note also that the BHAC98 model for 8 Gyr and the solar Z predicts the largest radit of all sets. despite predicting the lowest temperatures and luminosities.," Note also that the BHAC98 model for 8 Gyr and the solar $Z$ predicts the largest radii of all sets, despite predicting the lowest temperatures and luminosities."
+ The presented models suggest that the radius and temperature discrepancies may not be significant for older stars., The presented models suggest that the radius and temperature discrepancies may not be significant for older stars.
+ This seems to be supported by the recent discovery of à 0.88 + 0.86 M. evolved eclipsing binary in the famous globular cluster (Thompsonetal.2010)., This seems to be supported by the recent discovery of a 0.88 + 0.86 $_\odot$ evolved eclipsing binary in the famous globular cluster \citep{tho09}.
+. Several sets of theoretical models succeeded to fit the observed radii and bolometric luminosities of this binary components with a single isochrone., Several sets of theoretical models succeeded to fit the observed radii and bolometric luminosities of this binary components with a single isochrone.
+ The estimated age was 11.3 Gyr and |Fe/H|2-0.70 was assumed., The estimated age was 11.3 Gyr and [Fe/H]=-0.70 was assumed.
+ Considering the similar masses of the ASAS-04 components we may expect that the almost perfect fits of¢>10 Gyr isochrones are plausible., Considering the similar masses of the ASAS-04 components we may expect that the almost perfect fits of $t>10$ Gyr isochrones are plausible.
+ From the discussion above. we can deduce the age of ASAS-04 to be 5-12 Gyr and the metal abundance between 0.008 and 0.02 with the ranges of 8—I1 Gyr and Z from 0.012 to ~0.018 be the most probable ones.," From the discussion above, we can deduce the age of ASAS-04 to be $5 - 12$ Gyr and the metal abundance between 0.008 and 0.02 with the ranges of $8 - 11$ Gyr and $Z$ from 0.012 to $\sim0.018$ be the most probable ones."
+ This makes, This makes
+lo give where Q(r) is the total energv input to the wind.,to give where $Q(r)$ is the total energy input to the wind.
+ Since we are interested in the terminal velocity of the outflow. we choose a point above the heating shell where the energy has reached its steady state value where the energy is constant in Figure 3.. top panel) and integrate outwards using the energy ancl Mach number at this point to solve (11)) as an initial value problem.," Since we are interested in the terminal velocity of the outflow we choose a point above the heating shell where the energy has reached its steady state value where the energy is constant in Figure \ref{fig:nrmdoten}, top panel) and integrate outwards using the energy and Mach number at this point to solve \ref{eq:machno}) ) as an initial value problem."
+ Note that in [act the terminal velocity is determined by the (constant) value of the Bernoulli energy above, Note that in fact the terminal velocity is determined by the (constant) value of the Bernoulli energy above
+number of bins ranging from 2 to 10.,number of bins ranging from 2 to 10.
+" Roughly. one gets an idea of the distribution of the planet abundances with 6=3. but one can realistically only start talking about a ""planetary mass function"" for 625."," Roughly, one gets an idea of the distribution of the planet abundances with $b \geq 3$, but one can realistically only start talking about a “planetary mass function” for $b \geq 5$."
+" While a planetary mass-radius-separation function yi)(ip.rp.ανAd.Z.7) depending on the stellar mass. metallicity. and age involves 6 parameters. less detailed ++parameter functions are e.g. the planetary mass-separation in,,ny.e:AL.Z) or mass radius function snasUpryM,Z) depending on stellar mass and metallicity. or a planetary mass-radius-separation function depending on stellar mass only. and 2-parameter functions would e.g. be the planetary mass function «τηνCn:M,) depending on stellar mass only. or the planetary mass-separation function ια(1.(1) irrespective of the stellar properties."," While a planetary mass-radius-separation function $\varphi_{m_\rmn{p},r_\rmn{p},a}(m_\rmn{p},r_\rmn{p},a;M_\star,Z,\tau)$ depending on the stellar mass, metallicity, and age involves 6 parameters, less detailed 4-parameter functions are e.g. the planetary mass-separation $\varphi_{m_\rmn{p},a}(m_\rmn{p},a;M_\star,Z)$ or mass radius function $\varphi_{m_\rmn{p},r_\rmn{p}}(m_\rmn{p},r_\rmn{p};M_\star,Z)$ depending on stellar mass and metallicity, or a planetary mass-radius-separation function depending on stellar mass only, and 2-parameter functions would e.g. be the planetary mass function $\varphi_{m_\rmn{p}}(m_\rmn{p};M_\star)$ depending on stellar mass only, or the planetary mass-separation function $\varphi_{m_\rmn{p},a}(m_\rmn{p},a)$ irrespective of the stellar properties."
+ We now have a total sample of about 450 planets orbiting stars other than the Sun. where it took about 10 years to detect the first 150. then about 3 years to detect the next 150. and then just about | year to detect the equal number of 150.," We now have a total sample of about 450 planets orbiting stars other than the Sun, where it took about 10 years to detect the first 150, then about 3 years to detect the next 150, and then just about 1 year to detect the equal number of 150."
+ Table 1. shows how long campaigns with a constant detection rate of 150 planets per year would have to last in order ο obtain the respective functions with desired accuracies., Table \ref{tab:nplanets} shows how long campaigns with a constant detection rate of 150 planets per year would have to last in order to obtain the respective functions with desired accuracies.
+ Right now. the collected data allow to measure. [-parameter 'unctions. find the basic structure structure ()ο 10) of 2-parameter ‘unctions. see basic trends (>= 3) in 4-parameter functions. and some hint on the dependency of the planet abundance on further »uameters.," Right now, the collected data allow to measure 1-parameter functions, find the basic structure structure $b \geq 10$ ) of 2-parameter functions, see basic trends $b \geq 3$ ) in 4-parameter functions, and some hint on the dependency of the planet abundance on further parameters."
+" With 150 planets per year. or more realistically. a fair ""actor of this rate. rough ideas (> 2:5) of 4-parameter planetary mass functions (5 5) and an indication of trends (b. 3) ‘or 6-parameter planetary mass functions are obtainable within oreseeable time frames. but the numbers call for more aggressive searches."," With 150 planets per year, or more realistically, a fair factor of this rate, rough ideas $b \geq 5$ ) of 4-parameter planetary mass functions $b \geq 5$ ) and an indication of trends $b \geq 3$ ) for 6-parameter planetary mass functions are obtainable within foreseeable time frames, but the numbers call for more aggressive searches."
+" While stars with and without planets have been distinguished by referring to the fraction f,CV...Z.7.O) of stars that host planets and defining the differential planetary mass-radius-orbit function «(Πριr0:M,Z.7.02) to relate to these only. a further statistic is the distribution of the number of planets amongst all planetary systems."," While stars with and without planets have been distinguished by referring to the fraction $f_\rmn{p}(M_\star,Z,\tau,\Omega)$ of stars that host planets and defining the differential planetary mass-radius-orbit function $\varphi(m_\rmn{p},r_\rmn{p},a,\varepsilon; M_\star, Z, \tau,\Omega)$ to relate to these only, a further statistic is the distribution of the number of planets amongst all planetary systems."
+ With multiplicity indices c; that denote the fraction of planetary systems containing & planets. where the planetary mass-radius-orbit function can be decomposed as where In general. all ji(mmyryec:Ad.Z.7.0) may be different.," With multiplicity indices $\zeta_k$ that denote the fraction of planetary systems containing $k$ planets, where the planetary mass-radius-orbit function can be decomposed as where In general, all $\hat{\varphi}_k(m_\rmn{p},r_\rmn{p},a,\varepsilon; M_\star, Z, \tau,\Omega)$ may be different."
+ Together with the multiplicity indices ος. one would be left with an infinite number of parameters.," Together with the multiplicity indices $\zeta_k$, one would be left with an infinite number of parameters."
+" This however can be meaningfully avoided by adopting a functional dependence of ος and 4, on k that is described by a small finite number of parameters.", This however can be meaningfully avoided by adopting a functional dependence of $\zeta_k$ and $\hat{\varphi}_k$ on k that is described by a small finite number of parameters.
+" In particular. one might want to distinguish stars with a single planets to multiple-planet systems. described by Qj (with l Gh gilmore:M.Z.T.QO). and In fact. ὁ have argued that there is evidence for 4, being different from yfunut-"," In particular, one might want to distinguish stars with a single planets to multiple-planet systems, described by $\zeta_1$ (with $\zeta_\rmn{mult} = 1-\zeta_1$ ), $\hat{\varphi}_1(m_\rmn{p},r_\rmn{p},a,\varepsilon; M_\star, Z, \tau,\Omega)$, and In fact, \citet{Wright} have argued that there is evidence for $\hat{\varphi}_1$ being different from $\hat{\varphi}_\rmn{mult}$."
+ The assessment of planetary multiplicity however poses a huge challenge for properly interpreting the observational data. given that our knowledge of the absence of further planets in observed systems is quite limited.," The assessment of planetary multiplicity however poses a huge challenge for properly interpreting the observational data, given that our knowledge of the absence of further planets in observed systems is quite limited."
+ If Hot Jupiters are considered lonely. whereas Neptune-mass planets are frequently found in multiple systems (22).. how much does this have to be attributed to the faet that observational techniques that report Hot Jupiters are insensitive to less massive planets. whereas if the sensitivity extends down to lower masses. other such planets are spotted rather easily?," If Hot Jupiters are considered lonely, whereas Neptune-mass planets are frequently found in multiple systems \citep{Mayor:abundance,HARPS:abundance2}, how much does this have to be attributed to the fact that observational techniques that report Hot Jupiters are insensitive to less massive planets, whereas if the sensitivity extends down to lower masses, other such planets are spotted rather easily?"
+ It is intriguing to see that observations of transit timing variations led to the suggestion of the presence of a 15 Earth-mass planet in the WASP-3 system (2). that was already Known to host a Hot Jupiter (2).., It is intriguing to see that observations of transit timing variations led to the suggestion of the presence of a 15 Earth-mass planet in the WASP-3 system \citep{MacPlanet} that was already known to host a Hot Jupiter \citep{WASP3}.
+ Planets reported by microlensing in particular cannot be claimed to be the only ones in the system. they were just the only ones that revealed their presence during a transient event.," Planets reported by microlensing in particular cannot be claimed to be the only ones in the system, they were just the only ones that revealed their presence during a transient event."
+ ? explicitly found that the acquired data do not exclude the presence of gas-giant planets at any separation orbiting the lens star, \citet{390further} explicitly found that the acquired data do not exclude the presence of gas-giant planets at any separation orbiting the lens star
+Generally speaking. with these data it 1s not possible to firmly separate carbon-rich from oxygen-rich stars among our AGB candidates.,"Generally speaking, with these data it is not possible to firmly separate carbon-rich from oxygen-rich stars among our AGB candidates."
+ However. for such metal-poor galaxies we would expect to find carbon-rich stars at colors Jy—Koz1.5 (e.g.?.andreferencestherein). and a few stars with these colors are indeed present in all of our target galaxies.," However, for such metal-poor galaxies we would expect to find carbon-rich stars at colors $J_0-K_0\gtrsim1.5$ \citep[e.g.][and references therein]{kang06}, and a few stars with these colors are indeed present in all of our target galaxies."
+ We check whether our stellar samples contain dust enshrouded AGB stars., We check whether our stellar samples contain dust enshrouded AGB stars.
+ This kind of objects are extremely faint or undetected in the optical. very red at NIR wavelengths and thus not easily detectable in the J-band because of incompleteness effects in our observations.," This kind of objects are extremely faint or undetected in the optical, very red at NIR wavelengths and thus not easily detectable in the $J$ -band because of incompleteness effects in our observations."
+ For example. ? consider a sample of ~40 stellar clusters with a range of ages and metallicities in the Small and Large Magellanic Clouds.," For example, \citet{vanloon05} consider a sample of $\sim40$ stellar clusters with a range of ages and metallicities in the Small and Large Magellanic Clouds."
+ They find a total of about 30 dust enshrouded AGB stars in ~20 young and intermediate-age clusters., They find a total of about 30 dust enshrouded AGB stars in $\sim20$ young and intermediate-age clusters.
+ These stars are found at /-K>2.5. and have metallicities higher than [Fe/H|—0.9 dex.," These stars are found at $J-K>2.5$, and have metallicities higher than $=-0.9$ dex."
+ However. for clusters with ages and metallicities comparable to our target galaxies. no dust enshrouded AGB stars were detected.," However, for clusters with ages and metallicities comparable to our target galaxies, no dust enshrouded AGB stars were detected."
+ We thus do not expect a significant number of dust enshrouded stars to be present in our target galaxies., We thus do not expect a significant number of dust enshrouded stars to be present in our target galaxies.
+ We search for stars that have good K-band measurement but no J- counterpart. and find two such objects in CenA-dEI. none in $GC1319.1-4216 and one in ESO269-066.," We search for stars that have good $K$ -band measurement but no $J$ -band counterpart, and find two such objects in CenA-dE1, none in SGC1319.1-4216 and one in ESO269-066."
+ Of the mentioned sources. in CenA-dE] one is found slightly outside the limiting radius. while the second is close to the center but has a good measurement only for one of the two K-bands: in ESO269-066 the dust enshrouded candidate also has à bac measurement in one of the two bands.," Of the mentioned sources, in CenA-dE1 one is found slightly outside the limiting radius, while the second is close to the center but has a good measurement only for one of the two $K$ -bands; in ESO269-066 the dust enshrouded candidate also has a bad measurement in one of the two bands."
+ We thus mention that these are candidates but could just as well be unresolved background galaxies (see previous Sect.)., We thus mention that these are candidates but could just as well be unresolved background galaxies (see previous Sect.).
+ We can also look for additional AGB candidates by considering variability. which is an intrinsic. characteristic of luminous AGB stars.," We can also look for additional AGB candidates by considering variability, which is an intrinsic characteristic of luminous AGB stars."
+ For all of the target galaxies we have at least two observations 1n the K-band. so we use the difference between the stellar magnitudes at different epochs as a_ variability indicator.," For all of the target galaxies we have at least two observations in the $K$ -band, so we use the difference between the stellar magnitudes at different epochs as a variability indicator."
+ For a long period variable star. the typical maximum magnitude difference is ~O.1 to ~1.5 mag in the K-band. and the period is on the order of ~10°? days (seeforex-ample?.andreferences therein)..," For a long period variable star, the typical maximum magnitude difference is $\sim0.1$ to $\sim1.5$ mag in the $K$ -band, and the period is on the order of $\sim10^{2-3}$ days \citep[see for example][and references therein]{rejkuba03}."
+ We should thus expect to see variations of a few tens of a magnitude at most. given the observing timescales for our targets (see Tab. 2)).," We should thus expect to see variations of a few tens of a magnitude at most, given the observing timescales for our targets (see Tab. \ref{infonir}) )."
+ For CenA-dEI. there are three observations in the K-band due to one repeated observation.," For CenA-dE1, there are three observations in the $K$ -band due to one repeated observation."
+ There are 36 days between the first and the last one (see Tab. 2)).," There are 36 days between the first and the last one (see Tab. \ref{infonir}) ),"
+ which are barely enough to put a lower limit on the number of possible long-period variables., which are barely enough to put a lower limit on the number of possible long-period variables.
+ We check whether there are variations between the different K-band observations. but find none.," We check whether there are variations between the different $K$ -band observations, but find none."
+ We then also check the whole combined list of sources. looking for stars that display a magnitude variation of more than 3 times the combined photometric errors of the individual measurements.," We then also check the whole combined list of sources, looking for stars that display a magnitude variation of more than 3 times the combined photometric errors of the individual measurements."
+ We find two additional variable sources that lie just below the lower limits of the AGB selection boxes. and thus include them in our AGB candidates list.," We find two additional variable sources that lie just below the lower limits of the AGB selection boxes, and thus include them in our AGB candidates list."
+ However. when checking them or the images we find that their profiles look like those of barely resolved background galaxies.," However, when checking them on the images we find that their profiles look like those of barely resolved background galaxies."
+ In Fig., In Fig.
+ 10. (upper panel) we display the A-band magnitude difference between the seconc and the third epochs. since these are the ones with better seeing. for all the candidate AGB stars except one. because it has a bad measurement in the second K-band observation.," \ref{variab} (upper panel) we display the $K$ -band magnitude difference between the second and the third epochs, since these are the ones with better seeing, for all the candidate AGB stars except one, because it has a bad measurement in the second $K$ -band observation."
+ Shown (1 green) are also the two likely background galaxies., Shown (in green) are also the two likely background galaxies.
+ The K-band observations of SGCI319.1-4216 were take 57 days apart. and three of the AGB candidates. display variability (blue dots in the central panel of Fig. 10)).," The $K$ -band observations of SGC1319.1-4216 were taken 57 days apart, and three of the AGB candidates display variability (blue dots in the central panel of Fig. \ref{variab}) )."
+ Whe considering the entire sample. two stars that lie just leftwards of the NIR selection box. and are found inside the optical selectio box. are indeed variables exhibiting a luminosity change by more than 307 (green symbols).," When considering the entire sample, two stars that lie just leftwards of the NIR selection box, and are found inside the optical selection box, are indeed variables exhibiting a luminosity change by more than $3\sigma$ (green symbols)."
+ We add the two latter to the number of candidate AGB stars for SGCI319.1-4216 (anc report them in the electronic version of Tab. 3))., We add the two latter to the number of candidate AGB stars for SGC1319.1-4216 (and report them in the electronic version of Tab. \ref{agb_list}) ).
+ Also in this, Also in this
+1851 and NGC 1904.,1851 and NGC 1904.
+" These clusters were selected for being located at r=16.7 and 18.8 kpc from the galactic center, respectively."," These clusters were selected for being located at r=16.7 and 18.8 kpc from the galactic center, respectively."
+" Compared to the other globular clusters studied so far, NGC 1851 and NGC 1904 are approximately twice as distant from the Milky Way center."," Compared to the other globular clusters studied so far, NGC 1851 and NGC 1904 are approximately twice as distant from the Milky Way center."
+" Thus are experiencing a tidal heating, proportional to , about one order of magnitude smaller, making its effects negligible."," Thus are experiencing a tidal heating, proportional to $r^{-3}$, about one order of magnitude smaller, making its effects negligible."
+" The initial selection of targets was based on their color, as derived from the analysis of ESO Imaging Survey frames and ESO 2.2m Wide Field Imager data."," The initial selection of targets was based on their color, as derived from the analysis of ESO Imaging Survey frames and ESO 2.2m Wide Field Imager data."
+" Targets have been selected requiring color difference from the cluster main sequence V-1I«0.05 and V—I<0.1, and apparent magnitude of 19>m18 and 19>m17, respectively for NGC 1851 and 1904."," Targets have been selected requiring color difference from the cluster main sequence $V-I<0.05$ and $V-I<0.1$, and apparent magnitude of $19>m>18$ and $19>m>17$, respectively for NGC 1851 and 1904."
+ The cut in luminosity was made close to the base of the giant branch to probe the cluster stellar population in a well populated region to ensure good probability to find cluster members at large distances from the cluster., The cut in luminosity was made close to the base of the giant branch to probe the cluster stellar population in a well populated region to ensure good probability to find cluster members at large distances from the cluster.
+" Indeed, according to Milky Way stellar population models Vanhollebeke,Groenewegen,andGirardi(2009) we expect a contamination of only 0.029 and 0.118 stars per arcmin squared in the selected color-luminosity range."," Indeed, according to Milky Way stellar population models \cite{Vanhollebeke09} we expect a contamination of only 0.029 and 0.118 stars per arcmin squared in the selected color-luminosity range."
+ With this surface density we expect a contamination of about, With this surface density we expect a contamination of about
+matter to [low from the voids towards the surrounding galaxy walls. implying coneentric shells of matter may collide.,"matter to flow from the voids towards the surrounding galaxy walls, implying concentric shells of matter may collide."
+ We therefore consider a model of a void surrounced bv àn overdense region., We therefore consider a model of a void surrounded by an overdense region.
+" Within voids. due to the lower amount of matter than in the homogeneous background. the curvature of the space is negative. thus the explicit forms of mass (M) and curvature (expressed by the function. ZZ) are where Ady is the mass in the corresponding volume of the homogeneous universe lie. Aly=(4706Ac)2Porat? and Pots is the background density at the last scattering instant]. Al,=sAbaο 77. MS=03 kpe. a=12 kpc."," Within voids, due to the lower amount of matter than in the homogeneous background, the curvature of the space is negative, thus the explicit forms of mass $M$ ) and curvature (expressed by the function $E$ ) are where $M_0$ is the mass in the corresponding volume of the homogeneous universe [i.e. $M_0 = (4 \pi G /3c^2) \rho_{b,ls} r^3$ and $\rho_{b,ls}$ is the background density at the last scattering instant], $M_1 = 8 M_2 a^{-3} {\rm e}^{-3/2}$ , $M_2 = -0.3$ kpc, $a = 12$ kpc."
+" where E,=4Esbeto ES=Ll1ο, ὁ= kpe."," where $E_1 = 4 E_2 b^{-2} {\rm e}^{-1}$, $E_2 = -1.1 \times 10^{-5}$, $b = 10.9$ kpc."
+ lt should. be noted. that other models οἱ voids are also possible even ones which do not evolve. from initial rarclactions but from condensation. ef. Mustapha&Lellaby (2001).," It should be noted that other models of voids are also possible – even ones which do not evolve from initial rarefactions but from condensation, cf. \citet{MH01}."
+.. However. this particular void. model was chosen because it develops. as we will show. a shell crossing singularity.," However, this particular void model was chosen because it develops, as we will show, a shell crossing singularity."
+ As can bee seen for ro24 the mass distribution as well as the curvature is the same as in the homogeneous FLAW mioclels., As can bee seen for $r>24$ the mass distribution as well as the curvature is the same as in the homogeneous FLRW models.
+ These functions were used as an initial condition specified. at the last scattering instant., These functions were used as an initial condition specified at the last scattering instant.
+ The initial density distribution for these moccels is very close to the form given in the first panel of Fig L.., The initial density distribution for these models is very close to the form given in the first panel of Fig \ref{fig1}.
+ One can see here that the void region extends from Iz1.5Alper. and is surrounded by the galaxy wall which has à density up to twice the value of the void.," One can see here that the void region extends from $R\approx 1.5$, and is surrounded by the galaxy wall which has a density up to twice the value of the void."
+ We start the evolution of both. the pressurc-frec LemaittreTolman and the Lemaittre models from the same profile of mass and curvature clistributions., We start the evolution of both the pressure-free Lemaîttre–Tolman and the Lemaîttre models from the same profile of mass and curvature distributions.
+ The only discrepaney between these models is with the equation of state. which was chosen to be of a polvtropic form The polvtropic index is chosen to be η=3/2 which is the case of à mono-atomic gas.," The only discrepancy between these models is with the equation of state, which was chosen to be of a polytropic form The polytropic index is chosen to be $n =3/2$ which is the case of a mono-atomic gas."
+ This equation of state is à good approximation to deseribe degenerate star cores. giant gaseous planets. or even for rocky planets;," This equation of state is a good approximation to describe degenerate star cores, giant gaseous planets, or even for rocky planets."
+ Thus. although realistic conditions within high-density regions inside walls might lead to a more complicated dependence of pressure. this simple polvtropie equation of state can be treated as a σου first approximation to the problem considered in thisLeller.," Thus, although realistic conditions within high-density regions inside walls might lead to a more complicated dependence of pressure, this simple polytropic equation of state can be treated as a good first approximation to the problem considered in this."
+ The constant. Ix for the LemaitreTolman model. which is pressure-free. is A=0 and for Lemaittre moccl is chosen to be A=1.98Ott m? s. (," The constant K for the Lemaîttre–Tolman model, which is pressure-free, is $K=0$ and for Lemaîttre model is chosen to be $K = 1.98 \times 10^{14}$ $^2$ $^2$. ("
+Sec. 71) ,Sec. \ref{evolution}) )
+and A—Los107 n So(Sec. ??))., and $K = 1.08 \times 10^{14}$ $^2$ $^2$ (Sec. \ref{acousticosc}) ).
+ These are very high values., These are very high values.
+ For comparison the ratio of standard pressure of ai (pare=101.325 kPa) to its density (pi;=1.202 m) αἱ O° C is approximately equal to 7.84101 mes?) , For comparison the ratio of standard pressure of air $p_{air} = 101.325$ kPa) to its density $\rho_{air} = 1.292$ $^3$ ) at $^{\circ}$ C is approximately equal to $7.84 \times 10^4$ $^2$ $^2$.
+Sue values were chosen in order to better depict the influence of pressure gradients on the evolution of matter., Such values were chosen in order to better depict the influence of pressure gradients on the evolution of matter.
+ Llowever. even if such very stiff. equations of state are. emploved. their impact on the evolution is visible only when density eradients become large.," However, even if such very stiff equations of state are employed, their impact on the evolution is visible only when density gradients become large."
+ Phus. the incorporation of this gradient of pressure mostly. affects only regions where the shell crossing singularities would occur.," Thus, the incorporation of this gradient of pressure mostly affects only regions where the shell crossing singularities would occur."
+ The algorithm which is used to calculate the evolution in the LemaittreTolman model is the same asthe one used, The algorithm which is used to calculate the evolution in the Lemaîttre–Tolman model is the same asthe one used
+and z>8.0 compared to that of Pen99.,and $z > 8.0$ compared to that of Pen99.
+ We note that the error in our method decreases steadily with redshift approaching «0.014 at z=1100., We note that the error in our method decreases steadily with redshift approaching $< 0.014$ at $z=1100$.
+" In comparison, for high redshifts, Pen99 error always stays ~ and does not decrease appreciably."," In comparison, for high redshifts, Pen99 error always stays $\sim$ and does not decrease appreciably."
+ A contour plot of AE based on the method of Pen99 with various z and Qa is shown in figure 2.., A contour plot of $\Delta E$ based on the method of Pen99 with various $z$ and $\Omega_\Lambda$ is shown in figure \ref{ContourPen}. .
+ Relatively complicated distribution of variations in the AE can be seen for the parameter space characterized by z and Qa., Relatively complicated distribution of variations in the $\Delta E$ can be seen for the parameter space characterized by $z$ and $\Omega_\Lambda$.
+" However, a contour plot of AE for our method, whichis shown in Figure 3,, 5.."," However, a contour plot of $\Delta E$ for our method, whichis shown in Figure \ref{ContourWick}, \ref{HistoPenWick}."
+consider the combination of three dillerent. requirements on them: Note that over the Galaxy. as a whole. the requirement on 0 only allects the distribution in 0 and the requirement on J only alfects the distribution in J the two distributions can be thought of asindependent?.,"consider the combination of three different requirements on them: Note that over the Galaxy as a whole, the requirement on $\bolth$ only affects the distribution in $\bolth$ and the requirement on $\bolJ$ only affects the distribution in $\bolJ$ – the two distributions can be thought of as."
+. It is only because of the finite survey. volume. and therefore the finite range of 0 for which stars with for a given J will be observed. that. the J condition significantly allects the observed 9 distribution (and vice versa)," It is only because of the finite survey volume, and therefore the finite range of $\bolth$ for which stars with for a given $\bolJ$ will be observed, that the $\bolJ$ condition significantly affects the observed $\bolth$ distribution (and vice versa)."
+" In Figure 10. I plot the density of the GCS stars as a function of JL, and J).", In Figure \ref{fig:actions} I plot the density of the GCS stars as a function of $J_\phi$ and $J_r$.
+ The density of stars in my phase- model is also plotted. for comparison.," The density of stars in my phase-mixed model is also plotted, for comparison."
+" For a given J,, there is a minimum 4,νι lor stars to reach the Solar neighbourhood. which can be thought of as a mininiun epievelic amplitude for a given guiding centreraclius*."," For a given $J_\phi$ there is a minimum $J_r=J_{r,min}$ for stars to reach the Solar neighbourhood, which can be thought of as a minimum epicyclic amplitude for a given guiding centre."
+. This is the cause of the near-parabolie lower boundary scen in ligure 10.., This is the cause of the near-parabolic lower boundary seen in Figure \ref{fig:actions}.
+ Phe Pleiades and Sirius moving groups can be clearly. seen as small overdensities in this plot., The Pleiades and Sirius moving groups can be clearly seen as small overdensities in this plot.
+" The Lyacles moving group is seen as a rather more spread out overdensity al à range of J,. around. J;,0.97.75,4. tending towards slightly lower J, with increasing ο."," The Hyades moving group is seen as a rather more spread out overdensity at a range of $J_r$ , around $J_\phi=0.97J_{\phi,0}$, tending towards slightly lower $J_\phi$ with increasing $J_r$ ."
+" The dotted ancl dashed. lines in Figure LO are 2:1 OLI and ILR lines respectively. these are lines along which 20,(J)|OF)=20, and 20,(J)Q,(J)=20, respectively. for cillerent values of £3. the perturber pattern speed. chosen such that the resonance lines reach J;=0 at J,ο.1. or Lido."," The dotted and dashed lines in Figure \ref{fig:actions}
+ are 2:1 OLR and ILR lines respectively, these are lines along which $2\Omega_\phi(\bolJ)+\Omega_r(\bolJ)=2\Omega_p$ and $2\Omega_\phi(\bolJ)-\Omega_r(\bolJ)=2\Omega_p$ respectively, for different values of $\Omega_p$, the perturber pattern speed, chosen such that the resonance lines reach $J_r=0$ at $J_\phi=0.9,\,1$, or $1.1J_{\phi,0}$."
+" Changing the value of Ον moves the resonance lines in -/,,. but does not significantly alter their &eracient in this range of J."," Changing the value of $\Omega_p$ moves the resonance lines in $J_\phi$, but does not significantly alter their gradient in this range of $\bolJ$."
+ The Livacdes overdensity seems o lie around a Lindblad. resonance line. but this could. be either an OLR or LR bine it was this [act which lead SLO to claim this was an Lindblad resonance. but that one needed o investigate the distribution in angle to determine which one.," The Hyades overdensity seems to lie around a Lindblad resonance line, but this could be either an OLR or ILR line – it was this fact which lead S10 to claim this was an Lindblad resonance, but that one needed to investigate the distribution in angle to determine which one."
+ Other resonances the 3:1 or 4:1 OLR. or ILI lines would appear very similar on Figure 10.. though the 2:1 Η ine is the furthest from the vertical.," Other resonances – the 3:1 or 4:1 OLR or ILR lines – would appear very similar on Figure \ref{fig:actions}, though the 2:1 ILR line is the furthest from the vertical."
+ It is also worth noting hat the slope of the various resonance lines is sensitive to he Galactic potential in a logarithmic potential (of the kind. used by SLO). the gradients of the 2:1 OLR and ILH ines in this part of J-space are nearly identical.," It is also worth noting that the slope of the various resonance lines is sensitive to the Galactic potential – in a logarithmic potential (of the kind used by S10), the gradients of the 2:1 OLR and ILR lines in this part of $\bolJ$ -space are nearly identical."
+ Lt may »f possible to use the slope of resonance lines in action space to provide information about the Galactic potential by comparing them to observed dynamical substructure. but that is bevond the scope of this study.," It may be possible to use the slope of resonance lines in action space to provide information about the Galactic potential by comparing them to observed dynamical substructure, but that is beyond the scope of this study."
+" ‘To explore the expected distribution of stars in the Solar neighbourhood associated with a resonance. E consider a rrelated to the phase-mixecl uusecl previously. adjusted to include a resonant component: where fy is the distribution. function. described. in Section 3.. €' is a constant chosen such that the resonant component contributes S percent of the stars observed. in the Solar neighbourhood. and Jos is a function of J, and is chosen such that LOAdae)dΠοdies)=const. lor J.=0. andl 0,,,; is a function of ϐ,, and is chosen such that IB,s|me,=const."," To explore the expected distribution of stars in the Solar neighbourhood associated with a resonance, I consider a related to the phase-mixed used previously, adjusted to include a resonant component: where $f_0$ is the distribution function described in Section \ref{sec:num}, , $C$ is a constant chosen such that the resonant component contributes $8$ percent of the stars observed in the Solar neighbourhood, and $J_{\phi,res}$ is a function of $J_r$ and is chosen such that $l\Omega_r(J_r,J_{\phi,res})+m\Omega_\phi(J_r,J_{\phi,res}) 
+= const$, for $J_z=0$, and $\theta_{r,res}$ is a function of $\theta_\phi$ and is chosen such that $l\theta_{r,res}+m\theta_\phi = const$."
+" The values Ay. and Ap... give the width of the resonance peak around the exact resonance lines in JJ, and 8,.. respectively."," The values $\Delta_{J,res}$ and $\Delta_{\theta,res}$ give the width of the resonance peak around the exact resonance lines in $J_\phi$ and $\theta_r$, respectively."
+" One could. equally. describe the width in action or angle in terms of a spread in J, or 6,, respectively. but for convenience I have chosen to describe it in terms of the coordinates with the e&reater ranges of values in these data."," One could, equally, describe the width in action or angle in terms of a spread in $J_r$ or $\theta_\phi$ respectively, but for convenience I have chosen to describe it in terms of the coordinates with the greater ranges of values in these data."
+ Phe width Avy. is ellectivelv a width in frequency about the pattern speed. of the perturber.," The width $\Delta_{J,res}$ is effectively a width in frequency about the pattern speed of the perturber."
+ 1n the toy models E show here E take Avy.=0.014. Noes=0.3.," In the toy models I show here I take $\Delta_{J,res}=0.01J_{\phi,0}$ , $\Delta_{\theta,res}=0.3$."
+ ] consider two toy models. cach designed. to. produce models with an overdensitv in phase-space in a similar volume to that where the ναός moving group is found (but not tuned to produce abest fit). one corresponding to an OLR (/=I. m= 2) and one corresponding to an ILI (= —]l1.m —2).," I consider two toy models, each designed to produce models with an overdensity in phase-space in a similar volume to that where the Hyades moving group is found (but not tuned to produce abest fit), one corresponding to an OLR $l=1$, $m=2$ ) and one corresponding to an ILR $l=-1$, $m=2$ )."
+" For the OLR model. E take 6...|26,= 1.9. and for the ILIt mocel 06,,,;|24,=13."," For the OLR model, I take $\theta_{r,res}+2\theta_\phi=-1.9$ , and for the ILR model $-\theta_{r,res}+2\theta_\phi=1.3$."
+" In the OLI case L take JuC],=0)0.975.459. and in the ILB Case Ayal.m0)—03985.,,n."," In the OLR case I take $J_{\phi,res}(J_r=0)=0.975J_{\phi,0}$, and in the ILR case $J_{\phi,res}(J_r=0)=0.985J_{\phi,0}$."
+" Figure LL shows contour plots of the density in the 8, 9,, plane of the OLR. and LLL models. and plots of 6,|n6,, (as in Figure 6)) restricted to 5»=#2 in the interests of brevity."," Figure \ref{fig:IOLR_mod_cont}
+ shows contour plots of the density in the $\theta_r$, $\theta_\phi$ plane of the OLR and ILR models, and plots of $\theta_r+n\theta_\phi$ (as in Figure \ref{fig:meat}) ) restricted to $n=\pm2$ in the interests of brevity."
+ Both the ILIt and OLI. models reproduce some of the features of the HIvades overdensity., Both the ILR and OLR models reproduce some of the features of the Hyades overdensity.
+ In. both cases the overdensity in angle space is somewhat triangular in shape. like the EIvades overdensity. rather than following a single line as one would expect if only the condition on angle (eq. 1))," In both cases the overdensity in angle space is somewhat triangular in shape, like the Hyades overdensity, rather than following a single line as one would expect if only the condition on angle (eq. \ref{eq:res}) )"
+ was relevant., was relevant.
+ In both cases the overdensity in angle is strong for the two cases mn=x2. as well for other values of n (not shown).," In both cases the overdensity in angle is strong for the two cases $n=\pm2$, as well for other values of $n$ (not shown)."
+" In an elfort to explain the structure of the overdensity in the 6,.4,, plane. the upper panels of Figure 11. also show the lines 4,=6)... for the two models. and lines corresponding to the condition on J."," In an effort to explain the structure of the overdensity in the $\theta_r$, $\theta_\phi$ plane, the upper panels of Figure \ref{fig:IOLR_mod_cont} also show the lines $\theta_r=\theta_{r,res}$ for the two models, and lines corresponding to the condition on $\bolJ$."
+" The latter are found by taking the condition that 1,=ενω.) Cor d,mhayesENpas OL Id=done.F2N es) ancl determining the two possible values of@ that a star with these actions would have at the Sun's position in the relevantpart of phase. space. lower values of J, correspond to smaller (ic. closer tozero) values of 6,,."," The latter are found by taking the condition that $J_\phi=J_{\phi,res}(J_r)$ (or $J_\phi=J_{\phi,res}\pm\Delta_{J,res}$ or $J_\phi=J_{\phi,res}\pm2\Delta_{J,res}$ ) and determining the two possible values of$\bolth$ that a star with these actions would have at the Sun's position – in the relevantpart of phase space, lower values of $J_r$ correspond to smaller (i.e. closer tozero) values of $\theta_\phi$ ."
+" ""his gives a sense of the two competing cllects which (in addition to the general selection ellects illustratedin Figure 5)) determine the shape of the overdensity in", This gives a sense of the two competing effects which (in addition to the general selection effects illustratedin Figure \ref{fig:mod_0_cont}) ) determine the shape of the overdensity in
+Additional position errors may have occurred because source OS (ol M98) is only 50 aresee away (ie the ENIM of the oll- PSE).,Additional position errors may have occurred because source 98 (of M98) is only 50 arcsec away (ie the FWHM of the off-axis PSF).
+ Alternatively. the counterpart may be a variable AGN. undetected in the WRI and with Re24.5 mag at the epoch of the optical imaging.," Alternatively, the counterpart may be a variable AGN, undetected in the HRI and with $>$ 24.5 mag at the epoch of the optical imaging."
+ Alrough this source is outside the complete survey ares. it is potentially of interest.," Although this source is outside the complete survey area, it is potentially of interest."
+ It is listed by MOS as a |ank field (ie no counterpart with 1123 mae)., It is listed by M98 as a blank field (ie no counterpart with $<$ 23 mag).
+ A possible extremely distant. cluster of. galaxies lies 20 aresee south of the PSPC position., A possible extremely distant cluster of galaxies lies 20 arcsec south of the PSPC position.
+ The likely brightest. cluster galaxy (BCG) has 10223 mag. giving an estimated redshift zz L where the redshift has been estimated from the BCC magnitude by extrapolating the relation of Vikhlinin (1998).," The likely brightest cluster galaxy (BCG) has $\ga$ 23 mag, giving an estimated redshift $\ga$ 1.3, where the redshift has been estimated from the BCG magnitude by extrapolating the relation of Vikhlinin (1998)."
+ Llowever. the detection in the LIRL data of a X-ray source coincident with a R=23 mag galaxy 9 arcsec south of the PSPC position suggests that this galaxy is the counterpart. and not the intra-cluster medium of the possible clistant cluster.," However, the detection in the HRI data of a X-ray source coincident with a R=23 mag galaxy 9 arcsec south of the PSPC position suggests that this galaxy is the counterpart, and not the intra-cluster medium of the possible distant cluster."
+ This source is also outside the complete survey area but we list it here because a compact group of three galaxies of Reels mag lie within the PSPC error box., This source is also outside the complete survey area but we list it here because a compact group of three galaxies of $\approx$ 18 mag lie within the PSPC error box.
+ The redshifts are unknown. but based on the BCC: magnitude probably lie in jo range Z=0.25-0.3.," The redshifts are unknown, but based on the BCG magnitude probably lie in the range z=0.25-0.3."
+ This source is potentially an extremely distant. cluster., This source is potentially an extremely distant cluster.
+ lt is listed. by MOS as a blank field., It is listed by M98 as a blank field.
+ A galaxy of R=23.1 mage surrounded by several fainter galaxies lies within the PSPC error circle., A galaxy of R=23.1 mag surrounded by several fainter galaxies lies within the PSPC error circle.
+ Near infra-red Ix band imaging (Newsam 1997) shows that many of the galaxies of very red: the brightest has R-h=4.2+0.4 and. R-L=1.6 (to be compared with R-h=2.6 anc R-L=0.6 for zero redshift’ ellipticals)., Near infra-red K band imaging (Newsam 1997) shows that many of the galaxies of very red; the brightest has $\pm$ 0.4 and R-I=1.6 (to be compared with R-K=2.6 and R-I=0.6 for zero redshift ellipticals).
+ If this galaxy is the BCG of a cluster. then the redshift estimated from its Ro band magnitude is 221.3. and from its Ix18.9 magnitude. involving a less uncertain extrapolation. “1.5 ," If this galaxy is the BCG of a cluster, then the redshift estimated from its R band magnitude is $\ga$ 1.3, and from its K=18.9 magnitude, involving a less uncertain extrapolation, $\approx$ "
+strounding the black hole is that of a standard accretion disk. this does not come about.,"surrounding the black hole is that of a standard accretion disk, this does not come about."
+ We now consider two cases of accretion flow. which are not standard thin accretion disks. in which να advection of poloidal Seld is more likely to occur.," We now consider two cases of accretion flow, which are not standard thin accretion disks, in which inward advection of poloidal field is more likely to occur."
+ Iu au acvection-dominated accretion flow (ADAF). the basic idea is that energy released in the accretion process is not radiated locally but is. rather. retained by the fluid as internal energv aud advected iuto the hole (see. for exaniple. Naravan Yi 1995: aud the review by Sveusson 1998).," In an advection-dominated accretion flow (ADAF), the basic idea is that energy released in the accretion process is not radiated locally but is, rather, retained by the fluid as internal energy and advected into the hole (see, for example, Narayan Yi 1995; and the review by Svensson 1998)."
+ As fay as the present discussion is concerned. tle major difference between this kind of accretion flow aud the standard disk is that the disk is geometrically thick in the seuse that JP~HR.," As far as the present discussion is concerned, the major difference between this kind of accretion flow and the standard disk is that the disk is geometrically thick in the sense that $H\sim
+R$."
+ The accretion is driven bv viscous processes with a~1. aud heuce with eg~ey.," The accretion is driven by viscous processes with $\alpha\sim1$, and hence with $v_R\sim v_\phi$."
+ What this implies (see Section 3.2) is that the invard How velocity is comparable in maeuitucde to the outward diffusiou velocity for a poloidal field threading the disk., What this implies (see Section 3.2) is that the inward flow velocity is comparable in magnitude to the outward diffusion velocity for a poloidal field threading the disk.
+ This meaus that there could iu principle be some non-ieelieible radial advection of poloidal flux., This means that there could in principle be some non-negligible radial advection of poloidal flux.
+ While it Πο]! ο possible to set up a steady coufleuration in which inward advection of poloidal flux is balanced at cach radius x outward diffusion. there is no reason to expect that the field threading the hole (which is in any case generated x currents in the disk) can significantly exceed the field hreacding the iuuer disk.," While it might be possible to set up a steady configuration in which inward advection of poloidal flux is balanced at each radius by outward diffusion, there is no reason to expect that the field threading the hole (which is in any case generated by currents in the disk) can significantly exceed the field threading the inner disk."
+ It is evident that if we wish to produce significant advection of poloidal flux to the πιο disk regions it is necessary to ensure that the radial inflow velocity iu the disk exceeds the radial diffusive outflow rate of poloidal field., It is evident that if we wish to produce significant advection of poloidal flux to the inner disk regions it is necessary to ensure that the radial inflow velocity in the disk exceeds the radial diffusive outflow rate of poloidal field.
+ Since this cannot be done uxiug a standard disk iu which the inflow is duc to outward diffusion of angular momentum through the disk. it follows that we ποσα to look for other mechanisnis for outward transport of aueular 1uonientuni.," Since this cannot be done using a standard disk in which the inflow is due to outward diffusion of angular momentum through the disk, it follows that we need to look for other mechanisms for outward transport of angular momentum."
+ It the disk is selteravitatius. ax is thought to occur in the carly stages of protostellar disks. aud inu the outer regions of disks around galactic nuclei. then: non-axisviunietrie instabilities can eive rise fo significant outward trausport of angular momentum (Paczvisski LOTS: Boss 198 E:Authouv Carlbere 1988: Lin Priugle 1987. 1990: Sellwood Liu 1989: Laughlin. IKorchaei- Adams 1997).," It the disk is self-gravitating, as is thought to occur in the early stages of protostellar disks, and in the outer regions of disks around galactic nuclei, then non-axisymmetric instabilities can give rise to significant outward transport of angular momentum (Paczyńsski 1978; Boss 1984; Anthony Carlberg 1988; Lin Pringle 1987, 1990; Sellwood Lin 1989; Laughlin, Korchagin, Adams 1997)."
+ Since such a process is not driven by lvdromaguetic oeiustabilities it is conceivable that the magnetic Praudtl nuuber στ be quite different frou unity. and that significant nmsvurd transport of poloidal field might be able to take place.," Since such a process is not driven by hydromagnetic instabilities it is conceivable that the magnetic Prandtl number might be quite different from unity, and that significant inward transport of poloidal field might be able to take place."
+ Although the inner regions of disks around black holes either in ACN or in QCalactie binaries are not usually considered to be seltberavitating. there müeght be au iuterestiug exception here if one considers the disk generated in the dynamical iguptiou of a neutron star by a black hole. which occurs oe1 sone models for x-ray bursts (Rasio 1996: Móssziros Rees 1997: Paczvüsski 1998).," Although the inner regions of disks around black holes either in AGN or in Galactic binaries are not usually considered to be self-gravitating, there might be an interesting exception here if one considers the disk generated in the dynamical disruption of a neutron star by a black hole, which occurs in some models for $\gamma$ -ray bursts (Rasio 1996; Mésszárros Rees 1997; Paczyńsski 1998)."
+" It has been argued by a nmuuber of authors DDBludford Pavue 1982: ιο Norman 1986: Kounigl 1989: Pelleticr ιο 1992: Lovelace. Romanova, Contopolous 1993) that a imaguetically driven disk wind might be the main mechanisii by which excess angular momentum is removed from disk material. and so nüght be the main mechanisina that drives au accretion disk."," It has been argued by a number of authors Blandford Payne 1982; Pudritz Norman 1986; Könnigl 1989; Pelletier Pudritz 1992; Lovelace, Romanova, Contopolous 1993) that a magnetically driven disk wind might be the main mechanism by which excess angular momentum is removed from disk material, and so might be the main mechanism that drives an accretion disk."
+ Were again the inflow velocity can in principle significantly exceed any outward diffusion rate for poloidal field. especially if the poloidal ficld is strong enough to suppress the Balbus-Wawley iustabilitv.," Here again the inflow velocity can in principle significantly exceed any outward diffusion rate for poloidal field, especially if the poloidal field is strong enough to suppress the Balbus-Hawley instability."
+ If sich a 1nechanisu were able to give rise to a steady state. it would be necessary to appeal to some process (such as the interchange iustabilitv: Spruit Taain 1990: Lubow Spruit 1995: Sprit. Stelle. Papaloizou 1995) that counterbalauces the steady inward drageing of poloidal field and allows outward diffusion of field to occur.," If such a mechanism were able to give rise to a steady state, it would be necessary to appeal to some process (such as the interchange instability; Spruit Taam 1990; Lubow Spruit 1995; Spruit, Stehle, Papaloizou 1995) that counterbalances the steady inward dragging of poloidal field and allows outward diffusion of field to occur."
+ Thus. as in Section L2.1. it is envisaged that a steady poloidal field configuration is set up. with mad advection and outward diffusion producing a balance aud a steady eradieut in the poloidal feld.," Thus, as in Section 4.2.1, it is envisaged that a steady poloidal field configuration is set up, with inward advection and outward diffusion producing a balance and a steady gradient in the poloidal field."
+ However. for the reasons discussed above. there is no reason to expect such physical processes to eive rise to a poloidal feld threading the hole that is significantly cubanced over the poloidal field threading the imuer disk.," However, for the reasons discussed above, there is no reason to expect such physical processes to give rise to a poloidal field threading the hole that is significantly enhanced over the poloidal field threading the inner disk."
+ Iu addition. the idea that such a steady balance cau be set up at all has been brought into question (Lovelace. Romanova. Newman 1991: Lubow. Papaloizou. Pringle 1991b: Agapitou Papaloizon 1995).," In addition, the idea that such a steady balance can be set up at all has been brought into question (Lovelace, Romanova, Newman 1994; Lubow, Papaloizou, Pringle 1994b; Agapitou Papaloizou 1998)."
+ The main vont here is that the process of wind removal of angular nuoiientunmi occurs locally and directly at each radius in the disk., The main point here is that the process of wind removal of angular momentum occurs locally and directly at each radius in the disk.
+ Au annulus in the disk which succeeds in getting rid of angular momentum to a wind docs rot require the presence of neighbouring annuli to do so., An annulus in the disk which succeeds in getting rid of angular momentum to a wind does not require the presence of neighbouring annuli to do so.
+ Thus different anuuli which manage to dispose of heir angular ποιο in this wav are to some extent independeut dynamical entities., Thus different annuli which manage to dispose of their angular momentum in this way are to some extent independent dynamical entities.
+" Furthermore. cfiicicut removal of angular momentum from a particular annulus cads to mwurd movement. outward beudiug of poloidal ficld lines. and consequently euhauced wind outflow. enhanced removal of augulu ποιοτα, and further iuflow (Lubow. Papaloizou. Pringle 1991b). ("," Furthermore, efficient removal of angular momentum from a particular annulus leads to inward movement, outward bending of poloidal field lines, and consequently enhanced wind outflow, enhanced removal of angular momentum, and further inflow (Lubow, Papaloizou, Pringle 1994b). ("
+This instability may be tempered by the fact that stroug bending of field) lines impedes au outflow bv iakine the disk sub-Neplerian: Ogilvie Livio 1998.),This instability may be tempered by the fact that strong bending of field lines impedes an outflow by making the disk sub-Keplerian; Ogilvie Livio 1998.)
+ ILowever. even if such unstable wind driven accretion occurs (aud at least in the disks i cataclysinic variables there is evidence that it does not. Livio 1997). there is no particular reason to suppose that at any stage the streugth of the poloidal field threading the hole is slenificantly ereater that the streneth of the poloidal field threading the imuer disk. except possibly for brief dvuamical interludes.," However, even if such unstable wind driven accretion occurs (and at least in the disks in cataclysmic variables there is evidence that it does not, Livio 1997), there is no particular reason to suppose that at any stage the strength of the poloidal field threading the hole is significantly greater that the strength of the poloidal field threading the inner disk, except possibly for brief dynamical interludes."
+" Thus. here again. it seenmis difficult to set up a credible picture iu which electromagnetic extraction of spin enerev from the hole donünates iu a steady. or even a time-averaged κακο, over electromagnetic extraction of spin cucrey frou the disk material."," Thus, here again, it seems difficult to set up a credible picture in which electromagnetic extraction of spin energy from the hole dominates in a steady, or even a time-averaged sense, over electromagnetic extraction of spin energy from the disk material."
+ Dlaudford Zuajek (1977) noted that if the poloidal magnetic field threading the black hole is comparable iu, Blandford Znajek (1977) noted that if the poloidal magnetic field threading the black hole is comparable in
+to fit several indices simultaneously could. spuriously. be interpreted as an indication of nou-solar abundance ratios.,"to fit several indices simultaneously could, spuriously, be interpreted as an indication of non-solar abundance ratios."
+ Also. uiisiatches between observational data aud model SSPs are often taken as evidence for the presence of components which are not fully accounted for iu the models. such as nissiue (or extreme) stellar evolutionary stages or a composite population.," Also, mismatches between observational data and model SSPs are often taken as evidence for the presence of components which are not fully accounted for in the models, such as missing (or extreme) stellar evolutionary stages or a composite population."
+" demonstrated that. for SSPs. the Fe5106 iudex traces Fe only. whilst [MgFo]. traces total uctallicity, Z (as first noted by 73)."," demonstrated that, for SSPs, the Fe5406 index traces Fe only, whilst [MgFe] traces total metallicity, $Z$ (as first noted by )."
+ These two iudices. iu. conibiuation therefore provide an estimate of the level of a-eulianceient. or simply whether a population has non-solu abundance ratios.," These two indices, in combination therefore provide an estimate of the level of $\alpha$ -enhancement, or simply whether a population has non-solar abundance ratios."
+ Figure L shows the Fe5106|AleFe} plane with lines from the DaSTI 11 Gyr (coustaut age) scaled-solar aud a-cuhanced models. jomed at »oiuts of approximately equal |Fe/T] (uote that at constant Z. [Fo/II| is lower for a-euhauced nodels than the corresponding scaled-solar ones).," Figure \ref{fig4} shows the Fe5406–[MgFe] plane with lines from the BaSTI 14 Gyr (constant age) scaled-solar and $\alpha$ -enhanced models, joined at points of approximately equal [Fe/H] (note that at constant $Z$ , [Fe/H] is lower for $\alpha$ -enhanced models than the corresponding scaled-solar ones)."
+" Overplotted ire the results from the whole-isochrouc tests for increased 7,4, (LOO). inereased ogg (0.25 dex) aud iucreased/decreased. [Fo/TI] (0.15 dex)."," Overplotted are the results from the whole-isochrone tests for increased $T_{eff}$ (100K), increased $g$ (0.25 dex) and increased/decreased [Fe/H] (0.15 dex)."
+ Iu this diagraiu one expects that ασ deviation i abundance ratios will move poiuts iorizontallv. ie. any degree of a-enhancenmeut uoves points from the scaled-solar liue on the eft. towards the a-cnhanced line on the right (see Figure 9 of uote that lines of different age are completely degenerate iu this diagram).," In this diagram one expects that any deviation in abundance ratios will move points horizontally, i.e. any degree of $\alpha$ -enhancement moves points from the scaled-solar line on the left, towards the $\alpha$ -enhanced line on the right (see Figure 9 of – note that lines of different age are completely degenerate in this diagram)."
+ Hore it can be seen that altering auy of +the atinosplieric paralucters simply moves the scaled-solar SSP points along the scaled-solar line. changing the inferred |Fe/HI| but not altering the inferred abundance ratios for other elements to nou-solar ratios.," Here it can be seen that altering any of the atmospheric parameters simply moves the scaled-solar SSP points along the scaled-solar line, changing the inferred [Fe/H] but not altering the inferred abundance ratios for other elements to non-solar ratios."
+ This is an important point to notice siuce individual clement abunudances. inclhudius several a clemenuts. appear to alter substantially. as demonstrated by the measured offsets in the various line iudices listed in Tables 1. aud 2..," This is an important point to notice since individual element abundances, including several $\alpha$ elements, appear to alter substantially, as demonstrated by the measured offsets in the various line indices listed in Tables \ref{tab:4Gtests} and \ref{tab:14Gtests}."
+ Iu couclusion. we urge caution against the over-interpretation of stellay population parameters Toni line iudex data. m terms of the tuferred scaled-solar or non-caled-solu abuudauce ratios and also the inferred. presence of a composite »opulation. especially when nmultipledudex or full- fitting methods are eiiploved.," In conclusion, we urge caution against the over-interpretation of stellar population parameters from line index data, in terms of the inferred scaled-solar or non-scaled-solar abundance ratios and also the inferred presence of a composite population, especially when multiple-index or full-SED fitting methods are employed."
+ We find that. or SSPs. Fe5106 in combination with |MgFCO| oxovides the iost robust indication of non-solar abundance ratios.," We find that, for SSPs, Fe5406 in combination with [MgFe] provides the most robust indication of non-solar abundance ratios."
+ We remind the reader that our results potentially impact on all SPS methods. whether fitting functions or full SEDs are eniploved.," We remind the reader that our results potentially impact on all SPS methods, whether fitting functions or full SEDs are employed."
+ Measured offsets for 23 commonly used diagnostic line indices are provided. aud we encourage the user to deteriuine the overall iupact ou their observational data and preferred fitting method.," Measured offsets for 23 commonly used diagnostic line indices are provided, and we encourage the user to determine the overall impact on their observational data and preferred fitting method."
+ We thaul the anonviuous referee for a constructive report aud some useful suggestions which helped to put our results in context., We thank the anonymous referee for a constructive report and some useful suggestions which helped to put our results in context.
+ SALP. would like to express heartfelt thanks to Elaine Siith-Freeiian for many useful diseussious aud for providing the initial motivation to dothis work., S.M.P. would like to express heartfelt thanks to Elaine Smith-Freeman for many useful discussions and for providing the initial motivation to dothis work.
+ S.ALP. acknowledges financial support from the Scicuce Techuoloey Facilities Council (STFC) through a Postdoctoral Researcl Fellowship., S.M.P. acknowledges financial support from the Science Technology Facilities Council (STFC) through a Postdoctoral Research Fellowship.
+The current study was limited. not only by the small FOV of TIP-IL which nowadays has a slit twice as long as during our observing campaign in 2005. but also by the observational gap on July 4th.,"The current study was limited, not only by the small FOV of TIP-II, which nowadays has a slit twice as long as during our observing campaign in 2005, but also by the observational gap on July 4th."
+ Magnetic field extrapolations could undoubtedly shed more light on clarifying the magnetic structure of this filament., Magnetic field extrapolations could undoubtedly shed more light on clarifying the magnetic structure of this filament.
+ It is now crucial to carry out more multrwavelength measurements. as the one presented in here. with higher cadence and bigger FOVs to fit the pieces of the puzzle together. in order to fully understand the origin. evolution and magnetie topology of AR filaments.," It is now crucial to carry out more multiwavelength measurements, as the one presented in here, with higher cadence and bigger FOVs to fit the pieces of the puzzle together, in order to fully understand the origin, evolution and magnetic topology of AR filaments."
+ In particular. continuous vector magnetograms of active regions together with simultaneous imaging of the corona should be able to prove/disprove the proposed scenario for the last stages of their evolution.," In particular, continuous vector magnetograms of active regions together with simultaneous imaging of the corona should be able to prove/disprove the proposed scenario for the last stages of their evolution."
+ The instrument suite on board the NASA/SDO satellite is the best candidate for such an study., The instrument suite on board the NASA/SDO satellite is the best candidate for such an study.
+ , 
+times larger than that monitored by current high-: SN Ia searches down to the same limiting magnitude 00z20) (Perlmutter 1999).,times larger than that monitored by current $z$ SN Ia searches down to the same limiting magnitude $m_B \approx 20$ ) (Perlmutter 1999).
+" The simple fireball model (rie. a spherically symmetric relativistic external shock wave expanding into a homogeneous medium) has been ""in and out of the hospital"" for months. but notices of its death appear to be premature."," The simple fireball model (i.e., a spherically symmetric relativistic external shock wave expanding into a homogeneous medium) has been “in and out of the hospital” for months, but notices of its death appear to be premature."
+" This amazes me. given the wealth of complexities one can easily imagine in the fireball itself and in its environment (Mésszárros 1999),"," This amazes me, given the wealth of complexities one can easily imagine in the fireball itself and in its environment (Mésszárros 1999)."
+ If the simple relativistic fireball model (or even more complex variants of it) suffice to explain burst afterglows (see Figure 5). much can be learned. including the energy of the fireball per unit solid angle. the ratio of the energy in the magnetic field to that in relativistic electrons. and the density of the external medium into which the fireball expands (Wijers Galama 1999; van Paradijs 1999: Lamb. Castander Reichart 1999).," If the simple relativistic fireball model (or even more complex variants of it) suffice to explain burst afterglows (see Figure 5), much can be learned, including the energy of the fireball per unit solid angle, the ratio of the energy in the magnetic field to that in relativistic electrons, and the density of the external medium into which the fireball expands (Wijers Galama 1999; van Paradijs 1999; Lamb, Castander Reichart 1999)."
+" It should be possible. in principle. to use the effects on the afterglow spectrum of extinction due to dust in the host galaxy and of absorption by the Lyman-o forest to determine the redshift of the burst itself. but so far. this goal has eluded modelers (see. e.g. Lamb. Castander Reichart 1999),"," It should be possible, in principle, to use the effects on the afterglow spectrum of extinction due to dust in the host galaxy and of absorption by the $\alpha$ forest to determine the redshift of the burst itself, but so far, this goal has eluded modelers (see, e.g, Lamb, Castander Reichart 1999)."
+" Currently. we are in the regime in terms of what we learn from each individual because. given the diversity of GRBs. GRB afterglows. and host galaxies. we have yet to sample the full ""phase space"" of afterglow or host galaxy properties."," Currently, we are in the regime in terms of what we learn from each individual because, given the diversity of GRBs, GRB afterglows, and host galaxies, we have yet to sample the full “phase space” of afterglow or host galaxy properties."
+" Still less have we sampled the full ""phase space"" of combinations of burst. afterglow. and host properties."," Still less have we sampled the full “phase space” of combinations of burst, afterglow, and host properties."
+ At the same time. we are in the stronglynon-linear regime. in terms of what we learn from each individual of a burst afterelow.," At the same time, we are in the strongly regime, in terms of what we learn from each individual of a burst afterglow."
+ The value of each astronomer's observation 1s enhanced by the observations made by all other astronomers., The value of each astronomer's observation is enhanced by the observations made by all other astronomers.
+ As we have heard from several speakers at this workshop. the amount of information that can be gleaned from a given afterelow depends greatly on the number of measurements that exist both simultaneously in time and in wavelength. from the radio through the millimeter. sub-millimeter. near-infrared. optical. and X-ray bands.," As we have heard from several speakers at this workshop, the amount of information that can be gleaned from a given afterglow depends greatly on the number of measurements that exist both simultaneously in time and in wavelength, from the radio through the millimeter, sub-millimeter, near-infrared, optical, and X-ray bands."
+ Furthermore. since the range of redshifts for the bursts (and therefore also their atterglows) is large. we cannot know in advance which bands will be crucial.," Furthermore, since the range of redshifts for the bursts (and therefore also their afterglows) is large, we cannot know in advance which bands will be crucial."
+ Thus simultaneous or near-simultaneous multi-wavelength observations of burst afterglows are essential. and therefore observations by as many observers as possible must be encouraged.," Thus simultaneous or near-simultaneous multi-wavelength observations of burst afterglows are essential, and therefore observations by as many observers as possible must be encouraged."
+ Finally. greater co-operation and co-ordination among observers is important. and should be facilitated. as has been done by setting up the invaluable service represented by the Gamma-Ray Burst Coordinate Network (GCN) (Barthelmy et al.," Finally, greater co-operation and co-ordination among observers is important, and should be facilitated, as has been done by setting up the invaluable service represented by the Gamma-Ray Burst Coordinate Network (GCN) (Barthelmy et al."
+ 1999)., 1999).
+ Star forming regions consist of a cluster of O/B stars that lie in and around a clumpy cloud of dust and gas., Star forming regions consist of a cluster of O/B stars that lie in and around a clumpy cloud of dust and gas.
+ We expect Ae>>1 for O/B stars embedded in the cloud. and Ay;~0 for O/B stars that have drifted out of the cloud and/or lie near the surface of the cloud and have expelled the gas and dust in their vicinity.," We expect $A_V >> 1$ for O/B stars embedded in the cloud, and $A_V \approx 0$ for O/B stars that have drifted out of the cloud and/or lie near the surface of the cloud and have expelled the gas and dust in their vicinity."
+ Thus the optical/UV spectrum of star forming regions is à sum of the spectra of many hot (blue) stars. some of which are embedded in the cloud. and therefore heavily extinguished. and some of which lie on the surface or around the cloud. and are therefore essentially un-extinguished.," Thus the optical/UV spectrum of star forming regions is a sum of the spectra of many hot (blue) stars, some of which are embedded in the cloud, and therefore heavily extinguished, and some of which lie on the surface or around the cloud, and are therefore essentially un-extinguished."
+ This composite spectrum is rather blue. and yields a value AVEzl when a single extinction curve Is fitted to it.," This composite spectrum is rather blue, and yields a value $A_V^{\rm eff}
+\approx 1$ when a single extinction curve is fitted to it."
+ The situation is very different when we consider an individual line-of-sight. as is appropriate for the afterglow of a GRB.," The situation is very different when we consider an individual line-of-sight, as is appropriate for the afterglow of a GRB."
+ If the GRB source lies outside and far away from any star-forming region. we expect Aversion<1: if the GRB source lies outside but near a star-forming region. we expect Αμπο1 about half the time and Ay!altcrelow>>L about half the time.," If the GRB source lies outside and far away from any star-forming region, we expect $A_V^{\rm afterglow} \lesssim 1$; if the GRB source lies outside but near a star-forming region, we expect $A_V^{\rm afterglow} \lesssim
+1$ about half the time and $A_V^{\rm afterglow} >> 1$ about half the time."
+" Finally. if the GRB source is embedded Π the star-forming; region.; we expect AVS""afterglow>>4]."," Finally, if the GRB source is embedded in the star-forming region, we expect $A_V^{\rm afterglow} >> 1$."
+ Thus. if GRB sources actually lie in star-forming regions. one would expect .1.o> (values of Ay~1030 are not uncommon for dense. cool molecular clouds in the Galaxy).," Thus, if GRB sources actually lie in star-forming regions, one would expect $A_V^{\rm afterglow} >> 1$ (values of $A_V \sim 10-30$ are not uncommon for dense, cool molecular clouds in the Galaxy)."
+ Is this consistent with what we see?, Is this consistent with what we see?
+ No., No.
+ However. this may not mean that GRB sources do not lie in star-forming regions.," However, this may not mean that GRB sources do not lie in star-forming regions."
+ The reason is that the soft X rays and the UV radiation from the GRB and its afterglow are capable. during the burst and immediately afterward. of vaporizing all of the dust in their path (Lamb Reichart 1999b).," The reason is that the soft X rays and the UV radiation from the GRB and its afterglow are capable, during the burst and immediately afterward, of vaporizing all of the dust in their path (Lamb Reichart 1999b)."
+ Thus the value of Appt that we measure may have nothing to do with the pre-existing value of the extinction through the star-forming region in which the burst source is embedded. but may instead reflect merely the extinction due to dust and gas in the disk of the host galaxy.," Thus the value of $A_V^{\rm afterglow}$ that we measure may have nothing to do with the pre-existing value of the extinction through the star-forming region in which the burst source is embedded, but may instead reflect merely the extinction due to dust and gas in the disk of the host galaxy."
+ The GRB. and its soft X-ray and UV afterglow. are also capable of tonizing gas in any envelope material expelled by the progenitor of the burst source and in the interstellar medium of the host galaxy.," The GRB, and its soft X-ray and UV afterglow, are also capable of ionizing gas in any envelope material expelled by the progenitor of the burst source and in the interstellar medium of the host galaxy."
+" This will produce Strómmgren spheres or very narrow cones (if the burst and its afterglow are beamed) in hydrogen. helium and various metals (Bisnovatyi-Kogan Timokhin 1998. Timokhin Bisnovatyi-Kogan 1999, Mésszárros 1999)."," This will produce Strömmgren spheres or very narrow cones (if the burst and its afterglow are beamed) in hydrogen, helium and various metals (Bisnovatyi-Kogan Timokhin 1998, Timokhin Bisnovatyi-Kogan 1999, Mésszárros 1999)."
+ Recombination of the tonized hydrogen eventually produces intense [CII]. [CIV]. [OVI] and [CIII] emission lines in the UV. and intense Πα and H./ emission lines in the optical.," Recombination of the ionized hydrogen eventually produces intense [CII], [CIV], [OVI] and [CIII] emission lines in the UV, and intense $\alpha$ and $\beta$ emission lines in the optical."
+ However. the line fluxes may still not be strong enough to be detectable at the large redshift distances of GRB host galaxies.," However, the line fluxes may still not be strong enough to be detectable at the large redshift distances of GRB host galaxies."
+ Interaction of the GRB and its soft X-ray afterelow with any envelope material expelled by the progenitor of the burst source and with the surrounding interstellar medium can also produce intense fluorescent iron line emission (see. e.g. Mésszárros 1999). but it is again difficult to see how the line flux could be large enough to be detectable or to explain the hints of a fluorescent iron emission line in the X-ray afterglows of GRB 980703 (Piro et al.," Interaction of the GRB and its soft X-ray afterglow with any envelope material expelled by the progenitor of the burst source and with the surrounding interstellar medium can also produce intense fluorescent iron line emission (see, e.g., Mésszárros 1999), but it is again difficult to see how the line flux could be large enough to be detectable or to explain the hints of a fluorescent iron emission line in the X-ray afterglows of GRB 980703 (Piro et al."
+ 1999) and GRB 980828 (Yoshida et al., 1999) and GRB 980828 (Yoshida et al.
+" 1999),", 1999).
+"this latter effect is commonly parameterized through electron acceleration parameters e, and &, both assumed to be on the order of ~107! (e.g., Bóttcher&Dermer (2010))).","this latter effect is commonly parameterized through electron acceleration parameters $\epsilon_e$ and $\zeta_e$, both assumed to be on the order of $\sim 10^{-1}$ (e.g., \cite{bd10}) )."
+" Those parameters are defined so that e, is the fraction of swept-up proton energy transferred to relativistic electrons, and Z, the fraction of swept-up electrons that are accelerated to relativistic energies."," Those parameters are defined so that $\epsilon_e$ is the fraction of swept-up proton energy transferred to relativistic electrons, and $\zeta_e$ the fraction of swept-up electrons that are accelerated to relativistic energies."
+" Specifically, if the electron acceleration process results in an electron spectral index of q>2, the cutoff is given by A A A A where [shock is the Lorentz factor of the (internal) forward or reverse shock resulting from the collision of two shells of relativistic ejecta, measured in the co-moving frame of the shocked material (Bóttcher&Dermer (2010)))."," Specifically, if the electron acceleration process results in an electron spectral index of $q > 2$, the low-energy cutoff is given by         where $\Gamma_{\rm shock}$ is the Lorentz factor of the (internal) forward or reverse shock resulting from the collision of two shells of relativistic ejecta, measured in the co-moving frame of the shocked material \cite{bd10}) )."
+" While Pghock is a strong function of the relative Lorentz factors of the colliding shells, e, and Z, are expected to depend on the efficiency of the generation of turbulent magnetic fields mediating the energy transfer between protons and electrons behind the shock fronts."," While $\Gamma_{\rm shock}$ is a strong function of the relative Lorentz factors of the colliding shells, $\epsilon_e$ and $\zeta_e$ are expected to depend on the efficiency of the generation of turbulent magnetic fields mediating the energy transfer between protons and electrons behind the shock fronts."
+" In Bóttcheretal.(2007), constraints on jet parameters could be derived from the estimated synchrotron peak flux in the 15 January 2006 SED, and a hint for hard lags among the optical (BVR) bands."," In \cite{bbj07}, constraints on jet parameters could be derived from the estimated synchrotron peak flux in the 15 January 2006 SED, and a hint for hard lags among the optical (BVR) bands."
+ The synchrotron peak flux was translated into a magnetic-field estimate according to Eq., The synchrotron peak flux was translated into a magnetic-field estimate according to Eq.
+ 4 in, 4 in
+ÀcTT74A.. respectively.,"$\lambda \sim 7774\,$, respectively."
+ All the targets were selected from the photometric catalog published by(2004).. choosing only relatively isolated objects.," All the targets were selected from the photometric catalog published by, choosing only relatively isolated objects."
+ Two different pointings were perlormed in order to sample the inner aud the most external regions of the cluster., Two different pointings were performed in order to sample the inner and the most external regions of the cluster.
+ Our ability to observe BSS in the inner wwas grea(lv restricted by Che crowding of bot stars and fibers., Our ability to observe BSS in the inner was greatly restricted by the crowding of both stars and fibers.
+ In each pointing ~30 BSS were observed. while ~20 fibers were used (o acquire skv spectra.," In each pointing $\sim 30$ BSS were observed, while $\sim 20$ fibers were used to acquire sky spectra."
+ Seventeen BSS were observed in both the pointings and were used to test the internal accuracy of the abundance measures., Seventeen BSS were observed in both the pointings and were used to test the internal accuracy of the abundance measures.
+ The total exposure time of each pointing was split in sub-exposures of about one hour each., The total exposure time of each pointing was split in sub-exposures of about one hour each.
+ In summary. for each of the two pointings we have obtained 3 spectra sampling the lines and 2 for theOr.," In summary, for each of the two pointings we have obtained 3 spectra sampling the lines and 2 for the."
+ By combining the sub-exposures we finally obtained mean spectra with a S/.V250 (per resolution element) for most of the selected BSS., By combining the sub-exposures we finally obtained mean spectra with a $S/N \ga 50$ (per resolution element) for most of the selected BSS.
+ Raw spectra were recuced in IRAF following the stanclared procedure., Raw spectra were reduced in IRAF following the standard procedure.
+ The task APALL was used to deline and extract the apertures and (o calibrate the one-dimensional spectra in wavelength. by adopting the dispersion solution derived by Th-Ar lamps accquired after each spectra.," The task APALL was used to define and extract the apertures and to calibrate the one-dimensional spectra in wavelength, by adopting the dispersion solution derived by Th-Ar lamps acquired after each spectra."
+ A full description of the data analvsis will be given in a forthcoming paper (Sabbi et al., A full description of the data analysis will be given in a forthcoming paper (Sabbi et al.
+ 2006. in preparation).," 2006, in preparation)."
+ Ilere we focus on the discussion of C and O abundances., Here we focus on the discussion of C and O abundances.
+ The analvsis of ihe chemical abundances was performed using the ROSA package1, The analysis of the chemical abundances was performed using the ROSA package.
+"9558), The equivalent width (EW) of each measurable line was measured by Gaussian fitüng of the line profile. adopting a relationship between EW anc EWIIM. as described in(2001): an iterative clipping average over a traction of the highest spectral points around each line was applied to define a local continuum."," The equivalent width (EW) of each measurable line was measured by Gaussian fitting of the line profile, adopting a relationship between EW and FWHM, as described in; an iterative clipping average over a fraction of the highest spectral points around each line was applied to define a local continuum."
+ Abundances were derived [rom the measured EW once appropriate atmospheric parameters have been adopted., Abundances were derived from the measured EW once appropriate atmospheric parameters have been adopted.
+" In. particular: (1) stellar temperatures (7~6600—8000 IXIN) were estimated by (he empirical relation log--—0.38(D—V)4-3.99. οριαπο from Che IL, temperatures of TO stars and of a few BSS. observed at high resolution (/2~ 40.000) with UVES2006): (2) gravites (logg~4.3. 4.8) were estimated Irom the BSSlocation in the Color Magnitude Diagram (CMD. see Fig. 5))."," In particular: (1) stellar temperatures $T\simeq 6600-8000$ K) were estimated by the empirical relation $\log T=-0.38 { (B-V)} +3.99$, obtained from the $_\alpha$ temperatures of TO stars and of a few BSS, observed at high resolution $R\sim 40,000$ ) with UVES; (2) gravities $\log g\simeq 4.3$ –4.8) were estimated from the BSSlocation in the Color Magnitude Diagram (CMD, see Fig. \ref{f2}) ),"
+" and (3) a value of was assumed for the microturbulence velocity,", and (3) a value of was assumed for the microturbulence velocity.
+ Finally. |Fe/I1I]|——0.67 has been adopted from(2004).," Finally, $=-0.67$ has been adopted from."
+.. The derived abundances of [O/Fe]| were corrected for departures from the local thermodynamic equilibrium (NLTE). following 0.345...," The derived abundances of [O/Fe] were corrected for departures from the local thermodynamic equilibrium (NLTE), following ."
+forward shock emission should be dominated at this time.,forward shock emission should be dominated at this time.
+" Thus. we expect that At 5000 s. from equation (12) and (13) and using the constrained value of E and 5 from the forward shock LOR5 ""-gez.∣−26g,4,Hz⊜⋯⋯⋯∏⋅∖∖⇁⊜∶↔⊺⊜⊓∕⋯∣∿↻∙⊃⊝∖ and i4 3.29.10Mey3 Hz. so the optical frequency mp (4.7«10' Hz) is expected to be in the regime rj;<opcMa and the flux density should decline with a temporal index of ~—2."," Thus, we expect that At 5000 s, from equation (12) and (13) and using the constrained value of $E$ and $n$ from the forward shock emission, we get $\nu_{{\rm mr}}\sim 6.39\times
+10^8\eta_3^2\epsilon_{{\rm er},-1}^2\epsilon_{{\rm Br},-1}^{1/2}$ Hz and $\nu_{{\rm cr}}\sim 3.29\times 10^{17}\epsilon_{{\rm Br},-1}^{-3/2}$ Hz, so the optical frequency $\nu_{{\rm opt}}$ $4.7\times 10^{14}$ Hz) is expected to be in the regime $\nu_{{\rm mr}} < \nu_{{\rm opt}} < \nu_{{\rm cr}}$ and the flux density should decline with a temporal index of $\simeq -2$."
+" If the radio frequency rui, (8.5«10? Hz) also lies in the same regime (Le. Ly<Maio My) at 4.67« 1009. the radio flux then relates with the optical flux as Combining. Eqs.(17) and (18). we obtain p«1.5."," If the radio frequency $\nu_{{\rm radio}}$ $8.5\times 10^9$ Hz) also lies in the same regime (i.e. $\nu_{{\rm mr}} < \nu_{{\rm radio}} < \nu_{{\rm cr}}$ ) at $4.67\times 10^5$ s, the radio flux then relates with the optical flux as Combining Eqs.(17) and (18), we obtain $p<1.5$."
+ Such à small p is inconsistent with the observed decay slope of the late-time optical and X-ray emission., Such a small $p$ is inconsistent with the observed decay slope of the late-time optical and X-ray emission.
+" As rayx£777, we note that at 4.67« IOs. r4,>Maio only if ja>|126!,€3 ."," As $\nu_{{\rm mr}}\propto \hat t^{-73/48}$, we note that at $4.67\times 10^5$ s, $\nu_{{\rm mr}}>\nu_{{\rm radio}}$ only if $\eta_3>112\epsilon_{{\rm er},-1}^{-1}\epsilon_{{\rm Br},-1}^{-1/4}$ ."
+ Since such an initial Lorentz factor ts too large. this spectral regime ts unlikely (see c.f.," Since such an initial Lorentz factor is too large, this spectral regime is unlikely (see c.f."
+ loka 2010)., Ioka 2010).
+ Radio flux from the reverse shock may be also affected by the synchrotron self absorption (SSA) of the radiating electrons., Radio flux from the reverse shock may be also affected by the synchrotron self absorption (SSA) of the radiating electrons.
+ The SSA frequency in the slow cooling case is (Wu et al., The SSA frequency in the slow cooling case is (Wu et al.
+" 20051) where co&15 Leis nearly a constant and X==""E e.is shock.the column density of electrons heated by the reverse "," 2005a) where $c_0\simeq 15$ is nearly a constant and $\Sigma=\frac{E/\eta
+m_pc^2}{4\pi R^2}$ is the column density of electrons heated by the reverse shock."
+"We find ri,<MaioMar requires Fsunνεο.Ἱεμνι210+ at 4.67 1008 according to equations (3). (19) and (20)."," We find $\nu_{{\rm mr}}<\nu_{{\rm radio}}<\nu_{\rm ar}$ requires $
+E_{54}n_{-3}\epsilon_{{\rm er},-1}\epsilon_{{\rm Br},-1}\ga 10^4$ at $4.67\times
+10^5$ s according to equations (3), (19) and (20)."
+ If we substitute Es;22.0 and 7.3=0.40. which are derived from the constraints by the forward shock emission. into this inequality. we obtain ερegiZ1.25«107. which is unlikely.," If we substitute $E_{54}=2.0$ and $n_{-3}=0.40$, which are derived from the constraints by the forward shock emission, into this inequality, we obtain $\epsilon_{{\rm er},-1}\epsilon_{{\rm Br},-1}\gtrsim 1.25\times10^4$, which is unlikely."
+ Therefore we conclude that SSA effect can not solve the problem that the reverse shock scenario for the optical emission overpredicts the radio emission., Therefore we conclude that SSA effect can not solve the problem that the reverse shock scenario for the optical emission overpredicts the radio emission.
+ It is usually assumed that the blast wave that produces the afterglow emission is nearly adiabatic Rees 1997; Sart et al., It is usually assumed that the blast wave that produces the afterglow emission is nearly adiabatic Rees 1997; Sari et al.
+ 1998). re. the total kinetic energy of the relativistic shock is a constant.," 1998), i.e. the total kinetic energy of the relativistic shock is a constant."
+ However. the energy loss of the blast wave could be significant under some circumstance.," However, the energy loss of the blast wave could be significant under some circumstance."
+ The radiation efficiency of the blast wave Is given by (Wu et al., The radiation efficiency of the blast wave is given by (Wu et al.
+" 200Sa) As we can see from the above equation. the energy loss is especially important in the ""fast-cooling"" case (1.8. νο) with a large ες."," 2005a) As we can see from the above equation, the energy loss is especially important in the ""fast-cooling"" case (i.e. $\nu_{\rm m} >\nu_{\rm c} $ ) with a large $\epsilon_{\rm e} $."
+ In this case. the decreasing blast wave energy at early times will result in a faster decay of the afterglow emission than the adiabatic case.," In this case, the decreasing blast wave energy at early times will result in a faster decay of the afterglow emission than the adiabatic case."
+ At late time. as 7» decreases with time. the radiation efficiency drops and the decay slope changes to the adiabatic case.," At late time, as $\frac{\nu_{\rm m} }{\nu_{\rm c} }$ decreases with time, the radiation efficiency drops and the decay slope changes to the adiabatic case."
+ We examine whether this scenario can explain the fast decay in the early optical and high-energy LAT emission., We examine whether this scenario can explain the fast decay in the early optical and high-energy LAT emission.
+ According to Huang et al. (, According to Huang et al. (
+1999) and Wu et al. (,1999) and Wu et al. (
+2005a). the isotropic-equivalent energy E of the blast wave evolve with time as The quantities describing the synchrotron spectrum in such a semi-radiative shock are similar to equations (3)-(5). except that the constant E in these equations should be replaced by a time-dependent E as deseribed by Eq. (,"2005a), the isotropic-equivalent energy $E$ of the blast wave evolve with time as The quantities describing the synchrotron spectrum in such a semi-radiative shock are similar to equations (3)–(5), except that the constant $E$ in these equations should be replaced by a time-dependent $E$ as described by Eq. ("
+22).,22).
+" So the synchrotron emission flux decays as F,«T0710070 for Vp<< iy."," So the synchrotron emission flux decays as $F_{{\rm \nu}}\propto
+T^{-3(p-1+\epsilon)/(4-\epsilon)}$ for $\nu_{{\rm mf}}<\nu <\nu_{{\rm cf}}$ ."
+ To explain the decay slope à=1.6 of the early optical emission in GRB 090902B with this model. we need €Z0.6.," To explain the decay slope $\alpha\ga 1.6$ of the early optical emission in GRB 090902B with this model, we need $\epsilon \ga 0.6$."
+ Thus we need To explain the late optical. X-ray and radio observations. we also need: #F!opt(10?s)- 0.01mdy. E!radio10's)20.03ml y. and FXGOO8)):2 0.2;y (radio observations starts at about 1005. we extrapolate it to 10?s).," Thus we need To explain the late optical, X-ray and radio observations, we also need: $\kappa F_{{\rm opt}}^{\rm f}(10^5{\rm s})\simeq 0.01$ mJy, $F_{{\rm radio}}^{\rm f}(10^5{\rm s})\simeq 0.03$ mJy, and $F_{{\rm X}}^{\rm f}(10^5{\rm
+s})\simeq 0.2 \mu$ Jy (radio observations starts at about $1.3\times
+10^5$ s, we extrapolate it to $10^5$ s)."
+" At 10°s. the two characteristic frequencies in the synchrotron spectrum are Myrc2.0«OPELτιεςHz MEand epjq7mous11LOVEEspPaPasPRapsέρωςΗΖ for e,20.6 and 4p22.2."," At $10^5$ s, the two characteristic frequencies in the synchrotron spectrum are $\nu_{{\rm mf}}\simeq 2.0\times 10^{12} E_{54}^{1/2}\epsilon_{{\rm Bf},-5}^{1/2}
+{\rm Hz}$ and $\nu_{{\rm cf}}\simeq 1.1\times
+10^{17}E_{54}^{-11/18}n_{-3}^{-10/9}\epsilon_{{\rm Bf},-5}^{-11/18} {\rm
+Hz}$ for $\epsilon_{\rm e} =0.6$ and $p=2.2$."
+ As rag. «Ix. Hio He in three different frequency regimes. we have three independent constraints on the shock parameters.," As $\nu_{{\rm opt}}$, $\nu_{{\rm X}}$, $\nu_{{\rm radio}}$ lie in three different frequency regimes, we have three independent constraints on the shock parameters."
+" Finally we obtain For these parameters. the deceleration time of the blast wave Is Tyo.& 120s. and we can obtain the blast wave kinetic energy at the deceleration time 744, according to Eq.(22)). Le. EsTa)&50."," Finally we obtain For these parameters, the deceleration time of the blast wave is $T_{\rm dec}\simeq 120$ s, and we can obtain the blast wave kinetic energy at the deceleration time $T_{\rm dec}$ according to \ref{ET}) ), i.e. $E_{54}(T_{\rm dec})\simeq 50$."
+ This energy is extraordinary large., This energy is extraordinary large.
+ With such a high isotropic energy. we the flux density in LAT band at the deceleration time to be expectFs(a)& Hy. which is one order of magnitude higher than the observed flux (—0.1/Jy).," With such a high isotropic energy, we expect the flux density in LAT band at the deceleration time to be $F_{\rm
+LAT}^{\rm f}(T_{\rm dec})\simeq 1\mu$ Jy, which is one order of magnitude higher than the observed flux $\simeq 0.1\mu {\rm Jy}$ )."
+ Therefore. we conclude that this model can not explain the broadband data of GRB 090902B. Jets from GRBs may have complex structure.," Therefore, we conclude that this model can not explain the broadband data of GRB 090902B. Jets from GRBs may have complex structure."
+ For the sake of calculation ease. the structured jet can be simplified as a two-component jet.," For the sake of calculation ease, the structured jet can be simplified as a two-component jet."
+" It assumes that the Jet consists of two components: a narrow component with a relatively small half- angle (4) and a large isotropic-equivalent energy in the center. and àwide component with a larger half- angle (£y) and a smaller isotropic-equivalent energy (hereafter. we use the superscripts/subscripts ""NT and W? represents the quantities of the narrow component and the"," It assumes that the jet consists of two components: a narrow component with a relatively small half-opening angle $\theta_N$ ) and a large isotropic-equivalent energy in the center, and awide component with a larger half-opening angle $\theta_W$ ) and a smaller isotropic-equivalent energy (hereafter, we use the superscripts/subscripts 'N' and 'W' represents the quantities of the narrow component and the"
+forms a central black hole of mass Ady.,forms a central black hole of mass $M_{\rm bh}$.
+" The black hole mass fraction kr=μιAdal, Is assumed to obey. a loe- probability distribution with logry=35 and σ=O44 (Ihbunanu Loeb L999b) These values roughly reflect the distribution of jack hole to bulge mass ratios found iu a sample of 36 local galaxies (Alagorrian et al.", The black hole mass fraction $r\equiv M_{\rm bh}/M_{\rm halo}$ is assumed to obey a log-Gaussian probability distribution with $\log r_0=-3.5$ and $\sigma=0.5$ (Haiman Loeb 1999b) These values roughly reflect the distribution of black hole to bulge mass ratios found in a sample of 36 local galaxies (Magorrian et al.
+ 1998) for a xuvonie nass fraction of ~(Οιου)z0.1., 1998) for a baryonic mass fraction of $\sim (\Omega_{\rm b}/\Omega_0)\approx 0.1$.
+" We further xoxtulate that each black hole cuits a timedepeucdent xoimetrie huuiuositvin proportion to its mass, L4—Myafg=MigLgaaexptt fy). where Lada=15x1U7Ay/AL.eres5 ds the Eddinetou luminosity. f is he time clapsed since the formation of the black hole. aud ty= lU)vr is the characteristic quasar lifetime."," We further postulate that each black hole emits a time–dependent bolometric luminosityin proportion to its mass, $L_{\rm q}\equiv M_{\rm bh}f_{\rm q}=M_{\rm bh}L_{\rm
+Edd} \exp(-t/t_0)$ , where $L_{\rm Edd}=1.5\times10^{38}~M_{\rm bh}/{\rm
+M_\odot}~{\rm erg~s^{-1}}$ is the Eddington luminosity, $t$ is the time elapsed since the formation of the black hole, and $t_0=10^6$ yr is the characteristic quasar lifetime."
+ Finally. we assuine that the shape of the emitted spectrum follows he mean spectrum of the quasar sample iu Elvis et al. (," Finally, we assume that the shape of the emitted spectrum follows the mean spectrum of the quasar sample in Elvis et al. ("
+1991) up to a photon cherey of LO keV. We extrapolate he spectrum up to ~50 keV. assuniug a spectral slope of a —U (or a photon iudex of -1).,"1994) up to a photon energy of 10 keV. We extrapolate the spectrum up to $\sim 50$ keV, assuming a spectral slope of $\alpha$ =0 (or a photon index of -1)."
+ This simple model was demonstrated to accurately reproduce the evolution of the optical luminosity function in the Bband (Pei 1995) at redshifts :22.2 (Wana Loeb 1998)., This simple model was demonstrated to accurately reproduce the evolution of the optical luminosity function in the B–band (Pei 1995) at redshifts $z\gsim 2.2$ (Haiman Loeb 1998).
+ Because our model mceorporates several sinplifiug assuiptious. we regard it as the minimal tov model which successtully reproduces the existing data.," Because our model incorporates several simplifying assumptions, we regard it as the minimal toy model which successfully reproduces the existing data."
+ If one of our input assunipfious was drastically violated and our model had failed to ft the observed LF. then a uodification of its basic ineredicuts would be needed.," If one of our input assumptions was drastically violated and our model had failed to fit the observed LF, then a modification of its basic ingredients would be needed."
+" Iu lusLetter, we focus on the predictions of this nüninmual nodel in auticipation of the forthcoming lauuch ofCXO: an investigation of a broader range of plausible tov nocels will be mace clsewhere."," In this, we focus on the predictions of this minimal model in anticipation of the forthcoming launch of; an investigation of a broader range of plausible toy models will be made elsewhere."
+" We adopt the concordance cosmology of Ostriker Steinhardt (1995). namely à AC DAL model with a tilted power spectrum (Qy.04.0,ioyn) (0.35. 0.65. 0.0L. 1.65. 08T. 0.96)."," We adopt the concordance cosmology of Ostriker Steinhardt (1995), namely a $\Lambda$ CDM model with a tilted power spectrum $\Omega_0,\Omega_\Lambda, \Omega_{\rm b},h,\sigma_{8h^{-1}},n$ )=(0.35, 0.65, 0.04, 0.65, 0.87, 0.96)."
+" Figure 1 shows the adopted spectrum of quasars. assunine a black hole mass μι=LOAD... placed at two different redshifts. 7,=Il aud 7.=6."," Figure \ref{fig:spectrum} shows the adopted spectrum of quasars, assuming a black hole mass $M_{\rm bh}=10^8{\rm M_\odot}$, placed at two different redshifts, $z_{\rm s}=11$ and $z_{\rm s}=6$."
+ Iu computing the intergalactic absorption. we included the opacity of both livdrogen and helium as well as the effect of electron scattering.," In computing the intergalactic absorption, we included the opacity of both hydrogen and helium as well as the effect of electron scattering."
+" We assumed that reionization occurred at z,=10 aud that at higher redshifts the IGM was homogeneous and fully neutral."," We assumed that reionization occurred at $z_{\rm
+r}=10$ and that at higher redshifts the IGM was homogeneous and fully neutral."
+" At lower redshifts. Qoioκτν, we included the livdrvogen opacity of the Ίνα forest given by Madau (1995). and extrapolated his fitting formulae for the evolution of the mmuberdensity of absorbers bevoud 2=5 when necessary."," At lower redshifts, $0<z<z_{\rm r}$, we included the hydrogen opacity of the $\alpha$ forest given by Madau (1995), and extrapolated his fitting formulae for the evolution of the numberdensity of absorbers beyond $z=5$ when necessary."
+" As Figure 1 shows. the muninuun black hole mass detectable by the 2.1019crestan3 flux limit ofCVO (see below) is Af,~LOSΔΙ. at +=10 and Ay,~2«10*Mat 2=5."," As Figure \ref{fig:spectrum} shows, the minimum black hole mass detectable by the $\sim 2\times 10^{-16}~{\rm erg~s^{-1}~cm^{-2}}$ flux limit of (see below) is $M_{\rm bh}\sim 10^8~{\rm M_\odot}$ at $z=10$ and $M_{\rm bh}\sim 2\times10^7~{\rm M_\odot}$ at $z=5$."
+ Iu our model. the correspouding halo masses are Mya~3.10H ALL. and Afi~6x1019ALL. respectively.," In our model, the corresponding halo masses are $M_{\rm halo}\sim 3\times10^{11}~{\rm M_\odot}$ , and $M_{\rm halo}\sim 6\times 10^{10}~{\rm M_\odot}$, respectively."
+ Although such massive halos are rare. their abundance is detectable in wide-field surveys.," Although such massive halos are rare, their abundance is detectable in wide-field surveys."
+ Note that an accurate determination of the spectrmu below ~0.1 κο could have provided an estimate of the reionization redshift τν., Note that an accurate determination of the spectrum below $\sim0.1$ keV could have provided an estimate of the reionization redshift $z_{\rm r}$.
+ Uitortunately. this spectral regime suffers from Galactic absorption (O'Flaherty and Jakobsen 1997) aud is outside the 0.1.GkeV detection band ofCYO.," Unfortunately, this spectral regime suffers from Galactic absorption (O'Flaherty and Jakobsen 1997) and is outside the 0.4–6keV detection band of."
+" Tn our model. the X.rav hiuuinosity fiction at a redshitt > fin MpeP(erg/s) 1| is eiven by a stun over halos that formed just before that redshift. ONENESS""TENarse where Ly is the observed ταν luuinosity in the iustruient* detection baud (0.5 3 keV for ROSAT and 0.1.6 keV for CNO): fx is the fraction of the quasar’s bolometric Wuninosity emitted in this baud: 4?N74Mdt is the black halo formation rate. eiven by a convolution of the PressSchechter halo mass function with equation (1)): and Af=?¢(-) fis the time elapsed from a cosmic time f until a redshift +."," In our model, the X–ray luminosity function at a redshift $z$ [in ${\rm
+Mpc^{-3}~(erg/s)^{-1}}$ ] is given by a sum over halos that formed just before that redshift, (L_X,z) where $L_X$ is the observed X–ray luminosity in the instrument's detection band $0.5$ $3$ keV for ROSAT and 0.4–6 keV for ); $f_{\rm X}$ is the fraction of the quasar's bolometric luminosity emitted in this band; $d^2N/dM_{\rm bh}dt$ is the black halo formation rate, given by a convolution of the Press–Schechter halo mass function with equation \ref{eq:scat}) ); and $\Delta t=t(z)-t$ is the time elapsed from a cosmic time $t$ until a redshift $z$."
+ Although our model was constructed so as to fit the observed optical LE. Figure 2. demonstrates that it is also in good aerecment with the data on the Xrax LF.," Although our model was constructed so as to fit the observed optical LF, Figure \ref{fig:LF} demonstrates that it is also in good agreement with the data on the X–ray LF."
+ This implies that the choice of quasar spectrum im our model is reasonable., This implies that the choice of quasar spectrum in our model is reasonable.
+ The solid curve in this fleure shows the prediction of equation (3)) at 2= 3.5. near the highest redshift where N-vay data is available.," The solid curve in this figure shows the prediction of equation \ref{eq:LF}) ) at $z=3.5$ , near the highest redshift where X-ray data is available."
+ The bottom curve corresponds to a cutoff in circular velocity for the host halos of ea50kns Lo which is introduced here in," The bottom curve corresponds to a cutoff in circular velocity for the host halos of $v_{\rm circ}\geq 50~{\rm km~s^{-1}}$ , which is introduced here in"
+Recent work has demonstrated that several [x gian stars exhibit. low-amplitude. long-term radial velocity (RY) variations with periods of several hundreds: of days (MeClure ct al.,"	Recent work has demonstrated that several K giant stars exhibit low-amplitude, long-term radial velocity (RV) variations with periods of several hundreds of days \nocite{mcc85} \nocite{irw89} \nocite{wal92} \nocite{hat93} (McClure et al."
+ 1985: Irwin. et al., 1985; Irwin et al.
+ 1989: Walker e al., 1989; Walker et al.
+ 1992: Ilatzes Cochran 1993)., 1992; Hatzes Cochran 1993).
+ The nature of these variations is presently unknown., The nature of these variations is presently unknown.
+ Racial pulsations can be excluded: as a cause since the period of the fundamenta racial mode is expected to be about a week. C, Radial pulsations can be excluded as a cause since the period of the fundamental radial mode is expected to be about a week. (
+Ehese stars also show short-term variability on timescales of a few davs which are due to radial or nonracial pulsations.),These stars also show short-term variability on timescales of a few days which are due to radial or nonradial pulsations.)
+ This eaves nonradial pulsations (NIA). rotational modulation w surface features. or low-niass companions as possible explanations for the long-period RY variations.," This leaves nonradial pulsations (NRP), rotational modulation by surface features, or low-mass companions as possible explanations for the long-period RV variations."
+ Ix giant stars iive large radii and low projected: rotational velocities so he period of rotation for these stars is expected. to. be several hundreds ofdays., K giant stars have large radii and low projected rotational velocities so the period of rotation for these stars is expected to be several hundreds of days.
+ Rotational modulation thus seenis o be a front-running hypothesis for explaining the lone-erm variabilitv., Rotational modulation thus seems to be a front-running hypothesis for explaining the long-term variability.
+ Lf the surface features are related to stellar activity (spots. plage. ete.)," If the surface features are related to stellar activity (spots, plage, etc.)"
+ then we should expect to find variations in the equivalent widths of lines. particularly ones ormed in the chromospheric.," then we should expect to find variations in the equivalent widths of lines, particularly ones formed in the chromospheric."
+ Indeed. Lambert (1987) found variations in the He E line in Arcturus with the same period (233 clavs) that was later found in the RV variations.," Indeed, Lambert (1987) \nocite{lam87}
+ found variations in the He I line in Arcturus with the same period (233 days) that was later found in the RV variations."
+ Larson et al. (, Larson et al. (
+1993) found evidence for the 545 day RW period in the equivalent width: variations of Ca IL in 3 Geminorum.,1993) found evidence for the 545 day RV period in the equivalent width variations of Ca II in $\beta$ Geminorum.
+ I thus seems that the long-period variability is consistent with rotational moculation. at least for Arcturus and 3 Geminorum.," It thus seems that the long-period variability is consistent with rotational modulation, at least for Arcturus and $\beta$ Geminorum."
+ The confirmation of rotational modulation as the source of the RY variability has. vet to be established. for Aldebaran (= a Tauri). a IX giant. with an HV. period of 643 clavs and a 2-4. (peak-to-peak) amplitude of (llatzes Cochran 1993: hereafter 11092).," 	The confirmation of rotational modulation as the source of the RV variability has yet to be established for Aldebaran $=$ $\alpha$ Tauri), a K giant with an RV period of 643 days and a $K$ (peak-to-peak) amplitude of (Hatzes Cochran 1993; hereafter HC93)."
+ The interesting aspect of this variability is that is seems to have been oesent ancl coherent. (same amplitude ancl phase) in RV measurcments spanning over 12 vears., The interesting aspect of this variability is that is seems to have been present and coherent (same amplitude and phase) in RV measurements spanning over 12 years.
+ IH surface structure is responsible for the RW variability of this star then it mus ος very. long-lived. which at first seems unlikely. but since nothing is known about surface structure on Ix. giants this ivpothesis cannot be summarilv rejected.," If surface structure is responsible for the RV variability of this star then it must be very long-lived, which at first seems unlikely, but since nothing is known about surface structure on K giants this hypothesis cannot be summarily rejected."
+ The lifetime anc cohereney of the long-period RV. variations in. Xldebaran would normally argue in favour of a low-mass companion., The lifetime and coherency of the long-period RV variations in Aldebaran would normally argue in favour of a low-mass companion.
+ After all. one would expect changes in the amplitude anc phase of the variations if they were due to a surface structure," After all, one would expect changes in the amplitude and phase of the variations if they were due to a surface structure"
+certain regions of our focal plane. leaving others relatively uncderpopulated.,"certain regions of our focal plane, leaving others relatively underpopulated."
+ We have simulated. the performance of our optimizations in real-life situations by making use of ealaxy mock catalogs extracted from the Bolshoi simulation (?).., We have simulated the performance of our optimizations in real-life situations by making use of galaxy mock catalogs extracted from the Bolshoi simulation \citep{Klypin2010}.
+ Important. for this work. the clustering properties of mock galaxies match those of real galaxies with gooc accuracy.," Important for this work, the clustering properties of mock galaxies match those of real galaxies with good accuracy."
+ The reason for using mock catalogs instead of real catalogs is that they provide us with more Iexibility when it comes to selecting dillerent samples with dilleren number densities., The reason for using mock catalogs instead of real catalogs is that they provide us with more flexibility when it comes to selecting different samples with different number densities.
+ We have performed a 2-D projection of a simulation box of 250 Mpe/h on a side situated at. redshift l onto our focal plane and used the circular. velocity of haloes as an empirical threshold. for selecting cilleren densities., We have performed a 2-D projection of a simulation box of 250 Mpc/h on a side situated at redshift 1 onto our focal plane and used the circular velocity of haloes as an empirical threshold for selecting different densities.
+ In order to allow for à fair comparison with the results presented in the previous sections. we have selecte catalogs with target-to-positioner ratios of ~ 0.5. 1. 3 ane 5.," In order to allow for a fair comparison with the results presented in the previous sections, we have selected catalogs with target-to-positioner ratios of $\sim$ 0.5, 1, 3 and 5."
+ Simulations were performed with 9 different realizations for each η. obtained by rotating the box In Fig. 6..," Simulations were performed with 9 different realizations for each $\eta$, obtained by rotating the box In Fig. \ref{fig:mocks},"
+ we show a portion of the focal plane of an instrument. represented by an array of patrol discs. in a sequence of the first tiles.," we show a portion of the focal plane of an instrument, represented by an array of patrol discs, in a sequence of the first tiles."
+ The over-plotted dots svmbolize he projected. positions of targets belonging to a mock galaxy catalog with yo3., The over-plotted dots symbolize the projected positions of targets belonging to a mock galaxy catalog with $\eta \sim 3$.
+ The first panel in this figure illustrates the initial situation where the entire sample of argets is to be observed., The first panel in this figure illustrates the initial situation where the entire sample of targets is to be observed.
+ Note that. as explained. above. argets on the focal plane of the spectrograph. rather than ing randomly distributed. accumulate in filaments.," Note that, as explained above, targets on the focal plane of the spectrograph, rather than being randomly distributed, accumulate in filaments."
+ The rest of the panels show the distribution of targets assigned o each tile by using the optimized. algorithm described in Section 4.., The rest of the panels show the distribution of targets assigned to each tile by using the optimized algorithm described in Section \ref{sec:draining}.
+ In. this example. ~75% of targets have been selected alter 3 The performance of our optimizations with mock galaxy catalogs is presented in the same format às in previous sections in Fable 3. and Fig. 7..," In this example, $\sim 75\%$ of targets have been selected after 3 The performance of our optimizations with mock galaxy catalogs is presented in the same format as in previous sections in Table \ref{tab:real} and Fig. \ref{fig:real}."
+ X direct consequence of the presence of clustering in our catalogs is that the fraction of targets assigned to each tile decreases significantly. as à comparison between Table 20 anc Table 3.7 demonstrates.," A direct consequence of the presence of clustering in our catalogs is that the fraction of targets assigned to each tile decreases significantly, as a comparison between Table \ref{tab:random} and Table \ref{tab:real} demonstrates."
+ ηνίαιν. the probability that a single positioner has to deal with several targets is now higher ancl. consequently. it becomes harder to move targets towards the first. tiles.," Trivially, the probability that a single positioner has to deal with several targets is now higher and, consequently, it becomes harder to move targets towards the first tiles."
+ As an example. with 51 and two tiles we could observe almost 94% of all targets in a random catalog. even whithout allowing for rotation of the focal plane.," As an example, with $\eta \sim 1$ and two tiles we could observe almost $94 \%$ of all targets in a random catalog, even whithout allowing for rotation of the focal plane."
+ In a real catalog we could only assign S54 of targets in. the same number of tiles., In a real catalog we could only assign $85 \%$ of targets in the same number of tiles.
+ Similarly. in areal catalog with 5 we would need 5in tiles to barely reach a completeness. of SOUL. at least S4 below our expectations from. random catalogs.," Similarly, in areal catalog with $\eta \sim 5$ we would need 5 tiles to barely reach a completeness of $80 \%$ , at least $8\%$ below our expectations from random catalogs."
+ Again. we refer the reader to Table 3. for the exact fractions of targets assigned with a random approach. with the draining algorithm alone and with the draining algorithm complemented. with rotation of the focal plane.," Again, we refer the reader to Table \ref{tab:real} for the exact fractions of targets assigned with a random approach, with the draining algorithm alone and with the draining algorithm complemented with rotation of the focal plane."
+ In order to analyze the performance of our optimizations as compared to à random approach we point the reader to Fig. 7.., In order to analyze the performance of our optimizations as compared to a random approach we point the reader to Fig. \ref{fig:real}.
+ This figure shows that the gain provided. by our optimizations as compare to à greedy approach when implemented in a real-life situation is consistent. with what we could infer from. random catalogs., This figure shows that the gain provided by our optimizations as compare to a greedy approach when implemented in a real-life situation is consistent with what we could infer from random catalogs.
+ X closer inspection. however. reveals that the gain provided by the draining algorithm alone is slightly smaller now. falling below 2%. whereas the gain achieved by combining this algorithm with a ROT optimization remains in the range of 56%.," A closer inspection, however, reveals that the gain provided by the draining algorithm alone is slightly smaller now, falling below $2\%$, whereas the gain achieved by combining this algorithm with a ROT optimization remains in the range of $5-6\%$."
+ I only slight rotations were allowed (ROT? optimization) we could still improve the eLllicieney of the process in 3.54.5%. Note that. as mentioned previously. the results shown in this work were obtained. ignoring Liber collisions. as the ellect of these depends on the geometry of the fiber positioning robot itself. and. therefore. cannot be extrapolated to any fiber-fed spectrograph of this kind.," If only slight rotations were allowed (ROT2 optimization) we could still improve the efficiency of the process in $3.5-4.5\%$ Note that, as mentioned previously, the results shown in this work were obtained ignoring fiber collisions, as the effect of these depends on the geometry of the fiber positioning robot itself, and, therefore, cannot be extrapolated to any fiber-fed spectrograph of this kind."
+ The tvpical number of these events obviously increases in mock ealaxy catalogs (and hence in real catalogs) due to the fact that objects are more clustered (the fraction of collisions is almost negligible in random catalogs)., The typical number of these events obviously increases in mock galaxy catalogs (and hence in real catalogs) due to the fact that objects are more clustered (the fraction of collisions is almost negligible in random catalogs).
+ Even in real catalogs the fraction of objects in conflict. remains small for SLOL: even for y=5., Even in real catalogs the fraction of objects in conflict remains small for SIDE: $\lesssim 1\%$ even for $\eta=5$.
+ However. an important advantage of the draining algorithm is that collisions can be solved optimally. and. additional methods. are not. needed.," However, an important advantage of the draining algorithm is that collisions can be solved optimally, and additional methods are not needed."
+ The results presented in Table 3.7 and Fig., The results presented in Table \ref{tab:real} and Fig.
+ 7 would. basically not change if collisions were taken into account., \ref{fig:real} would basically not change if collisions were taken into account.
+ Only very slight variations are expected in the very last tiles. which are not relevant in real-life observations.," Only very slight variations are expected in the very last tiles, which are not relevant in real-life observations."
+ The results. presented. in this section confirm. that. despite the physical restrictions of this state-of-the-art. fiber positioning robots. it is possible to optimize a survey strateew in a remarkable wav just by assigning targets to tiles and positioners conveniently.," The results presented in this section confirm that, despite the physical restrictions of this state-of-the-art fiber positioning robots, it is possible to optimize a survey strategy in a remarkable way just by assigning targets to tiles and positioners conveniently."
+ In addition. these results represent a strong support for allowing the focal plane of future instruments to rotate.," In addition, these results represent a strong support for allowing the focal plane of future instruments to rotate."
+ We have shown how even a slight rotation produces a remarkable optimization., We have shown how even a slight rotation produces a remarkable optimization.
+ In the next section we discuss on the implications of our results., In the next section we discuss on the implications of our results.
+ The results on the optimization of the fiber positioning process that we present in this work are valid for any focal plane consisting of an array of positioners as that described in Section 3.., The results on the optimization of the fiber positioning process that we present in this work are valid for any focal plane consisting of an array of positioners as that described in Section \ref{sec:robot}. .
+ To first order. therefore. our results are not dependent on the size ofthe tile or the number," To first order, therefore, our results are not dependent on the size ofthe tile or the number"
+0.5 truecia The outer regions of M31 have become an increasingly complicated field of study as it has become clear in recent vears that the role of accretion in halo formation is of considerable importance.,0.5 truecm The outer regions of M31 have become an increasingly complicated field of study as it has become clear in recent years that the role of accretion in halo formation is of considerable importance.
+ present star-count maps showing what appear to be extensive tidal disturbances iu the halo of M1. including iu the vicinity of the massive elobular cluster Cl.," \citet{fer02} present star-count maps showing what appear to be extensive tidal disturbances in the halo of M31, including in the vicinity of the massive globular cluster G1."
+ Reitzel&Cuhathakurta(2002.hereafterRCGU2) find evidence for a subtle stream-like feature. in an outer halo field located 19 Ispc from the ceuter in projection along the southeastern ninor-axis. using a combination of kinematics and metallicity measuremeuts of red ejut brauch (RGB) stars;," \citet[][hereafter RG02]{rei02} find evidence for a subtle stream-like feature, in an outer halo field located 19 kpc from the center in projection along the southeastern minor-axis, using a combination of kinematics and metallicity measurements of red giant branch (RGB) stars."
+ Cuhbathakurta&Reitzel(2002) confirm that this feature continues along the nünor axis in two inner halo fields located near the globular clusters C312 and C302. located 11 and 7 kpe from the uucleus of M31 respectively near the southeastern minor axis.," \citet{guh02}
+ confirm that this feature continues along the minor axis in two inner halo fields located near the globular clusters G312 and G302, located 11 and 7 kpc from the nucleus of M31 respectively near the southeastern minor axis."
+ ΑΟ5 two closest satellites M32 aud NGC 205 are known to be undergoing tidal stripping (Choi.Culiathaurta.&Johustou2002):: vet there is no definite proposal for a companion that night have been responsible for the large-scale streams seen iu the halo., M31's two closest satellites M32 and NGC 205 are known to be undergoing tidal stripping \citep*{cho02}; yet there is no definite proposal for a companion that might have been responsible for the large-scale streams seen in the halo.
+ The area around Cl is a particularly interesting field to study; as this object has been proposed to be the core of a tidallv-disrupted dwarf galaxy (Alevlanetal.2001).," The area around G1 is a particularly interesting field to study, as this object has been proposed to be the core of a tidally-disrupted dwarf galaxy \citep{mey01}."
+. Tf this is the case. one uüght expect to find the tidal debris surrouudiug the main body of the object with velocities and metallicites simular to Cl itself.," If this is the case, one might expect to find the tidal debris surrounding the main body of the object with velocities and metallicites similar to G1 itself."
+ In addition. the field around CU is expected to have roughly equal umubers of N31 halo aud disk stars so this eives us the opportunity to study the disk population of AD further out (0—31 kpe) than las been done to date.," In addition, the field around G1 is expected to have roughly equal numbers of M31 halo and disk stars \citep*{rei98,hod95} so this gives us the opportunity to study the disk population of M31 further out $r\sim34$ kpc) than has been done to date."
+ Ferguson&Johuson(2001). study a field 30 kpe from the nucleus of M31 along the NE major axis aud estimate a mean metallicity of [Fe/II|=0.7: they find that the population is mostly old (>8 Ctr) but there is evidence for an intermediate-age population as well., \citet{fer01} study a field 30 kpc from the nucleus of M31 along the NE major axis and estimate a mean metallicity of $\rm[Fe/H]\simeq-0.7$; they find that the population is mostly old $>8$ Gyr) but there is evidence for an intermediate-age population as well.
+ Iu a paper that accoupanics this work. Richetal.(2003) report evidence for an iutermediate-age population," In a paper that accompanies this work, \citet{ric03} report evidence for an intermediate-age population"
+The discovery of these two brown dwarf events is remarkable because each event had a hieh magnification. aud therefore required a very small distance of closest approach.,"The discovery of these two brown dwarf events is remarkable because each event had a high magnification, and therefore required a very small distance of closest approach."
+ Such evenis are (herelore rare., Such events are therefore rare.
+ Each of the two events therefore represents a large number of additional brown-dwarf events of short duration., Each of the two events therefore represents a large number of additional brown-dwarf events of short duration.
+ It would be difficult to use the detection of these two events (to lormulate a realistic estimate of the total number of short-duration brown-dwarl-lens events presently expected., It would be difficult to use the detection of these two events to formulate a realistic estimate of the total number of short-duration brown-dwarf-lens events presently expected.
+ Nevertheless. these detections add plausibility to (he estimates we have mace above. based on rate calculations and the combined OGLE and MOA detection rates.," Nevertheless, these detections add plausibility to the estimates we have made above, based on rate calculations and the combined OGLE and MOA detection rates."
+ Figure 1 demonstrates (hat brown-dwarf events with τε in the range of 8—16 days can take place ab distances greater than a hundred pe for velocities in excess of ~50. km /., Figure 1 demonstrates that brown-dwarf events with $\tau_E$ in the range of $8-16$ days can take place at distances greater than a hundred pc for velocities in excess of $\sim 50~$ km $^{-1}$.
+ Note in addition. that the total volume. hence the number of possible lenses ancl the rate of lensing by anv given population of lenses. increases with distance from us. (," Note in addition, that the total volume, hence the number of possible lenses and the rate of lensing by any given population of lenses, increases with distance from us. ("
+See. e.g.. DiStefano 2008a. 2003b for details.),"See, e.g., Stefano 2008a, 2008b for details.)"
+ Therefore the largest number of brown-dwarl lenses generating 8— 16-day events should have velocities 50 kms. + and be located at distances larger (han LOO pc., Therefore the largest number of brown-dwarf lenses generating $8-16$ -day events should have velocities $> 50~$ km $^{-1}$ and be located at distances larger than $100$ pc.
+ This is consistent with (he events observed to date., This is consistent with the events observed to date.
+ Consider a planetary svstem in which one planet serves as a lens. producing a event wilh a measured value of Te. Suppose Chat the star orbited by the planet lens is detected.," Consider a planetary system in which one planet serves as a lens, producing a short-duration event with a measured value of $\tau_E.$ Suppose that the star orbited by the planet lens is detected."
+ Suppose further. that a sequence of high-resolution measurements allows the geometric parallax. proper motion. aud Einstein angle of the star to be measured (see. e.g. Di Stefano 2009).," Suppose further, that a sequence of high-resolution measurements allows the geometric parallax, proper motion, and Einstein angle of the star to be measured (see, e.g., Di Stefano 2009)."
+" The combination of Dj and 9,5, produces a high-precision value ol the gravitational mass. M,. of the star."," The combination of $D_L$ and $\theta_{E,\ast}$ produces a high-precision value of the gravitational mass, $M_\ast,$ of the star."
+ In general. the stellar mass is estimated based on spectral ancl flix information.," In general, the stellar mass is estimated based on spectral and flux information."
+ A direct measurement of the gravitational mass allows stellar models to be tested., A direct measurement of the gravitational mass allows stellar models to be tested.
+ In some cases. it max be possible to conduct subsequent transit or raclial-velocity studies to measure the gravitational mass of the star in a second wax. i.e. bv studsyiug the orbit of the planet that served as a lens and/or the orbits of other planets.," In some cases, it may be possible to conduct subsequent transit or radial-velocity studies to measure the gravitational mass of the star in a second way, i.e., by studying the orbit of the planet that served as a lens and/or the orbits of other planets."
+ Thus. for some stars orbited by planet lenses. we may be able to compare (he gravitational ass measured via lensing with the gravitational mass measured via orbital dvnamics.," Thus, for some stars orbited by planet lenses, we may be able to compare the gravitational mass measured via lensing with the gravitational mass measured via orbital dynamics."
+ Up to this point in our discussion. the planet has plaved only a peripheral role: (1) it produced a photometric event that alerted us to the possibility of measuring astrometric lensing by (he star. aud (2) it alerted us to (he presence of at least one planet orbiting the star. (hereby motivating subsequent transit and/or radial-velocity studies.," Up to this point in our discussion, the planet has played only a peripheral role: (1) it produced a photometric event that alerted us to the possibility of measuring astrometric lensing by the star, and (2) it alerted us to the presence of at least one planet orbiting the star, thereby motivating subsequent transit and/or radial-velocity studies."
+ The lensing event can of course leach us a good deal about the planet., The lensing event can of course teach us a good deal about the planet.
+ First. if finite-source-size elfects are detected. (hen 05prance can be directly measured.," First, if finite-source-size effects are detected, then $\theta_{E,planet}$ can be directly measured."
+ The distance to the planet is. to high precision. the same as the distance to the central star.," The distance to the planet is, to high precision, the same as the distance to the central star."
+" The combination of 0,prone and Dj measures the mass of the planet."," The combination of $\theta_{E,planet}$ and $D_L$ measures the mass of the planet."
+ With the gravitational masses of both the planet and star measured. the mass ratio q can be computed.," With the gravitational masses of both the planet and star measured, the mass ratio $q$ can be computed."
+ If. in addition. the projected orbital separation. a. between the central star and planet lens is less than roughlv 3.5Hp. or if the event “repeats”. then q and a can both be estimated from a fit to the planet-lens light curve.," If, in addition, the projected orbital separation, $a,$ between the central star and planet lens is less than roughly $3.5\, R_{E,\ast},$ or if the event “repeats”, then $q$ and $a$ can both be estimated from a fit to the planet-lens light curve."
+ The value of 4 so measured can be checked [or consistency with the measured values of the planets and stars gravitational lüasses., The value of $q$ so measured can be checked for consistency with the measured values of the planet's and star's gravitational masses.
+predominantly methanol.,predominantly methanol.
+ It is not easy. to untangle the blending and derive useful inlormation from the resulting line prolile., It is not easy to untangle the blending and derive useful information from the resulting line profile.
+" Instead we opted to observe (wo positions south of IRc2 with declination offsets of aancd100"".. respectively."," Instead we opted to observe two positions south of IRc2 with declination offsets of and, respectively."
+ The two positions correspond to local maxima in the IC] map of OMC-L by Schilke.Phillips.&Wang(1995)., The two positions correspond to local maxima in the HCl map of OMC-1 by \citet{spw1995}.
+. The (0. 60) position also coincides with the CS elump LSI (Mundyetal.1986).. and the (ο. 1007)) position Orion S. Figure 4 show the HC] J=1—0 spectra for the remaining sources in (he sample.," The (0, ) position also coincides with the CS clump LS1 \citep{mundy1986}, and the (0, ) position Orion S. Figure \ref{fig:spectra4} show the HCl $J=1-0$ spectra for the remaining sources in the sample."
+ The HC] hvperfine structure. resulting from the interaction between the electric [ield and chlorine nuclear spin. splits the J=1—0 transition into three components. with the outer (vo separated by —6.35 and 8.22 rrespectivelv from the strongest middle component for LICL.," The HCl hyperfine structure, resulting from the interaction between the electric field and chlorine nuclear spin, splits the $J=1-0$ transition into three components, with the outer two separated by $-6.35$ and 8.22 respectively from the strongest middle component for ."
+ ForCL. the separations are —5.05 aand 6.45|... respectively.," For, the separations are $-5.05$ and 6.45, respectively."
+ Statistical weights for the upper level of the hyperfine transitions dictate optically thin line strength ratio of 1:3:2., Statistical weights for the upper level of the hyperfine transitions dictate optically thin line strength ratio of 1:3:2.
+ In the observed sample. the (hree-component hvperline pattern is clearly seen in many sources. but (he components often do not conform to the optically thin ratio. indicating deviations rom the optically (hin limit or differences in excitation between (he three components.," In the observed sample, the three-component hyperfine pattern is clearly seen in many sources, but the components often do not conform to the optically thin ratio, indicating deviations from the optically thin limit or differences in excitation between the three components."
+ Under the assumption that the hvperfine components share the same excitation and kinematic characteristics. lits to the line profiles can independently provide the optical depth and excitation temperature (Xwan&Scoville1975).," Under the assumption that the hyperfine components share the same excitation and kinematic characteristics, fits to the line profiles can independently provide the optical depth and excitation temperature \citep{kwan1975}."
+. Such a utility is readily available in CLASS. the spectral line data reduction package developed ancl maintained by IAM.," Such a utility is readily available in CLASS, the spectral line data reduction package developed and maintained by IRAM."
+ The HFS fit in CLASS provides four parameters: AT+Traine Vpsg. we the line width. ad Traine," The HFS fit in CLASS provides four parameters: $\Delta T\cdot\tau_{main}$, $V_{LSR}$, $w$ the line width, and $\tau_{main}$."
+ Tain Is the total opticaldepth from all hvperfine components of the line. and," $\tau_{main}$ is the total opticaldepth from all hyperfine components of the line, and"
+for calibration‘ of the interferometer. pnphase. enhancing‘? the coherence time of the visibilities ‘and ‘allowing© the detection of fainter target,"for calibration of the interferometer phase, enhancing the coherence time of the visibilities and allowing the detection of fainter targets."
+ The typical s.one sigma image sensitivity. cerived [rom these datasets is approximately 1: mJv/beam., The typical one sigma image sensitivity derived from these datasets is approximately 1 mJy/beam.
+ The angular resolution varies with source declination but is tvpically approximately LOO mas., The angular resolution varies with source declination but is typically approximately 100 mas.
+ ligure 2 shows a typical e-VLBI image of one. of the sources in Table 1.. all of which were unresolved on scales of mamas or less.," Figure \ref{fig:J031010-573041} shows a typical e-VLBI image of one of the sources in Table \ref{table:observed_targets}, all of which were unresolved on scales of mas or less."
+ ‘Table 3 lists the maximum angular size of cach source as measured [from the images using the MIIRLAD task IMETE., Table \ref{table:position_and_size} lists the maximum angular size of each source as measured from the images using the MIRIAD task IMFIT.
+ Xn c-VLBI position is also listed for each source along with the flux density measured [rom the e-VLBL (~ O.Laaresee GCGllzbeam) and ΑΟ ( l5aaresec beam) images., An e-VLBI position is also listed for each source along with the GHz flux density measured from the e-VLBI $\sim0.1$ arcsec beam) and AT20G $\sim15$ arcsec beam) images.
+ Phe uncertainty in the the eVLBI positions is dominated by the phase referencing of the calibration source. and is typically muimas in RA and DEC.," The uncertainty in the the e-VLBI positions is dominated by the phase referencing of the calibration source, and is typically mas in RA and DEC."
+ Note that the [ANT20€C. and eVLBI Iux-density measurements are not simultaneous and. were made up to three vears apart., Note that the AT20G and e-VLBI flux-density measurements are not simultaneous and were made up to three years apart.
+ The mean Lux ratio (ονοι vrsuc) is 0.90 with a standard: deviation of 0.22., The mean flux ratio $_{\rm VLBI}$ $_{\rm AT20G}$ ) is 0.90 with a standard deviation of 0.22.
+ Howe exclude the source 255450. which appears to be variable (see SS44.1.2 below). the mean [lux ratio rises to 0.94 and the standard deviation drops to 0.15.," If we exclude the source $-$ 255450, which appears to be variable (see 4.1.2 below), the mean flux ratio rises to 0.94 and the standard deviation drops to 0.18."
+ These results suggest that (i) the nearby AT20G GPS sources are compact. with ~90% of their GGlIA emission arising on scales smaller than," These results suggest that (i) the nearby AT20G GPS sources are compact, with $\sim$ of their GHz emission arising on scales smaller than"
+value of E/AM results in higher expausiou velocitics of the ejecta and hotter post-shock gas. and a lower equilibrium ionization parzuuoeter.,"value of $E/M$ results in higher expansion velocities of the ejecta and hotter post-shock gas, and a lower equilibrium ionization parameter."
+ However. the τοι! also decelerates more rapidly. causing Z to increase at a faster rate than for the η=12 case.," However, the remnant also decelerates more rapidly, causing $\Xi$ to increase at a faster rate than for the $n = 12$ case."
+ This ultimately lunders the formation of cool gas relative to the 12 case. and only a small amount is able to form (see Table 2)).," This ultimately hinders the formation of cool gas relative to the $n = 12$ case, and only a small amount is able to form (see Table \ref{tab:mass_cool_gas}) )."
+ Receutly. it is has become clear that there exists a class of type II supernova explosions which are under cuerectic(6.8... Zampier 2003)).," Recently, it is has become clear that there exists a class of type II supernova explosions which are under energetic, Zampieri \cite{Z2003}) )."
+ For two explosions examined in detail. Zampieri al.(2003)) found AZ=LIAL. and E060.9<10ere.," For two explosions examined in detail, Zampieri \cite{Z2003}) ) found $M \gtsimm 14 \Msol$ and $E \approx 0.6 - 0.9 \times 10^{51} \erg$."
+ We do uot expect remmauts with. such paraicters to evolve significantly differeutlv to our canonical models with Af=LOAL.. aud E=10ere., We do not expect remnants with such parameters to evolve significantly differently to our canonical models with $M = 10 \Msol$ and $E = 10^{51} \erg$.
+ We lave calculated svuthetic Lue profiles for ciission from the cooled eas, We have calculated synthetic line profiles for emission from the cooled gas.
+ The 2D axisviinietrie eric is rotated onto a 3D cartesian erid and the cussion from volume clemeuts containing cool gas was integrated udder the assmuption that it is optically thin and the volume emission rate varies as v7., The 2D axisymmetric grid is rotated onto a 3D cartesian grid and the emission from volume elements containing cool gas was integrated under the assumption that it is optically thin and the volume emission rate varies as $n^{2}$.
+ Since the gas is cool. thermal Doppler broadening is neslieible.," Since the gas is cool, thermal Doppler broadening is negligible."
+ A previous investigation (Dottorff 2000)) failed to reach aux strong couclisious concerning whether was favoured by observations. so it is not iucluded im our model.," A previous investigation (Bottorff \cite{BFBK2000}) ) failed to reach any strong conclusions concerning whether was favoured by observations, so it is not included in our model."
+ The cussion is blue- or redshifted according to the line of sight velocity of the eas., The emission is blue- or redshifted according to the line of sight velocity of the gas.
+ Absorption was also asstuned to be neelieible., Absorption was also assumed to be negligible.
+fixed to 1.32. as reported by Capectal.(2007). from the ddata. eives a very poor result (42/5 = 897/16).,"fixed to 1.32, as reported by \citet{grupe07} from the data, gives a very poor result $\chi^2/\nu$ = 897/46)."
+ The fit cau be improved by leaving the decay slope as a free parameter refley its... model2)," The fit can be improved by leaving the decay slope as a free parameter \\ref{lc_fits}, model 2)."
+ Thisresultsinasingledecaugslopeofos = ].I5x00.010. but the light curve still deviates sieuificautlv a later times from this slope. resulting iu an unacceptable y2/r = 100/15.," This results in a single decay slope of $\alpha_3$ = 0.01, but the light curve still deviates significantly at later times from this slope, resulting in an unacceptable $\chi^2/\nu$ = 400/45."
+ A broken power law fit to the entire late-tine Πο curve (nodel 3) reveals a break at about 2 Ms after the burst: iu contrast to the result of Capeetal.(2007).. in which we could fit the late-tine ddata with just one decay slope. the addition of the 2007 data requires a break in the ddata.," A broken power law fit to the entire late-time light curve (model 3) reveals a break at about 2 Ms after the burst; in contrast to the result of \citet{grupe07}, in which we could fit the late-time data with just one decay slope, the addition of the 2007 data requires a break in the data."
+ The late decay slope. αι = 1.85. isdriven by the last two oobservations. while the \? is also strongly affected by two very ligh data points at ~2 MIs and ~5 Als. Makine the assumption that these two high points are late-time N-rav flares unrelated to the afterelow of the external shock. we removed them frou futher fits (sce models Laud 5).," The late decay slope, $\alpha_4$ = 1.85, isdriven by the last two observations, while the $\chi^2$ is also strongly affected by two very high data points at $\sim2$ Ms and $\sim5$ Ms. Making the assumption that these two high points are late-time X-ray flares unrelated to the afterglow of the external shock, we removed them from further fits (see models 4 and 5)."
+" We then fit the data between LOO ks and 30 MIs with a broken power law Guodel6). obtaining a break of Diveak.3=1.01 MIs aud slopes of tlie(2132!dadete002 anda,=L61. ν"," We then fit the data between 100 ks and 30 Ms with a broken power law (model 6), obtaining a break time of $T_{\rm break,3} = 1.01^{+0.35}_{-0.22}$ Ms and slopes of $\alpha_3 = 1.32^{+0.02}_{-0.05}$ and $\alpha_4 = 1.61^{+0.10}_{-0.06}$."
+αThis fit is plotted as d Jine iin m rofxravjc.., This fit is plotted as the dashed line in \\ref{xray_lc}.
+ The last two oobservatious deviate from this fit. sugecsting a break at about a vear after the burst.," The last two observations deviate from this fit, suggesting a break at about a year after the burst."
+ Because these last two points lave very few counts (and consequently large nucertaimtics). a broken power law fit to the late-time heht curve cannot constrain either the break time or the late-time decay slope a5 uuless at least one parameter is fixed.," Because these last two points have very few counts (and consequently large uncertainties), a broken power law fit to the late-time light curve cannot constrain either the break time or the late-time decay slope $\alpha_5$ unless at least one parameter is fixed."
+ We therefore approached the question of a final break in steps., We therefore approached the question of a final break in steps.
+" A suele power law fit to the lieht curve for T>1.2 Ms refle pits. model7ygivesa, = 1.68400.08 aud ὧν = 12/15."," A single power law fit to the light curve for $T \geq 1.2$ Ms \\ref{lc_fits}, model 7) gives $\alpha_4$ = 0.08 and $\chi^2/\nu$ = 12/15."
+ Although this is already an acceptable fit. we investigated the possibility of a late-time break which is expected from CRB theory (c.f.ce.Zhangetal.2006for 30060).," Although this is already an acceptable fit, we investigated the possibility of a late-time break which is expected from GRB theory \citep[c.f. e.g.][]{zhang06, meszaros06}. ."
+" At first we fitted the light curve 1.2 Ms xTzx 35 Ms with a single power law Guodel SN) which results in a, = 1.61IM and uvy = 6/13.", At first we fitted the light curve for $1.2$ Ms $\leq T \leq$ 35 Ms with a single power law (model 8) which results in $\alpha_4$ = $1.61^{+0.07}_{-0.13}$ and $\chi^2/\nu$ = 6/13.
+ We then fitted a broken power law model to the eutire heht curve with 7> 1.2 Ms with αι fixed at 1.61 (the best-fit result when the last two oobservations are excluded: model δ in refley/ts) jtodeterminewhetherthedatarcquirea ver ylatebredk Adthetigittdonce fHe," We then fitted a broken power law model to the entire light curve with $T
+>$ 1.2 Ms with $\alpha_4$ fixed at 1.61 (the best-fit result when the last two observations are excluded; model 8 in \\ref{lc_fits}) ) to determine whether the data require a very late break in the light curve slope."
+gits. c inodelüygieesTisa) = 43/22. Ms and ο”.=1.65ML with 4?/»=6/11.," This fit \\ref{lc_fits}, model 9) gives $T_{\rm break,4}$ = $^{+4.2}_{-5.1}$ Ms and $\alpha_5=4.65^{+2.05}_{-1.34}$, with $\chi^2/\nu=6/14$."
+ displavs the contour plot between the final break time and the final slope., \\ref{contour} displays the contour plot between the final break time and the final slope.
+ It shows that they are still not swelbconstrained., It shows that they are still not well-constrained.
+ Although the best-fit breax time is [1 Ms (2007 November). a break as carly as 26 Ms. with a late-time decay slope of a5 = 2.5. is consistent with the data at the lo level.," Although the best-fit break time is 41 Ms (2007 November), a break as early as $\sim 26$ Ms, with a late-time decay slope of $\alpha_5$ = 2.5, is consistent with the data at the $1 \sigma$ level."
+ Iu addition to the broken power law fits with a sharp break. the late-time light. curve was also fitted by the sincothed double-broken power law model defined bx Beucrimannetal.(1999).," In addition to the broken power law fits with a sharp break, the late-time light curve was also fitted by the smoothed double-broken power law model defined by \citet{beuermann99}."
+. Tere we found decay. slopes a3=Ll.20400.07. ay=1.70400.05. and 0522. 1100.26 with break times at LO9ELTL.OL As and 20 Mx (fixed).," Here we found decay slopes $\alpha_3$ 0.07, $\alpha_4$ 0.05, and $\alpha_5$ 0.26 with break times at 1.01 Ms and 20 Ms (fixed)."
+ The smooth paramcter is fixed to 3.0 and 2.0 for the breaks at about 1 Ms and 20 Ms. respectively.," The smooth parameter is fixed to 3.0 and 2.0 for the breaks at about 1 Ms and 20 Ms, respectively."
+ This results in an acceptable fit with 42/» = 12/32., This results in an acceptable fit with $\chi^2/\nu$ = 42/32.
+ Possible interpretations of these temporal breaks are discussed iu rofdiscuss.., Possible interpretations of these temporal breaks are discussed in \\ref{discuss}.
+ Temporal breaks are often associated with spectral breaks (e.g.Sarictal.1998:MészárosotZhangetal.," Temporal breaks are often associated with spectral breaks \citep[e.g. ][]{sari98, meszaros98, zhang06}."
+2006)..— τοςayrprdisplagstheSwift NNReountrateandh, \\ref{lc_cr_hr} displays the XRT count rate and hardness ratio light curves for the interval between 100 ks and 5 Ms after the burst.
+ardnessratiolig, The hardness ratios are plotted segment by segment.
+htearces forthe itercalbet ac 0.3. after the break the spectruu hardens to IIR. ~0.15 with even larder values at later times.," While the hardness ratios before the break at 1 Ms after the burst are of order HR $\sim 0.3$ , after the break the spectrum hardens to HR $\sim 0.45$ with even harder values at later times."
+ The spectrum before the 1. Ms break cau be fitted with asingle : absorbed power law with: Vy-=(31nmU.2655)«102!102 cm7 and an cucrey spectral slope 4 = L1s00.11 qev = Bl/s2)," The spectrum before the 1 Ms break can be fitted with a single absorbed power law with $N_{\rm H} = (1.34^{+0.27}_{-0.25})\times
+10^{21}$ $^{-2}$ and an energy spectral slope $\beta_{\rm x}$ = 0.11 $\chi^2/\nu$ = 81/82)."
+ The spectrum after the 1 Ms break was also fitted by an absorbed single power law uodel., The spectrum after the 1 Ms break was also fitted by an absorbed single power law model.
+ Leaving the absorption column density as a free xwanmeter. however. results in an merease of the column deusitv. which does not secur plausible.," Leaving the absorption column density as a free parameter, however, results in an increase of the column density, which does not seem plausible."
+ Therefore we fixed the absorption column deusity to Ny=131«10?! cn7. the value. obtained before the break.," Therefore we fixed the absorption column density to $N_{\rm H} = 1.34\times 10^{21}$ $^{-2}$, the value obtained before the break."
+ This fit results in a slightly flatter cucrey spectral slope 3. = 8S9-EOO0.11., This fit results in a slightly flatter energy spectral slope $\beta_{\rm x}$ = 0.11.
+ These values were used in PIMMS to convert he AACTS-S count rates into the fluxes eiven in retsrav oyandplottedin re feraye., These values were used in PIMMS to convert the ACIS-S count rates into the fluxes given in \\ref{xray_log} and plotted in \\ref{xray_lc}.
+ The ddata μι be analyzed in the Poison Πιτ. coluplicating proper analysis of possible spectral various at vorv late times.," The data must be analyzed in the Poisson limit, complicating proper analysis of possible spectral variations at very late times."
+" Usine the Bayesian approach described by Parkctal.(2006).. we estimated the harduess ratios in the aud their uncertainties. both before aud after the break at 38 Ms. We obtain mean values of 0.39 for the 2007 Marchli-June data (before the final break: 38 counts total) and 0.80 for the very late data (after the final break: δ counts total). with confidence lanits of IIR—0.60 to O17 and TR= L00to 0.58 respectively,"," Using the Bayesian approach described by \citet{park06}, we estimated the hardness ratios in the and their uncertainties, both before and after the break at 38 Ms. We obtain mean values of $-0.39$ for the 2007 March-June data (before the final break; 38 counts total) and $-0.80$ for the very late data (after the final break; 8 counts total), with confidence limits of $= -0.60$ to $-0.17$ and $= -1.00$ to $-0.58$ , respectively."
+ Although this is a suggestion of spectral softening across the final break. wecannot exclude (at the sos LTbes fyestiloblity that the μανάτος» ratio is constant.," Although this is a suggestion of spectral softening across the final break, wecannot exclude (at the confidence level) the possibility that the hardness ratio is constant."
+ Note that due to the different energv bands and detector response matrices it is not possible to compare the aud: hharducss ratios directly., Note that due to the different energy bands and detector response matrices it is not possible to compare the and hardness ratios directly.
+ Even though GRB 0607209 was one of the brightest bursts detected in X-rays. it is uot the brightest oue so," Even though GRB 060729 was one of the brightest bursts detected in X-rays, it is not the brightest one so"
+"Using the K-S test, we find the likelihoods of 1x107+ and 0.016 that the OCDFs of the simulated HVSs are the same as those from observations for the first and second population of the detect HVSs, respectively.","Using the K-S test, we find the likelihoods of $1\times 10^{-4}$ and $0.016$ that the $\Theta$ CDFs of the simulated HVSs are the same as those from observations for the first and second population of the detect HVSs, respectively."
+" Therefore, we conclude that the detected HVSs are highly unlikely to be produced from the tidal breakup of isotropically distributed progenitorial binary stars."," Therefore, we conclude that the detected HVSs are highly unlikely to be produced from the tidal breakup of isotropically distributed progenitorial binary stars."
+" This further strengthens the conclusion obtained in Luetal. (2010),, i.e., the detected HVSs are probably originated from two thin disks with orientations similar to the CWS disk and the northern arm of the minispiral (or the warped outer part of the CWS disk) in the GC, respectively."," This further strengthens the conclusion obtained in \citet{Luetal10}, i.e., the detected HVSs are probably originated from two thin disks with orientations similar to the CWS disk and the northern arm of the minispiral (or the warped outer part of the CWS disk) in the GC, respectively."
+ Figure 13 shows the vCDFs for both the simulated HVSs (obtained from different models) and the observations., Figure \ref{fig:f13} shows the $v$ CDFs for both the simulated HVSs (obtained from different models) and the observations.
+" Our numerical simulations show that the vCDF is almost independent of the thickness of the disk(s) where the HVS progenitors are originated, but it does depend on how close the stellar binaries can approach the MBH and on the initial distribution of the semimajor axes of the stellar binaries."," Our numerical simulations show that the $v$ CDF is almost independent of the thickness of the disk(s) where the HVS progenitors are originated, but it does depend on how close the stellar binaries can approach the MBH and on the initial distribution of the semimajor axes of the stellar binaries."
+ Different models produce quite different vCDFs., Different models produce quite different $v$ CDFs.
+" In the “LP” models and the “UB” models, for example, more than of the resulted HVSs (with v92700kms!) have velocities larger than the maximum velocity of the detected HVSs (c1000km51), while the “RW” models can produce a steep vCDF which is quite similar to the observational ones."," In the “LP” models and the “UB” models, for example, more than of the resulted HVSs (with $v^{\infty}_{\rm ej}\ga700 \kms$ ) have velocities larger than the maximum velocity of the detected HVSs $\sim1000 \kms$ ), while the “RW” models can produce a steep $v$ CDF which is quite similar to the observational ones."
+" The models with log-normal distributions of ap; produce less HVSs at the high-velocity end because of the fraction of stellar binaries with small ap; (i.e., <0.3AU) is substantially smaller compared with that in those models with the Oppik law."," The models with log-normal distributions of $a_{\rm b,i}$ produce less HVSs at the high-velocity end because of the fraction of stellar binaries with small $a_{\rm b,i}$ (i.e., $\la 0.3\AU$ ) is substantially smaller compared with that in those models with the Öppik law."
+" For the first HVS population, our K-S tests find 2.7x1079 (7.9x107?) and 2.5x107 (1.3x1072) likelihoods that the vCDFs obtained from the “LP-1” (“UB-1”) model and the “LP-2” (“UB-2”) model are drawn from the same distribution as the observational ones, respectively, which suggests that the first HVS population is unlikely to be produced by either of the “LP” model and the “UB” model."," For the first HVS population, our K-S tests find $2.7\times 10^{-6}$ $7.9\times 10^{-3}$ ) and $2.5\times 10^{-4}$ $1.3\times
+10^{-2}$ ) likelihoods that the $v$ CDFs obtained from the “LP-1” (“UB-1”) model and the “LP-2” (“UB-2”) model are drawn from the same distribution as the observational ones, respectively, which suggests that the first HVS population is unlikely to be produced by either of the “LP” model and the “UB” model."
+" For the second HVS population, the K-S likelihood is 0.01 (0.05) and 0.06 (0.07) for the “LP-1” (“UB-1”) model and the “LP- (“UB-2”) model, respectively."," For the second HVS population, the K-S likelihood is $0.01$ $0.05$ ) and $0.06$ $0.07$ ) for the “LP-1” (“UB-1”) model and the ``LP-2'' (“UB-2”) model, respectively."
+" These numbers suggest that the second population is not likely to be produced by the “LP” or “UB” models with an initial ap,; distribution of 1/ay,;; but it may not be inconsistent with the “LP” (or “UB”) models with a log-normal distribution of ay, (though with limited statistics)."," These numbers suggest that the second population is not likely to be produced by the “LP” or “UB” models with an initial $a_{\rm b,i}$ distribution of $1/a_{\rm b,i}$ but it may not be inconsistent with the “LP” (or “UB”) models with a log-normal distribution of $a_{\rm b,i}$ (though with limited statistics)."
+" However, the vCDFs resulted from the “RW” models appear to be consistent with the observations as the K-S tests find the likelihoods of 0.52 and 0.43 (0.52) that the observational vCDFSs of the (0.13)first (second) HVS population are the same as that obtained from the “RW-1” model and the *RW-2"" model, respectively."," However, the $v$ CDFs resulted from the “RW” models appear to be consistent with the observations as the K-S tests find the likelihoods of 0.52 (0.13) and 0.43 (0.52) that the observational $v$ CDFs of the first (second) HVS population are the same as that obtained from the “RW-1” model and the ``RW-2'' model, respectively."
+ Adopting a different form of the Galactic potential may affect the estimation of the vCDF for the detected HVSs in Section 2 and the selection effects discussed in Section 4.1., Adopting a different form of the Galactic potential may affect the estimation of the $v$ CDF for the detected HVSs in Section 2 and the selection effects discussed in Section 4.1.
+" For example, if adopting a simple Galactic potential model as that described by Equation (8) in Kenyonetal.(2008), the ug? of the detected HVSs ranges from 850kms! to 1200kms! and the slope of the vCDF is slightly flatter than that estimated in Section 2."," For example, if adopting a simple Galactic potential model as that described by Equation (8) in \citet{Kenyon08}, the $v^{\infty}_{\rm ej}$ of the detected HVSs ranges from $850 \kms$ to $1200 \kms$ and the slope of the $v$ CDF is slightly flatter than that estimated in Section 2."
+" The simulated vCDFs from the models of “LP-1”, *LP-2"", “UB-1”, and “UB-2” are not likely to be consistent with the vCDF of the detected HVSs, while both the “RW-1” model and the “RW-2” model can produce a vCDF similar to that estimated for the detected HVSs according to the new Galactic potential."," The simulated $v$ CDFs from the models of “LP-1”, “LP-2”, “UB-1”, and “UB-2” are not likely to be consistent with the $v$ CDF of the detected HVSs, while both the “RW-1” model and the “RW-2” model can produce a $v$ CDF similar to that estimated for the detected HVSs according to the new Galactic potential."
+ Our main conclusion that the TBK mechanism can reproduce the detected vCDF made in this section is not affected by the choice of the Galactic potential (also see discussion in Sesanaetal.(2007) and Kenyon (2008)))., Our main conclusion that the TBK mechanism can reproduce the detected $v$ CDF made in this section is not affected by the choice of the Galactic potential (also see discussion in \citet{Sesana07} and \citet{Kenyon08}) ).
+" 'To close this section, we note here that the fraction of stellar binaries resulting in ejection of HVSs with properties similar to the detected ones is around"," To close this section, we note here that the fraction of stellar binaries resulting in ejection of HVSs with properties similar to the detected ones is around"
+a inagnetar progenitor. and hiehlights a discrepancy vetween the presence of RSCs in Wel aud the predictions of evolutionary models. which sugecst that tle most uunous RSCGs should evolve from siguificautlv lower isses.,"a magnetar progenitor, and highlights a discrepancy between the presence of RSGs in Wd1 and the predictions of evolutionary models, which suggest that the most luminous RSGs should evolve from significantly lower masses."
+ A first study of the binary fraction aunongst owoer-Iqnuinositvy ate-O II-III stars dno Wdl will be oreseuted. di ao subsequent paper iu this seres. but uaznyv binary svstenis are already available for follow-up study.," A first study of the binary fraction amongst lower-luminosity late-O II-III stars in Wd1 will be presented in a subsequent paper in this series, but many binary systems are already available for follow-up study."
+ These include short-period spectroscopic biuariesE32...W23003:: Paper D. eclipsing binaries within the WR. OD supereiant and main sequence populations (Bonanos2007).. and N-rav. aud racdio-selected collidius-wind binaries (Clarketal.2008:Dougherty2010).," These include short-period spectroscopic binaries,; Paper I), eclipsing binaries within the WR, OB supergiant and main sequence populations \citep{bonanos}, and X-ray and radio-selected colliding-wind binaries \citep{clark08, dougherty}."
+. Consideration of these data will allow further dvuamiucal constraints to be placed ou the progenitor masses of the evolved stars within Wdl as well as the general mass i1uninositv relation for stars in the upper reaches of the IIR diagram ae the post-MS patlisvavs they follow., Consideration of these data will allow further dynamical constraints to be placed on the progenitor masses of the evolved stars within Wd1 as well as the general mass luminosity relation for stars in the upper reaches of the HR diagram and the post-MS pathways they follow.
+" Moreover. hey will vield the first characterisation of the binary xoperties of a homogeneous population of massive stars. of critical importance for studies of both star aud cluster ornation and miuuerous hieh-energy phenomena such as ""üupernovae. enmlla-rav bursters and the formation of ugh mass X-ray binaries."," Moreover, they will yield the first characterisation of the binary properties of a homogeneous population of massive stars, of critical importance for studies of both star and cluster formation and numerous high-energy phenomena such as supernovae, Gamma-ray bursters and the formation of high mass X-ray binaries."
+where Tarsens (Teen) Is the true optical depth of the strong (weak) component. 71. 7» are the observed (apparent) optical depths and ἐν. {ο are the observed residual intensitics (normalized. by the local continuum. lu) al the same velocity for the strong and weak component of the doublet. respectively.,"where $\tau_{\rm strong}$ $\tau_{\rm weak}$ ) is the true optical depth of the strong (weak) component, $\tau_1$, $\tau_2$ are the observed (apparent) optical depths and $I_1$, $I_2$ are the observed residual intensities (normalized by the local continuum, $I_0$ ) at the same velocity for the strong and weak component of the doublet, respectively."
+ If the lines are optically thin. then from these exact expressions the relation can be obtained which relates the covering factor to the observed residual intensities (Llamann Ferlanc 1999).," If the lines are optically thin, then from these exact expressions the relation can be obtained which relates the covering factor to the observed residual intensities (Hamann Ferland 1999)."
+ 1n what follows we shall use the label # to indicate a covering factor value obtained by means ofintensities., In what follows we shall use the label $R$ to indicate a covering factor value obtained by means of.
+ All lines lor which an adjustment of the zero levelhas been performed with acre listecl in ‘Table 1..., All lines for which an adjustment of the zero levelhas been performed with are listed in Table \ref{tab:adj}.
+ The Zero level adjustment. determined withVPELIT.. as a fraction of the continuum is eiven in column 7 and its le error in column 8.," The zero level adjustment, determined with, as a fraction of the continuum is given in column 7 and its $1\sigma$ error in column 8."
+ The covering actor determined. from.4... {μι for these lines is given in column 9 of the Table.," The covering factor determined from, $f_R$, for these lines is given in column 9 of the Table."
+ In principle. since the zero evel adjustment is introduced to compensate for an excess continuum Hux. the value for this adjustment should be in agreement with 1fn.," In principle, since the zero level adjustment is introduced to compensate for an excess continuum flux, the value for this adjustment should be in agreement with $1-f_R$."
+ Ehe last column in Table 1 shows he value of 1fe For all entries. fj. in column 9.," The last column in Table \ref{tab:adj} shows the value of $1-f_R$ for all entries, $f_R$, in column 9."
+ In other words. the last column can be thought of as giving an aadjustment based. on the residual. intensities.," In other words, the last column can be thought of as giving an adjustment based on the residual intensities."
+ In one hen compares corresponding entries in columns 7 (actual adjustment used) and 10 (expected adjustment from residual intensities). one may notice that these agree to within leo for line 17. to within zle for lines 1. 4 and 15. and to within 20 or more for the rest of the lines.," If one then compares corresponding entries in columns 7 (actual adjustment used) and 10 (expected adjustment from residual intensities), one may notice that these agree to within $1\sigma$ for line 17, to within $\approxgt 1\sigma$ for lines 1, 4 and 15, and to within $2\sigma$ or more for the rest of the lines."
+ We call this an It-tvpe comparison (see below) anc we shall return to this point later., We call this an R-type comparison (see below) and we shall return to this point later.
+ Further. from Equations 2. and 3.2 it is clear that the ratio of the observed optical depths. 7;/7». depends both on f and on the true optical depth for cach component. and thus. for given b. on the column density.," Further, from Equations \ref{equ:et1} and \ref{equ:et2} it is clear that the ratio of the observed optical depths, $\tau_1/\tau_2$, depends both on $f$ and on the true optical depth for each component, and thus, for given $b$, on the column density."
+ This.fheorcticad result is illustrated in for a line with b=10 aand clilferent values of the column density and covering factor., This result is illustrated in for a line with $b=10$ and different values of the column density and covering factor.
+ lt can be seen hat for low column densities and/or high covering factors. this ratio is close tο 2.," It can be seen that for low column densities and/or high covering factors, this ratio is close to 2."
+ This corresponds to the case where an ADR. is not οservable., This corresponds to the case where an ADR is not observable.
+ Llowever. for given f and b there is a threshold in column density. above which the ratio will deviate significantly from 2 and. depending on the SNR in the region of the line. it will not be possible to fit the line with a model Voigt profile (which assumes a ratio of 2) without obtaining large resicluals.," However, for given $f$ and $b$ there is a threshold in column density, above which the ratio will deviate significantly from 2 and, depending on the SNR in the region of the line, it will not be possible to fit the line with a model Voigt profile (which assumes a ratio of 2) without obtaining large residuals."
+ This elect is seen in1: relative to 75. τι is less than what eexpects and. hence. the fit appears markedly below the 1548 Iline ancl above the 1550 Iline.," This effect is seen in: relative to $\tau_2$, $\tau_1$ is less than what expects and, hence, the fit appears markedly below the 1548 line and above the 1550 line."
+ This discrepancy is due to excess flux. getting. through atthe velocity where a particular doublet: is observed., This discrepancy is due to excess flux getting through atthe velocity where a particular doublet is observed.
+ Lf the truce optical depth and the covering factor are known. the observed. continuum normalized. intensity is given by Equations 2. or 9 3..," If the true optical depth and the covering factor are known, the observed, continuum normalized intensity is given by Equations \ref{equ:et1} or \ref{equ:et2}. ."
+ Thus. although the continuum," Thus, although the continuum"
+1984).. but. no previous X-ray spectral or timing analvsis studies have been carried out for it.,", but no previous X-ray spectral or timing analysis studies have been carried out for it."
+ In. order to ensure that these objects are not. intermediate: polars (APs) and to look for orbital ancl spin modulation in the data. the power spectra were calculated by using a Lomb- periodogram which is used for period analvsis of unevenly spaced data.," In order to ensure that these objects are not intermediate polars (IPs) and to look for orbital and spin modulation in the data, the power spectra were calculated by using a Lomb-Scargle periodogram which is used for period analysis of unevenly spaced data."
+ When searching over the [requeney range 0.00001.0.03 Lz. no significant periodicities were seen at the 99 per cent confidence level.," When searching over the frequency range 0.00001–0.03 Hz, no significant periodicities were seen at the 99 per cent confidence level."
+ We carried out X-ray spectral analysis in order to study the underlsing spectra of the source sample. ancl. ultimately. to caleulate the I[uxes and luminosities of the sources.," We carried out X-ray spectral analysis in order to study the underlying spectra of the source sample, and, ultimately, to calculate the fluxes and luminosities of the sources."
+ To employ. Gaussian statistics. the N-ray spectra were binned al 20 ct  with and then fitted in1996).," To employ Gaussian statistics, the X-ray spectra were binned at 20 ct $^{-1}$ with and then fitted in."
+. In CVs. the power source of X-ray. emission is known to be accretion onto the white dwarl.," In CVs, the power source of X-ray emission is known to be accretion onto the white dwarf."
+ Phe accreted material is shock-heated to high temperatures (ρω 10502003). and this material has to cool before settling onto the white cwarl surface.," The accreted material is shock-heated to high temperatures $\sim$ 10–50, and this material has to cool before settling onto the white dwarf surface."
+" Thus. the cooling gas [low is assumed to consist of a range of temperatures which vary from the hot shock temperature 75,4 to the empoerature of the optically thin cooling material which eventually settles onto the surface. of the white cwarl1997)."," Thus, the cooling gas flow is assumed to consist of a range of temperatures which vary from the hot shock temperature $kT_{max}$ to the temperature of the optically thin cooling material which eventually settles onto the surface of the white dwarf."
+. Thus. when fitting X-ray spectra of CVs. cooling How spectral models should. represent more ohvsically correct. picture of the cooling plasma. unlike single temperature. spectral models.," Thus, when fitting X-ray spectra of CVs, cooling flow spectral models should represent more physically correct picture of the cooling plasma, unlike single temperature spectral models."
+ Cooling Low naiocdels rave successfully been applied. to CV spectra in. previous studies by c.g. and., Cooling flow models have successfully been applied to CV spectra in previous studies by e.g. and.
+(2003).. In this view. the multi-temperature characteristic is our motivation for emphasizing the cooling Low mocel in the rest of this work.," In this view, the multi-temperature characteristic is our motivation for emphasizing the cooling flow model in the rest of this work."
+" The cdillerential emission measure GEMAE for an isobaric cooling How can be described. by where nm, is the mass of a proton. ji the mean molecular weight (~ 0.6). οςια) total emissivitv per volume in units oberg tem 7. M accretion rate. p. particle density. and & the Boltzmann constant."," The differential emission measure dEM/dT for an isobaric cooling flow can be described by where $m_{p}$ is the mass of a proton, $\mu$ the mean molecular weight $\sim$ 0.6), $\epsilon$ (T,n) total emissivity per volume in units of erg $^{-1}$ $^{-3}$, $\dot{M}$ accretion rate, $n$ particle density, and $k$ the Boltzmann constant."
+ The source of the N-rav. emission above the white dwarf illuminates the surface of the white dwarf and thus causes a reflection. which is seen as Fe hea iron Illuorescence line at 6.4 keV 1901).," The source of the X-ray emission above the white dwarf illuminates the surface of the white dwarf and thus causes a reflection, which is seen as Fe $\alpha$ iron fluorescence line at 6.4 keV ."
+. According to George Fabian. an infinite slab reflector subtencding a total solid angle of$2 = 2: where the X-ray source is located right above the slab. produces an equivalent width of up to  150 eV for the 6.4 keV. Fe Ίνα fluorescence line.," According to George Fabian, an infinite slab reflector subtending a total solid angle of $\Omega$ = $\pi$ where the X-ray source is located right above the slab, produces an equivalent width of up to $\sim$ 150 eV for the 6.4 keV Fe $\alpha$ fluorescence line."
+ Ehe equivalent width of the 6.4 keV. iron line depends on the total abundance of the reflector1997).. the inclination angle between the surface of the rellector and the observer's line of sight. and the photon index of the spectrum of the X-ray emission source2009).," The equivalent width of the 6.4 keV iron line depends on the total abundance of the reflector, the inclination angle between the surface of the reflector and the observer's line of sight, and the photon index of the spectrum of the X-ray emission source."
+. Even though we believe that the cooling How -tvpe multi-temperature model is the correct. description of the ηνμον of the cooling gas How in CVs. previous works have often used single temperature plasma moclels.," Even though we believe that the cooling flow -type multi-temperature model is the correct description of the physics of the cooling gas flow in CVs, previous works have often used single temperature plasma models."
+ Για». in order o compare the cllects of two cillerent spectral models on he spectral fit parameters. we fitted the spectra with 1) a single temperature optically thin thermal plasma moclel and 2) a cooling Low model (mkcflow)) whichJmews6. was originally. developed to cdeseribe the cooling flows in clusters of galaxies1988).. adding photoclectric absorption to both models.," Thus, in order to compare the effects of two different spectral models on the spectral fit parameters, we fitted the spectra with 1) a single temperature optically thin thermal plasma model and 2) a cooling flow model ) which was originally developed to describe the cooling flows in clusters of galaxies, adding photoelectric absorption to both models."
+ In order to investigate the equivalent. width of the 6.4 keV iron emission line. à Gaussian line was added. at 6.4 keV with a line width fixed at σξ 10 eV. The spectral fits did not necessarily require the 6.4 keV line. e.g. for SS Aur the ον = 0.96/629 when a Gaussian line at 6.4 keV. was not included.," In order to investigate the equivalent width of the 6.4 keV iron emission line, a Gaussian line was added at 6.4 keV with a line width fixed at $\sigma$ = 10 eV. The spectral fits did not necessarily require the 6.4 keV line, e.g., for SS Aur the $\chi^{2}_{\nu}$ $\nu$ = 0.96/629 when a Gaussian line at 6.4 keV was not included."
+ TheSusaku NIST and Χο. spectra were fitted simultaneously for each source as well as the GIS and SES spectra of Z Cam and. WZ See with the models mentioned above.," The XIS1 and XIS0,2,3 spectra were fitted simultaneously for each source as well as the GIS and SIS spectra of Z Cam and WZ Sge with the models mentioned above."
+ Some cata sets required additional components to improve the fits., Some data sets required additional components to improve the fits.
+ Three of the sources. IUE Cas. VS93 Seo and Z Cam. required. partial covering absorption model.pcfabs.. to reduce residuals in the [ow energv end (between ~ 0.62 keV).," Three of the sources, HT Cas, V893 Sco and Z Cam, required partial covering absorption model, to reduce residuals in the low energy end (between $\sim$ 0.6–2 keV)."
+ To reduce. residuals around 0.80. keV. in the SS Cvg spectrum. we added. a Gaussian line at 0:51 keV. with a line width of 0.24 keV letting the line energv and. width. both to vary free.," To reduce residuals around 0.80 keV in the SS Cyg spectrum, we added a Gaussian line at 0.81 keV with a line width of 0.24 keV letting the line energy and width both to vary free."
+ For U Gom. single absorbed: optically thin. thermal plasma model vielded a 7/79 = 2.23/403.," For U Gem, single absorbed optically thin thermal plasma model yielded a $\chi^{2}$ $\nu$ = 2.23/403."
+ Since thefit was not statistically satisfactory. we added a second. optically thin," Since thefit was not statistically satisfactory, we added a second optically thin"
+which we actually use for computationalefficiency?.,which we actually use for computational.
+". Finally, to complete the full specification of our chemical model, we need to estimate the column density of the molecular gas, Ny,, for the self-shielding factor given by Equation (ATi)."," Finally, to complete the full specification of our chemical model, we need to estimate the column density of the molecular gas, $N_\H2$, for the self-shielding factor given by Equation \ref{eq:sh2}) )."
+" Unfortunately, we cannot simply use the Sobolev approximation to derive Ny, similar to the column density of dust in Equation (A9)), because H» absorption is concentrated in separate absorption lines and is sensitive to the internal velocity dispersion inside molecular clouds."," Unfortunately, we cannot simply use the Sobolev approximation to derive $N_\H2$ similar to the column density of dust in Equation \ref{eq:sd}) ), because $\H2$ absorption is concentrated in separate absorption lines and is sensitive to the internal velocity dispersion inside molecular clouds."
+" These velocities are unresolved in our simulations, but can greatly reduce the self-shielding of molecular gas."," These velocities are unresolved in our simulations, but can greatly reduce the self-shielding of molecular gas."
+" Dust, on the other hand, absorbs UV radiation in continuum and is thus not affected by velocity distribution of the gas."," Dust, on the other hand, absorbs UV radiation in continuum and is thus not affected by velocity distribution of the gas."
+" Therefore, we introduce the following simple ansatz for the effective column density Ny, for Equation (ATI), where L, is the velocity coherence length of the molecular hydrogen inside molecular clouds."," Therefore, we introduce the following simple ansatz for the effective column density $N_\H2$ for Equation \ref{eq:sh2}) ), where $L_c$ is the velocity coherence length of the molecular hydrogen inside molecular clouds."
+" Since we cannot deduce this quantity from observations or other calculations, we treat it as another parameter of our model."," Since we cannot deduce this quantity from observations or other calculations, we treat it as another parameter of our model."
+" With the expressions for the shielding factors above, the only two parameters of our model are C, and L,."," With the expressions for the shielding factors above, the only two parameters of our model are $C_\rho$ and $L_c$."
+ These parameters can only be determined by comparing the simulation results to the observational data., These parameters can only be determined by comparing the simulation results to the observational data.
+" As the primary data sets used to calibrate the model, we use the measurements of atomic and molecular gas surface densities in nearby spirals from and measurements of gas fractions along the lines of sight to individual stars for atomic and molecular gas in the Milky Way and Magellanic Clouds(???)."," As the primary data sets used to calibrate the model, we use the measurements of atomic and molecular gas surface densities in nearby spirals from and measurements of gas fractions along the lines of sight to individual stars for atomic and molecular gas in the Milky Way and Magellanic Clouds."
+. We calibrate the two parameters of the model: the clumping factor and the molecular coherence length L.., We calibrate the two parameters of the model: the clumping factor $C_\rho$ and the molecular coherence length $L_c$.
+" We find, however, that there is no unique best-fit set of parameters."," We find, however, that there is no unique best-fit set of parameters."
+" Instead, any combinationC, of these two parameters that satisfy the constraint provides an acceptable fit to the observational constraints."," Instead, any combination of these two parameters that satisfy the constraint provides an acceptable fit to the observational constraints."
+" As an example, we show on the left panel of Figure[AT]] fits to the measurements (averaged over all galaxies they observed) for three combinations of the parameters L, and C,."," As an example, we show on the left panel of Figure \ref{fig:asjust} fits to the measurements (averaged over all galaxies they observed) for three combinations of the parameters $L_c$ and $C_\rho$."
+" In general, higher clumping factors result in the lower atomic contents at high surface densities, but the trend is too weak to be of any statistically significant constraining power."," In general, higher clumping factors result in the lower atomic contents at high surface densities, but the trend is too weak to be of any statistically significant constraining power."
+" As a fiducial set of parameters we choose the combination L,=0.3pc and C,=30.", As a fiducial set of parameters we choose the combination $L_c=0.3\dim{pc}$ and $C_\rho=30$.
+" This choice provides a marginally better overall fit to the observations, and is also consistent with estimates of the gas clumping factor deep inside molecular clouds(?)."," This choice provides a marginally better overall fit to the observations, and is also consistent with estimates of the gas clumping factor deep inside molecular clouds."
+". The fiducial value of C, is somewhat larger than the estimates of the clumping factor from numerical simulations of turbulent molecular clouds, C,=eo, where σἹηρ31—1.5 is the dispersion of the lognormal density distribution inside the clouds."," The fiducial value of $C_\rho$ is somewhat larger than the estimates of the clumping factor from numerical simulations of turbulent molecular clouds, $C_\rho = e^{\sigma_{\ln\rho}^2}$, where $\sigma_{\ln\rho}\approx 1 - 1.5$ is the dispersion of the lognormal density distribution inside the clouds."
+" However, the value of C,=10, which was used in and is more consistent with the numerical simulations of turbulent molecular clouds would provide an almost equally good to the existing observations, if it is used with L,~1pc."," However, the value of $C_\rho=10$, which was used in and is more consistent with the numerical simulations of turbulent molecular clouds would provide an almost equally good to the existing observations, if it is used with $L_c\approx1\dim{pc}$."
+" Any sub-cell model would be of limited value, if it was only applicable to a narrow range of numerical resolutions."," Any sub-cell model would be of limited value, if it was only applicable to a narrow range of numerical resolutions."
+" In order to test the range of spatial resolutions over which our model performs robustly, we have re-run a subset of our test simulations,"," In order to test the range of spatial resolutions over which our model performs robustly, we have re-run a subset of our test simulations,"
+We uote that dynamical [friction still operates in circtumstauces where mass accretion is [rustratecd.,We note that dynamical friction still operates in circumstances where mass accretion is frustrated.
+ For example. a wind-emitting star moving through a eas cloud experiences mass loss rather than ass gain.," For example, a wind-emitting star moving through a gas cloud experiences mass loss rather than mass gain."
+ Cloud gas inpacting the wind upstream is arrested or refracted in a bowshock. as analytically calculated by Wilkin(1996).," Cloud gas impacting the wind upstream is arrested or refracted in a bowshock, as analytically calculated by \citet{w96}."
+. Downstream. the wind forms a supersonic jet.," Downstream, the wind forms a supersonic jet."
+" As long as the upstream staudolL radius of the shock lies within r, aud the dowustreanr jet is relatively uarrow. the [ar-fielcl perturbatious are close to what we have obtained. aud equation (66)) lor F still applies."," As long as the upstream standoff radius of the shock lies within $r_s$ and the downstream jet is relatively narrow, the far-field perturbations are close to what we have obtained, and equation \ref{eqn:bondifric}) ) for $F$ still applies."
+ When the object is actually able to accept eas freely. dynamical friction arises in two pliysically distinct ways.," When the object is actually able to accept gas freely, dynamical friction arises in two physically distinct ways."
+ First. there is the gravitational tug from the wake.," First, there is the gravitational tug from the wake."
+ Secoucl. momenttun is transferre directly to the object by gas falling onto it.," Second, momentum is transferred directly to the object by gas falling onto it."
+ Our findiug that these two forces sum to AZV is at least roughly consistent with simulations.," Our finding that these two forces sum to ${\dot M}\,V$ is at least roughly consistent with simulations."
+ In a numerical study directed primarily at the mass accretion issue. Rullert(1996) explicitly determined both force coutributious ou accretors of various size inary gas.," In a numerical study directed primarily at the mass accretion issue, \citet{r96} explicitly determined both force contributions on accretors of various size in a gas."
+ For aud0.6.. the sitmulation ended before the flow reache steady-state (see his Fig.," For and, the simulation ended before the flow reached steady-state (see his Fig."
+ 2)., 2).
+ After initial transients clied out. the gravitational drag was steady unti fe13PFpg. where the Bondi-Hoyle time /pj is races.," After initial transients died out, the gravitational drag was steady until $t\approx 13\ t_{\rm BH}$, where the Bondi-Hoyle time $t_{\rm BH}$ is $r_{\rm acc}/c_{\rm s}$."
+ Therealter. this force component declined for the rest of the integration.," Thereafter, this force component declined for the rest of the integration."
+ At the end of the simulation =32 μι). the sum of the eravitationa t," At the end of the simulation $t = 32\ t_{\rm BH}$ ), the sum of the gravitational drag and momentum accretion forces was $1.2\ {\dot M}\,V$."
+he two forces quickly leveled off. with a sum equal to 1.137V.," For and the same Mach number, the two forces quickly leveled off, with a sum equal to $1.4\ {\dot M}\,V$."
+ However. this simulation rau ouly until /=LO/py. so it is not clear whether the gravitational drag would Lave later declined. as iu the first case.," However, this simulation ran only until $t = 10\ t_{\rm BH}$, so it is not clear whether the gravitational drag would have later declined, as in the first case."
+ Following historical precedeut. we lave restricted our investigation to an isothermal gas.," Following historical precedent, we have restricted our investigation to an isothermal gas."
+ For an isentropic gas with 5>l. it seems likely that the friction force will still be given by AZV. as loug as the aceretor is moving subsonically.," For an isentropic gas with $\gamma > 1$, it seems likely that the friction force will still be given by $\dot{M}\, V$, as long as the accretor is moving subsonically."
+ Verifving this equality analytically would require a perturbation study analogous to the present one., Verifying this equality analytically would require a perturbation study analogous to the present one.
+ We leave such a project for future investigators., We leave such a project for future investigators.
+ Again the current body of numerical studies is in broad accord with our expectation., Again the current body of numerical studies is in broad accord with our expectation.
+ determined the total friction force on an accretor moving through a 5=5/3 gas., \citet{r94} determined the total friction force on an accretor moving through a $\gamma=5/3$ gas.
+ For and 3=0.6. the friction force was 1.1AV at |=70fq.," For and $\beta=0.6$, the friction force was $1.1\ \dot{M}\, V$ at $t=70\ t_{\rm BH}$."
+ For and the same Mach nuuber. the flow liad not achieved steady state by /=19fpy.," For and the same Mach number, the flow had not achieved steady state by $t=19\ t_{\rm BH}$."
+ The total force was 1.8MV. atthis time. but was falling rapidly.," The total force was $1.8\ \dot{M}\, V$ atthis time, but was falling rapidly."
+ A future project of interest would be to redo these M.imulations over a range of 5- aud 5-values. running the simlations long enough until a true steady state is reached.," A future project of interest would be to redo these simulations over a range of $\beta$ - and $\gamma$ -values, running the simulations long enough until a true steady state is reached."
+ For the more general isentropic case. AZ cau uo longer be approximated by equation (52)).," For the more general isentropic case, $\dot{M}$ can no longer be approximated by equation \ref{eqn:intermdot}) )."
+ Iustead the value of M at a given V. decreases with higher 5-values. as shown aualytically by [or V=0. aud as seen in the simulations of Bullert(199[.1995.1996) [or accretors moving relative to the background gas.," Instead the value of $\dot{M}$ at a given $V$ decreases with higher $\gamma$ -values, as shown analytically by \citet{b52} for $V=0$, and as seen in the simulations of \citet{r94,r95,r96} for accretors moving relative to the background gas."
+ Isentropie [lows are less compressible than isothermal ones. so the wake will be less cdeuse.," Isentropic flows are less compressible than isothermal ones, so the wake will be less dense."
+ As a result. the friction force will also be lower. presumably by tle same amount as the accretion rate AL.," As a result, the friction force will also be lower, presumably by the same amount as the accretion rate $\dot{M}$ ."
+ Iu the present investigation. we have been unable to tease apart analytically the two force," In the present investigation, we have been unable to tease apart analytically the two force"
+he mereecdl product is a MS star with the stellar material completely mixed aud (b) uo mass is lost rou the system during the mereer process.,the merged product is a MS star with the stellar material completely mixed and (b) no mass is lost from the system during the merger process.
+ The no mass-loss assumptiou is based on the results ol SPH stimulations of MS-MS mereers(e.g... Sillsetal.2001)) but such calculations vield ouly a united amount of mixiug.," The no mass-loss assumption is based on the results of SPH simulations of MS-MS mergers, \citealt{sil2001}) ) but such calculations yield only a limited amount of mixing."
+ The rejuvenated age of the mereecd MS star is determiued clepeuding ou the amount of unburut hydrogen fuel gained by the hydrogen-burniug core as a result of the uixiug., The rejuvenated age of the merged MS star is determined depending on the amount of unburnt hydrogen fuel gained by the hydrogen-burning core as a result of the mixing.
+ Iun. case of a mass trausler across a MS-MS binary. clistinetion is mace between the cases when the original accretor MS star las a radiative or a couvective core.," In case of a mass transfer across a MS-MS binary, distinction is made between the cases when the original accretor MS star has a radiative or a convective core."
+ For a convective core. the core grows with the eain of mass and mixes with the uuburut bydrogen fuel so that the accreting MS star appears vounger.," For a convective core, the core grows with the gain of mass and mixes with the unburnt hydrogen fuel so that the accreting MS star appears younger."
+ For the case of a radiative core. the [Traction of the hydrogen burut in the hydrogeun-buruing core remains nearly unallectecl by the gain of mass so that the ellective age ol the MS star decreases.," For the case of a radiative core, the fraction of the hydrogen burnt in the hydrogen-burning core remains nearly unaffected by the gain of mass so that the effective age of the MS star decreases."
+ The ellective age is cleterminecl so as to keep the elapsed fraction of its MS lifetime unchauged., The effective age is determined so as to keep the elapsed fraction of its MS lifetime unchanged.
+ The initial cluster mass of A4(0)zzLO?AL. in our models correspouds to :N(0)=170667 stars., The initial cluster mass of $M_{cl}(0) \approx 10^5\Ms$ in our models corresponds to $N(0) = 170667$ stars.
+ To our knowledge. direct N-body computations with such a large μιαος of stars. where the clusters are fully inass-segregated aud all the massive stars are in binaries. are beiug reported [or the first time.," To our knowledge, direct N-body computations with such a large number of stars, where the clusters are fully mass-segregated and all the massive stars are in binaries, are being reported for the first time."
+ We evolve f initial models with the above NV(Q). generated using different raucom uumber seeds. until zz3 Myr.," We evolve 4 initial models with the above $N(0)$, generated using different random number seeds, until $\approx 3$ Myr."
+ We take this age as au upper limit of the age of R136 (Crowther 2010)..," We take this age as an upper limit of the age of R136 \citep{crw2010,pz2010}."
+" We do all the computations on “NVIDIA [80 GTX"" GPU platforms.", We do all the computations on “NVIDIA 480 GTX” GPU platforms.
+ From the above computations. we trace the bodies that are ejected from the clusters curing their evolution (within z3 Myr).," From the above computations, we trace the bodies that are ejected from the clusters during their evolution (within $\approx 3$ Myr)."
+ We consider a single-star/binary/imultiplet to be a runaway member [rom its host cluster if it is found moving away from the cluster beyoud R>10 pe distance from the cluster's center of density., We consider a single-star/binary/multiplet to be a runaway member from its host cluster if it is found moving away from the cluster beyond $R>10$ pc distance from the cluster's center of density.
+ Although our primary focus is on the runaway VMS5s. we consider the whole mass spectrum of ejected stars as well.," Although our primary focus is on the runaway VMSs, we consider the whole mass spectrum of ejected stars as well."
+ Fig., Fig.
+ 2 shows the projected suapshots of the ruuaways with inasses AL>38M... combined from the 1 computatious. at /—1 Myr and 3 Myr evolutionary times.," \ref{fig:ejsnap} shows the projected snapshots of the runaways with masses $M>3\Ms$, combined from the 4 computations, at $t=1$ Myr and 3 Myr evolutionary times."
+ It cau be seen that by /23 Myr there are a siguificant. uumber of fast runaway VMSs with total (3-dimeusional) velocities upto se300 kins  anc a lew fast VMSs are already present at /—1 Myr., It can be seen that by $t=3$ Myr there are a significant number of fast runaway VMSs with total (3-dimensional) velocities upto $\approx 300$ km $^{-1}$ and a few fast VMSs are already present at $t=1$ Myr.
+ All these runaways are on their MSs aud heuce are OB stars., All these runaways are on their MSs and hence are OB stars.
+ We note that the vast majority of the massive ejected members are sinele stars — only 2 of the massive ejecta [rom our computations are found in hard binaries., We note that the vast majority of the massive ejected members are single stars — only 2 of the massive ejecta from our computations are found in hard binaries.
+ We shall discuss the multiplicity properties of the ejected stellar population iu detail in a future paper., We shall discuss the multiplicity properties of the ejected stellar population in detail in a future paper.
+" It. is worthwhile to note that although our adapted canonical IME has a 19041, upper limit. our models vield single-star runaways with masses upto zz220M... within /«3 Myr. cousicerably"," It is worthwhile to note that although our adapted canonical IMF has a $150\Ms$ upper limit, our models yield single-star runaways with masses upto $\approx 250\Ms$ within $t<3$ Myr, considerably"
+"thin disc by 30—70 km s! (Gilmore, Wyse Norris 2002), with Viag=50 km s! usually taken as a canonical value (see also Vallenari et al.","thin disc by $30-70$ km $^{-1}$ (Gilmore, Wyse Norris 2002), with $V_{\rm
+ lag}=50$ km $^{-1}$ usually taken as a canonical value (see also Vallenari et al."
+ 2006)., 2006).
+" On the other hand, in external galaxies the difference in rotation velocity of thin and thick discs appears to be diverse and strongly dependent on galaxy mass, with lower mass galaxies having larger differences between thin and thick velocities (Yoachim Dalcanton 2008)."," On the other hand, in external galaxies the difference in rotation velocity of thin and thick discs appears to be diverse and strongly dependent on galaxy mass, with lower mass galaxies having larger differences between thin and thick velocities (Yoachim Dalcanton 2008)."
+" In terms of velocity dispersions, our results are in relatively good agreement with the observations of the Milky Way."," In terms of velocity dispersions, our results are in relatively good agreement with the observations of the Milky Way."
+ Vallenari et al. (, Vallenari et al. (
+2006) estimated the velocity dispersion ellipsoid of the thin and thick discs of our Galaxy.,2006) estimated the velocity dispersion ellipsoid of the thin and thick discs of our Galaxy.
+" Their results for the thick disc are: (σ-,σφ,σε)~(74+11,507,38X7T) km s! at the solar radius."," Their results for the thick disc are: $(\sigma_r, \sigma_\phi, \sigma_z) \sim (74\pm 11, 50\pm 7, 38\pm 7)$ km $^{-1}$ at the solar radius."
+" For the thin disc, they estimate the velocity ellipsoid dividing stars into four stellar age bins, finding (07,04,02)~(25—34,2032,1018)kms! (with errors < 15%)."," For the thin disc, they estimate the velocity ellipsoid dividing stars into four stellar age bins, finding $(\sigma_r, \sigma_\phi, \sigma_z) \sim (25-34, 20-32, 10-18)\,{\rm
+ km\, s}^{-1}$ (with errors $\lesssim 15\%$ )."
+" As described above, the eight simulated galaxies are diverse and show a wide range in velocity dispersions that agree relatively well with these results (although we note that simulated galaxies are not expected to resemble the Milky Way in detail)."," As described above, the eight simulated galaxies are diverse and show a wide range in velocity dispersions that agree relatively well with these results (although we note that simulated galaxies are not expected to resemble the Milky Way in detail)."
+" Finally, we note that the velocity structure of the simulated stellar components is complex, in general having important asymmetries."," Finally, we note that the velocity structure of the simulated stellar components is complex, in general having important asymmetries."
+ Fig., Fig.
+ 9 shows 2D face-on maps of tangential velocity for theeight simulations (including both disc and spheroid stars) within the inner 30 kpc., \ref{vtita_maps_stars} shows 2D face-on maps of tangential velocity for theeight simulations (including both disc and spheroid stars) within the inner $30$ kpc.
+" From these plots we can read off the velocity structure of simulated galaxies, the sizes of bulges and discs, and we can also observe bar patterns, particularly in Aq-C-5, Aq-E-5 and Aq-G-5 (see also the next section)."," From these plots we can read off the velocity structure of simulated galaxies, the sizes of bulges and discs, and we can also observe bar patterns, particularly in Aq-C-5, Aq-E-5 and Aq-G-5 (see also the next section)."
+" Aq-A-5, Aq-C-5, Aq-D-5 and Aq-E-5 have the largest tangential velocities, as expected since these are the most massive galaxies and therefore have higher circular velocities."," Aq-A-5, Aq-C-5, Aq-D-5 and Aq-E-5 have the largest tangential velocities, as expected since these are the most massive galaxies and therefore have higher circular velocities."
+" In the case of the disc starts to dominate at a relatively large radius, showing a ring-like structure."," In the case of the disc starts to dominate at a relatively large radius, showing a ring-like structure."
+" As we discuss below, the absence of a disc in Aq-F-5 is evident from this figure, and also the small disc component of Aq-H-5."," As we discuss below, the absence of a disc in Aq-F-5 is evident from this figure, and also the small disc component of Aq-H-5."
+ We find that the dynamical structure of discs found for our simulations is also present in the lower resolution runs., We find that the dynamical structure of discs found for our simulations is also present in the lower resolution runs.
+" In all cases, the older stars define thicker discs compared to the younger populations."," In all cases, the older stars define thicker discs compared to the younger populations."
+" In particular, Aq-E-5 and Aq-E-6b show very good agreement, as shown in Fig. 10.."," In particular, Aq-E-5 and Aq-E-6b show very good agreement, as shown in Fig. \ref{disk_dynamics_resolution}."
+" In this case, differences in tangential velocities and velocity dispersions are always lower than 5%."," In this case, differences in tangential velocities and velocity dispersions are always lower than $5\%$."
+" For the youngest stars (left-hand panel), tangential velocities are lower for Aq-E-6b than for Aq-E-5."," For the youngest stars (left-hand panel), tangential velocities are lower for Aq-E-6b than for Aq-E-5."
+" Velocity dispersion are, regardless of stellar age, larger for the lower resolution run."," Velocity dispersion are, regardless of stellar age, larger for the lower resolution run."
+ We detect more significant differences for Aq-C-6 and AqE-6 with respect to Aq-C-5 and Aq-E-5 respectively., We detect more significant differences for Aq-C-6 and Aq-E-6 with respect to Aq-C-5 and Aq-E-5 respectively.
+" In these cases, the tangential velocities are typically 15% lower in the low resolution runs, while velocity dispersions are ~50% larger."," In these cases, the tangential velocities are typically $15\%$ lower in the low resolution runs, while velocity dispersions are $\sim 50\%$ larger."
+ These results show that low resolution runs can artificially boost the degree of disc heating., These results show that low resolution runs can artificially boost the degree of disc heating.
+" In this section, we discuss the structure of the inner and outer spheroids of our simulated galaxies."," In this section, we discuss the structure of the inner and outer spheroids of our simulated galaxies."
+ In Figs., In Figs.
+" and 12,, we show maps of surface mass density for inner spheroids (up to 0.5X ropt) and outer spheroids (up to"," \ref{maps_innerspheroid} and \ref{maps_outerspheroid}, , we show maps of surface mass density for inner spheroids (up to $0.5\times
+r_{\rm opt}$ ) and outer spheroids (up to"
+inter-order gaps.,inter-order gaps.
+ Figure 11 shows a possible detection of one line., Figure \ref{f_n_8216_synthesis} shows a possible detection of one line.
+ The N abundance is log e(N)<7.1 but this might properly be considered an upper limit., The N abundance is $\log\epsilon$ $\leq 7.1$ but this might properly be considered an upper limit.
+" 0.2 cm The gf-values for the permitted and forbidden lines are taken from Wiese, Fühhr Deter (1996)."," 0.2 cm The $gf$ -values for the permitted and forbidden lines are taken from Wiese, Fühhr Deter (1996)."
+" On the McDonald spectrum, the forbidden oxygen lines at 5577A,, 6300 aand 6363 aare detected."," On the McDonald spectrum, the forbidden oxygen lines at 5577, 6300 and 6363 are detected."
+ Weak permitted lines of RMT 10 near aare also analyzed., Weak permitted lines of RMT 10 near are also analyzed.
+ Figure 12 shows the best-fitting synthetic spectrum for the O16156 rregion., Figure \ref{f_o_synthesis} shows the best-fitting synthetic spectrum for the 6156 region.
+ Forbidden and permitted lines give a very similar abundance., Forbidden and permitted lines give a very similar abundance.
+" The strong triplet at 7774 aand the 8446 feature give an abundance about 1.5 dex higher abundance, a difference attributed to non-LTE effects."," The strong triplet at 7774 and the 8446 feature give an abundance about 1.5 dex higher abundance, a difference attributed to non-LTE effects."
+ 0.2 cm The g f-values are taken from the database., 0.2 cm The $gf$ -values are taken from the database.
+ Five lines were suitable for abundance analysis from the McDonald spectrum (Table 3)., Five lines were suitable for abundance analysis from the McDonald spectrum (Table 3).
+ RMT and 6 give somewhat different results but we assign all lines the 4same weight., RMT 4 and 6 give somewhat different results but we assign all lines the same weight.
+deviation of the amisotropic amplitude posteriors should vary with multipoles ἐν as done by Grocuchoom&Evilk-sen (2009).,"deviation of the anisotropic amplitude posteriors should vary with multipoles $\ell$, as done by \cite{groeneboom:2008b}."
+ Before performing a full-scale analysis of sinulated polarized Planck data. we wish to validate our code.," Before performing a full-scale analysis of simulated polarized Planck data, we wish to validate our code."
+ We therefore simulate a low-resolution ως—32 map with E- mode data included., We therefore simulate a low-resolution $N_{\textrm{side}}=32$ map with E- mode data included.
+ Assuming an anisotropic amplitude of gy.=1.0. we perform both a brute-force and a ietropolis-hastings analysis of a full-xky map with uo beam nor noise.," Assuming an anisotropic amplitude of $g_*=1.0$, we perform both a brute-force and a metropolis-hastings analysis of a full-sky map with no beam nor noise."
+ The resulting posteriors for the TT-case and the TT|TEEE-case are shown in Figure. 1.., The resulting posteriors for the TT-case and the TT+TE+EE-case are shown in Figure \ref{fig:simulated_posteriors}.
+ It is worth to note that the posterior is more narrow when inchiding polarization data. as there is more data available.," It is worth to note that the posterior is more narrow when including polarization data, as there is more data available."
+ A typical posterior of the estimated direction n together with the input TT|EE ACW-sigual is secu in Fieure 5.., A typical posterior of the estimated direction $n$ together with the input TT+EE ACW-signal is seen in Figure \ref{fig:res_asymmetric}.
+ We now consider a Plauck simmlation., We now consider a Planck simulation.
+ We first simulate a temperature-only ACW-anisotropic map with, We first simulate a temperature-only ACW-anisotropic map with
+2009).,.
+". In this paper the vectors v. n,. and ng are delined in the spacecraft coordinate svstem."," In this paper the vectors $\bf v$ , $\bf n_{_A}$ , and $\bf n_{_B}$ are defined in the spacecraft coordinate system."
+" The errors of line-oEsight (LOS) vector. An, and n. will produce à pseudo-dipole difference signal The dipoles for each observations have to be removed from the raw data before because their intensities are roughly 10 to 20 times greater than those of the CMD anisolropies."," The errors of line-of-sight (LOS) vector, $\Delta \bf{n}_{_A}$ and $\Delta \bf{n}_{_B}$, will produce a pseudo-dipole difference signal The dipoles for each observations have to be removed from the raw data before map-making because their intensities are roughly 10 to 20 times greater than those of the CMB anisotropies."
+ A small error of antenna direction will produce an error in. predicted dipole intensity and (hen cause a pseudo-dipole signal in the resulting CMD map more noticeably through the Doppler dipole subtraction ., A small error of antenna direction will produce an error in predicted dipole intensity and then cause a pseudo-dipole signal in the resulting CMB map more noticeably through the Doppler dipole subtraction .
+ For example. a LOS error of ~7. just about a hall-pixel in the WMAP resolution. can consequently catse the dipole signal to be deviated by dx. which can not be ignored compared to the very weak CAIB signal.," For example, a LOS error of $\sim7'$, just about a half-pixel in the WMAP resolution, can consequently cause the dipole signal to be deviated by $\mu$ K, which can not be ignored compared to the very weak CMB signal."
+ An asvnehronous betweenthe attitude and differential data can also produce the signal., An asynchronous betweenthe attitude and differential data can also produce the pseudo-dipole signal.
+" The WALAP mission uses two separate clocks for the attitude data and science data respectively,", The WMAP mission uses two separate clocks for the attitude data and science data respectively.
+ Therelore. if (here is a small constant timing error. (here will be a constant direction difference between the “observed pixel” and the true pixel.," Therefore, if there is a small constant timing error, there will be a constant direction difference between the ""observed pixel"" and the true pixel."
+ This has the same ellect as a constant LOS error in spacecraft coordinates (Liu.Niong&Li2010)., This has the same effect as a constant LOS error in spacecraft coordinates \citep{liu10}.
+. Another possible source of pseudo-dipole signal in released WALAP maps is the sidelobe signal contamination., Another possible source of pseudo-dipole signal in released WMAP maps is the sidelobe signal contamination.
+ Like all radio telescopes. the WAIAP antennas have both main beam response and sidelobe response.," Like all radio telescopes, the WMAP antennas have both main beam response and sidelobe response."
+ The WAIAP antenna sidelobe response was described by Barnesοἱal.(2003) and the corresponding data file is publiclyavailable?., The WMAP antenna sidelobe response was described by \citet{barnes03} and the corresponding data file is publicly.
+. The data files are fits Format fall skv maps in spacecraft coordinates in which the sidelobe responses are given in normalized gain C. where the normalization rule is that the summation of all gains [or one antenna (including the main beam) equals to AN. the nmunber of pixels in themap.," The data files are fits format full sky maps in spacecraft coordinates in which the sidelobe responses are given in normalized gain $G$, where the normalization rule is that the summation of all gains for one antenna (including the main beam) equals to $N$, the number of pixels in themap."
+Thus for. eachdifferential⋅⊳⋅ observation.⋅ the recorded difference⋅⊳ signal⋅ is⋅ 357⋏∖↓↽↿4(6;—GP)T;/pua- N.,"Thus for eachdifferential observation, the recorded difference signal is $\sum_{i=0}^{N-1} (G_i^A-G_i^B)T_i/N$ ."
+ Let, Let
+"Here we use as additional constraints Li observations that we performed for a large sample of field red giant stars (subgiant, RGB, and early-AGB stars) with metallicities around solar.","Here we use as additional constraints Li observations that we performed for a large sample of field red giant stars (subgiant, RGB, and early-AGB stars) with metallicities around solar."
+ All sample stars have Hipparcos parallaxes so that their mass and evolutionary status could be relatively well determined (Charbonnel et al., All sample stars have Hipparcos parallaxes so that their mass and evolutionary status could be relatively well determined (Charbonnel et al.
+ in preparation)., in preparation).
+" In Fig.14 they are distinguished with respect to their mass (less or more massive than 2 M, in the left and right panels respectively).", In \ref{fig:Lagardeetal_posterIAU268_figb} they are distinguished with respect to their mass (less or more massive than 2 $_{\odot}$ in the left and right panels respectively).
+" Let us consider first the stars with initial masses lower than 2 Mc, whose Li properties are compared with predictions for the 1.5 and 2 Me models (left panel of Fig.14))."," Let us consider first the stars with initial masses lower than 2 $_{\odot}$, whose Li properties are compared with predictions for the 1.5 and 2 $_{\odot}$ models (left panel of \ref{fig:Lagardeetal_posterIAU268_figb}) )."
+ The theoretical Li behaviour is relatively straightforward., The theoretical Li behaviour is relatively straightforward.
+" On the main sequence and on the early-RGB, rotation-induced mixing leads to stronger Li depletion than in the standard case (compare e.g. the red curve with the black one); in this mass range indeed standard models predict no Li depletion on the main sequence and a N(Li) of the order of 1.5 at the end of the first dredge-up, which is at odds with the data."," On the main sequence and on the early-RGB, rotation-induced mixing leads to stronger Li depletion than in the standard case (compare e.g. the red curve with the black one); in this mass range indeed standard models predict no Li depletion on the main sequence and a N(Li) of the order of 1.5 at the end of the first dredge-up, which is at odds with the data."
+" After the end of the first dredge-up (Teff ~ 4800 K), the theoretical Li abundance remains temporarily constant as the convective envelope withdraws in mass."," After the end of the first dredge-up (Teff $\sim$ 4800 K), the theoretical Li abundance remains temporarily constant as the convective envelope withdraws in mass."
+" When thermohaline mixing becomes efficient (Teff ~ 4200 K), the theoretical Li abundance drops again in drastic manner (while it would stay constant in the standard case)."," When thermohaline mixing becomes efficient (Teff $\sim$ 4200 K), the theoretical Li abundance drops again in drastic manner (while it would stay constant in the standard case)."
+" After the star has reached the RGB tip its effective temperature increases (up to ~ 4800 K) as it settles on the clump, before decreasing again when the star starts climbing the early-AGB."," After the star has reached the RGB tip its effective temperature increases (up to $\sim$ 4800 K) as it settles on the clump, before decreasing again when the star starts climbing the early-AGB."
+ The second dredge-up that occurs then leads to a final decrease of N(Li)., The second dredge-up that occurs then leads to a final decrease of N(Li).
+" On this graph we do not plot the Li increase that is predicted to occur during the TP-AGB phase at a Teff of ~ 3200 K due to thermohaline mixing, and which is discussed in 4.1.2."," On this graph we do not plot the Li increase that is predicted to occur during the TP-AGB phase at a Teff of $\sim$ 3200 K due to thermohaline mixing, and which is discussed in 4.1.2."
+" As can be seen in Fig.14,, the present predictions are in perfect agreement with the data all along the evolutionary sequence and explain very well the upper limits observed for the brightest sample giant stars."," As can be seen in \ref{fig:Lagardeetal_posterIAU268_figb}, the present predictions are in perfect agreement with the data all along the evolutionary sequence and explain very well the upper limits observed for the brightest sample giant stars."
+ The observed Li dispersion at a given effective temperature reflects dispersion in the initial rotation velocity and in the initial stellar mass???)., The observed Li dispersion at a given effective temperature reflects dispersion in the initial rotation velocity and in the initial stellar mass.
+". The case of the more massive stars, whose Li observational behaviour is compared to predictions for the 2.5 and 2.7 Me models (right panel of Fig.14)) is even more simple."," The case of the more massive stars, whose Li observational behaviour is compared to predictions for the 2.5 and 2.7 $_{\odot}$ models (right panel of \ref{fig:Lagardeetal_posterIAU268_figb}) ) is even more simple."
+" In these objects indeed no thermohaline mixing occurs on the too short RGB, and rotation-induced mixing alone explains very well the data."," In these objects indeed no thermohaline mixing occurs on the too short RGB, and rotation-induced mixing alone explains very well the data."
+" In all the models that we have computed along the TP-AGB, non negligible fresh lithium production is obtained, although"," In all the models that we have computed along the TP-AGB, non negligible fresh lithium production is obtained, although"
+"the observed bispectrum in Eq. (17)), CtCes,","the observed bispectrum in Eq. \ref{eq:Dfnlanalyt}) ), $B^\mathrm{loc}_{\ell_1\ell_2\ell_3} /
+C_{\ell_1}C_{\ell_2}C_{\ell_3}$,"
+" which rapidly increases with multipole Bi°%,0,/Ce,as the product of spectra decreases more quickly than the bispectrum.", which rapidly increases with multipole as the product of spectra decreases more quickly than the bispectrum.
+ This leads to a 1/C dependence in squeezed configurations and to a 1/C? dependence in equilateral configurations., This leads to a $1/C_\ell$ dependence in squeezed configurations and to a $1/C_\ell^2$ dependence in equilateral configurations.
+" When the observed bispectrum is associated to CMB signal alone, its decrease cancels the increase of the weights so that the sum in Eq. (17))"," When the observed bispectrum is associated to CMB signal alone, its decrease cancels the increase of the weights so that the sum in Eq. \ref{eq:Dfnlanalyt}) )"
+ converges., converges.
+" Conversely, the sum diverges when the observed bispectrum is associated with a non-CMB signal and does not decrease with 6 as fast as the CMB."," Conversely, the sum diverges when the observed bispectrum is associated with a non-CMB signal and does not decrease with $\ell$ as fast as the CMB."
+ The bias Aff is maximal at 30 GHz and rapidly decreases with frequency., The bias $\Delta f_\mathrm{NL}^\mathrm{RAD}$ is maximal at 30 GHz and rapidly decreases with frequency.
+ It slightly increases again at the two highest frequencies following the amplitude of the bispectrum in temperature units which is plotted in the upper panel of Fig. 4.., It slightly increases again at the two highest frequencies following the amplitude of the bispectrum in temperature units which is plotted in the upper panel of Fig. \ref{fig:amplRAD}.
+ The relative error of ΔΙΑ for Ímax=700 is of the order of independently of the frequency.," The relative error of $\Delta
+f_\mathrm{NL}^\mathrm{RAD}$ for $\ell_{\mathrm{max}}=700$ is of the order of independently of the frequency."
+ It amounts to for €max=2048., It amounts to for $\ell_{\mathrm{max}}=2048$.
+ These errors bars were computed with simulations using the catalog of sources present in Sehgal et al., These errors bars were computed with simulations using the catalog of sources present in Sehgal et al.
+"’s maps As shown in Table 1,, masking sources above the ERCSC flux limit proves very efficient to significantly decrease the radio contamination to /wr at all the frequencies.","'s maps As shown in Table \ref{fnlradtable}, masking sources above the ERCSC flux limit proves very efficient to significantly decrease the radio contamination to $f_\mathrm{NL}$ at all the frequencies."
+" At a Planck-like resolution, £44.=2048, the bias AfR#? is reduced below unity above 150 GHz."," At a Planck-like resolution, $\ell_\mathrm{max}= 2048$, the bias $\Delta
+f_\mathrm{NL}^\mathrm{RAD}$ is reduced below unity above 150 GHz."
+ It is of the order of Planck's expected error bars at 90 GHz., It is of the order of Planck's expected error bars at 90 GHz.
+ At 30 GHz the bias is still important., At 30 GHz the bias is still important.
+" The bias due to IR sources Aft, is always negative, see Table 2.."," The bias due to IR sources $\Delta f_\mathrm{NL}^\mathrm{IR}$ is always negative, see Table \ref{fnlirtable}."
+" As a matter of fact, we have shown that the IR bispectrum peaks in squeezed configurations just like the CMB bispectrum and these configurations thus dominate the sum in Eq. (17))."," As a matter of fact, we have shown that the IR bispectrum peaks in squeezed configurations just like the CMB bispectrum and these configurations thus dominate the sum in Eq. \ref{eq:Dfnlanalyt}) )."
+" Moreover, in the squeezed limit the CMB bispectrum is negative while the IR bispectrum is positive."," Moreover, in the squeezed limit the CMB bispectrum is negative while the IR bispectrum is positive."
+" For the same reason as for radio sources, the bias AfMi blows up at high multipoles."," For the same reason as for radio sources, the bias $\Delta
+f_\mathrm{NL}^\mathrm{IR}$ blows up at high multipoles."
+" This is particularly important at a Planck-like resolution, max=2048, where primordial NG tests will need to carefully handle the contamination by IR sources."," This is particularly important at a Planck-like resolution, $\ell_\mathrm{max}=2048$, where primordial NG tests will need to carefully handle the contamination by IR sources."
+ The IR sources emission plummets at radio frequencies so that Afl is completely negligible below 220 GHz.," The IR sources emission plummets at radio frequencies so that $\Delta
+f_\mathrm{NL}^\mathrm{IR}$ is completely negligible below 220 GHz."
+ It becomes of the order of Planck's error bars at 277 GHz and it reaches WMAP’s central values for fri at 350 GHz., It becomes of the order of Planck's error bars at 277 GHz and it reaches WMAP's central values for $f_\mathrm{NL}$ at 350 GHz.
+ The relative error of Aft ranges between 6 and from 148 to 350 GHz for €max=700., The relative error of $\Delta f_\mathrm{NL}^\mathrm{IR}$ ranges between 6 and from 148 to 350 GHz for $\ell_{\mathrm{max}}=700$.
+ It ranges between 3 and for €max=2048. (, It ranges between 3 and for $\ell_{\mathrm{max}}=2048$. (
+These error bars were computed analytically with the weak NG approximation — see Appendix ??)) At higher frequencies the IR contamination to the bispectrum is likely larger but the contamination from our Galaxy needs to be taken into account as well.,These error bars were computed analytically with the weak NG approximation – see Appendix \ref{appendix:wngvar}) ) At higher frequencies the IR contamination to the bispectrum is likely larger but the contamination from our Galaxy needs to be taken into account as well.
+Using observations of the and transitions ofII2CO.. we have successfully. coustrained the spatial deusities of a sample of ealactic star-forming reeious.,"Using observations of the and transitions of, we have successfully constrained the spatial densities of a sample of galactic star-forming regions."
+ Both trausitious were observed toward 15 sources with relative ease. requiring an average of 17 min of iuteeration time aud resulting in ouly 2 nonudoetectious in the J=1 transition.," Both transitions were observed toward 18 sources with relative ease, requiring an average of 17 min of integration time and resulting in only 3 nondetections in the $J=4$ transition."
+ Accurate of the spatial deusitv η) were mnade“---Dicamus for 13 objects and useful lanits were placed on the une a colmbination of Large Velocity Gradieut (LVG) aud Loca Thermodvuamic Equilibrimm (LTE) analyses., Accurate measurements of the spatial density )] were made for 13 objects and useful limits were placed on the remainder using a combination of Large Velocity Gradient (LVG) and Local Thermodynamic Equilibrium (LTE) analyses.
+ Molecular ivdrogeu densities iu the range of ane ortho-formaldehyde colin densities per unit Lue width )etween and are foun or most sources. m general agreement with previous neasurements.," Molecular hydrogen densities in the range of and ortho-formaldehyde column densities per unit line width between and are found for most sources, in general agreement with previous measurements."
+ Detailed analyses of the advantages and limitations o this «ο... technique have also been provided., Detailed analyses of the advantages and limitations to this densitometry technique have also been provided.
+ deusitometry proves to be best suited to objects with z 100 EK. above which the LVG models come relatively independent of kineticο teniperaturoe.," densitometry proves to be best suited to objects with $\gtrsim$ 100 K, above which the LVG models become relatively independent of kinetic temperature."
+ Compared with the similarly utilized aud ransitious. the 7—3 aud | A--doublets provide higher spatial resolution aud scusitivity to hot. deuse material. which makes them more efficieut probes of spatial density iu wnolecular cores.," Compared with the similarly utilized and transitions, the $J=3$ and 4 -doublets provide higher spatial resolution and sensitivity to hot, dense material, which makes them more efficient probes of spatial density in molecular cores."
+ However. beam widths comparable o the auticipated source sizes make source structure considerations iniportant.," However, beam widths comparable to the anticipated source sizes make source structure considerations important."
+ Since nappiig mneasurenments inve vet to be conducted for these trausifious. the correlation between the spatial extent traced bv the J—35 and | K-doublets aud that of other deuse gas racers Is uncertain. naking spatial assuniptious based on past nieasuremoeuts problematic.," Since mapping measurements have yet to be conducted for these transitions, the correlation between the spatial extent traced by the $J=3$ and 4 -doublets and that of other dense gas tracers is uncertain, making spatial assumptions based on past measurements problematic."
+ This work serves as asuccessful proof-of-concept for he J=3/J| K-doublet densitometry echuique. adding a useful new diagnostic to the study of deuse molecular cuvironiments.," This work serves as asuccessful proof-of-concept for the $J=3/J=4$ -doublet densitometry technique, adding a useful new diagnostic to the study of dense molecular environments."
+ The Green Bank Telescope (CBT) staff was characteristically helpful and coutributed significantly o the success of our observing program., 	 The Green Bank Telescope (GBT) staff was characteristically helpful and contributed significantly to the success of our observing program.
+ The authors also thank the referee for several conuuenuts that ereatlv chhanced this work., The authors also thank the referee for several comments that greatly enhanced this work.
+ P. L M. would like to hank Tarold Dutuer aud the NRAOQ REU students of 2009 for their continued support., P. I. M. would like to thank Harold Butner and the NRAO REU students of 2009 for their continued support.
+ Funding for his project was provided bv the NSF through the NRAO REU Program (award No., Funding for this project was provided by the NSF through the NRAO REU Program (award No.
+ 0755390) aud vo NRÀO though an undereraduate internship. Fucilities::, 0755390) and by NRAO though an undergraduate internship. :
+ , 
+corona (e.g. see 9105. or refer to (he density profile in Figure 1 to be discussed in 833).,"corona (e.g. see SI05, or refer to the density profile in Figure 1 to be discussed in 3)."
+" Roughly speaking. for ο=3xI0! km/s. lupSUH, when n>10°em ."," Roughly speaking, for $v_b=3\times 10^4$ km/s, $l_{\rm mfp} \lesssim H_\rho$ when $n>10^9$ $^{-3}$."
+ Therefore. we anticipate the inwarcd-propagating electron beam {ο heat up ambient media from the lower corona to the upper chromosphere.," Therefore, we anticipate the inward-propagating electron beam to heat up ambient media from the lower corona to the upper chromosphere."
+" This allows us to assume that the incoming electron beam starts dissipating lrom μας=1.148. (lower corona) to ry,=1.0012. (upper chromosphere).", This allows us to assume that the incoming electron beam starts dissipating from $r_{\rm max} =1.1R_{\odot}$ (lower corona) to $r_{\rm min}=1.001R_{\odot}$ (upper chromosphere).
+ We model the energy flux of the beam to decrease inwardly according to where fis a parameter (hat describes (he spacial distribution of the heating., We model the energy flux of the beam to decrease inwardly according to where $k$ is a parameter that describes the spacial distribution of the heating.
+" =1 corresponds to constant volumetric heating: namely, the heating rate per unit mass is higher in (he upper region (i.e. corona) than in the lower region (i.e. chromosphere)."," $k=1$ corresponds to constant volumetric heating; namely, the heating rate per unit mass is higher in the upper region (i.e. corona) than in the lower region (i.e. chromosphere)."
+ II 7=0.1. the heating rate per unit mass is more unilormlwv distributed.," If $k= 0.1$, the heating rate per unit mass is more uniformly distributed."
+ In (his work. we adopt =0.1 to resemble (he situation of a constant beam-heatine rate per unit mass.," In this work, we adopt $k=0.1$ to resemble the situation of a constant beam-heating rate per unit mass."
+" We also assume (hat the momentum flux of the beam J, clissipates in (he same manner as that described bv eq. (5))", We also assume that the momentum flux of the beam $P_b$ dissipates in the same manner as that described by eq. \ref{eq:F_b_diss}) )
+ for the energy flux., for the energy flux.
+" This implies that although the beam velocity ορ does not decay in our dissipation model. the beam density p, declines. meaning that more and more beam electrons have been transformed into thermal electrons as the beam progresses downwids in (he dissipation region."," This implies that although the beam velocity $v_b$ does not decay in our dissipation model, the beam density $\rho_b$ declines, meaning that more and more beam electrons have been transformed into thermal electrons as the beam progresses downwards in the dissipation region."
+ The beam heating al the footpoint of one open [fiekl line occurs when the planet's nagnetosphere is crossing the field line., The beam heating at the footpoint of one open field line occurs when the planet's magnetosphere is crossing the field line.
+ ence. by means of eq. (1)).," Hence, by means of eq. \ref{eq:R_mp}) ),"
+ the beam heating proceeds on the timescale The stellar atmosphere model is based on the 1-D magnetohydrocdvuamic simulation with radiative cooling and thermal conduction in an open flux tube (5105)., the beam heating proceeds on the timescale The stellar atmosphere model is based on the 1-D magnetohydrodynamic simulation with radiative cooling and thermal conduction in an open flux tube (SI05).
+ In the original simulation. the heating is given bv the nonlinear dissipation of the Alfvénn waves excited by (he eranulations at the photosphere.," In the original simulation, the heating is given by the nonlinear dissipation of the Alfvénn waves excited by the granulations at the photosphere."
+" In (he case of D,=1 G. we set arms average amplitude «dome1.8 km/s al the photosphere. which is estimated from (he scaling with the surface convective flux (Suzuki 2007) [rom the Sun (S105)."," In the case of $B_*=1$ G, we set a rms average amplitude $<dv_{\perp}> \sim 1.8$ km/s at the photosphere, which is estimated from the scaling with the surface convective flux (Suzuki 2007) from the Sun (SI05)."
+" In the case of the stronger field D,=5 G. (the larger rms velocity fluctuation <de2273.6 km/s is used for the experiment. which"," In the case of the stronger field $B_*=5$ G, the larger rms velocity fluctuation $<dv_{\perp}>
+\sim 3.6$ km/s is used for the experiment, which"
+Ultra-high energv cosmic ravs (ULLIECIS) are energetic particles (2107 eV) that must originate in the most powerlul particle accelerators in. the Universe.,Ultra-high energy cosmic rays (UHECRs) are energetic particles $> 10^{19}$ eV) that must originate in the most powerful particle accelerators in the Universe.
+ These extreme events have fascinatecl scientists from. the time of their ciscovery in. 1962 (Linsley19039).., These extreme events have fascinated scientists from the time of their discovery in 1962 \citep{linsley}.
+. Since then. speculation about their origin has Llourishecl (Pierpaoli&Ferrar2005:Ghisellinietal.2008:Cuesta&Prada 2009).," Since then, speculation about their origin has flourished \citep{pierpaoli,ghise,cuesta}."
+. Theoretically. active galactic nucle: (AGN) have long been awored as the best candidates for particle acceleration to hese extreme energies (Ginzburg&Svrovatskii1964:Lillas 1984).," Theoretically, active galactic nuclei (AGN) have long been favored as the best candidates for particle acceleration to these extreme energies \citep{ginzburg,hillas}."
+. Unfortunately. no firm astrophysical association has »en established.," Unfortunately, no firm astrophysical association has been established."
+ Possibly the most intriguing suggestion of an association between VIECRs and AGN was put forward by he collaboration (Abrahametal.2004). that ound a possible correlation between the arrival direction of 27 events collected in their Southern Observatory ancl the »osition of nearby ας75 Alpe) ACN from the Verónn-C'etty catalog (Abrahamctal.2007)., Possibly the most intriguing suggestion of an association between UHECRs and AGN was put forward by the collaboration \citep{abraham} that found a possible correlation between the arrival direction of 27 events collected in their Southern Observatory and the position of nearby $\leq 75$ Mpc) AGN from the Verónn-Cetty catalog \citep{abraham2}.
+. The correlation. remains out. has weakened. slightly. with the inclusion of additional ULLECTtU events (for a grand total of 58) collected between 2004 January and 2009 March (Abrahamctal.2009)., The correlation remains but has weakened slightly with the inclusion of additional UHECR events (for a grand total of 58) collected between 2004 January and 2009 March \citep{abraham3}.
+. eonetheless. questions remain about the likelihood of the reported correlation (Abbasietal.2008a).," Nonetheless, questions remain about the likelihood of the reported correlation \citep{abbasi2}."
+. Lt is expected that 1 AGN are responsible for ULIECας. these must be preferentially placed. within a distance of 100 Mpe (Abrahametal.2007).," It is expected that if AGN are responsible for UHECRs, these must be preferentially placed within a distance of 100 Mpc \citep{abraham2}."
+ Interactions with Cosmic Microwave: Background. (CALB) photons should starve UllEClis arriving [rom larger distances through the Cireisen-Zatsepin-Ixuzmin (GZIx) ΓΣ 1966:Zatsepin&Ixuzmin 1966)..," Interactions with Cosmic Microwave Background (CMB) photons should starve UHECRs arriving from larger distances through the Greisen-Zatsepin-Kuzmin (GZK) effect \citep{greisen,zatse}."
+ In. fact. a combination of recent observations appear to corroborate the existence of a suppression of ULIECISs above 45107 eV. consistent with a GZlx horizon (Abbasictal.2008b:Abrahamet 2008h).," In fact, a combination of recent observations appear to corroborate the existence of a suppression of UHECRs above $4 \times 10^{19}$ eV consistent with a GZK horizon \citep{abbasi,abraham4}."
+. It is important to note that recent results from theAuger collaboration point to a transition to heavier Composition with increasing energies (Unger2007).., It is important to note that recent results from the collaboration point to a transition to heavier composition with increasing energies \citep{unger}.
+ If indeed iron nuclei are the dominant component of CLUECRs. the large deviation angles expected. for heavy. nuclei would render the detection of astrophysical counterparts even the nearest ones nearly impossible (Abrahametal.2008a).," If indeed iron nuclei are the dominant component of UHECRs, the large deviation angles expected for heavy nuclei would render the detection of astrophysical counterparts even the nearest ones nearly impossible \citep{abraham5}."
+. llere. we present a cross-correlation study of ας and the recently released: LLarge Area Telescope First Source Catalog (1EGL) in the 100 MeV. to 100 GeV energy range CXbdoetal.2010a).," Here, we present a cross-correlation study of UHECRs and the recently released Large Area Telescope First Source Catalog (1FGL) in the 100 MeV to 100 GeV energy range \citep{abdo}."
+. In contrast with previous studies. our cross-correlation analvsis is unique in that the ceatalog comprises a wide range of particle accelerators with gamma-ray production above LOO MeV directed in the Earth's direction.," In contrast with previous studies, our cross-correlation analysis is unique in that the catalog comprises a wide range of particle accelerators with gamma-ray production above $100$ MeV directed in the Earth's direction."
+ The Large Area Telescope on board. olfers à major improvement in sensitivity over previous GeV. detectors (Atwood 2007).., The Large Area Telescope on board offers a major improvement in sensitivity over previous GeV detectors \citep{atwood}. .
+ In its survey observation, In its survey observation
+of Perhuan et al. (,of Perlman et al. (
+"1996). also did not show a broad IL, coniponoent.",1996) also did not show a broad $\rm H_{\alpha}$ component.
+ Following Goodrich et al. (, Following Goodrich et al. (
+1990). the non-detection of a sieuificaut broad Ta euission liue suggests an extinction value of the Broad Line Region (BLR) of at least Ay = 3.7.,"1994), the non-detection of a significant broad $\alpha$ emission line suggests an extinction value of the Broad Line Region (BLR) of at least $\rm A_V$ = 3.7."
+ Near-infrared spectra of LES 1927|651 are required to estimate the upper limit of the extinction of the BLR by the detection of broad emission line componcuts at longer wavelength. e.g. at Dr 5 or Bra.," Near-infrared spectra of 1ES 1927+654 are required to estimate the upper limit of the extinction of the BLR by the detection of broad emission line components at longer wavelength, e.g. at Br $\gamma$ or $\alpha$."
+ The Ay values discussed above refer to the stm of the Galactic aud intrinsic extinction., The $\rm A_V$ values discussed above refer to the sum of the Galactic and intrinsic extinction.
+ The Galactic Ay value in the direction of LES 1927|65 Lis 0.291 as found iun NED., The Galactic $\rm A_V$ value in the direction of 1ES 1927+654 is 0.291 as found in NED.
+ Therefore. the corrected Ay for the Narrow Line Region iu the distant ACN is at most 1.71 (or slightly lower if we could correct for the underlying absorption line at /).," Therefore, the corrected $\rm A_V$ for the Narrow Line Region in the distant AGN is at most 1.71 (or slightly lower if we could correct for the underlying absorption line at $\beta$ )."
+ The rapid. giant and persistent X-ray variability. of LES 1927|651 detected with ROSAT aud coufirmed by Chandra point to a type P ACN classification. ic. we have a direct view to the central region.," The rapid, giant and persistent X-ray variability of 1ES 1927+654 detected with ROSAT and confirmed by Chandra point to a type 1 AGN classification, i.e. we have a direct view to the central region."
+ The optical spectra. however. give a Sevtert 2 classification with some contribution from the host ealaxyv.," The optical spectra, however, give a Seyfert 2 classification with some contribution from the host galaxy."
+ There secus to be a senificaut discrepancy (factor ~ 6) between the lower lait of the Ay value of 3.7 for the BLR. estimated from the optical spectrmu. auc the ιαπα Ay of 0.58 determined from the N-ray spectu.," There seems to be a significant discrepancy (factor $\sim$ 6) between the lower limit of the $\rm A_V$ value of 3.7 for the BLR, estimated from the optical spectrum, and the maximum $\rm A_V$ of 0.58 determined from the X-ray spectrum."
+ We note that if the dust is optically thin for N-vavs the intrinsic X-ray Ay cau not ereatlv exceed the observed value of 0.58. otherwise we would observe much strouger absorption edges as well as a lieher cuereyv soft X-ray cutoff.," We note that if the dust is optically thin for X-rays the intrinsic X-ray $\rm A_V$ can not greatly exceed the observed value of 0.58, otherwise we would observe much stronger absorption edges as well as a higher energy soft X-ray cutoff."
+ In this section we will discuss the possible scenarios resulting iu this unique colmbination of X-ray and optical properties., In this section we will discuss the possible scenarios resulting in this unique combination of X-ray and optical properties.
+ One possible explanation for the apparent disagreeineut i N-rav and optical properties ids an extremely underbhuuiuous. or even absent. BLR.," One possible explanation for the apparent disagreement in X-ray and optical properties is an extremely underluminous, or even absent, BLR."
+ The bhuuimositv derivec for the putative broad Πα component (see Fig., The luminosity derived for the putative broad $\alpha$ component (see Fig.
+ 10) is about LOeres.JH. whereas the total huninosity iu the optical B-haud is about 107erest.," 10) is about $\rm 10^{38} erg s^{-1}$, whereas the total luminosity in the optical $B$ -band is about $\rm 10^{43}\ erg\ s^{-1}$."
+ The typical ratio of cluission line to total B-baud huninosity iu Sevtert ealaxies is approximately 1 to 5 per ceut (Netzer. private commnimnuication). thus we would expect a line luninosity of about 104ere«| iu LES 19271651.," The typical ratio of emission line to total B-band luminosity in Seyfert galaxies is approximately 1 to 5 per cent (Netzer, private communication), thus we would expect a line luminosity of about $\rm 10^{41}\ erg\ s^{-1}$ in 1ES 1927+654."
+ This may be sugeestive of an underhuninous BLR., This may be suggestive of an underluminous BLR.
+ The true optical extinction might be larecr than that derived frou the optical spectrum., The true optical extinction might be larger than that derived from the optical spectrum.
+ It is possible hat. for exinuple. the dust erains in the absorbing matter could be individually thick to N-ravs (cf.," It is possible that, for example, the dust grains in the absorbing matter could be individually thick to X-rays (cf."
+ Fircuau 1971)., Fireman 1974).
+ In this case. the soft N-rav spectrum cutoff does not chanec with increasing BAy values. o. it remains constant in the rav observations.," In this case, the soft X-ray spectrum cutoff does not change with increasing $\rm A_V$ values, i.e. it remains constant in the X-ray observations."
+ Iu addition. we could have a large selfblauketiug factor. in which case the absorbed hWmuuinositv from the BLR Ιστ be even larger than the value derived in the case where the dust is optically thin to N-ravs.," In addition, we could have a large self-blanketing factor, in which case the absorbed luminosity from the BLR might be even larger than the value derived in the case where the dust is optically thin to X-rays."
+ We also can not exclude that LES 1927|651 exhibits a higher dust ο gas ratio than that of about 0.01 derived. from interstellar dust to gas ratios (CGorensteiu 1975)., We also can not exclude that 1ES 1927+654 exhibits a higher dust to gas ratio than that of about 0.01 derived from interstellar dust to gas ratios (Gorenstein 1975).
+into three peaks. and that the 3.44 mm source detected by Hunteretal.(1999) actually consists of five millimeter sources.,"into three peaks, and that the 3.4 mm source detected by \citet{hunter1999} actually consists of five millimeter sources."
+" Regarding the molecular outflow emission in. this region. Zhangetal.(2007) identify three molecular outflows in CO((2-1) and 65-5,. one of them coinciding with an ooutflow and the well-collimated SiO jet detected by Hunteretal. (1999)."," Regarding the molecular outflow emission in this region, \citet{zhang2007} identify three molecular outflows in (2–1) and $6_{5}$ $5_{4}$, one of them coinciding with an outflow and the well-collimated SiO jet detected by \citet{hunter1999}."
+" All three outflows appear to originate from the dust condensation in a region of about 3"".", All three outflows appear to originate from the dust condensation in a region of about $3''$.
+ All this information indicates that active star formation in clustered mode ts taking place in 55142., All this information indicates that active star formation in clustered mode is taking place in 5142.
+ The dense gas emission in this region has been studied with single-dish telescopes inNH3..CS.. HCN. HCO™..CH;OH.. and (Verdes-Montenegroetal.1989:Estalella1993:Hunteretal.1995.1999;Cesaroni 1999)).," The dense gas emission in this region has been studied with single-dish telescopes in, HCN, and \citealt{verdes-montenegro1989,estalella1993,hunter1995,hunter1999,cesaroni1999}) )."
+ In particular. the dense gas emission traced by the mmolecule has been observed with high angular resolution using the VLA (Zhangetal.2002).," In particular, the dense gas emission traced by the molecule has been observed with high angular resolution using the VLA \citep{zhang2002}."
+ The high angular resolution eemission consists of a central and compact core associated with the dust condensation. harboring at least. three intermediate/high-mass young stars. surrounded by fainter ccores located in a more extended structure with no signs of stellar activity (no maser nor molecular outflow emission associated) indicating that the region may harbor cores in different evolutionary stages.," The high angular resolution emission consists of a central and compact core associated with the dust condensation, harboring at least three intermediate/high-mass young stars, surrounded by fainter cores located in a more extended structure with no signs of stellar activity (no maser nor molecular outflow emission associated) indicating that the region may harbor cores in different evolutionary stages."
+ The presence of cores contaming massive star(s) together with cores with no star-formation activity makes this region a good choice to study how the rratio behaves in high-massregions.., The presence of cores containing massive star(s) together with cores with no star-formation activity makes this region a good choice to study how the ratio behaves in high-mass.
+ In this paper we report CARMA observations of the continuum emission at 3.2 mm and the dense gas traced by .((1-0) toward 55142., In this paper we report CARMA observations of the continuum emission at 3.2 mm and the dense gas traced by (1–0) toward 5142.
+ The paper layout ts as follows: in .22 we summarize our observations and the data reduction process., The paper layout is as follows: in 2 we summarize our observations and the data reduction process.
+ In .33 we present the main results for the continuum and mmolecular line emission., In 3 we present the main results for the continuum and molecular line emission.
+ In .J4 we analyze the molecular emission of several species by computing their column density maps., In 4 we analyze the molecular emission of several species by computing their column density maps.
+ In .55 we show the mai results of the chemical model to qualitatively reproduce the abundances of the region., In 5 we show the main results of the chemical model to qualitatively reproduce the abundances of the region.
+ Finally. ins .66 we discuss our findings. and we list the main conclusions ins 77," Finally, in 6 we discuss our findings, and we list the main conclusions in 7."
+", The Combined Array for Research in Millimeter-wave (CARMA) was used to observe the 3.2 mm continuum and the (1-0) emission toward AFGL 5142.", The Combined Array for Research in Millimeter-wave (CARMA) was used to observe the 3.2 mm continuum and the (1-0) emission toward AFGL 5142.
+ CARMA consists of six 10 m and nine 6 m antennas located at 2200 meters elevation at Cedar Flat in the Inyo Mountains of California., CARMA consists of six 10 m and nine 6 m antennas located at 2200 meters elevation at Cedar Flat in the Inyo Mountains of California.
+ The observations were carried out on 2007 February 4 and March 11 using the array in the C configuration with 14 antennas in the array., The observations were carried out on 2007 February 4 and March 11 using the array in the C configuration with 14 antennas in the array.
+" The projected baselines ranged from 26 to 370 m. The phase center was set at &=0530""48:02, 0=-33747'54""747."," The projected baselines ranged from 26 to 370 m. The phase center was set at $\alpha=05^{\rm h}30^{\rm m}48\rm \fs02$, $\delta=+33\degr47\arcmin54\farcs47$."
+" The FWHM of the primary beam at the frequency of the observations was 132"" for the 6 m antennas and 77"" for the 10 m antennas.", The FWHM of the primary beam at the frequency of the observations was $''$ for the 6 m antennas and $''$ for the 10 m antennas.
+ System temperatures were around 250 K during both The digital correlator was configured to observe simultaneously the continuum emission and the (1-0) group of the hyperfine transitions (93.176331. GHz. in the lower sideband).," System temperatures were around 250 K during both The digital correlator was configured to observe simultaneously the continuum emission and the (1–0) group of the hyperfine transitions (93.176331 GHz, in the lower sideband)."
+ The continuum data were recorded in, The continuum data were recorded in
+"26,891 secure (95% or greater confidence) redshifts, while the VVDS and zCOSMOS data sets include only","$26,891$ secure $95\%$ or greater confidence) redshifts, while the VVDS and zCOSMOS data sets include only"
+The starless cores owe their reputation as the sites of future star formation to the similarities they share with neighboring cores already in the pains of stellar birth.,The starless cores owe their reputation as the sites of future star formation to the similarities they share with neighboring cores already in the pains of stellar birth.
+" Small (tenths of pe). dense (1g,~10? to 109 em?) dark clouds of a few solar masses. the starless cores contain no infrared sources above the sensitivity level of the IRAS satellite (about 0.1 L.. at the distance of Taurus)."," Small (tenths of pc), dense $n_{\rm H_2} \sim 10^3$ to $10^6$ $^{-3}$ ) dark clouds of a few solar masses, the starless cores contain no infrared sources above the sensitivity level of the IRAS satellite (about 0.1 $_\odot$ at the distance of Taurus)."
+ Mixed in with a population of cores. embedded infrared protostars. and young T-Tauri stars. the starless cores seem on the verge of issue. (Myersetal.1983:Bergin&Tafalla2007 ).," Mixed in with a population of cores, embedded infrared protostars, and young T-Tauri stars, the starless cores seem on the verge of issue. \citep{MyersLinkeBenson1983, MyersBenson1983, BensonMyers1989, Beichman1986,
+DiFrancesco2007, Ward-Thompson2007,BerginTafalla2007}."
+. The starless cores are remarkable in another sense., The starless cores are remarkable in another sense.
+ They may be the structures in the interstellar medium in quasi-static equilibriumonly with lifetimes longer than their sound-crossing times (Beichmanetal.1986;Jessop&Ward-Thompson2000:Kirketal.2005:Ward-Thompson 2007).," They may be the only structures in the interstellar medium in quasi-static equilibrium with lifetimes longer than their sound-crossing times \citep{Beichman1986,JessopWT2000,KirkWTAndre2005,Ward-Thompson2007}."
+. Whereas all larger-scale clouds are better described as transient structures within a turbulent cascade (Fieldetal.2008).. the starless cores show observed densities that are consistently matched (Ward-etal.2005) by the truncated solutions of the isothermal Lane-Emden equation. re.. the Bonnor-Ebert (BE) spheres (Bonnor1956).," Whereas all larger-scale clouds are better described as transient structures within a turbulent cascade \citep{Field2008}, the starless cores show observed densities that are consistently matched \citep{Ward-Thompson1994,Andre1996,Ward-Thompson1999,Bacmann2000,Evans2001,AlvesLadaLada2001,Tafalla2002,Kandori2005,KirkWTAndre2005}
+ by the truncated solutions of the isothermal Lane-Emden equation, i.e., the Bonnor-Ebert (BE) spheres \citep{Bonnor1956}."
+. These are self-gravitating spheres supported by thermal energy and bounded by an external pressure., These are self-gravitating spheres supported by thermal energy and bounded by an external pressure.
+ The source of the external pressure appears to depend upon the star forming region. though its necessity is well documented.," The source of the external pressure appears to depend upon the star forming region, though its necessity is well documented."
+ Within the Pipe nebula. the core energetics are consistent with a constant external pressure supplied by hotter. more rareified gas (Ladaetal.2008).," Within the Pipe nebula, the core energetics are consistent with a constant external pressure supplied by hotter, more rareified gas \citep{Lada2008}."
+. In other regions. such as Lupus (Teixeiraetal.2005).. the cores appear to be surrounded by lower density molecular gas with observed molecular line widths =|] kms.," In other regions, such as Lupus \citep{Teixeira2005}, the cores appear to be surrounded by lower density molecular gas with observed molecular line widths $\geq 1$ kms."
+ Here and in similar regions small-scale (smaller than the observing beam) supersonic turbulence in the low-density molecular gas around the cores may supply the confining pressure., Here and in similar regions small-scale (smaller than the observing beam) supersonic turbulence in the low-density molecular gas around the cores may supply the confining pressure.
+ While a continuous extension to the larger-scale ISM prompts the question. of whether the cores are individual dynamical entities. a distinct boundary ts evident in several observations.," While a continuous extension to the larger-scale ISM prompts the question of whether the cores are individual dynamical entities, a distinct boundary is evident in several observations."
+ First. maps of dust absorption in the Taurus region (Bacmannetal.2000) often show sharp edges defining the core boundaries.," First, maps of dust absorption in the Taurus region \citep{Bacmann2000} often show sharp edges defining the core boundaries."
+ Second. around at least one core... a boundary is defined by a sharp transition (within less than an observing beam width) between subsonic and supersonic turbulence in the molecular gas (Pinedaetal.2010).," Second, around at least one core, a boundary is defined by a sharp transition (within less than an observing beam width) between subsonic and supersonic turbulence in the molecular gas \citep{Pineda2010}."
+. Third. a boundary is a necessary feature in models of radiative cooling and UV heating (Evansetal.2001:Shirley2002:Zuc-etal.2004;Keto&Field2005) that use the basic BE structure and accurately predict the observed gas and dust temperatures (Ward-Thompsonetal.2002:Pagani2003.2004:Crapsietal. 2007).. the observed variation in the excitation of the CO molecule across the starless cores (Berginetal.2006:Schneeetal.2007;Keto&Caselli 2008).. and the observed brightness and profiles of molecular lines (Keto&Caselli2010).," Third, a boundary is a necessary feature in models of radiative cooling and UV heating \citep{ Evans2001,ShirleyEvansRawlings2002,Zucconi2001,StamatellosWhitworth2003,Goncalves2004,KetoField2005}
+ that use the basic BE structure and accurately predict the observed gas and dust temperatures \citep{Ward-Thompson2002,Pagani2003,Pagani2004,Crapsi2007}, the observed variation in the excitation of the CO molecule across the starless cores \citep{Bergin2006,Schnee2007,KetoCaselli2008}, and the observed brightness and profiles of molecular lines \citep{KetoCaselli2010}."
+. The predictive success of models based on BE spheres (including a core boundary) implies that the starless cores are indeed individual dynamie entities in quasi-static equilibrium. balancing the forces of their own self-gravity. internal energy. and the larger-scale ISM.," The predictive success of models based on BE spheres (including a core boundary) implies that the starless cores are indeed individual dynamic entities in quasi-static equilibrium, balancing the forces of their own self-gravity, internal energy, and the larger-scale ISM."
+ In some cases. this equilibrium may be unstable. leading the core to contract. though their internal pressure ts sufficiently high to prevent free-fall.," In some cases, this equilibrium may be unstable, leading the core to contract, though their internal pressure is sufficiently high to prevent free-fall."
+ The hydrodynamic equations that describe the equilibrium configuration of both stars and the starless cores allow for perturbations in the form of sonic waves that may persist for many crossing times., The hydrodynamic equations that describe the equilibrium configuration of both stars and the starless cores allow for perturbations in the form of sonic waves that may persist for many crossing times.
+ Such oscillations are observed in the Sun. other main sequence stars. and white dwarfs. and given the turbulent state of the larger-scale ISM. should," Such oscillations are observed in the Sun, other main sequence stars, and white dwarfs, and given the turbulent state of the larger-scale ISM, should"
+The Dedeliud tri-axial ellipsoids are an example of non-axisviunietric. but stationary solutions witlin Newtonian gravity.,"The Dedekind tri-axial ellipsoids are an example of non-axisymmetric, but stationary solutions within Newtonian gravity."
+ Duc to internal motions. hey are. m fact. stationary du an inertial frame.," Due to internal motions, they are, in fact, stationary in an inertial frame."
+ When addressing the question of whether or not stationary. but non-axisviunetrie solutions are oossible within General Relativity. this property uakes the Dedekind ellipsoids a natural choice upon which to base one's considerations.," When addressing the question of whether or not stationary, but non-axisymmetric solutions are possible within General Relativity, this property makes the Dedekind ellipsoids a natural choice upon which to base one's considerations."
+ It was. ii part. with this question in mind that 77? iuned their atteutious to the post-Newtonian (DN) approximation of the Dedekind ellipsoids.," It was, in part, with this question in mind that \citet{CE74, CE78_CE78} turned their attentions to the post-Newtonian (PN) approximation of the Dedekind ellipsoids."
+ Iu a paper from the same series. 7 µας already considered the axisvuuuetrce huit of the PN σιacobi ellipsoids at length and was able to show that it coincides with a certain PN Alaclaurin spheroid (just as thei Newtomlan counterparts coincide at the point of bifurcation).," In a paper from the same series, \citet{Chandrasekhar67c_C67} had already considered the axisymmetric limit of the PN Jacobi ellipsoids at length and was able to show that it coincides with a certain PN Maclaurin spheroid (just as their Newtonian counterparts coincide at the point of bifurcation)."
+ This is related to the fact that the PN figures were chosen to rotate uniforiiulv., This is related to the fact that the PN figures were chosen to rotate uniformly.
+ Ou the other haud. the PN velocity field chosen in ?/— excludes the possibility of uniform rotation in the axisvuuuetricliuit although it is possible iu the axisvunuetricc," On the other hand, the PN velocity field chosen in \cite{CE78_CE78} excludes the possibility of uniform rotation in the axisymmetric although it is possible in the axisymmetric."
+"ase, This restriction seenis neither natural nor advisable iu the context of trving to settle the question as to the existence of relativistic. nou-axisviunuetrie. stationary solutions."," This restriction seems neither natural nor advisable in the context of trying to settle the question as to the existence of relativistic, non-axisymmetric, stationary solutions."
+ The naive expectation is that tle axisvuumetric PN Dedekind ellipsoids contain the PN Alaclaurin spheroids in the axisvuuuetric luit (up to arbitrary order)., The ve expectation is that the axisymmetric PN Dedekind ellipsoids contain the PN Maclaurin spheroids in the axisymmetric limit (up to arbitrary order).
+ Tn this article. we beein in 77 by examining the axisviuuetric case of a ecneralization to the solution prescuted in ?..," In this article, we begin in \ref{axisymm_point} by examining the axisymmetric case of a generalization to the solution presented in \citet{CE78_CE78}."
+ We proceed in 3 to consider a (continous) lit to axisvuuuctry., We proceed in \ref{axisymm_limit} to consider a (continuous) limit to axisymmetry.
+ In ?7 the connection to the post-Newtouian Maclamin spheroids is examiunect., In \ref{discussion} the connection to the post-Newtonian Maclaurin spheroids is examined.
+ We cousider a generalization of the PN Dedekiud ellipsoids preseuted in 7? (veferred to. from here on in as CETS) in which we add post-Newtoman terms to the velocity., We consider a generalization of the PN Dedekind ellipsoids presented in \citet{CE78_CE78} (referred to from here on in as CE78) in which we add post-Newtonian terms to the velocity.
+ We comply with the notation used in CET78 aud refer the reader to the definitions there for the various quautities., We comply with the notation used in CE78 and refer the reader to the definitions there for the various quantities.
+there. forcing much of the enerey out at both longer aud shorter waveleugths.,"there, forcing much of the energy out at both longer and shorter wavelengths."
+ The dashed line iu Figure 12 shows the integrated flux as a function of wavelength for secondary. eclipse (blue dashed line. Figure 12)).," The dashed line in Figure \ref{energy} shows the integrated flux as a function of wavelength for secondary eclipse (blue dashed line, Figure \ref{energy}) )."
+ The flux from the planet iu the 1-15 juu rauge accounts for of the planct’s enmergeut enerev. wlüch makes this an especially Huportant waveleneth rauge for determining the atinospheric properties of €J136b.," The flux from the planet in the 1-15 $\mu$ m range accounts for of the planet's emergent energy, which makes this an especially important wavelength range for determining the atmospheric properties of GJ436b."
+ It is interesting to note the increased variability in the cuitted flix from the planet as a function of orbital position in the « solar case compared with the L< solar case in Figures 11 and 12.., It is interesting to note the increased variability in the emitted flux from the planet as a function of orbital position in the $\times$ solar case compared with the $\times$ solar case in Figures \ref{fluxes} and \ref{energy}.
+ The flux emitted from the 50. solar case at secondary eclipse (blue line. Figure 13) is lacking in many of the predoiinaut spectral features secu at other orbital phases. due toa shallower day-sicle temperature eradicut.," The flux emitted from the $\times$ solar case at secondary eclipse (blue line, Figure \ref{fluxes}) ) is lacking in many of the predominant spectral features seen at other orbital phases, due to a shallower day-side temperature gradient."
+" The absorption features due to CIT, are much. weaker at secondary eclipse indicating a reduction in CIT, abundance on the dav-side compared to the iigbht-de which is seen during transit (orauge line. Figure 11))."," The absorption features due to $_4$ are much weaker at secondary eclipse indicating a reduction in $_4$ abundance on the day-side compared to the night-side which is seen during transit (orange line, Figure \ref{fluxes}) )."
+ Observing the flux emitted. from (100 as a fiction of wavelength at several different poiuts along its orbit could reveal a great deal about its overall chemical composition., Observing the flux emitted from GJ436b as a function of wavelength at several different points along its orbit could reveal a great deal about its overall chemical composition.
+ The atmospheric models presented here are not oulv useful for exploring circulation regimes. chemistry. and radiative transfer. but can also provide insight into current observations aud help euide future observations of €GJ136b.," The atmospheric models presented here are not only useful for exploring circulation regimes, chemistry, and radiative transfer, but can also provide insight into current observations and help guide future observations of GJ436b."
+ Figure 10 also juchides tle availableSpifzer secondary eclipse mncasurehicuts from ?.., Figure \ref{light_curves} also includes the available secondary eclipse measurements from \citet{ste10}.
+" In. all metallicity cases, our predicted planct/star flux ratio falls short of the measured 2.6. 5.5. 8.0. 16.0. and 210 pau values and is higher than the observed 1.5 44 value."," In all metallicity cases, our predicted planet/star flux ratio falls short of the measured 3.6, 5.8, 8.0, 16.0, and 24.0 $\mu$ m values and is higher than the observed 4.5 $\mu$ m value."
+ However. it is useful to uote that < solar model plauct/star fux ratio comes inuch closer to matching the observations ναι the Ls solar model. which could lint at a lüeh metallicity as already sugeested by ?..," However, it is useful to note that $\times$ solar model planet/star flux ratio comes much closer to matching the observations than the $\times$ solar model, which could hint at a high metallicity as already suggested by \citet{ste10}."
+ Wiel metallicity solutions (030 « solar)are also favored bv 7. to match the 8 jaa observatious of GJ136b from ?.., High metallicity solutions $\sim$ $\times$ solar)are also favored by \citet{spi10} to match the 8 $\mu$ m observations of GJ436b from \citet{demi07}. .
+ The predicted plauct/star flux ratios im the 50. & solar case are within two sigma of the values measured by ?. at all bandpasses except 3.6 fan. Iun comparing our predicted planet/star flux ratios with those observed bv ο it is niportant to remember that we have not altered the teniperature or chemistry iu our models in an attempt to match observations., The predicted planet/star flux ratios in the 50 $\times$ solar case are within two sigma of the values measured by \citet{ste10} at all bandpasses except 3.6 $\mu$ m. In comparing our predicted planet/star flux ratios with those observed by \citet{ste10} it is important to remember that we have not altered the temperature or chemistry in our models in an attempt to match observations.
+ The interplay between the equilibriuni chemistry mining ratios. the absorption and emission of flix. aud the atmospheric dyaianuics dictates the temperature structure and cimerecut spectrum.," The interplay between the equilibrium chemistry mixing ratios, the absorption and emission of flux, and the atmospheric dynamics dictates the temperature structure and emergent spectrum."
+ As pointed out in ? it is likely that disequilibrimim chemistry plays a strong role in the atimosplicre of CJ0)., As pointed out in \citet{ste10} it is likely that disequilibrium chemistry plays a strong role in the atmosphere of GJ436b.
+" They suggest a CO/CIE, ratio that is many orders of magnitude larger than what one would predict from an equilibrimm chemistry model.", They suggest a $_4$ ratio that is many orders of magnitude larger than what one would predict from an equilibrium chemistry model.
+" Euhancement of CO at the expense of CTL, is well known in the atiuosplieres of the solar system's giant planets and brown dwarts, but not to such an extreme degree."," Enhancement of CO at the expense of $_4$ is well known in the atmospheres of the solar system's giant planets and brown dwarfs, but not to such an extreme degree."
+ Additionally. they suggest that the photochemical destruction of CIT; is Óuportaut to further lessen this molecule’s importance in the planet's atinospliere.," Additionally, they suggest that the photochemical destruction of $_4$ is important to further lessen this molecule's importance in the planet's atmosphere."
+ The carbon chemistry of an atmosphere will strongly affect flux iicasurements dn the 3.6. L5. 5.5.and Wm yma baudpasses.," The carbon chemistry of an atmosphere will strongly affect flux measurements in the 3.6, 4.5, 5.8,and 8.0 $\mu$ m bandpasses."
+" As shown in Figure ll and discussed in Section ?7.. lowering the amount of CIE, in the atmosphere will increase the eiiergeut flux from the planet in the 3.6 aud 8.0 gan baudpasses."," As shown in Figure \ref{fluxes} and discussed in Section \ref{spec}, lowering the amount of $_4$ in the atmosphere will increase the emergent flux from the planet in the 3.6 and 8.0 $\mu$ m bandpasses."
+ Higher order hydrocarbons produced by mixiug aud photochemistry (7). such as Colo. Coll). and Collg. are also strong absorbers throughout the near- and mid-infrared.," Higher order hydrocarbons produced by mixing and photochemistry \citep{zah09_2}, such as $_2$ $_2$, $_2$ $_4$, and $_2$ $_6$, are also strong absorbers throughout the near- and mid-infrared."
+ Tucreasing the amount of CO». along with CO. in the atmosphere will decrease the eiuicrseut fiux frou. the planet in the £5 pau baudpass while at the sanie time iucroasius the cuiergeut fux at waveleneth ou either side of this bandpass. including the 3.6 aud 5.5 gma baudpasses.," Increasing the amount of $_2$, along with CO, in the atmosphere will decrease the emergent flux from the planet in the 4.5 $\mu$ m bandpass while at the same time increasing the emergent flux at wavelength on either side of this bandpass, including the 3.6 and 5.8 $\mu$ m bandpasses."
+ Althoughour models do not include disequilibriuu chemistry. they can provide au important coustraiut for disequilibrimu chemistry modelsnamely. estimates of dynamical timescales aud vertical mixing rates.," Although our models do not include disequilibrium chemistry, they can provide an important constraint for disequilibrium chemistry models—namely, estimates of dynamical timescales and vertical mixing rates."
+" Disequilibrium chemistry is expected in all reeious of the atinosphere where dvuauiic timescales. 745,,. areshorter than clicmical tiniescales. Tobenm- "," Disequilibrium chemistry is expected in all regions of the atmosphere where dynamic timescales, $\tau_{dyn}$ , areshorter than chemical timescales, $\tau_{chem}$ ."
+The chynaic timescale is elven sinplv by where Lo ds the relevant leneth scale iu the horizontal⋅ or vertical. direction. aud Wis2. horizoutal: or vertical⋅ wind⋅ speed., The dynamic timescale is given simply by where $L$ is the relevant length scale in the horizontal or vertical direction and $V$ is horizontal or vertical wind speed.
+ For one-dimensional⋅ photochemical, For one-dimensional photochemical
+ For one-dimensional⋅ photochemical⋅, For one-dimensional photochemical
+a single point source was fitted to the visibility data for PKS 1718-649.,a single point source was fitted to the visibility data for PKS 1718–649.
+ The antenna phases were self-calibrated with a 10 s solution interval. and the model-fitting procedure was repeated.," The antenna phases were self-calibrated with a 10 s solution interval, and the model-fitting procedure was repeated."
+ The resulting flux density was recorded after three more iterations of self-calibration and model fitting., The resulting flux density was recorded after three more iterations of self-calibration and model fitting.
+ The two frequency bands were processed separately until this point., The two frequency bands were processed separately until this point.
+ The mean of the 8896 MHz time series was smaller than the mean of the 8640 MHz time series. as expected from the previously measured radio spectrum of the source (Tingay et 11997).," The mean of the 8896 MHz time series was smaller than the mean of the 8640 MHz time series, as expected from the previously measured radio spectrum of the source (Tingay et 1997)."
+ We increased the 8896 MHz flux densities by1.4%.. and then averaged the results of the two bands at each epoch.," We increased the 8896 MHz flux densities by, and then averaged the results of the two bands at each epoch."
+ The final light curve is shown in Figure 2a., The final light curve is shown in Figure 2a.
+ The error bars span the difference between the results of the two bands (after having normalized the 8896 MHz data)., The error bars span the difference between the results of the two bands (after having normalized the 8896 MHz data).
+ These error bars give an indication of the statistical noise in each measurement., These error bars give an indication of the statistical noise in each measurement.
+ The fluctuations of PKS 1718-649 are small. with an RMS variation of 29 mJy. or of the mean flux density of 4.04 Jy.," The fluctuations of PKS 1718–649 are small, with an RMS variation of 29 mJy, or of the mean flux density of 4.04 Jy."
+" The local sidereal times of the epochs were strongly clustered around two values. 14:00 and 23:00 (hereafter. “early” and ""]ate). because of the preference for particular hour angles mentioned in 2.."," The local sidereal times of the epochs were strongly clustered around two values, 14:00 and 23:00 (hereafter, “early” and “late”), because of the preference for particular hour angles mentioned in \ref{sec:design}."
+ In Figure 2a. the open symbols represent data from the early LSTs. and the filled symbols represent data from the late LSTs.," In Figure 2a, the open symbols represent data from the early LSTs, and the filled symbols represent data from the late LSTs."
+ The different symbol shapes represent different array configurations., The different symbol shapes represent different array configurations.
+ The reason for encoding this, The reason for encoding this
+study by virtually every. N-ray. satellite during the past decade.,study by virtually every X-ray satellite during the past decade.
+ Defined by their optical emission-line properties (FWIIM Ibs<2000 km | and ii| A5007/IL2 <3) NLSIs are different [rom type 2 Sevfert galaxies. which generally have ΟΠΗ>3 (Osterbrock Poegge 1985: Goodrich 1989).," Defined by their optical emission-line properties (FWHM $\beta < 2000$ km $^{-1}$ and ] $\lambda$ $\beta < 3$ ) NLS1s are different from type 2 Seyfert galaxies, which generally have $\beta > 3$ (Osterbrock Pogge 1985; Goodrich 1989)."
+ NLSIs usually have strong permitted lines of11. resembling the well-known prototvpe of their class. I Zw 1.," NLS1s usually have strong permitted lines of, resembling the well-known prototype of their class, I Zw 1."
+ The exireme N-rax properties ofNarrow-line Sevler( 1 galaxies are now well estabelished., The extreme X-ray properties of Narrow-line Seyfert 1 galaxies are now well estabelished.
+ Thev frequently show a strong soft excess component. their hard X-ray. spectrum tends (o be steeper (han in similar broad-lime Sevfert 1 galaxies. ancl thev show enhanced X-rav variability (e.g.. Leighlv 1999a.b and relerences (herein).," They frequently show a strong soft excess component, their hard X-ray spectrum tends to be steeper than in similar broad-line Seyfert 1 galaxies, and they show enhanced X-ray variability (e.g., Leighly 1999a,b and references therein)."
+ The most promising explanation Lor (his behavior is that NLSIs have a higher mass accretion rate wilh respect to the Eddington value (han ordinary Sevlert galaxies with broad optical lines (Laor2000)., The most promising explanation for this behavior is that NLS1s have a higher mass accretion rate with respect to the Eddington value than ordinary Seyfert galaxies with broad optical lines \citep{laor00}.
+. This result is potentially very important: since AGN are believed in general to be powered by accretion. study of objects with the highest accretion rate may help us understand AGN accretion in general.," This result is potentially very important: since AGN are believed in general to be powered by accretion, study of objects with the highest accretion rate may help us understand AGN accretion in general."
+ Ark 564 (2—0.02468-:0.00007: HIuchra. Vogelev. Geller 1999) is an interesting object from the point of view that shows a peculiar emission line-like feature near 1 keV in low resolution spectra.," Ark 564 $=$ $\pm$ 0.00007; Huchra, Vogeley, Geller 1999) is an interesting object from the point of view that shows a peculiar emission line-like feature near 1 keV in low resolution spectra."
+ This emission line-like feature has been reported [rom various observations performed byROSAT. andBeppoSAN (Brandtetal.1994:Leiehly1999b:Turner.George.etal. 2001).," This emission line-like feature has been reported from various observations performed by, and \citep {bra94, lei99b, tur99, com01}."
+. Turner et al. (, Turner et al. (
+1999) modeled the feature with a strong (equivalent width: EW~70 eV) and relatively broad (6—0.16 keV) line at n keV. assuming (hat spectra components are Galactic absorption. a power-law and a Gaussian line in (he energy band of 0.65.0 and 7.510 keV. Using the same data and the different continuummodel (Galactic absorbed a power-law plus a black body). Leighly (1999b) obtained weaker CEW=30 eV) line at η keV with o of 0.1 eV. Comastrietal.(2001) also reported that the mode dependent EW is as low as 2050 eV with the same model as Leighlvs from independent analvses.,"1999) modeled the feature with a strong (equivalent width: $\sim$ 70 eV) and relatively broad $\sigma$ =0.16 keV) line at $^{+0,02}_{-0.04}$ keV, assuming that spectral components are Galactic absorption, a power-law and a Gaussian line in the energy band of 0.6–5.0 and 7.5–10 keV. Using the same data and the different continuummodel (Galactic absorbed a power-law plus a black body), Leighly (1999b) obtained weaker $=30$ eV) line at $^{+0.024}_{-0.031}$ keV with $\sigma$ of 0.1 eV. \citet {com01} also reported that the model dependent EW is as low as 20–50 eV with the same model as Leighly's from independent analyses."
+ In theBeppoS observation. a line-like feature around 1 keV is clearly visible (Comastrietal.2001).," In the observation, a line-like feature around 1 keV is clearly visible \citep {com01}."
+. Still the origin of (his feature was far [rom being unambiguously identified., Still the origin of this feature was far from being unambiguously identified.
+ In order to identify (his feature. we carried out an observation with the IWhieh Energy Transmission Grating Spectrometer (ILETGS).," In order to identify this feature, we carried out an observation with the High Energy Transmission Grating Spectrometer (HETGS)."
+ Observed with higher energy resolution. some possible origins for this feature. such as blends of narrow emission lines. can be investigated.," Observed with higher energy resolution, some possible origins for this feature, such as blends of narrow emission lines, can be investigated."
+ If. we can identilv (his feature. it may. enable us to investigate (he physical conditions such as ionization state of the surrounding matter of the AGN.," If we can identify this feature, it may enable us to investigate the physical conditions such as ionization state of the surrounding matter of the AGN."
+" Ark 564 is one of the brightest narrow-line Sevlert 1: galaxies in the hard X-ray band and the 2.10 keV [αν is a few times Hergem7s J|, so a erating observation using Chendrais clearly feasible."," Ark 564 is one of the brightest narrow-line Seyfert 1 galaxies in the hard X-ray band and the 2–10 keV flux is a few times $^{-11}~{\rm erg~cm^{-2}~s^{-1}}$ , so a grating observation using is clearly feasible."
+is mareinally significant evidence that +A has à low spectral index compared to the cloud.,is marginally significant evidence that 4A has a low spectral index compared to the cloud.
+ There is also strong evidence that the region associated with the 4X. outflow as a low spectral index., There is also strong evidence that the region associated with the 4A outflow has a low spectral index.
+ Various possible causes of this were discussed., Various possible causes of this were discussed.
+ “Phere is no evidence suggesting tha he emission is optically thick at either. wavelength., There is no evidence suggesting that the emission is optically thick at either wavelength.
+ The apparentlv low spectral index of 4X was dependent on emperature assumptions. but the outllow region shoulc not be systematically cooler than the surrounding cloud. so this is unlikely to provide an explanation.," The apparently low spectral index of 4A was dependent on temperature assumptions, but the outflow region should not be systematically cooler than the surrounding cloud, so this is unlikely to provide an explanation."
+ We considere 1e possibility that molecular line emission could. manifes unself as a region of low spectral index by contaminating the DTlux density., We considered the possibility that molecular line emission could manifest itself as a region of low spectral index by contaminating the flux density.
+ Based on the measured. line strengths. [rom BSDGALA. this explanation seemed. unlikely to be the sole ause. although it could constitute a partial cause.," Based on the measured line strengths from BSDGMA, this explanation seemed unlikely to be the sole cause, although it could constitute a partial cause."
+ Le was ierefore concluded that the low spectral index of 4A was most likely an indication of cust grain growth in the dense circumstellar disks., It was therefore concluded that the low spectral index of 4A was most likely an indication of dust grain growth in the dense circumstellar disks.
+ Εις grain growth then appears to be more advanced in 4X than in 4D or 4C. We have seen an area of low spectral index running through LRASLA. in the position of the molecular outLow seen by XSDGAMALA.," This grain growth then appears to be more advanced in 4A than in 4B or 4C. We have seen an area of low spectral index running through IRAS4A, in the position of the molecular outflow seen by BSDGMA."
+ This ridge is seen on both sides of IRASLA. and is clearly distinct from IRAS 4€ to the Northeast.," This ridge is seen on both sides of IRAS4A, and is clearly distinct from IRAS 4C to the Northeast."
+ We argued in Section 5 that contamination from molecular line Hus is probably not sullicient to explain this., We argued in Section \ref{lines} that contamination from molecular line flux is probably not sufficient to explain this.
+ We suggest the possibility that dust. from the vicinity of IRASLA is being swept up and entrained in the outflow., We suggest the possibility that dust from the vicinity of IRAS4A is being swept up and entrained in the outflow.
+ ‘This suggests models in which the driving mechanism for the outflow is a jet embedded in the protostellar core itself (see e.g. Masson Chernin 1993: Chernin et al 1994: Raga Cabrit. 1993)-," This suggests models in which the driving mechanism for the outflow is a jet embedded in the protostellar core itself (see e.g. Masson Chernin 1993; Chernin et al 1994; Raga Cabrit, 1993)."
+ AM£ this explanation is correct. the main issues for outllow theories would seem to be that the driving mechanism should not be so violent that dust &rains. are destroved in larec numbers. and that significant amounts of dust material should occupy the outflow cavity after the main jet working surfaces have passed through.," If this explanation is correct, the main issues for outflow theories would seem to be that the driving mechanism should not be so violent that dust grains are destroyed in large numbers, and that significant amounts of dust material should occupy the outflow cavity after the main jet working surfaces have passed through."
+ The JCMT is operated by the Joint Astronomy Centre. on behalf of the Wis Particle Physies and Astronomy Research Council. the Netherlands Organization lor Pure Rescarch. and the National Research Council of Canada.," The JCMT is operated by the Joint Astronomy Centre, on behalf of the UK Particle Physics and Astronomy Research Council, the Netherlands Organization for Pure Research, and the National Research Council of Canada."
+ The authors thank James Deane and Jane Greaves for helpful ciseussions ancl advice., The authors thank James Deane and Jane Greaves for helpful discussions and advice.
+of the hiehly reddened CA. ~ 8.18.9) CO stu. DCls. is. displaved in Fig. l..,"of the highly reddened $A_{\rm v}$ $\sim$ 8.4–8.9) G9 star, DC48, is displayed in Fig. \ref{zams}."
+ We tested all of the TTS for Nav. variability using the methods described in Wambarvanetal.(1999)., We tested all of the TTS for X-ray variability using the methods described in \cite{ham99}.
+. The oulv E Tauri Star which showed N-vav variability is the newly identified star RXJ0255.5|2005 that was detected both in the RASS and in theROSAT pointed observation and flared diving the pointed observation (see the light curve displayed in Fig. 6))., The only T Tauri Star which showed X-ray variability is the newly identified star RXJ0255.5+2005 that was detected both in the RASS and in the pointed observation and flared during the pointed observation (see the light curve displayed in Fig. \ref{lc}) ).
+ The peak N-rav count rate duiug the flare increased by ore than a factor of 6 from the pre-flare count rate., The peak X-ray count rate during the flare increased by more than a factor of 6 from the pre-flare count rate.
+ Although: we do no have, Although we do no have
+The results of the above procedure applied to our four quasar fields are presented in Table 10.,The results of the above procedure applied to our four quasar fields are presented in Table 10.
+ In rows 1-2 we list the R.A. anc Deel., In rows 1-2 we list the R.A. and Decl.
+ corrections to our measured. PAL to account for the rotation of the plane of the LMC. aud in rows 3-1. the correspondiug corrected PAL values. in equatorial coordinates. as viewed by an observer located at the center of the LAIC.," corrections to our measured PM to account for the rotation of the plane of the LMC, and in rows 3-4, the corresponding corrected PM values, in equatorial coordinates, as viewed by an observer located at the center of the LMC."
+ Ln rows 5-8 we give calculated PAL values relative to the GRE. both in equatorial aud. galactic coordinates.," In rows 5-8 we give calculated PM values relative to the GRF, both in equatorial and galactic coordinates."
+ These values correspond. to the LMC's PAL as seen by an observer located at the Sun. with the contributious to the PM. from the peculiar solar motion aud from the LSR’s motion. removecl.," These values correspond to the LMC's PM as seen by an observer located at the Sun, with the contributions to the PM, from the peculiar solar motion and from the LSR's motion, removed."
+ In rows 9-11 we give the IL. O aud Z components of the space velocity in a rectangular cartesian coordinate system centered ou the LMC (as cefinecl by Schweitzer et al.," In rows 9-11 we give the $\Pi$ , $\Theta$ and $Z$ components of the space velocity in a rectangular cartesian coordinate system centered on the LMC (as defined by Schweitzer et al.,"
+ 1995. for the Sculptor dSphli).," 1995, for the Sculptor dSph)."
+ The Ll component is parallel to the projection outo the Galactic plane of the radius vector rou the center of the Galaxy to the center of the LNIC. aud is positive when it poiuts racdially away form the Galactic center.," The $\Pi$ component is parallel to the projection onto the Galactic plane of the radius vector from the center of the Galaxy to the center of the LMC, and is positive when it points radially away form the Galactic center."
+ The Θ component is perpeucicular to the LL component. parallel to he Galactic plane. aud poiuts iu the direction of rotation of the Galactic disk.," The $\Theta$ component is perpendicular to the $\Pi$ component, parallel to the Galactic plane, and points in the direction of rotation of the Galactic disk."
+ The Z component »oiuts in the direction of the Galactic uorth pole., The $Z$ component points in the direction of the Galactic north pole.
+ These three componeuts are [ree from the Suns »eculiar motion aud LSR motion., These three components are free from the Sun's peculiar motion and LSR motion.
+" ln rows 12-13 we give the LNICs radial aud transverse space velocities. as seen by an hypothetical observer located at the center of the Galaxy. aud at rest witli 'espect to the Calactic All of the above calculations were carried out asstuning a distance of 50.1 kpe of the LMC from the Sun. a distance of 8.5 kpe of the Sun from the Galactic center. a 220 kins | circular velocity oL the LSR aud a peculiar velocity of the Sun relative to the LSR of (i. vow.) = (—10.5.25. 7.17) km | (Dehnen Binney 1998). These components are positive if u. pointsracially away [roin the Galactic center. v... points in the direction of Galactic rotation aud w.. is directed towards the Galactic north Although the matter was not addressed here. the values presented in table 10 can be used to determine the orbit of the LMC aud therefore study possible past aud future interactions of the LMC with other Local Group If we assume that the LMC is eravitationally bound to. aud in an elliptical orbit. arouud the Galaxy. aud that the mass of the Galaxy is coutaied within 50 kpc of the galactic center. we cau iuake an estimate of the lower limit of its mass through the expression: where r, is the Γλ1Οs apogalacticon distance aud rige dts present distauce."," In rows 12-13 we give the LMC's radial and transverse space velocities, as seen by an hypothetical observer located at the center of the Galaxy, and at rest with respect to the Galactic All of the above calculations were carried out assuming a distance of 50.1 kpc of the LMC from the Sun, a distance of 8.5 kpc of the Sun from the Galactic center, a 220 km $^{-1}$ circular velocity of the LSR and a peculiar velocity of the Sun relative to the LSR of $_{\sun}$ $_{\sun}$ $_{\sun}$ ) = $-$ 10,5.25,7.17) km $^{-1}$ (Dehnen Binney 1998), These components are positive if $_{\sun}$ pointsradially away from the Galactic center, $_{\sun}$ points in the direction of Galactic rotation and $_{\sun}$ is directed towards the Galactic north Although the matter was not addressed here, the values presented in table 10 can be used to determine the orbit of the LMC and therefore study possible past and future interactions of the LMC with other Local Group If we assume that the LMC is gravitationally bound to, and in an elliptical orbit, around the Galaxy, and that the mass of the Galaxy is contained within 50 kpc of the galactic center, we can make an estimate of the lower limit of its mass through the expression: where $_{\rm a}$ is the LMC's apogalacticon distance and $_{\rm LMC}$ its present distance."
+" Forr, = 300 kpc (Lin et al.", For$_{\rm a}$ = 300 kpc (Lin et al.
+ 1995) we obtain Me: values of : (8.2 = 1.3). (9.04 1.6). (3.0 + ," 1995) we obtain $\rm M_{G}$ values of : (8.2 $\pm$ 1.3), $\pm$ 1.6), (3.0 $\pm$ "
+we will address the case of rotating winds which iav be applicable to classical Be stars as these stars do secu to have simall 1ias-lIoss rates.,we will address the case of rotating winds which may be applicable to classical Be stars as these stars do seem to have small mass-loss rates.
+ JMDP is supported by a PPARC postdoctoral research assistantship., JMP is supported by a PPARC postdoctoral research assistantship.
+ The referee. Dr. S.P. Osvocki is thauked for constructive suggestionsOO which improved this paper.," The referee, Dr. S.P. Owocki is thanked for constructive suggestions which improved this paper."
+Sagittarius2.. A (Ser. A) is. an extremely bright. radio. source situated. at theCentre.,Sagittarius A (Sgr A) is an extremely bright radio source situated at the.
+" At its. core is- Ser . i.a region. which. harbours a supermassive. black hole with. a mass o[23 10"" ML Genzel 2000)."," At its core is Sgr $^*$, a region which harbours a supermassive black hole with a mass of 2–3 $\times 10^6$ $_{\sun}$ Genzel 2000)."
+ Many reviews have acdressed the variety of structures. which are apparent on a wide. range of spatial. scales. and the complex web of interactions which characterise this region Yusel-Zadeh 2000: FalekeFe 1999: Mezger.Mezeer. Duschl: Zvlkalk 1996).," Many reviews have addressed the variety of structures, which are apparent on a wide range of spatial scales, and the complex web of interactions which characterise this region Yusef-Zadeh 2000; Falcke 1999; Mezger, Duschl Zylka 1996)."
+ sThe Ser: A complex consists. of ⋅↴Ser A West⇁ and, The Sgr A complex consists of Sgr A West and.
+ Ser ;XWest: includess Ser A‘. XI.a threc-armii spiral-likespir: structure (the mini-spiral) in orbit around Ser A. and the central star cluster (H5 16).," Sgr AWest includes Sgr $^*$, a three-arm spiral-like structure (the mini-spiral) in orbit around Sgr $^*$, and the central star cluster (IRS 16)."
+ On the plane of the sky. the radio source encompasses Ser A West and has à non-thermal shell-like structure (Ekersetal.1975).," On the plane of the sky, the radio source encompasses Sgr A West and has a non-thermal shell-like structure \cite{Ekers1975}."
+ This morphology has been explained in terms of a supernova remnant (SNR) SNR COO10.0: Jones 1974: Ekers 1983: Creen 20013., This morphology has been explained in terms of a supernova remnant (SNR) SNR G0.0+0.0; Jones 1974; Ekers 1983; Green 2001).
+ Llowever. alternative and niore exotic interpretations have also been. proposed. this is the remnant of an outllow triggered by an explosion in Ser A*.," However, alternative and more exotic interpretations have also been proposed, this is the remnant of an outflow triggered by an explosion in Sgr $^*$."
+ ]t is certainly the case that if has a direct physical link to Sgr X West and Ser X. then its study is particularly interesting. in. the context of. past activityun of the central supermassive. black hole.," It is certainly the case that if has a direct physical link to Sgr A West and Sgr $^*$, then its study is particularly interesting in the context of past activity of the central supermassive black hole."
+ ;But equally well. if.n is simply an SNR. then its study should lead us to a better unclerstancling of the special interstellar environment of theCentre region.," But equally well, if is simply an SNR, then its study should lead us to a better understanding of the special interstellar environment of the region."
+ Phe non-thermal shell of is elongated nearly parallel to the Galactic plane with an overall size scale of 275 6 pew inLor a 8.0 kpe distance)., The non-thermal shell of is elongated nearly parallel to the Galactic plane with an overall size scale of $\times$ 5 $\times$ 6 $^2$ for a 8.0 kpc distance).
+| Lt is surrounded by ⇁⊳a dust ringTne (AlezgerwoopοἱOfal.+1980).OKC and.κ in.: projection.went] overlaps with the giant molecular clouds 0.07 (the 50 km  molecular cloud: Serabvn. Lacy Achtermann 1992) and 0.08 (the 20 kms P cloud: Mezger 1986).," It is surrounded by a dust ring \cite{Mezger1989} and, in projection, overlaps with the giant molecular clouds $-$ $-$ 0.07 (the 50 km $^{-1}$ molecular cloud; Serabyn, Lacy Achtermann 1992) and $-$ $-$ 0.08 (the `20 km $^{-1}$ ' cloud; Mezger 1986)."
+ Phe morphology. of the dust ring and the cloud 0.07. as well as the detection of coincident OL maser emission. (Yusel-Zadeh 1996. 1999). stronely suggests that the non-thermal shell physically interacts with the dust ring and. cloud.," The morphology of the dust ring and the cloud $-$ $-$ 0.07, as well as the detection of coincident OH maser emission (Yusef-Zadeh 1996, 1999), strongly suggests that the non-thermal shell physically interacts with the dust ring and cloud."
+ Dased on this observational result. Yusef-Zadeh Morris (1987) proposed a model in which a supernova occurred inside the molecular cloud. ancl created the shell.," Based on this observational result, Yusef-Zadeh Morris (1987) proposed a model in which a supernova occurred inside the molecular cloud and created the shell."
+ More recently Yusef-Zadeh (2000) have outlined a picture in which Ser A West is embedded in the frontmost. region, More recently Yusef-Zadeh (2000) have outlined a picture in which Sgr A West is embedded in the frontmost region
+The synthetic spectrum from slit position 6 reproduces the rec-shiftecl velocity ellipse. which is present in the sspectrum.,"The synthetic spectrum from slit position 6 reproduces the red-shifted velocity ellipse, which is present in the spectrum."
+ The svnthetic spectrum. extracted: from. slit position S successfully reproduces. the incomplete. blue-shifted velocity cllipse., The synthetic spectrum extracted from slit position 8 successfully reproduces the incomplete blue-shifted velocity ellipse.
+ In general. the svnthetic spectra extracted. from. the model viewed at 15 (Fig. 5))," In general, the synthetic spectra extracted from the model viewed at $^{\circ}$ (Fig. \ref{modelcomp}) )"
+ resemble the observed velocity trends in the NEP longslit spectra more closely than the 107 model., resemble the observed velocity trends in the NTT longslit spectra more closely than the $^{\circ}$ model.
+ At this inclination. the svnthetie spectrum from slit position 3 rellects the tilted structure observed in the sspectrum and the svnthetie spectrum from slit. position 4 demonstrates a greater olfset between the bright red- and bluc-shifted filaments. which is in close agreement with the sspectrum.," At this inclination, the synthetic spectrum from slit position 3 reflects the tilted structure observed in the spectrum and the synthetic spectrum from slit position 4 demonstrates a greater offset between the bright red- and blue-shifted filaments, which is in close agreement with the spectrum."
+ Only the synthetic spectrum from slit. position 6 fails to convincingly model the observed spectrum at an inclination of 15: the synthetic spectrum does not form a complete rec-shiftect velocity ellipse as seen in the NTT spectrum., Only the synthetic spectrum from slit position 6 fails to convincingly model the observed spectrum at an inclination of $^{\circ}$; the synthetic spectrum does not form a complete red-shifted velocity ellipse as seen in the NTT spectrum.
+ Overall. both ancl models satisfactorily reproduce the observed. images and. spectra. with the nmiocel providing a better fit to the spectroscopy and the model a better fit to the imagery.," Overall, both and models satisfactorily reproduce the observed images and spectra, with the model providing a better fit to the spectroscopy and the model a better fit to the imagery."
+ We. therefore. favour a nebular inclination between15.," We, therefore, favour a nebular inclination between."
+ Comparison of both models with the observations shows no evidence for deviation from homologous expansion (Llubble-type How) nor for anv additional turbulent broadening in the nebular shell., Comparison of both models with the observations shows no evidence for deviation from homologous expansion (Hubble-type flow) nor for any additional turbulent broadening in the nebular shell.
+ The modelling clearly proves that Sp 1 exhibits an hourglass-like morphology with a well-defined. waist - no other morphology. would. be consistent with both the spectroscopy and imagery. presented here., The modelling clearly proves that Sp 1 exhibits an hourglass-like morphology with a well-defined waist - no other morphology would be consistent with both the spectroscopy and imagery presented here.
+ The racial velocities produced by both the aand sspatio-kinematical models accurately rellect those observed in the longslit spectra., The radial velocities produced by both the and spatio-kinematical models accurately reflect those observed in the longslit spectra.
+ This suggests that the adopted maximum expansion velocities of 55 for the ellipsoids and 20. for the faint. surrounding sphere are νον σου approximations.," This suggests that the adopted maximum expansion velocities of 55 for the ellipsoids and 20 for the faint, surrounding sphere are very good approximations."
+ “Phe heliocentric systemic velocity of he model PN is. IS+5Lo somewhat at odds with he value of 3143 ddetermined. by 2...," The heliocentric systemic velocity of the model PN is $-18 \pm 5$, somewhat at odds with the value of $-31\pm3$ determined by \citet{meatheringham88}."
+ No reasonable explanation could. be ound for this discrepancy - spectra. from slits 3 ane 5. which cross the nebular centre. were collapsed ancl then itted with a gaussian profile following the method. of ? determining a value consistent with that taken from the spatio-kinematical modelling.," No reasonable explanation could be found for this discrepancy - spectra from slits 3 and 5, which cross the nebular centre, were collapsed and then fitted with a gaussian profile following the method of \citet{meatheringham88} determining a value consistent with that taken from the spatio-kinematical modelling."
+ To ensure that this was not an instrumental issue. the svstemic velocity was confirmed using two UCLISS spectra (Section 2.2)). again following the same method used by 2..," To ensure that this was not an instrumental issue, the systemic velocity was confirmed using two UCLES spectra (Section \ref{sec:AAT}) ), again following the same method used by \citet{meatheringham88}."
+ The angular extent of the svnthetic spectra is consistent with the observed. spectra. which suggests the adopted ecometrical dimensions in the model are also very good approximations.," The angular extent of the synthetic spectra is consistent with the observed spectra, which suggests the adopted geometrical dimensions in the model are also very good approximations."
+ The angular distance between the ends of both lobes is  ((Fig., The angular distance between the ends of both lobes is $\sim$ (Fig.
+ 4bb)., \ref{ellipse_model}b b).
+ Adopting a distance to Sp 1 of 1.5 kpe (?) gives a physical distance between the ends of both lobes of 0.5 parsees and a kinematical age of ~STOO vears., Adopting a distance to Sp 1 of 1.5 kpc \citep{sabbadin86} gives a physical distance between the ends of both lobes of 0.5 parsecs and a kinematical age of $\sim$ 8700 years.
+ Ligh spatial and spectral resolution long-slit sspectra have been obtained. from the Fine Ring Nebula. Sp 1l.," High spatial and spectral resolution long-slit spectra have been obtained from the Fine Ring Nebula, Sp 1."
+ These spectra. together with deep narrow-band imagery. have been used to derive a spatio-kinematic mocel of the nebula proving its. bipolar nature.," These spectra, together with deep narrow-band imagery, have been used to derive a spatio-kinematic model of the nebula proving its bipolar nature."
+" The spatio- model of Sp 1 fits with the classical ""butterfly? morphology for a PN defined in the classification scheme of ?..", The spatio-kinematical model of Sp 1 fits with the classical `butterfly' morphology for a PN defined in the classification scheme of \citet{balick87}.
+ The svmmetry axis of the nebula is inclined. almost iong the line of sight (10727:157)., The symmetry axis of the nebula is inclined almost along the line of sight $10\degr\geq i \geq 15\degr$ ).
+ A Hubble-tvpe How is assumed. with an equatorial expansion velocity of ~25 JL. consistent with typical PNe expansion. velocities.," A Hubble-type flow is assumed with an equatorial expansion velocity of $\sim$ 25, consistent with typical PNe expansion velocities."
+ Vhe heliocentric velocity of the PN was found to be 155+., The heliocentric velocity of the PN was found to be $-18 \pm 5$.
+ The kinematical age of Sp 1. at a cistance of 1.5 kpe (ο). is found to be “S700 vears - well within the range of typical PN ages.," The kinematical age of Sp 1, at a distance of 1.5 kpc \citep{sabbadin86}, is found to be $\sim$ 8700 years - well within the range of typical PN ages."
+ Sp 1. does exhibit an exceptional morphology amongst the known sample of PNe with close-binary central stars., Sp 1 does exhibit an exceptional morphology amongst the known sample of PNe with close-binary central stars.
+ “Phe bipolar structure of the main nebula is fairly prevalent amongst other PNe with binary central stars that have also been subjected to. detailed: spatio-kinematical modelling (??)..," The bipolar structure of the main nebula is fairly prevalent amongst other PNe with binary central stars that have also been subjected to detailed spatio-kinematical modelling \citep{jones11a,jones11b}."
+ Phe faint sphere of material which surrounds the bipolar shell. referred to in the text as a “halo”. does not encompass the entire inner nebula. appearing mainlv outside the waist of the bipolar shell.," The faint sphere of material which surrounds the bipolar shell, referred to in the text as a “halo”, does not encompass the entire inner nebula, appearing mainly outside the waist of the bipolar shell."
+ It is. therefore. possible that the “halo” is more akin to an equatorial enhancement (such as the one in seen in Lar 4. 7)) than to the extended haloes found in some PNe (?)..," It is, therefore, possible that the “halo” is more akin to an equatorial enhancement (such as the one in seen in HaTr 4, \citealp{tyndall11b}) ) than to the extended haloes found in some PNe \citep{corradi03}."
+" Equatorial enhancements have been identified as prevalent in the morphologies of PNe with close-binary central stars. which could be indicative that the formation of the ""halo"" in Sp Lis in some way connected with the shedcding of the CE (as is suspected. for more markedly toroidal structures. 7) 2... ?.. ?. ? ?22))"," Equatorial enhancements have been identified as prevalent in the morphologies of PNe with close-binary central stars, which could be indicative that the formation of the “halo” in Sp 1 is in some way connected with the shedding of the CE (as is suspected for more markedly toroidal structures, \citealp{miszalski09b}) \citealp{mitchell07b}, \citealp{jones10b}, \citealp{hillwig10}, \citealp{huckvale11} \citealp{tyndall11a,tyndall11b})"
+" Equatorial enhancements have been identified as prevalent in the morphologies of PNe with close-binary central stars. which could be indicative that the formation of the ""halo"" in Sp Lis in some way connected with the shedcding of the CE (as is suspected. for more markedly toroidal structures. 7) 2... ?.. ?. ? ?22))."," Equatorial enhancements have been identified as prevalent in the morphologies of PNe with close-binary central stars, which could be indicative that the formation of the “halo” in Sp 1 is in some way connected with the shedding of the CE (as is suspected for more markedly toroidal structures, \citealp{miszalski09b}) \citealp{mitchell07b}, \citealp{jones10b}, \citealp{hillwig10}, \citealp{huckvale11} \citealp{tyndall11a,tyndall11b})"
+"highlighting the proximity of the downflows and the chromospheric enhancements, it does not provide information about whether the spatial correspondence between the two endures with time.","highlighting the proximity of the downflows and the chromospheric enhancements, it does not provide information about whether the spatial correspondence between the two endures with time."
+" In order to separate the short lived enhancements from the persistent ones, event maps were constructed from the time sequence of Ca filtergrams acquired the SP scans in the following manner."," In order to separate the short lived enhancements from the persistent ones, event maps were constructed from the time sequence of Ca filtergrams acquired before/during the SP scans in the following manner."
+" First, a small before/duringquiet Sun region was selected to determine the time averaged chromospheric intensity for each of the sequences."," First, a small quiet Sun region was selected to determine the time averaged chromospheric intensity for each of the sequences."
+ This value was then used to normalize the intensity of the individual filtergrams., This value was then used to normalize the intensity of the individual filtergrams.
+ A histogram of the intensities in the AR was derived to determine a suitable threshold value., A histogram of the intensities in the AR was derived to determine a suitable threshold value.
+" From the trailing part of the histogram, threshold values of 0.9, 1.0 and 2.0 were chosen for the three ARs respectively."," From the trailing part of the histogram, threshold values of 0.9, 1.0 and 2.0 were chosen for the three ARs respectively."
+" Using these values, a binary image was created for each individual image in the time sequence, where all pixels with an intensity above the threshold were set to one and the rest to zero."," Using these values, a binary image was created for each individual image in the time sequence, where all pixels with an intensity above the threshold were set to one and the rest to zero."
+" The binary maps were then added in time, yielding at the end, a map with pixels having values indicating the number of chromospheric events."," The binary maps were then added in time, yielding at the end, a map with pixels having values indicating the number of chromospheric events."
+ 'The resulting event maps are shown in Figure 8.., The resulting event maps are shown in Figure \ref{event}.
+ The map constructed for AR. 10923 displays two large patches on the penumbral filaments at the site of the downflows., The map constructed for AR 10923 displays two large patches on the penumbral filaments at the site of the downflows.
+ They are labeled as BE in the top panel of Figure 7.., They are labeled as BE in the top panel of Figure \ref{ca-vel}.
+ The counts at those locations indicate that brightness enhancements persisted for nearly one-third of the 1 hr sequence., The counts at those locations indicate that brightness enhancements persisted for nearly one-third of the 1 hr sequence.
+" There are no obvious signatures of the MJs, reflecting the transient nature and weakness of these events."," There are no obvious signatures of the MJs, reflecting the transient nature and weakness of these events."
+ The strongest and long-lived chromospheric brightenings observed in AR 10953 are confined to a region between the two downflowing patches., The strongest and long-lived chromospheric brightenings observed in AR 10953 are confined to a region between the two downflowing patches.
+ AR 10953 produced several chromospheric enhancements from April 29 to the end of May 1 (Louisetal.2008;Shimizuetal. 2009).," AR 10953 produced several chromospheric enhancements from April 29 to the end of May 1 \citep{Rohan2008,Shimizu2009}."
+. Some of them were co-spatial with supersonic downflows observed in the light bridge nearly 10 hours earlier than the downflows reported here (Louisetal. The 2009).., Some of them were co-spatial with supersonic downflows observed in the light bridge nearly 10 hours earlier than the downflows reported here \citep{Rohan2009}. .
+bottom panel of Figure 8 shows the event map corresponding to AR 11029., The bottom panel of Figure \ref{event} shows the event map corresponding to AR 11029.
+ The Ca filtergrams were acquired during the SP scan with alow cadence of 5, The Ca filtergrams were acquired during the SP scan with alow cadence of 5
+CL and [XP cases. or a modified: non-standard svnchrotron spectrum model should work better.,"CI and KP cases, or a modified non-standard synchrotron spectrum model should work better."
+ In order to explain such disagreement between the predicted. anc observed. X-ray spectrum in AIST jet. knots. several possible solutions to this problem. are proposed.," In order to explain such disagreement between the predicted and observed X-ray spectrum in M87 jet knots, several possible solutions to this problem are proposed."
+ For example. if considering svnchrotron cooling as well as electron. acceleration processes in a non-uniform magnetic field. Bicknell&Beeclman(1996) suggest. it ds. possible to achieve a larger break in spectral index than standard CL model. ic. Ae70.5.," For example, if considering synchrotron cooling as well as electron acceleration processes in a non-uniform magnetic field, \citet{Bicknell96}
+ suggest it is possible to achieve a larger break in spectral index than standard CI model, i.e. $\Delta\alpha>0.5$."
+ With a similar scenario. Honda&Llonda(2007) obtain à modified. N-rav. spectral index consistent with observations of MIST knot A. During the preparation of this paper. we also notice the latest result of Liu&Shen(2007).," With a similar scenario, \citet{Honda07} obtain a modified X-ray spectral index consistent with observations of M87 knot A. During the preparation of this paper, we also notice the latest result of \citet{Liu07}."
+. μον propose a moclified Ο model to explain the racdio-to-N-rav continua in six knots of AIST jet., They propose a modified CI model to explain the radio-to-X-ray continua in six knots of M87 jet.
+ Considering the thin acceleration. region (i.e. shock front) locating at the immediately upstream. of the, Considering the thin acceleration region (i.e. shock front) locating at the immediately upstream of the
+"a representative sample of 63 galaxies with ΛιW1.5x10A, and spatially-resolved kinematics derived from FLAMES/GIRAFFE observations at the VLT.","a representative sample of 63 galaxies with $M_{stellar} \geq 1.5\times
+10^{10} M_\odot$ and spatially-resolved kinematics derived from FLAMES/GIRAFFE observations at the VLT."
+ They found that of z~0.6 intermediate-mass galaxies (1.e.. the progenitors of local spirals). have chaotic velocity fields inconsistent with expectations from pure rotating disks.," They found that of $\sim$ 0.6 intermediate-mass galaxies (i.e., the progenitors of local spirals), have chaotic velocity fields inconsistent with expectations from pure rotating disks."
+ Subsequent analyzes suggest that most of them are likely associated with major mergers (Hammer 2009b)., Subsequent analyzes suggest that most of them are likely associated with major mergers \citep{hammer09b}.
+. Mergers are also found to be a good driver for the large scatter seen in the distant Tully-Fisher relation (Puechetal.2010a:Covingtonetal.2010) .," Mergers are also found to be a good driver for the large scatter seen in the distant Tully-Fisher relation \citep{puech09b,covington10}."
+ Actually. the remarkable co-evolution of the morphology. kinematics. star formation density. metal density. and stellar mass density. all could find a common and natural explanation in the frame of the spiral rebuilding scenario. according to which 50 to of present-day spiral disks were rebuilt after a major merger since ΖΞ1. as proposed by Hammeretal.(200.," Actually, the remarkable co-evolution of the morphology, kinematics, star formation density, metal density, and stellar mass density, all could find a common and natural explanation in the frame of the spiral rebuilding scenario, according to which 50 to of present-day spiral disks were rebuilt after a major merger since z=1, as proposed by \cite{hammer05}."
+M To this respect. the Milky Way appears to be exceptional. w a remarkable quite past history compared to other local spiral galaxies (Hammeretal.2007).," To this respect, the Milky Way appears to be exceptional, with a remarkable quite past history compared to other local spiral galaxies \citep{hammer07}."
+. There is now a growing body of evidence suggesting that disk rebuilding indeed took place at zxl. which makes it as a viable driver for galaxy evolution at these epochs.," There is now a growing body of evidence suggesting that disk rebuilding indeed took place at $\leq$ 1, which makes it as a viable driver for galaxy evolution at these epochs."
+ On the theoretical side. numerical simulations showed how gas can be expelled in tidal tailsduring such mergers. and how it can be subsequently re-accreted (Barnes2002).," On the theoretical side, numerical simulations showed how gas can be expelled in tidal tailsduring such mergers, and how it can be subsequently re-accreted \citep{barnes02}."
+. This re-accreted gas is expected to cool down and form new stars. re-building a new disk around a spheroidal remnant (Springel&Hernquist2005:Robertsonetal.2006).. which might correspond to morphologies as late as Sb galaxies (Lotzetal.2008:Hopkinsetal. 2009c).," This re-accreted gas is expected to cool down and form new stars, re-building a new disk around a spheroidal remnant \citep{springel05, robertson06}, which might correspond to morphologies as late as Sb galaxies \citep{lotz08,hopkins09c}."
+. Recent theoretical developments jive shed light on the underlying process. which appears to be surely gravitational (Hopkinsetal.2009).," Recent theoretical developments have shed light on the underlying process, which appears to be purely gravitational \citep{hopkins09}."
+. The requirement for a major merger to rebuild a new disk depends mainly on the gas fraction during the final coalescence. which needs to be at east (Robertsonetal.2006).," The requirement for a major merger to rebuild a new disk depends mainly on the gas fraction during the final coalescence, which needs to be at least \citep{robertson06}."
+. Cosmological simulations are now also producing such re-processed disks at z«l. although the role of cosmological gas accretion is not totally understood. but aking this process into account could result in a lower gas fraction hreshold for rebuilding a new disk (Governatoetal.2009).," Cosmological simulations are now also producing such re-processed disks at $<$ 1, although the role of cosmological gas accretion is not totally understood, but taking this process into account could result in a lower gas fraction threshold for rebuilding a new disk \citep{governato08}."
+. On he observational side. first examples of rebuilt disks were recently detected at z~0.6 (Puechetal.2009:Hammer2009)... and he auto-consistency of the disk rebuilding process starts being investigated. both theoretically (Hopkinsetal.2009b:al.2009:Hopkinset20€S and observationally (HammeretCompanyetal.2010," On the observational side, first examples of rebuilt disks were recently detected at $\sim$ 0.6 \citep{puech09,hammer09}, and the auto-consistency of the disk rebuilding process starts being investigated, both theoretically \citep{hopkins09b,stewart09,hopkins09c} and observationally \citep{hammer09b,kannappan09,bundy09,huertas10}."
+) If the spiral rebuilding scenario appears to achieve encouraging successes in describing galaxy evolution at zl. some points still need to be investigated.," If the spiral rebuilding scenario appears to achieve encouraging successes in describing galaxy evolution at $<$ 1, some points still need to be investigated."
+ In particular. the impact of the expected numerous mergers on the survival of thin disks is still deb: (Toth&OstrikerMNHopkinsetal.2008:Pur- 2009b).," In particular, the impact of the expected numerous mergers on the survival of thin disks is still debated \citep{toth92,hopkins08,purcell09,moster09b}."
+. Furthermore. Luminous InfraRed Galaxies (LIRGs) account for about of the star formation density reported at z&l (Hammeretal.2005).. but their morphologies reveal that only half of 220.5 LIRGs are compatible with mergers. while in the other half appear to be spiral (Melbourneetdensityal.2005).," Furthermore, Luminous InfraRed Galaxies (LIRGs) account for about of the star formation density reported at $\leq$ 1 \citep{hammer05}, but their morphologies reveal that only half of $>$ 0.5 LIRGs are compatible with mergers, while those in the other half appear to be spiral \citep{melbourne05}."
+.oM thoseLIRGsMarcillacetal.(2006). showed that the large number at these epochs suggests that they could experience between two and four star formation bursts until 2Ξ0. with typical timescales of «0.1 Gyr. which is not consistent with a simply continuous star formation history.," \cite{marcillac06} showed that the large number density of LIRGs at these epochs suggests that they could experience between two and four star formation bursts until z=0, with typical timescales of $\sim$ 0.1 Gyr, which is not consistent with a simply continuous star formation history."
+ They concluded that minor mergers. tidal interactions. or gas accretion remain plausible triggering mechanisms in distant LIRGs harboring a spiral morphology.," They concluded that minor mergers, tidal interactions, or gas accretion remain plausible triggering mechanisms in distant LIRGs harboring a spiral morphology."
+ Interestingly. of local LIRGs are barred. which could play a role in regulating star formation in such objects (Wangetal.2006).," Interestingly, of local LIRGs are barred, which could play a role in regulating star formation in such objects \citep{wang06}."
+. At higher redshifts (.e.. z> 1). the co-moving density of galaxy appears to be dominated by clumpy irregular galaxies (Elmegreenetal.2007).," At higher redshifts (i.e., $>$ 1), the co-moving density of galaxy appears to be dominated by clumpy irregular galaxies \citep{elmegreen07}."
+". Interest for such objects dates back to Cowieetal.(1995)... who first noticed the unusual aspects of some high- galaxies. dubbed as “chain galaxies"" and described as “linearly organized giant star-forming regions""."," Interest for such objects dates back to \cite{cowie95}, who first noticed the unusual aspects of some high-z galaxies, dubbed as “chain galaxies” and described as “linearly organized giant star-forming regions”."
+ The large occurrence of blue star-forming knots in less edge-on objects was also later recognized as a general and intriguing feature of distant galaxies. orobably linked to an early phase in the formation of local spiral galaxies (Cowieetal.1995:vandenBergh 1996).. and were atter referred to as “clump clusters” by Elmegreenetal.(2004).," The large occurrence of blue star-forming knots in less edge-on objects was also later recognized as a general and intriguing feature of distant galaxies, probably linked to an early phase in the formation of local spiral galaxies \citep{cowie95,vandenbergh96}, , and were latter referred to as “clump clusters” by \cite{elmegreen04}."
+. The improved spatial resolution provided by the HST/ACS re-invigorated the interest for these objects. which were all suggested o be different incarnations of the same underlying population viewed along different inclination angles (O'Neiletal.2000:Elmegreenetal. 2004b).," The improved spatial resolution provided by the HST/ACS re-invigorated the interest for these objects, which were all suggested to be different incarnations of the same underlying population viewed along different inclination angles \citep{oneil00,elmegreen04b}."
+". Both kind of objects are therefore often referred to as “clumpy galaxies"".", Both kind of objects are therefore often referred to as “clumpy galaxies”.
+" They are found to be typically made of 5-10 kpe-sized clumps with stellar masses ~107"" M. G.e.. «100 times more massive than the largest star complexes in present-day spiral galaxies). and they typically account for one third of the total galaxy emission (Elmegreen&Elmegreen2005)"," They are found to be typically made of 5-10 kpc-sized clumps with stellar masses $\sim 10^{7-9}$ $_\odot$ (i.e., $\sim$ 100 times more massive than the largest star complexes in present-day spiral galaxies), and they typically account for one third of the total galaxy emission \citep{elmegreen05}."
+ Such clumps are thought to be linked to the formation of disks. as suggested by the increase of the inter-clump surface density. and the decrease of the mass surface density contrast between the clumps and the inter-clump regions. when going from clumpy galaxies with no evident inter-clump emission to clumpy galaxies with faint rec disks. and spiral galaxies (Elmegreenetal.2009b).," Such clumps are thought to be linked to the formation of disks, as suggested by the increase of the inter-clump surface density, and the decrease of the mass surface density contrast between the clumps and the inter-clump regions, when going from clumpy galaxies with no evident inter-clump emission to clumpy galaxies with faint red disks, and spiral galaxies \citep{elmegreen09b}."
+. What is driving the formation of these clumps has been the subject of many attentions during the past decade., What is driving the formation of these clumps has been the subject of many attentions during the past decade.
+ Numerica simulations suggested that clumps might originate from the loca gravitational instability of very gas-rich disks of young galaxies (Noguchi1998:Immelietal.2004).," Numerical simulations suggested that clumps might originate from the local gravitational instability of very gas-rich disks of young galaxies \citep{noguchi98,immeli04}."
+.. Due to their large masses. he clumps would experience strong dynamical friction and spira owards the galaxy center within a few Gyr. which might lead to the ormation of a bulge (Noguchi1999:Elmegreenetal.2008.2009).. as well as a stellar disk through strong stellar scattering (BournaudetMal.t2009b).," Due to their large masses, the clumps would experience strong dynamical friction and spiral towards the galaxy center within a few Gyr, which might lead to the formation of a bulge \citep{noguchi99,elmegreen08,elmegreen09}, as well as a thick stellar disk through strong stellar scattering \citep{bournaud09b}."
+. ονSimulated clumps are found to show woperties similar to (Immelietal.20040:Bournaudetal. 2007).," Simulated clumps are found to show properties similar to observations \citep{immeli04b,bournaud07}."
+. The lifetime of these clumps is so short and the Traction of clumpy galaxies at zl so high. that making the clumpy phase a long-term phenomenon requires a continuous and rapid resh supply of cold gas in order to feed the disk and regenerate new clumps (Dekeletal. 2009b)..," The lifetime of these clumps is so short and the fraction of clumpy galaxies at $>$ 1 so high, that making the clumpy phase a long-term phenomenon requires a continuous and rapid fresh supply of cold gas in order to feed the disk and regenerate new clumps \citep{dekel09b}. ."
+ Theoretical developments indeed suggested that early galaxy formation is fed by cold streams venetrating through dark matter halos (Dekeletal.2009)., Theoretical developments indeed suggested that early galaxy formation is fed by cold streams penetrating through dark matter halos \citep{dekel09}.
+. These cold streams are expected to maintain a dense disk that can undergo gravitational fragmentation into several giant clumps (Dekeletal. 2009b)., These cold streams are expected to maintain a dense disk that can undergo gravitational fragmentation into several giant clumps \citep{dekel09b}.
+. This possible link between high-z clumpy galaxies and the cosmological context was strengthened both by recent cosmological numerical simulations (Agertzetal.2009:Ceverinoetal. 20001. and semi-analytic models (Khochfar&Silk2009).," This possible link between high-z clumpy galaxies and the cosmological context was strengthened both by recent cosmological numerical simulations \citep{agertz09,ceverino09}, and semi-analytic models \citep{khochfar08}."
+. Alternatively. it was proposed that clumps could also result from on-going mergers or interactions (Taniguchi&Sh- ," Alternatively, it was proposed that clumps could also result from on-going mergers or interactions \citep{taniguchi01,overzier08,dimatteo08}. ."
+Discriminating between the merger and fragmentation scenarii is not straightforward because it requires high-resolution integral feld spectroscopy in high-z galaxies., Discriminating between the merger and fragmentation scenarii is not straightforward because it requires high-resolution integral field spectroscopy in high-z galaxies.
+ Indeed. as stated by Noguchi (1999).. the most straightforward and powerful test for discriminating between the two hypothesis is to examine the Kinematicsof the clumps: in the mergerscenario. a random," Indeed, as stated by \cite{noguchi99}, , the most straightforward and powerful test for discriminating between the two hypothesis is to examine the kinematicsof the clumps: in the mergerscenario, a random"
+of narrow dust lanes and emission-line gas in nearby galaxies.,of narrow dust lanes and emission-line gas in nearby galaxies.
+ The technique is based on the Richardson-Lucy (R-L) image restoration algorithm. which uses the PSF to identify structure to enhance as part of the deconvolution process.," The technique is based on the Richardson-Lucy (R-L) image restoration algorithm, which uses the PSF to identify structure to enhance as part of the deconvolution process."
+ Mathematicallv. a structure map 5 is defined as: ∖∖⇁↥∐↲↕⋅≼↲∫↕⋟∖⊽⊔∐↲∪↕⋅↕≸↽↔↴↕↕⋯⊔∐↓≀↕↴≸≟↩⋅∫∍↕⋟∖⊽⊔∐↲↕∐⋟∖⊽⊔⋅∏∐∐↲↕∐↕↴⊳↔⊲⇪⋖⋡≺∢∪∐⋟∖⊽⊔⋅∏≺∢∩↲≼⇂∖∖⇁↕⊔↥⊔∐↲↴⊺↕∐⋡∖↽↴⊺↕∐↓ ⋟∖⊽∩∐∖∖⊽≀↧↴↕⋅≼↲∪↓≯↕∖⊽∏⊳∖⇁↥≪↽∖↽∐∪∪↳↽⊋∩∩⊥⋝⋅⋅∫∍∣↕⊳∖⇁⊔∐↲∏⋅≀↧↴∐⊳∖⇁↕↽≻∪⊳∖⊽≼↲∪↓⋟⊔∐↲↕↴⊳↔⊲⇪≀↧↴∐≼⇂∶↽↕↕⋝∖⊽⊔∐↲≺∢∪∐∖↽∪↥∏∐∪∐ ∪↕↽≻≼↲↕⋅≀↧↴↥∪↕⋅≼⋝⊳∖⇁≼↲≼," Mathematically, a structure map $S$ is defined as: where $I$ is the original image, $P$ is the instrument PSF \citep[constructed 
+with the TinyTim software of][]{krist04}, $P^t$ is the transpose of the PSF and $\otimes$ is the convolution operator \citep[see][]{pogge02}."
+↲↥↴∪≸≟≸≟≼↲≪↽∖↽⇀∖↕≀↧↴↕⋅∐∐↕∃∩∪∃↕⋝⋅⋅ ↴∏∐↲⋟∖⊽⊔⋅∏≺∢⊓∐⋅≼↲∐⋯↕↽≻∩↲≺∢∐∐↕≺⇂⋯↲↕⋅≼↲≺∢∪∖↽≼↲↕⋅⋟∖⊽⋟∖⊽↕∐↓∐≀↧↴↕⋅↕∐↓⋟∪↕⋅∐↓≀↕↴∐∪∐↥∪⊔∐↲∐↓∪↕⋅≼↲⊔⋅≀↧↴≼∐∐∪↕⋯↥∐∐↲⊔↥⋯⇂⋟∖⊽ ∪↓⋟≺∢∪∐⊔⋅≀↧⊔∖⊽↥≼↲∐∐≀↧↴∐≺∢≼↲∐∐↲↕∐⋅⋟∖⊽∏≺∢∐≀↧⊍∖⊽⇂∎↓⊔↕∐≸↽↔↴≀↧↴∐≺⇂⋟∖⊽∏∣↽≻⊔⋅≀↕↴≺∢∐∐≸≟≼↲∐↕↕↽≻∐≺∢≀↧↴↥↕⋟," The structure map technique recovers similar information to the more traditional methods of contrast enhancement, such as fitting and subtracting elliptical isophotes and construction of color maps, but offers several advantages."
+∖⊽∪↕↽≻∐∪∩↲⋟∖⊽≀↧↴∐≼⇂≺∢∪∐⋟∖⊽∏⋅∏≺∢∐∪∐ ∪↓⋟≺∢∪↥∪↕⋅∐↓≀↧↴↕↽≻⋟∖⇁⋅∣↽≻⋯∪∐⋡≼↲↕⋅⋟∖⊽⋟∖⊽≼↲∖↽≼↲↕⋅≀↧↴↥≀↧↴≼⇂∖↽≀↧↴↕∐≀↧↴≸↽↔↴≼↲⋟∖⊽⋅⊟≻↕⋅≼↲⇀↸≀↕↴∐↓↕↽≻↥≼↲⋅⊔∐↲∐⊔↕∐↖≺↽↔↴≀↧↴∐≺⇂⋟∖⇁∏∣↽≻⊔⋅≀↧↴≺∢∐∪∐∪↓⋟ ≼↲∐↕↕↽≻⋟∖⊽≼," For example, the fitting and subtraction of ellipses can lead to artificial features in the presence of strong brightness discontinuities or isophotal twists in the images."
+↲⋟∖⊽≺∢≀↧↴∐↥≼↲≀↧↴≺⇂↥∪≀↧↴↕⋅↥∐∎↓≺∢↕≀↧↴↥↓⋟≼↲≀↧↴��∐⋅≼↲⋟∖⊽↕∐⊔∐↲↕↽≻↕⋅≼↲⋟∖⊽≼↲∐≺∢≼↲∪↓⋟⋟∖⊽∏⋅∪∐↖≺≟∣↽≻↕⋅↕≸↽↔↴↥∐∐≼↲⋟∖⋱∖⊽≼∐⋟," In the case of color maps, there are simply not as many galaxies with observations in two filters as in one."
+∖⊽≺∢∪∐∐∐∏↕∐≼↲⋟∖⊽∪↕⋅ ↕⋝∖⊽∪↕↽≻∐∪↥≀↧↴↥↥∖∖⊽↕⊳∖⊽↥⊳∖⊽↕∐⊔∐↲↥∐⋯≸≟≼↲⋝∖⊽⋅↕∐⊔∐↲≺∢≀↧↪∖⇁≼↲∪↓⋟≺∢∪↥∪↕⋅∐↓≀↧↴↕↽≻⊳∖⊽⋅⊔∐↲↕⋅≼↲≀↧↴↕⋅≼↲⊳∖⊽↕∐↓↕, Possible mismatches between the PSFs in the two bands could also introduce artificial features in color maps.
+↽≻↥∡∖↽∐∪↥≀↧↴⊳∖⊽∐↓≀↧↴∐∡∖↽ ≸↽↔↴≀↧↴↥≀↧↴⇀↸↕≼↲⊳∖⊽∖∖↽↥⊔↥∪∣↽≻⊳∖⊽≼↲↕⋅∖↽≀↧↴∐∪∐⋝∖⊽↕∐↥∖∖↽∪⇂∎↓∐≼↲↕⋅⋝∖⊽≀↧⊔∖⊽↕∐∪∐≼↲⋅↥↴∪⋝∖⊽⊳∖⇁↕∣↽≻↥≼↲↕∐↕⊳∖⇁∐↓≀↧↴∩∢∐≼↲⊳∖⇁∣↽≻≼↲↥∖∖↽≼↲≼↲∐⊔∐↲↥↴⊳↔⊲⊟∖⊽↕∐ ⊔∐↲↥∖∖⇁∪∣↽≻≀↧↴∐≺⇂⋟∖⊽≺∢∪∏∐≀↕⊥∖⊽∪↕∐⊔⋅⋯⇂∏≺∢≼↲≀↕↴↕⋅↥∐∎↓≺∢↕≀↧↴↥↓⋟≼↲≀↕↴⊓∐⋅≼↲⋟∖⊽↕∐≺∢∪↥∪↕⋅∐↓≀↧↴↕↽≻⋟∖⊽⋅↴⊺∐≼↲⋟∖⊽≼↲↕↽≻↕⋅∪∣↽≻↥≼↲∐↓⋟∖⊽≀↧↴↕⋅≼↲≀↧↴∐ avoided with the use of structure maps. although some spurious features can still appear.," These problems are all avoided with the use of structure maps, although some spurious features can still appear."
+ For exaniple. overexposed pixels can produce dark rings that mimic absorption features. while dead pixels can appear as false bright spots in the images.," For example, overexposed pixels can produce dark rings that mimic absorption features, while dead pixels can appear as false bright spots in the images."
+ Nevertheless. such cases were easily identified and did not affect our analysis.," Nevertheless, such cases were easily identified and did not affect our analysis."
+ After constructing structure maps. we then visually inspected each galaxy to determine if dust structures were present. and measured (he spatial extent of anv dust.," After constructing structure maps, we then visually inspected each galaxy to determine if dust structures were present and measured the spatial extent of any dust."
+ We also noted the presence of emission features Chat seem to be nuclear stellar disks (see discussion below) and measured the projected racial extent of these features., We also noted the presence of emission features that seem to be nuclear stellar disks (see discussion below) and measured the projected radial extent of these features.
+ οἱαπο maps for all our galaxies are shown in Figures 3. to 1.., Structure maps for all our galaxies are shown in Figures \ref{fig-stmap1} to \ref{fig-stmap7}.
+ Figure 3. present structure maps of each earlv-tvpe (7x 0). AGN host next to its inactive. control galaxy for the matched sample.," Figure \ref{fig-stmap1} present structure maps of each early-type $T \leq 0$ ), AGN host next to its inactive, control galaxy for the matched sample."
+ Each galaxy pair has the active galaxy on the left and (he inactive ealaxy on (he right., Each galaxy pair has the active galaxy on the left and the inactive galaxy on the right.
+ These (wo figures illustrate the most striking result of this paper. namely that all of the earlv-tvpe active galaxies possess circumnuclear dust while only (7 of 26) of the earlv-tvpe inactive galaxies possess cireumnuclear dust.," These two figures illustrate the most striking result of this paper, namely that all of the early-type active galaxies possess circumnuclear dust while only (7 of 26) of the early-type inactive galaxies possess circumnuclear dust."
+ Figure 2. demonstrates (hat (his result also holds when we include the extended sample. specifically all of the early-tvpe active galaxies have circumnuclear dust and only (9 of 34) of the inactive galaxies in the extended sample have circumnuclear dust.," Figure \ref{fig-stmap4} demonstrates that this result also holds when we include the extended sample, specifically all of the early-type active galaxies have circumnuclear dust and only (9 of 34) of the inactive galaxies in the extended sample have circumnuclear dust."
+ This result also holds if we exclude lenticulars anc, This result also holds if we exclude lenticulars and
+ This result also holds if we exclude lenticulars ancl, This result also holds if we exclude lenticulars and
+(2000a.b) results in (ely.Buse29) = (7.6x10P. 0.35. 2.07. 1.77) for 3=146 and AnaBuneteS)=(02x10P. 145. 2.13. 142) for d—LT for LT<z348. (,"(2000a,b) results in $(A_{\rm HM},B_{\rm HM},z_{c},S)$ = $7.6\times 10^{-13}$, 0.35, 2.07, 1.77) for $\beta=1.46$ and $(A_{\rm HM},B_{\rm HM},z_{c},S)=(3.2\times 10^{-13}$, 1.45, 2.13, 1.42) for $\beta=1.7$ for $1.7 < z < 3.8$. ("
+5) Allowing for a varving equivalent width threshold across each QSO spectrum results in consistently higher values of (59) than are Found [rom the constant threshold treatments.,5) Allowing for a varying equivalent width threshold across each QSO spectrum results in consistently higher values of $J(\nu_{0})$ than are found from the constant threshold treatments.
+ At:>ld. the variable threshold solution is not well-constrained.," At $z > 1$, the variable threshold solution is not well-constrained."
+ Jackknile experiments indicate (hat (his is due the objects 0743-6719 and 0302-2223. namely the highest. ως} absorption lines in each of their spectra. (," Jackknife experiments indicate that this is due the objects 0743-6719 and 0302-2223, namely the highest $\omega(z)$ absorption lines in each of their spectra. ("
+6) Allowing for a cosmology in which (Q4.4)=(0.3.0.7). rather than (1..0.),"6) Allowing for a cosmology in which $(\Omega_{\rm M},\Omega_{\Lambda}) = (0.3,0.7)$, rather than (1.,0.)"
+ has no significant effect on the value of (274) derived from these data. (, has no significant effect on the value of $J(\nu_{0})$ derived from these data. (
+7) The z«1 result is in agreement with the range of values of the mean intensity of the hvdrogen-ionizing background allowed by a variety of local estimates. including [a imaging and modeling of galaxy IHE disk truncations.,"7) The $z < 1$ result is in agreement with the range of values of the mean intensity of the hydrogen-ionizing background allowed by a variety of local estimates, including $\alpha$ imaging and modeling of galaxy HI disk truncations."
+ To within the uneertainty in the measurement. (his result agrees wilh the one previous proximity effect measurement of the low redshift UV background (IXF93).," To within the uncertainty in the measurement, this result agrees with the one previous proximity effect measurement of the low redshift UV background (KF93)."
+ These results are consistent wilh calculated models based upon the integrated emission from QSOs alone (11396) and with models which include both QSOs and starburst galaxies (Shull et 11999)., These results are consistent with calculated models based upon the integrated emission from QSOs alone (HM96) and with models which include both QSOs and starburst galaxies (Shull et 1999).
+ The uncertainties do not make a distinction between these (wo models possible. (, The uncertainties do not make a distinction between these two models possible. (
+8) The results presented here tentatively confir the IGM evolution scenario provided by large scale hydrodynamie simulations (Davé et 11999).,8) The results presented here tentatively confirm the IGM evolution scenario provided by large scale hydrodynamic simulations (Davé et 1999).
+ This scenario. which is successful in describing many observed properlies of the low redshilt IGM. is dependent upon an evolving J(44) which decreases from 2= 2(0z=0.," This scenario, which is successful in describing many observed properties of the low redshift IGM, is dependent upon an evolving $J(\nu_{0})$ which decreases from $z = 2$ to $z = 0$."
+ However. the low redshift UV background required to match the observations of the evolution of the Ly-a. forest line density is larger than found from the data. indicating Chat structure formation is plaving a role in this evolution as well.," However, the low redshift UV background required to match the observations of the evolution of the $\alpha$ forest line density is larger than found from the data, indicating that structure formation is playing a role in this evolution as well."
+ Our results and the work of others are summarized in Figure 16.., Our results and the work of others are summarized in Figure \ref{fig:allzcomp}.
+ We find some evidence of evolution in (19). though it appears that even larger data sets. especially ab z«1 and/or improved proximity effect ionization models will be required to improve the significance.," We find some evidence of evolution in $J(\nu_{0})$, though it appears that even larger data sets, especially at $z < 1$ and/or improved proximity effect ionization models will be required to improve the significance."
+ The authors thank (he anonymous releree for a careful review of the paper and for helpful sugeestions., The authors thank the anonymous referee for a careful review of the paper and for helpful suggestions.
+ This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute. of Technology. under contract with the National Aeronautics aud Space Acdministration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
+ This project was supported bv STScI erant AAR-05785.02-944 and STSel grant GGO 0660601954. JS acknowledges support of the National Science Foundation Graduate Research Fellowship and the Zonta Foundation Amelia Earhart Fellowship., This project was supported by STScI grant AR-05785.02-94A and STScI grant GO 066060195A. JS acknowledges support of the National Science Foundation Graduate Research Fellowship and the Zonta Foundation Amelia Earhart Fellowship.
+ JS. JD. and MM received. financial support [rom NSF grant AST-9617060D. AD acknowledges support from NASA Contract," JS, JB, and MM received financial support from NSF grant AST-9617060B. AD acknowledges support from NASA Contract"
+The photometry of GS 000 variable stars in the LMC was performed. by the OGLE team and recently published by Zeebrun et al. (,The photometry of 68 000 variable stars in the LMC was performed by the OGLE team and recently published by Żeebruń et al. (
+2001).,2001).
+ Unfortunately. cue to the [act that OGLE searches for. microlensing phenomena. the vast majoritv of their cata is collected in. the £-band.," Unfortunately, due to the fact that OGLE searches for microlensing phenomena, the vast majority of their data is collected in the $I$ -band."
+ Adcditionalls. the exposure times are too short to produce very high quality data on RI Lyr variables.," Additionally, the exposure times are too short to produce very high quality data on RR Lyr variables."
+ Phe photometry described. by Zeebruá οἱ al. (, The photometry described by Żeebruń et al. (
+2001) was collected in 21 fields and for most of them the number of V-band frames was arounel 30-40.,2001) was collected in 21 fields and for most of them the number of $V$ -band frames was around 30-40.
+ Only four fields. namely SCR. 8€. SC4 and SC. were observed. over 50 times in V and thus the," Only four fields, namely SC2, SC3, SC4 and SC5, were observed over 50 times in $V$ and thus the"
+ ↕↓↕↓⊰↓⊰∟∙∖⇁↓⋅↿∖∺∠∢⋅⋠⊔∐↓≤⋗⊤≼∩↦↓⊰∐≺∶∢⋅⊔↓∖∺∢↗⊔⇂∪↓⋅⊳∺∠⋖⊾⊀⊔∐⊾∖↽↼∣⊔↓⋅≼⋱∖⊲↓↳ ⇉∪∪⊤⊐↦⇀∖∠≺⊲∙∖⇁⋏∙≟↿∖,"in RR Lyr \citep{rrl}, RR Gem \citep{rrgII}, , XZ Cyg \citep{xzc} and CZ Lac \citep{czl}."
+∟⋜↧≺⊲↓⊔∙∖⇁∠⋖↗⋅∢⊾⇂⋜↧↓⊳⇉∪⊔∃⋜⋯∠⊔⊲∠∟⋯∼↿∖∺∢↗⊔⇂∪↓⋅ ⋖⋅⇂⋜↧↓⊳⇉∪∐∩↦↓⊔⋃∢⊾⊳∖⋖⋅⊳∖⇂⋜⊔⋅⊳∖↿↓↕∢⋅⊳∖↿↓⋅⋖⊾⊔⋏∙≟↿↓↥⊳∖∪⇂⋅↿↓↕∢⊾⋜⋯↓↓≻∐⊓⊔⇂∢⋅ ⋜⋯∠⇂↿↓↥⋖⋅↓≻↓⋯⊳∖∢⋅⊔↓⋯⇂⊔↓⋜∐⊀↓∪⊔⊳∖⋜↧⊳∖∖∖⊽∢⋅⇂⋜↧⊳∖⇂↓↕⋖⋅↓≻⇂⇂↓≻⋜↧⇂↕∢≱↓↥⋜↧↓↥∠⇂ he modulation periods vary., In these stars the strengths of the amplitude and the phase modulations as well as the pulsation and the modulation periods vary.
+ These changes are. probably connected either to multi-periodicity of the modulation or o abrupt and continuous changes of the pulsation period., These changes are probably connected either to multi-periodicity of the modulation or to abrupt and continuous changes of the pulsation period.
+ In Table 2 we summarize the detected changes in the pulsation and the modulation. periods for all those variables. which rave data at least for two dillerent epochs.," In Table \ref{bllong} we summarize the detected changes in the pulsation and the modulation periods for all those variables, which have data at least for two different epochs."
+ We note. however. hat in those cases when the observed. period changes have »en determined Lor several epochs (c.g. XZ Dra. XZ (νο. UV Dra. RY UMa). no strict relation between the pulsation- and the mocdulation-period changes of the variables holds.," We note, however, that in those cases when the observed period changes have been determined for several epochs (e.g. XZ Dra, XZ Cyg, RW Dra, RV UMa), no strict relation between the pulsation- and the modulation-period changes of the variables holds."
+ The connection between the period. changes reflects. some endency rather than an exact relation., The connection between the period changes reflects some tendency rather than an exact relation.
+ Vherelore. results obtained from two-epoch data have to be taken with caution.," Therefore, results obtained from two-epoch data have to be taken with caution."
+ The columns of Table 2. give the pulsation and. the modulation periods. the absolute values of their changes calculated from the periods given in Table 1. or taken from he literature. the normalized relative period variations. and the relative period. and frequcney variations for seven AI5 and six field RRab stars. listed in order of increasing »ulsation period.," The columns of Table \ref{bllong} give the pulsation and the modulation periods, the absolute values of their changes calculated from the periods given in Table \ref{bltab} or taken from the literature, the normalized relative period variations, and the relative period and frequency variations for seven M5 and six field RRab stars, listed in order of increasing pulsation period."
+ IEExamining these data we failed to find any regular pattern. but some general trends are evident.," Examining these data we failed to find any regular pattern, but some general trends are evident."
+ First. it seems that in general. there is an anticorrelation between he directions of the pulsation and the modulation. period changes. Le. à negative sign of c1 is more common han a positive one.," First, it seems that in general, there is an anticorrelation between the directions of the pulsation and the modulation period changes, i.e. a negative sign of ${{\rm d}P_{\mathrm{Bl}}/{\rm d}P_0}$ is more common than a positive one."
+" Secondly. it can be noticed. that the relative period variations (P5 /d4/,) and their normalized values tend to have larger absolute values for variables with onger pulsation periods than for shorter period ones."," Secondly, it can be noticed that the relative period variations ${{\rm d}{P_{\mathrm{Bl}}}/{{\rm d} {P_{0}}}}$ ) and their normalized values tend to have larger absolute values for variables with longer pulsation periods than for shorter period ones."
+ The pulsation-period variation of one of the Blazhko stars. V56. can be described: with a long. periodic/evclic variation.," The pulsation-period variation of one of the Blazhko stars, V56, can be described with a long, periodic/cyclic variation."
+ This gives a unique possibility to test the binary origin of the long-period O—C' variation., This gives a unique possibility to test the binary origin of the long-period $O-C$ variation.
+ Lf this were caused by the light-time elfect the pulsation and modulation periods should have to show similar variation. because NPP is the function of the elements of the orbit for any periodic signal (Coutts1971a).," If this were caused by the light-time effect the pulsation and modulation periods should have to show similar variation, because $\Delta P/P$ is the function of the elements of the orbit for any periodic signal \citep{co71}."
+. The observed. normalised relative period variation of the pulsation and the modulation periods of V56 is. however. much larger than |. it is SLO contracdicting the orbital origin of the evelic. long-period variation.," The observed normalised relative period variation of the pulsation and the modulation periods of V56 is, however, much larger than 1, it is 840 contradicting the orbital origin of the cyclic, long-period variation."
+ However. keeping in mind that random variations in the modulation periods may override the supposed svnchronous variations of the pulsation and the modulation periods. this result is not in fact conclusive.," However, keeping in mind that random variations in the modulation periods may override the supposed synchronous variations of the pulsation and the modulation periods, this result is not in fact conclusive."
+ Although Blazhko stars are usually numerous in. globular clusters. which are rich in RR Lyrae stars. we hardly know anvthing about their properties because of a lack of data suitable. for studying them.," Although Blazhko stars are usually numerous in globular clusters, which are rich in RR Lyrae stars, we hardly know anything about their properties because of a lack of data suitable for studying them."
+ Results have been restricted to deriving the Blazhko periods. but even they have been successfully obtained only in a few cases. e.&. in M3 (Benkóetal. 2006).," Results have been restricted to deriving the Blazhko periods, but even they have been successfully obtained only in a few cases, e.g. in M3 \citep{m3}."
+. In. this paper the Πλας population of M5 are reviewed using all the available photometric data of the variables. in order to identify the Blazhko stars and to study their behaviour.," In this paper the RRab population of M5 are reviewed using all the available photometric data of the variables, in order to identify the Blazhko stars and to study their behaviour."
+ Modulation periods were determined for 13 out of the 20 Blazhko stars that were identified., Modulation periods were determined for 13 out of the 20 Blazhko stars that were identified.
+ The long time base of the observations also made it possible to record changes in the observed. properties of some of the Blazhko stars., The long time base of the observations also made it possible to record changes in the observed properties of some of the Blazhko stars.
+ The most important results of this study can be summarized as follows: Finally. we note that. the possibilities of. studying elobular-cluster. Blazhko stars are still unexploited.," The most important results of this study can be summarized as follows: Finally, we note that the possibilities of studying globular-cluster Blazhko stars are still unexploited."
+ Although these. studies would neec relatively large observational ellorts. the homogeneity of such samples would be helpful for determining what properties of the It Lyrac stars predestine them to become Blazhko variables.," Although these studies would need relatively large observational efforts, the homogeneity of such samples would be helpful for determining what properties of the RR Lyrae stars predestine them to become Blazhko variables."
+ The constructive. helpful comments of the referee. Ixatrien Ixolenberg are much appreciated.," The constructive, helpful comments of the referee, Katrien Kolenberg are much appreciated."
+ Lhe financial support of OTIAX erant Ix-068626 is acknowledged., The financial support of OTKA grant K-068626 is acknowledged.
+ €. Clement thanks the Natural Science and Engineering Council of Canada for financial support., C. Clement thanks the Natural Science and Engineering Council of Canada for financial support.
+ Zs., Zs.
+" H. thanks the ""Lendüllet program ofthe Hungarian Academy of Sciences for supporting his work.", H. thanks the `Lendüllet' program ofthe Hungarian Academy of Sciences for supporting his work.
+where Foo. Fog flew and qnos are the observed. and corrected values of the elfective racius (in arcsec) and. of the mean ellective surface brightness (in magaresec27). respectively. and eb is the major-to-minor axis ratio.,"where $r_{e,o}$ $r_{e,c}$ , $\mu_{e,o}$ and $\mu_{e,c}$ are the observed and corrected values of the effective radius (in arcsec) and of the mean effective surface brightness (in $\rm mag~arcsec^{-2}$ ), respectively, and $a/b$ is the major-to-minor axis ratio."
+" Total magnitudes are corrected for internal extinction as follows (CGavazzi Doselli 1996): where nrp, and nmep, are the observed: ancl corrected values. whatever the Hubble-tvpe is."," Total magnitudes are corrected for internal extinction as follows (Gavazzi Boselli 1996): where $m_{T,o}$ and $m_{T,c}$ are the observed and corrected values, whatever the Hubble-type is."
+ In fact. Gavazzi 3oselli cid not. find. anv morphological dependence of the corrections of total magnitudes for internal extinction in the near-IHi. pass-bands. as opposed to the optical ones.," In fact, Gavazzi Boselli did not find any morphological dependence of the corrections of total magnitudes for internal extinction in the near-IR pass-bands, as opposed to the optical ones."
+ Fig., Fig.
+ 1l shows the distribution of the sample galaxies in a three-dimensional fold-out of. the ποπ &-space., 1 shows the distribution of the sample galaxies in a three-dimensional fold-out of the near-IR $\kappa$ -space.
+ Llere we group the sample galaxies in ellipticals (19) plus lenticulars (SO). SO0a galaxies. anc giant spirals (SaSc) plus late-spirals/la/yBCDs (Scd.lfBCD) and generic irregular/peculiar galaxies.," Here we group the sample galaxies in ellipticals (E) plus lenticulars (S0), S0a galaxies, and giant spirals (Sa–Sc) plus late-spirals/Im/BCDs (Scd–Im/BCD) and generic irregular/peculiar galaxies."
+ I2|80. 50a. SabufBCD galaxies plus generic irregular/peculiar galaxies are represented with filled circles. asterisks ancl empty circles. respectively. in Fig.," E+S0, S0a, Sa–Im/BCD galaxies plus generic irregular/peculiar galaxies are represented with filled circles, asterisks and empty circles, respectively, in Fig."
+ 1., 1.
+ We assume characteristic cumulative uncertainties OW;=ÓWoOsaOL. consistent with typical uncertainties in the observables (ef," We assume characteristic cumulative uncertainties $\delta \kappa_1 = \delta \kappa_2 = \delta \kappa_3 = 0.1$, consistent with typical uncertainties in the observables (cf."
+ Sect., Sect.
+ 3)., 3).
+ The magnitudes of these 10 errors are reproduced in each panel of Fig., The magnitudes of these 1 $\rm \sigma$ errors are reproduced in each panel of Fig.
+ 1., 1.
+ We also reproduce the increase in all the three &s- coordinates due to the increase ofthe dynamical mass when the potential contribution of rotational velocity to the total kinetic energy of the ellipticals is taken into account (cl., We also reproduce the increase in all the three $\kappa$ -space coordinates due to the increase of the dynamical mass when the potential contribution of rotational velocity to the total kinetic energy of the ellipticals is taken into account (cf.
+ Sect., Sect.
+ 2)., 2).
+ Llere we need to say that the impact of the extinction corrections on the distribution of the late-tvpe galaxies in the near-H. &-3pace is small but the fact that the average elfective mass-to-light ratio of these galaxies increases bv when extinction corrections are applied., Here we need to say that the impact of the extinction corrections on the distribution of the late-type galaxies in the near-IR $\kappa$ -space is small but the fact that the average effective mass-to-light ratio of these galaxies increases by when extinction corrections are applied.
+ In the next subsections we analyse the galaxy distribution in each projection of the &-space as a function of morphology., In the next subsections we analyse the galaxy distribution in each projection of the $\kappa$ -space as a function of morphology.
+ Fherefore. in Fig.," Therefore, in Fig."
+ 2. we reproduce the three-dimensional fold-out. of the near-H1t. &-space for these four eroups of Hubble types: We adopt this classification from BBEN (except for the eeneric irregular/peculiar galaxies)," 2, we reproduce the three-dimensional fold-out of the near-IR $\kappa$ -space for these four groups of Hubble types: We adopt this classification from BBFN (except for the generic irregular/peculiar galaxies)."
+ In each panel of Fig., In each panel of Fig.
+" 2. we reproduce the Fundamental Plane relation in the &, a plane (see Sect."," 2, we reproduce the Fundamental Plane relation in the $\kappa_1$ $\kappa_3$ plane (see Sect."
+ 4.1.1). the distribution of the I2]SO galaxies. within £1σ from the mean. inthe &o 5; plane (see Sect.," 4.1.1), the distribution of the E+S0 galaxies, within $\rm \pm 1~\sigma$ from the mean, in the $\kappa_2$ $\kappa_3$ plane (see Sect."
+" 4.2.1) and the borderline of the ""Zone of IExcelusion in the &j πο plane (see Sect.", 4.2.1) and the borderline of the “Zone of Exclusion” in the $\kappa_1$ $\kappa_2$ plane (see Sect.
+ 1.3.1)., 4.3.1).
+ The distribution. of the elliptical and lenticular galaxies in the 5j 5 projection of the &-space confirms that the dvnanücal mass-to-near-Hi lisht ratio increases with the dynamical mass (ef, The distribution of the elliptical and lenticular galaxies in the $\kappa_1$ $\kappa_3$ projection of the $\kappa$ -space confirms that the dynamical mass-to-near-IR light ratio increases with the dynamical mass (cf.
+ PDdC€) as found in the optical (Bender. δισοι Faber 1992: Burstein et al.," PDdC) as found in the optical (Bender, Burstein Faber 1992; Burstein et al."
+ 1995: DDEN)., 1995; BBFN).
+ This result. was expected from the existence of the Fundamental Plane relation in its canonical notation (Pahre. de Carvalho Djorgovski 19908: S98).," This result was expected from the existence of the Fundamental Plane relation in its canonical notation (Pahre, de Carvalho Djorgovski 1998; S98)."
+" A linear fit to the data which minimizes the dispersion in both axes gives: The 4ope of the L-band FP (in &-space notation) that we derive corresponds to the sealing relations Al./L,xALOΟΣΕ απά ALL,LOPLEOS qu ", A linear fit to the data which minimizes the dispersion in both axes gives: The slope of the H-band FP (in $\kappa$ -space notation) that we derive corresponds to the scaling relations $M_e/L_e \propto M_e^{0.296 \pm 0.027}$ and $M_e/L_e \propto L_e^{0.421 \pm 0.055}$.
+(he K-band PDdC obtained. AMμυο ancl AL/L.OxLUCDYZTUQU1S-- '. τ, In the K-band PDdC obtained $M_e/L_e \propto M_e^{0.147 \pm 0.011}$ and $M_e/L_e \propto L_e^{0.172 \pm 0.013}$ .
+οThe originel5 o⋅ the significantenlD cliscrepancyR between these two sets of near-Lt scaling relations is twofold: the incompleteness bias and dillerences in data analysis., The origin of the significant discrepancy between these two sets of near-IR scaling relations is twofold: the incompleteness bias and differences in data analysis.
+ First we note that both these analyses rest on uncomplete samples. so that the derived. sealing relations are biased.," First we note that both these analyses rest on uncomplete samples, so that the derived scaling relations are biased."
+ The PDC sample is more than twice in size than ours. containing 251 elliptical ancl lenticular galaxies. selected wimarily from 13 nearby rich clusters (but. also from loose eroups and the general field). whose distribution in redshift velocity brackets the redshift velocities of Virgo and Coma. he main contributors of objects in our study.," The PDdC sample is more than twice in size than ours, containing 251 elliptical and lenticular galaxies, selected primarily from 13 nearby rich clusters (but also from loose groups and the general field), whose distribution in redshift velocity brackets the redshift velocities of Virgo and Coma, the main contributors of objects in our study."
+ According to he study of the incompleteness bias by Teerikorpi (1987. 1900). we may allow that. in our sample. the fraction of low-luminosity (and. low-mass) objects is skew towards ugher luminosities (ancl masses) than in the PDdC sample. so that a steeper relation between mass-to-light ratio and mass may be expected.," According to the study of the incompleteness bias by Teerikorpi (1987, 1990), we may allow that, in our sample, the fraction of low-luminosity (and low-mass) objects is skew towards higher luminosities (and masses) than in the PDdC sample, so that a steeper relation between mass-to-light ratio and mass may be expected."
+ Nonetheless. svstematic differences in the cata analysis between us anc PDdC€ probably have a larger impact on the obtained sealing relations.," Nonetheless, systematic differences in the data analysis between us and PDdC probably have a larger impact on the obtained scaling relations."
+ While the seeing corrections adopted by us and PDdC are consistent (cf., While the seeing corrections adopted by us and PDdC are consistent (cf.
+ Pahre 1999). the elfective radii and mean surface brightnesses are derived in two dillerent ways.," Pahre 1999), the effective radii and mean surface brightnesses are derived in two different ways."
+ We obtain these effective parameters from. isophotal. elliptical surface photometry (ef," We obtain these effective parameters from isophotal, elliptical surface photometry (cf."
+ Sect., Sect.
+ 3.1). while PDdC determined them from circular aperture photometry. in agreement with previous optical studies.," 3.1), while PDdC determined them from circular aperture photometry, in agreement with previous optical studies."
+ As discussed by Pahre (1999). the isophotal estimate of the elfective radius is. on average. slightlv [larger than its circular. aperture estimate. and. conversely.the isophotalestimate of the elfective intensity is lower than the corresponding circular aperture estimato.," As discussed by Pahre (1999), the isophotal estimate of the effective radius is, on average, slightly larger than its circular aperture estimate, and, conversely,the isophotalestimate of the effective intensity is lower than the corresponding circular aperture estimate."
+ This was already proposed. by SOS in. order to. explain the diserepaney. between the H-band and. the Ix-band. EP, This was already proposed by S98 in order to explain the discrepancy between the H-band and the K-band FP
+2 absorption edges/lines to the power law continuum.,2 absorption edges/lines to the power law continuum.
+" A narrow emission line at 5.59+0.07 keV (rest frame 6.42+0.07 keV) was only marginally significant in the combined data set, with an EW of 754-50 eV and a small improvement in the fit from x? of 269/228 to x? of 264/226."," A narrow emission line at $\pm$ 0.07 keV (rest frame $\pm$ 0.07 keV) was only marginally significant in the combined data set, with an EW of $\pm$ 50 eV and a small improvement in the fit from $\chi^2$ of 269/228 to $\chi^2$ of 264/226."
+" An absorption line at 6.8840.05 keV (rest frame 7.91+0.05 keV), of width σ.»100 eV and EW 360+50 eV, produced a larger improvement, to x? of 236/223."," An absorption line at $\pm$ 0.05 keV (rest frame $\pm$ 0.05 keV), of width $\sigma$$\sim$ 100 eV and EW $\pm$ 50 eV, produced a larger improvement, to $\chi^2$ of 236/223."
+" Finally, an absorption edge at Τ.6--0.2 keV (rest frame 8.7+0.2 keV), optical depth 7 of 0.30.1, further improved the fit to x? of 228/221."," Finally, an absorption edge at $\pm$ 0.2 keV (rest frame $\pm$ 0.2 keV), optical depth $\tau$ of $\pm$ 0.1, further improved the fit to $\chi^2$ of 228/221."
+ We defer further discussion of these features to Section 5.3., We defer further discussion of these features to Section 5.3.
+" The marginal detection of low energy absorption in the original oobservation of 0918+2117M (Wilkes 22002) was a surprise, given that iis a highly reddened object."," The marginal detection of low energy absorption in the original observation of M (Wilkes 2002) was a surprise, given that is a highly reddened object."
+" On the basis of the follow-up oobservations it now seems that the intermediate flux level oobservation may have been confused by the presence of a soft emission component, seen more clearly in the low flux state oobservation in 2003."," On the basis of the follow-up observations it now seems that the intermediate flux level observation may have been confused by the presence of a soft emission component, seen more clearly in the low flux state observation in 2003."
+" We now find, in the 2005 high state spectrum, clear evidence for low energy absorption corresponding to a line-of-sight column Ng-4x10?! ?."," We now find, in the 2005 high state spectrum, clear evidence for low energy absorption corresponding to a line-of-sight column $_{H}$$\sim$ $\times 10^{21}$ $^{-2}$."
+" With the low ionisation parameter and a 'normal dust to gas ratio, that absorbing column is consistent with the additional optical reddening (i.e. above that of an unobscured QSO; Barkhouse and Hall 2001) of J-K;~0.4 seen in 09184-2117M. Our analysis suggests that the spectral change between the 2003 and 2005 oobservations was dominated by a change in the X-ray power law continuum, with absorption in a column of Ny~4x 10?! cm""? of ‘cold’ or weakly ionised gas present in both cases."," With the low ionisation parameter and a `normal' dust to gas ratio, that absorbing column is consistent with the additional optical reddening (i.e. above that of an unobscured QSO; Barkhouse and Hall 2001) of $_{s}$$\sim$ 0.4 seen in M. Our analysis suggests that the spectral change between the 2003 and 2005 observations was dominated by a change in the X-ray power law continuum, with absorption in a column of $_{H}$$\sim$ $\times 10^{21}$ $^{-2}$ of `cold' or weakly ionised gas present in both cases."
+" As noted in W05, inclusion of such absorption in modelling the 2003 sspectrum, in addition to strong cold reflection, would allow the underlying continuum to assume a ‘normal’ photon index, in the range predicted by Comptonisation models (eg Haardt 11997)."," As noted in W05, inclusion of such absorption in modelling the 2003 spectrum, in addition to strong cold reflection, would allow the underlying continuum to assume a `normal' photon index, in the range predicted by Comptonisation models (eg Haardt 1997)."
+" Variability of the primary power law continuum has been identified as the dominant component in X-ray spectral changes in several well studied AGN (Fabian and Vaughan 2003, Pounds 22004)."," Variability of the primary power law continuum has been identified as the dominant component in X-ray spectral changes in several well studied AGN (Fabian and Vaughan 2003, Pounds 2004)."
+ It appears that iis a further example of that situation., It appears that is a further example of that situation.
+" While we have no direct information on the location of the absorbing gas, it is probably not very close to the central source unless in sufficiently dense clouds to avoid evaporation."," While we have no direct information on the location of the absorbing gas, it is probably not very close to the central source unless in sufficiently dense clouds to avoid evaporation."
+" However, the intermediate type 1.5 classification of ssuggests some absorbing matter is located at a radius comparable with the BLR."," However, the intermediate type 1.5 classification of suggests some absorbing matter is located at a radius comparable with the BLR."
+ Future observations to check for variability of the absorbing column and/or ionisation parameter should clarify that issue., Future observations to check for variability of the absorbing column and/or ionisation parameter should clarify that issue.
+ It is interesting to speculate on the nature of the soft X-ray, It is interesting to speculate on the nature of the soft X-ray
+from OB associations.,from OB associations.
+ PDRs exposed to weak FUV radiation fields have been studied in a few sources using i]] emission (e.g.?2?) but never observed with velocity—resolved m]] before Herschel/HIFI.," PDRs exposed to weak FUV radiation fields have been studied in a few sources using ] emission \citep[e.g.][]{Maezawa1999,Bensch06,Pineda07a} but never observed with velocity–resolved ] before Herschel/HIFI."
+" ? suggested that most of the m]] in our Galaxy originates from a clumpy medium exposed to a FUV field of yo 10-5, which is larger than the upper limit determined for the majority of the observed components."," \citet{Cubick2008} suggested that most of the ] in our Galaxy originates from a clumpy medium exposed to a FUV field of $\chi_0=10^{1.8}$ , which is larger than the upper limit determined for the majority of the observed components."
+" Note, however, that their model does not consider emission arising from diffuse clouds."," Note, however, that their model does not consider emission arising from diffuse clouds."
+" From our observed LOSs, we find that about of the total detected m]] emission arises from low-column density regions (without significant 1900 emission; ??)) while is emitted from dense PDRs."," From our observed LOSs, we find that about of the total detected ] emission arises from low–column density regions (without significant $^{13}$ CO emission; \citealt{Langer2010,Velusamy2010}) ) while is emitted from dense PDRs."
+" Nevertheless, the moderate FUV field predicted by ? might suggest that the predominance of low-FUV radiation field regions observed in our limited sample (covering less than of the entire survey) might hold for the entire Galaxy."," Nevertheless, the moderate FUV field predicted by \citet{Cubick2008} might suggest that the predominance of low–FUV radiation field regions observed in our limited sample (covering less than of the entire survey) might hold for the entire Galaxy."
+ We found two regions where our analysis suggests high densities (>105 ccm?) and strong FUV fields (between Xo=10 and 10°).," We found two regions where our analysis suggests high densities $>10^5$ $^{-3}$ ) and strong FUV fields (between $\chi_{\rm 0} =
+10^4$ and $10^6$ )."
+ These regions are likely associated with massive star formation., These regions are likely associated with massive star formation.
+ This conclusion is a result of the elevated u]]/CO ratio observed towards these regions., This conclusion is a result of the elevated ]/CO ratio observed towards these regions.
+ This identification suggests that the n]]/CO ratio is good tracer of the location of massive star formation regions in the galaxy. 1], This identification suggests that the ]/CO ratio is good tracer of the location of massive star formation regions in the galaxy. ]
+] observations will therefore provide an alternative method to determine the distribution of massive star forming regions in the galaxy (e.g.?).., observations will therefore provide an alternative method to determine the distribution of massive star forming regions in the galaxy \citep[e.g.][]{Bronfman2000}.
+ Note that velocity resolved observations are crucial for the interpretation of the m]]/CO ratio., Note that velocity resolved observations are crucial for the interpretation of the ]/CO ratio.
+ In our observed LOSs we have found velocity components showing [Cu]] emission but no CO as well as components showing CO but no rj]., In our observed LOSs we have found velocity components showing ] emission but no CO as well as components showing CO but no ].
+ Velocity unresolved observations would have given a distorted value of the m]]/CO ratio that would result in an incorrect interpretation of the physical conditions of the line-emitting gas., Velocity unresolved observations would have given a distorted value of the ]/CO ratio that would result in an incorrect interpretation of the physical conditions of the line–emitting gas.
+thickuess does not depeud ou the upstream magnetic field streugth.,thickness does not depend on the upstream magnetic field strength.
+ From the examples shown iu Fig., From the examples shown in Fig.
+ 2 and 3.. we can see that although the zeroeth-orcler solution follows the geueral behavior of C shocks. it is not accurate for strougly-1uaguetized cases (Fig. 3)).," \ref{ana_thick1} and \ref{ana_thick2}, we can see that although the zeroeth-order solution follows the general behavior of C shocks, it is not accurate for strongly-magnetized cases (Fig. \ref{ana_thick2}) )."
+ Compared with the dataset of exact solutions discussed in previous section. the RMS value of (Lexact—Leerocih)/Lexact is 0.355. and the range of (Lesaet—Leerveth}/Lexact is —0.8 to 0.28.," Compared with the dataset of exact solutions discussed in previous section, the RMS value of $(L_\mathrm{exact}-L_\mathrm{zeroeth})/L_\mathrm{exact}$ is $0.355$ , and the range of $(L_\mathrm{exact}-L_\mathrm{zeroeth})/L_\mathrm{exact}$ is $-0.8$ to $0.28$."
+ The depeudence ou the velocity. tou density. aud. collision coellicient in Equation (31)) cau be uuderstood in terms of the drag force between ious aud neutrals.," The dependence on the velocity, ion density, and collision coefficient in Equation \ref{thick_ana}) ) can be understood in terms of the drag force between ions and neutrals."
+ The total momentum flux iu ueutrals enteringH the shock isH poco., The total momentum flux in neutrals entering the shock is $\rho_0{v_0}^2$.
+> The mean drag forceH per volume isη ~apopi;gco.," The mean drag force per volume is $\sim\alpha\rho_0\rho_{i,0}v_0$."
+ The ration ofn these quantities. which is the characteristic distance over which momentum excliaugetakes place. is This depenclence is similar to Equation (3.12) in Draine&Mclxee(1993) ifthe Alfvéun speed in the fluid is similar to the upstream velocity. ονey.," The ratio of these quantities, which is the characteristic distance over which momentum exchangetakes place, is This dependence is similar to Equation (3.12) in \cite{1993ARA&A..31..373D} if the Alfvénn speed in the fluid is similar to the upstream velocity, $v_\mathrm{A}\sim v_0$."
+ Although they obtained an estimate usine different assumptions aud approximations. the basic idea that the momentum trausfer rate determines the shock thickuess is similar.," Although they obtained an estimate using different assumptions and approximations, the basic idea that the momentum transfer rate determines the shock thickness is similar."
+" To obtain a more accurate estimate of the C shock thickness. we return to the clillerential equation (9a)) for neutral momentum flux. makine use of Equation (21)) aud the ionization equilibrium condition r;=ry> We integratee this equation. usiugOm constant. values on the rielt-lauc-siclee where the minimum value of r,/rg can be derived explicitly[rom Equation (16)) as This yieldsa quadratic for rj as a fuuction of vr:"," To obtain a more accurate estimate of the C shock thickness, we return to the differential equation \ref{momN}) ) for neutral momentum flux, making use of Equation \ref{constD}) ) and the ionization equilibrium condition $r_i = r_n^{1/2}$ , We integrate this equation, using constant values on the right-hand-side where the minimum value of $r_n/r_B$ can be derived explicitlyfrom Equation \ref{rB}) ) as This yieldsa quadratic for $r_n$ as a function of $x$ :"
+errors.,errors.
+ The 2dEPGIBS has a brighter A by 0.05 mae. and a brighter pj by 0.85 mag ," The 2dFGRS has a brighter $M^*$ by 0.05 mag, and a brighter $\mu_e^*$ by 0.85 mag $^{-2}$."
+‘Taking into account the elfects of bulges and inclination as mentioned above. the 2dGIU distribution has become 0.15 mag brighter than the de Jong Lacey distribution and has a brighter pi by 0.3 mag ο7.," Taking into account the effects of bulges and inclination as mentioned above, the 2dFGRS distribution has become 0.15 mag brighter than the de Jong Lacey distribution and has a brighter $\mu_e^*$ by 0.3 mag $^{-2}$."
+ Considering that [ate-tvpe galaxies tend to be fainter and lower surface brightness. these results appear fully consistent.," Considering that late-type galaxies tend to be fainter and lower surface brightness, these results appear fully consistent."
+" Our value of 3,I can be converted. to a. Iuminositv- ≱∖⋯↓⋖⋅⊳∖⊀↓∠⋖⋅⋏∙≟↓⋅⋯⋠⊓⋅⊔↿⊰∣∣∶∪⋅∶∫≻≼≨∪∶∶∪⋅⋈⊔⊳⇀∖↓↿↓↥∪"," Our value of $\beta_{\mu}$ can be converted to a luminosity-scale size gradient $\beta_{r_e}=-0.360
+\pm 0.004$."
+⊔⋏∙≟↓↕↿↓↥⊲↓⊳∖ ∠⇂⊲↓∐⋅⋖⋅↓⋅⊳∖∐⋅∪⊔↓↿↓∐⋅∠⇂⋖⋅↼∣∪⊔⋏↳⊓∖↽∟⋯∼∢⊾∙∖⇁⊔∪⋈∩∖⇁⋜↧⇂⋯⋅⊳⊀↓⊔∩⋅↓⋅∢⊾⊳∖↿⊀↓⊔⋏∙≟↓∙∖⇁⊳ it agrees more closely. with their theoretical. prediction. of Ei (see de Jong Lacey 2000).," Although this differs from the de Jong Lacey (2000) value, interestingly, it agrees more closely with their theoretical prediction of $\beta_{r_e}=-\frac{1}{3}$ (see de Jong Lacey 2000)."
+" One possible reason for the variation in 3),i nav be a correlation between ≼∙∪⇂∪⊔↓⋅⋜⋯∠⋜↧∣⋡≱∖∪↓⋯∢⊾⊔↓⋜↧⋏∙≟⊔", One possible reason for the variation in $\beta_{\mu}$ may be a correlation between colour and absolute magnitude.
+⊀↓⊓⊔⇂⋖⊾⊳∐↓⋜⋯↿∪⊔∢⊾, Blanton et al. (
+"⇂⋜↧↓⊳⊔↭⊔∐⊔∠⇂ a strong correlation. between. (g*ry colour ancl M, brighter ealaxies are redder. fainter galaxies are blucr.","2001) find a strong correlation between $(g^*-r^*)$ colour and $M_{r^*}$: brighter galaxies are redder, fainter galaxies are bluer."
+ Alaking estimates of the colour-magnitude correlation from Fig., Making estimates of the colour-magnitude correlation from Fig.
+ 13 of Blanton et al. (, 13 of Blanton et al. (
+"2001). we find that Al,»=S7545(go8)αι δες.","2001), we find that $M_{r^*}=-8.75\pm_2^4(g^*-r^*)
+-14.88\pm_2^1$ ."
+ Using a mean b;5=11 (calculated from Fukugita. Shimasaku Ichikawa 1995 and Alacldox. Efstathiou Sutherland 1990). and assuming that the additional colour term. (b;p)=(qqry we calculate that the expected ἐξ0.36+0.23.," Using a mean $b_j-r^*=1.1$ (calculated from Fukugita, Shimasaku Ichikawa 1995 and Maddox, Efstathiou Sutherland 1990), and assuming that the additional colour term $\Delta(b_j-r^*)=\Delta(g^*-r^*)$ we calculate that the expected $\beta_{\mu,r^*}=0.36\pm0.23$."
+ The value estimated. from Fig., The value estimated from Fig.
+ 10 of Blanton et al. (, 10 of Blanton et al. (
+2001) is «των= O.50405. ,"2001) is $\beta_{\mu,r^*}=0.50\pm_{0.1}^{0.2}$ ."
+Thus the +! band: Iuminositv-surface. brightness correlation appears steeper than the b; band. Iuminosity-surface brightness correlation., Thus the $r^*$ band luminosity-surface brightness correlation appears steeper than the $b_j$ band luminosity-surface brightness correlation.
+ A similar colour-magnituce correlation in (b;—£) could explain the discrepancy between our result and de Jong Lacey (2000) result., A similar colour-magnitude correlation in $(b_j-I)$ could explain the discrepancy between our result and de Jong Lacey (2000) result.
+ The general good overall agreement between these substantially cifferent surveys is an important vinelication of both results., The general good overall agreement between these substantially different surveys is an important vindication of both results.
+ Cross et al. (, Cross et al. (
+2001) has extended the de Jong Lacey conclusions to the full range of galaxy types with M€16.,2001) has extended the de Jong Lacey conclusions to the full range of galaxy types with $M<-16$.
+ However. the different values obtained. for the luminosity surface brightness correlation may rellect a colour or morphologically dependent Iuminositv-surface brightness correlation.," However, the different values obtained for the luminosity surface brightness correlation may reflect a colour or morphologically dependent luminosity-surface brightness correlation."
+ Dlanton et al. (, Blanton et al. (
+2001) seem to have found similar results. but have not fitted a [function or tabulated their results.,"2001) seem to have found similar results, but have not fitted a function or tabulated their results."
+ As for the Schechter function it is trivial to calculate the luminosity density. j. by integrating the product. of the BBF and the luminosity over the complete range of surface brightness ancl absolute magnitucle.," As for the Schechter function it is trivial to calculate the luminosity density, j, by integrating the product of the BBF and the luminosity over the complete range of surface brightness and absolute magnitude."
+ ‘The solution is the same as the solution to the integral obtained [rom the Schechter function., The solution is the same as the solution to the integral obtained from the Schechter function.
+" When caleulated using the best fit parameters. the value of the luminosity density. j,,=(2.160.14)-1075EL. Mpcju"," When calculated using the best fit parameters, the value of the luminosity density, $j_{b_j}=(2.16\pm0.14)\times10^8h\,L_{\odot}$ $^{-3}$."
+ In Blanton et al. (, In Blanton et al. (
+2001). the Sloan team get a 40% ugher value for the luminosity density in the 6; filter than he 2dbORS team.,"2001), the Sloan team get a $40\%$ higher value for the luminosity density in the $b_j$ filter than the 2dFGRS team."
+ Does this mean that 2dCUS is missing some galaxies. or at least underestimating their luxes?," Does this mean that 2dFGRS is missing some galaxies, or at least underestimating their fluxes?"
+ For a start. the values of AZ are consistent. suggesting that th surveys are correcting magnitudes properly.," For a start the values of $M^*$ are consistent, suggesting that both surveys are correcting magnitudes properly."
+ However he measurement of ó* is over higher in Blanton ct al. (, However the measurement of $\phi^*$ is over higher in Blanton et al. (
+2001).,2001).
+ X more recent paper (Yasuda ct al., A more recent paper (Yasuda et al.
+ 2001) revises he SDSS luminosity density of Blanton et al.," 2001) revises the SDSS luminosity density of Blanton et al.,"
+ to js;=2.45t0.21107h? 3., to $j_{b_j}=2.43\pm0.21\times10^{-2}h^3$ $^{-3}$.
+ ιν revision is based on a fit o the galaxy number counts. suggesting that the Blanton et al.," This revision is based on a fit to the galaxy number counts, suggesting that the Blanton et al."
+ region was overdense by30%., region was overdense by.
+. The revised ó. value (=:.2.05£0.12oU.4551024875 By 7) is now consistent. with our measurement of ó*=2.060.00«10.24? 7. in the 6; band.," The revised $\phi^*$ value $\phi^*=2.05 \pm 0.12 
+^{+0.66}_{-0.28}\times10^{-2}h^3$ $^{-3}$ ) is now consistent with our measurement of $\phi^*=2.06 \pm 0.09\times10^{-2}h^3$ $^{-3}$, in the $b_j$ band."
+ Given this revised. value of 6°. the Blanton result. is still higher. but Blanton used a colour term b;=B.:L5(D VV). whereas the correct colour term for the ADM. tested using ELS data is b;=DO.28(2V) (Peacock. private communication).," Given this revised value of $\phi^*$, the Blanton result is still higher, but Blanton used a colour term $b_j=B-0.35(B-V)$ , whereas the correct colour term for the APM, tested using EIS data is $b_j=B-0.28(B-V)$ (Peacock, private communication)."
+ When these two factors are taken into account. the luminosity densities are entirely consistent.," When these two factors are taken into account, the luminosity densities are entirely consistent."
+ As demonstrated in Cross ct al. (, As demonstrated in Cross et al. (
+2001) the peak of the luminosity density lies well inside the selection boundaries.,2001) the peak of the luminosity density lies well inside the selection boundaries.
+" When the function is integrated over the range —24ςAlox 15.5. 20.1.«qn.s241. the value. obtained is Jo,=244LhkL. ? as compared to the summed BBD which gives: jp,=(241£0.20).105L. %."," When the function is integrated over the range $-24<M<-15.5$ , $20.1<\mu_e<24.1$, the value obtained is $j_{b_j}=2.14\times10^8h\,L_{\odot}$ $^{-3}$ as compared to the summed BBD which gives: $j_{b_j}=(2.11\pm0.20)\times10^8h\, L_{\odot}$ $^{-3}$."
+ These are the approximate selection. boundaries. so the correspondence is excellent.," These are the approximate selection boundaries, so the correspondence is excellent."
+ Unless the distribution shows an upturn outside the selection boundary the 280ECIUS data have uncovered over of the local B-bancl luminosity. clensitv., Unless the distribution shows an upturn outside the selection boundary the 2dFGRS data have uncovered over of the local B-band luminosity density.
+ In Cross et al. (, In Cross et al. (
+2001). Fig.,"2001), Fig."
+ 1 showed the variation in the LE às measured from a number of recent redshift surveys., 1 showed the variation in the LF as measured from a number of recent redshift surveys.
+ Lt was postulated that the large variation was due to surface brightness selection effects., It was postulated that the large variation was due to surface brightness selection effects.
+ In this section the variation of the luminosity function. with the limiting detection isophote is explored ancl compared to the range of published luminosity functions., In this section the variation of the luminosity function with the limiting detection isophote is explored and compared to the range of published luminosity functions.
+ Phroughout. an exponential surface brightness profile and an Einstein-de Sitter cosmology are assumed.," Throughout, an exponential surface brightness profile and an Einstein-de Sitter cosmology are assumed."
+ ‘To caleulate a derived luminosity function. we start with our fit to the BBE.," To calculate a derived luminosity function, we start with our fit to the BBF."
+ We take into account the overestimate in? by 0.55 mag 7. derivedin Appendix A. and use a value of po=22.45 mag 2> Note that in Cross et al. (," We take into account the overestimate in $\mu_e^*$ by 0.55 mag $^{-2}$ , derivedin Appendix A, and use a value of $\mu_e^*=22.45$ mag $^{-2}$ [Note that in Cross et al. ("
+2001) we used a less sophisticated: method to estimate,2001) we used a less sophisticated method to estimate
+" Typ, < 109 (?)..", $_{eff}$ $<$ $^6$ \citep{paquette86}.
+ 2107 material(???)..," $\gg$ $^{8}$ \citep{dupuis93b,dupuis92,dupuis93a}."
+ <70 problematic(???).. (??).. (2).. binaries(?)..," $<$ \citep{aannestad93,zuckerman98,zuckerman03}. \citep{koester05,koester06}. \citep{alcock86}. \citep{zuckerman87,koester97,holberg97,debes02,gianninas04}. \citep{zuckerman03}."
+By including the effect of the external electric field. the mean value of the LPAI Irequencey can be expressed as In (he absence of the electric field the dependence of the LPAI critical frequency on the surface potential of the quark star is (eppay)V. while in the presence of the electric field we have a logarithmic dependence on (he quark star potential.,"By including the effect of the external electric field, the mean value of the LPM frequency can be expressed as In the absence of the electric field the dependence of the LPM critical frequency on the surface potential of the quark star is $\left\langle \omega _{LPM}\right\rangle \sim V_{q}^{4}$, while in the presence of the electric field we have a logarithmic dependence on the quark star potential."
+ The dependence of (he mean LPAI critical frequency (eypay) on the quark star surface potential V;. in the low temperature limit and with ancl without the presence of an external electric field is represented in Fig.," The dependence of the mean LPM critical frequency $\left\langle \omega _{LPM}\right\rangle $ on the quark star surface potential $V_q$, in the low temperature limit and with and without the presence of an external electric field is represented in Fig."
+ 4., 4.
+ In order to calculate the spectral distribution of the intensity of the electromagnetic radiation emitted by an electron in the dense laver of the electrosphere of a quark star. by taking into account the effect of the multiple collisions and of the electric field. we use Eqs. (17—-21)).," In order to calculate the spectral distribution of the intensity of the electromagnetic radiation emitted by an electron in the dense layer of the electrosphere of a quark star, by taking into account the effect of the multiple collisions and of the electric field, we use Eqs. \ref{fin1}- \ref{fin2}) )."
+ The variation of the probability of emission of a photon by an electron moving in the electrosphere of the quark star is represented. for two sets of values of the quark star surface electric potential. in Fie.," The variation of the probability of emission of a photon by an electron moving in the electrosphere of the quark star is represented, for two sets of values of the quark star surface electric potential, in Fig."
+ 5., 5.
+ The mean value of the total intensity of the radiation. (Z(4)) is represented in Fig.," The mean value of the total intensity of the radiation, $\left\langle I(\omega )\right\rangle $ is represented in Fig."
+ 6., 6.
+" An important parameter describing the radiation properties of an electron gas is the plasma frequency «jy. defined as wy,=y4zenc(z.T)/p(2.T). and which is related to the medium permitüvityv 2(w) bv (ae)£21—wfar. 122ey (Jackson1999).."," An important parameter describing the radiation properties of an electron gas is the plasma frequency $\omega _p$, defined as $\omega _{p}=\sqrt{4\pi en_{e}(z,T)/\mu _{e}(z,T)}$, and which is related to the medium permittivity $\varepsilon \left( \omega
+\right) $ by $\varepsilon \left( \omega \right) \approx 1-\omega
+_{p}^{2}/\omega ^{2}$, $\omega >>\omega _{p}$ \citep{Ja75}."
+ The polarization of the electrosphere of the quark stars also strongly influences the radiation processes. and photons with frequency w<oy cannot escape from the electrosphere.," The polarization of the electrosphere of the quark stars also strongly influences the radiation processes, and photons with frequency $\omega <\omega _{p}$ cannot escape from the electrosphere."
+ The variation of the plasma [recuencey of the quark star electrosphere is represented. as à function of distance. in Figs.," The variation of the plasma frequency of the quark star electrosphere is represented, as a function of distance, in Figs."
+ f., 7.
+ The variation of the emissivity of the electrosphere of the quark star. by fully taking into the effect of the electric field on the suppression of the radiation due to the multiple scaltering effects. is presented in Figs.," The variation of the emissivity of the electrosphere of the quark star, by fully taking into the effect of the electric field on the suppression of the radiation due to the multiple scattering effects, is presented in Figs."
+ 8., 8.
+ :The total energv flux of. the electrosphere is: obtained: as FLJ=(1/z)5JiUXOd: ," The total energy flux of the electrosphere is obtained as $F_{\gamma }^{(ee)}=(1/\pi )\int_{0}^{\infty }\varepsilon _{\gamma
+}^{(ee)}dz $."
+:The variation of (he energy flux as a function of the temperature of the quark star surface is presented. for different values of the surface electric potential Vj. in Fig.," The variation of the energy flux as a function of the temperature of the quark star surface is presented, for different values of the surface electric potential $V_{I}$, in Fig."
+ 9., 9.
+ since the proposal of strange quark stars. much effort has been devoted to find major," Since the proposal of strange quark stars, much effort has been devoted to find major"
+[Fe/H] and only differs for their [Na/O] ratio.,[Fe/H] and only differs for their [Na/O] ratio.
+" This is most clearly shown in Figure 3,, where we plotted separately the (V, B-V) diagrams for the metal-poor groups (left panel), and for those of intermediate metallicity (right panel)."," This is most clearly shown in Figure \ref{f:fig3}, where we plotted separately the (V, B-V) diagrams for the metal-poor groups (left panel), and for those of intermediate metallicity (right panel)."
+" From these figures the presence of a correlation between colours along the RGB and stellar metallicity is apparent, as expected from evolutionary models."," From these figures the presence of a correlation between colours along the RGB and stellar metallicity is apparent, as expected from evolutionary models."
+" We may quantify this dependence, to produce a metallicity index which depends on the colours of the stars along the RGB."," We may quantify this dependence, to produce a metallicity index which depends on the colours of the stars along the RGB."
+ The usefulness of this index will be discussed later., The usefulness of this index will be discussed later.
+" In order to define it, we started by fitting a cubic polynomial through the RGB of the two metal-poor groups."," In order to define it, we started by fitting a cubic polynomial through the RGB of the two metal-poor groups."
+ We then measured the (B-V) offset A(B-V) between each star and this reference sequence., We then measured the (B-V) offset $\Delta$ (B-V) between each star and this reference sequence.
+" Finally, we fit a bilinear relation between [Fe/H] (dependent variable), and A(B-V) and V (independent variables)."," Finally, we fit a bilinear relation between [Fe/H] (dependent variable), and $\Delta$ (B-V) and V (independent variables)."
+" The best fit parameters from this relation allows to have for each star a parameter, that we call [m/H], which on average is identical to [Fe/H], but that obviously can be different from it for each individual star (and even for each group)."," The best fit parameters from this relation allows to have for each star a parameter, that we call [m/H], which on average is identical to [Fe/H], but that obviously can be different from it for each individual star (and even for each group)."
+ The average value of [m/H] for each group is also listed in Table 1.., The average value of [m/H] for each group is also listed in Table \ref{t:tab1}.
+" Figure 4 shows the correlation between [m/H] and [Fe/H] for the whole sample, as well as separately for the metal-poor,"," Figure \ref{f:fig4} shows the correlation between [m/H] and [Fe/H] for the whole sample, as well as separately for the metal-poor,"
+at least —3]aus ddetermunes a variation of of the Lue width.,at least $\sim$ 3 determines a variation of of the line width.
+ The latter is negligible to our purposes., The latter is negligible to our purposes.
+ Another explanation for the observed line broadening at high euerey might be an increase ofoptical depth for the lieh-cnerey transitions., Another explanation for the observed line broadening at high energy might be an increase of optical depth for the high-energy transitions.
+ It is possible to demoustrate that for an optical depth 7 at the line ceutre. the ratio between the line FWIIAL Ac. and the corresponding value in the optically thin limit. Ac (i.e. the intrinsic line width). is eiven by the expression From this. one can see that 7=1 is needed to obtain οz1.5. ie. the wast ratio seen iu the top panel of Fig. 11..," It is possible to demonstrate that for an optical depth $\tau$ at the line centre, the ratio between the line FWHM, $\Delta v$, and the corresponding value in the optically thin limit, $\Delta v_0$ (i.e. the intrinsic line width), is given by the expression From this, one can see that $\tau\simeq4$ is needed to obtain $\Delta v/\Delta v_0\simeq1.5$, i.e. the maximum ratio seen in the top panel of Fig. \ref{pend}."
+ As already discussed at leugth in Sects., As already discussed at length in Sects.
+" 3,2. and 2.1.1. such a large value iust refer to a very πια region. which cannot contribute to the ecluission observed by us on a much larger scale."," \ref{stempden}
+ and \ref{stau}, such a large value must refer to a very small region, which cannot contribute to the emission observed by us on a much larger scale."
+ Therefore. line broadening is uulikelv to be due to optical depth. effects.," Therefore, line broadening is unlikely to be due to optical depth effects."
+ Iu conclusion. we believe that the most plausible explanation for Fig.," In conclusion, we believe that the most plausible explanation for Fig."
+ ll ois that the velocity dispersion increases with cucrey of the transition and hence towards he centre of the chuups., \ref{pend} is that the velocity dispersion increases with energy of the transition and hence towards the centre of the clumps.
+ Various iechaisius cau explain this treud: keplerian rotation. iufall motions. aud urbuleuce caused by winds/outflows from the eiibedtaed. voung higl-imass stars.," Various mechanisms can explain this trend: keplerian rotation, infall motions, and turbulence caused by winds/outflows from the embedded, young high-mass stars."
+ Although the present observations cannot discriminate among these. it is worth noting hat the infall hvpothesis is consistent with the scenario xoposed in Sects.," Although the present observations cannot discriminate among these, it is worth noting that the infall hypothesis is consistent with the scenario proposed in Sects."
+ L1 and 1.2.2., \ref{sstab} and \ref{sdens}.
+" We have used the IRAM, 3020. aud FCRAO 1bin telescopes to observe a sample of 12 UC iregious in the J=65. 87. aud 1312 transitions ofCol..."," We have used the IRAM 30-m and FCRAO 14-m telescopes to observe a sample of 12 UC regions in the $J$ =6–5, 8–7, and 13–12 transitions of."
+ Emission has been detected in all of the objects and the following results have been obtained:, Emission has been detected in all of the objects and the following results have been obtained:
+"the consideration made by Llovel (2002)) in the sense of the universe as being similar to a quantum computer. where the big number 10"" plavs the role of the maximum number of elementary operations made by the universe on 10"" bits 2002)).","the consideration made by Lloyd ) in the sense of the universe as being similar to a quantum computer, where the big number $10^{120}$ plays the role of the maximum number of elementary operations made by the universe on $10^{90}$ bits )."
+ Our estimate has gone further. (he number of bits is found to be LO’? 20LIbb).," Our estimate has gone further, the number of bits is found to be $10^{122}$ b)."
+ We reinforce the statement of He and Ala (2011)):gravilq., We reinforce the statement of He and Ma ):.
+" The Schwarzschild radius R=26M/c? of a black hole may be combined with the Compton wavelength A=h/Ale to give 2011)): where { is the Planck's length /,=(Gh/ey'’?. G the gravitational constant. the Planck's constant and ο the speed of light."," The Schwarzschild radius $R = 2GM/c^2$ of a black hole may be combined with the Compton wavelength $\bar{\lambda} = \hbar /Mc$ to give ): where $l_p$ is the Planck's length $l_p = (Gh/c^3)^{1/2}$, $G$ the gravitational constant, $\hbar$ the Planck's constant and $c$ the speed of light."
+ Relation (1) states that the geonietric mean between R and A is of the order of Planck’s length., Relation (1) states that the geometric mean between $R$ and $\bar{\lambda}$ is of the order of Planck's length.
+ For a black hole. mass and length are proportional.," For a black hole, mass and length are proportional."
+ Since the Planck's units correspond to a quantum black hole: the eeneralization of the relation (1) gives as a result that all properties of a black hole. length A. mass AL and characteristic time £2/e. are related in (his way.," Since the Planck's units correspond to a quantum black hole; the generalization of the relation (1) gives as a result that all properties of a black hole, length $R$, mass $M$ and characteristic time $R/c$, are related in this way."
+ Ht means (hat anv black hole (2.M) has a black hole given by the relation (1). with properties of mass i and Selwarzschild radius A such that (1) is satisfied.," It means that any black hole $(R,M)$ has a black hole given by the relation (1), with properties of mass $m$ and Schwarzschild radius $\bar{\lambda}$ such that (1) is satisfied."
+ We have presented. and analvzed elsewhere 2010)). (he universe as a black hole.," We have presented, and analyzed elsewhere ), the universe as a black hole."
+ There have also been arguments in [avor of considering the universe. as a whole. to be a quantum object 2011)).," There have also been arguments in favor of considering the universe, as a whole, to be a quantum object )."
+ The combination of these (wo ideas gives, The combination of these two ideas gives
+relativistic particles would be much higher than the pressure of the maguetic field.,relativistic particles would be much higher than the pressure of the magnetic field.
+" Therefore either the confinement time of electrous in the knot would be eRef)-ὃς101RI,s. or else the oniknot would be inflating on the same timescale."," Therefore either the confinement time of electrons in the knot would be $t_{conf}^\prime
+\sim R/(c/\sqrt{3}) \simeq 2\times 10^{11} R_{kpc}\, \rm s$, or else the knot would be inflating on the same timescale."
+" This implies a comoving injection power in the kuot L'>5«WETSRy,ergs+. or a statiouary-fraue injection power/ D2L/>10!"" eres Ἐν"," This implies a comoving injection power in the knot $L^\prime \gtrsim 5\times 10^{44}
+W^{\prime}_{56}/R_{kpc} \,\rm erg \, s^{-1}$, or a stationary-frame injection power $\Gamma^2 L^\prime \gtrsim 10^{47}$ ergs $^{-1}$."
+ In fact. taking iuto account contribution of protons. Tavecchioetal.(2000) require a power of 3«107 cress1! to model the spectral energy distribution (SED) of knot WIx7.5 of PISS 0637-752 with the EC inodel.," In fact, taking into account contribution of protons, \citet{tav00} require a power of $3\times 10^{48}$ ergs $^{-1}$ to model the spectral energy distribution (SED) of knot WK7.8 of PKS 0637-752 with the EC model."
+ Although they areuc that this is consistent with total kinetic power of jets inferred to power the eiat radio lobes of ra10 galaxies. the largest jet power in the sample of Rawhlnes&Sauuders(1991). hardly exceeds 10 eres f.," Although they argue that this is consistent with total kinetic power of jets inferred to power the giant radio lobes of radio galaxies, the largest jet power in the sample of \citet{rs91} hardly exceeds $10^{47}$ ergs $^{-1}$."
+" This power also is 12 orders of maenitude ereater than the maximum peak οταν lunimosities inferred frou, EGRET observations of blazars. talàug iuto account the likely beaming factor of ~15 (Mukherjee&Chiang1999)."," This power also is 1–2 orders of magnitude greater than the maximum peak $\gamma$ -ray luminosities inferred from EGRET observations of blazars, taking into account the likely beaming factor of $\sim 1$ \citep{mc99}."
+. It should be noted here that in principal it is possible to satisfv the equipartition condition also for a given size R=Lkpe of the knot WIS. increasing the Doppler factor 6 further to ~20 (DerimeraudAtovan2001).," It should be noted here that in principal it is possible to satisfy the equipartition condition also for a given size $R=1 \,\rm kpc$ of the knot WK7.8, increasing the Doppler factor $\delta$ further to $ \simeq 20$ \citep{da04}."
+. This will also ininimize reducing) the power requirements for the EC model (CduisclliniaudCelotti2001:DenieraudAtovan 2001).," This will also minimize ) the power requirements for the EC model \citep{gc01,da04}."
+. However. this is possible oulv if the jet inclination angle is further decreased to <3°.," However, this is possible only if the jet inclination angle is further decreased to $\leq 3^\circ$."
+ Although this cannot be excluded for a particular source. such angles would be problematic to assunae for niauv sources. ax We discuss in Section...," Although this cannot be excluded for a particular source, such angles would be problematic to assume for many sources, as we discuss in Section 4."
+ Iu the EC/CAIBR aodel. N-vavs are produced w clectrous with στο~LO? in the stationary rae.," In the EC/CMBR model, X-rays are produced by electrons with $\gamma_{\rm C}\sim 10^3 $ in the stationary frame."
+ They cool ou the CAIBR only ou Cor nuescales το]=te of ({1)). unless their overall cooling is dominated by svuclirotron losses. which would then overproduce the observed racio Hux at low frequencies.," They cool on the CMBR only on Gyr timescales $t_{cool} \simeq t_{C}$ of \ref{tC}) ), unless their overall cooling is dominated by synchrotron losses, which would then overproduce the observed radio flux at low frequencies."
+ Since the photon target or the Compton process does not degrade at aly distance + fromthe parent ACN. the strong facing of the N-rav. knots and jet generally (outside the shots) with r iu many FR2 quasars. such as in SC 2732001).. PISS 1127-115 (Sictuieginowskaetal.2002).. or Pictor A (Wilsouetal.2001).. iuplies either a fast decrease of the total πο of electrons with 5~5c or a decline of the jet Doppler factor.," Since the photon target for the Compton process does not degrade at any distance $r$ from the parent AGN, the strong fading of the X-ray knots and jet generally (outside the knots) with $r$ in many FR2 quasars, such as in 3C 273, PKS 1127-145 \citep{siem02}, or Pictor A \citep{wys01}, implies either a fast decrease of the total number of electrons with $\gamma \sim
+\gamma_{\rm C}$ or a decline of the jet Doppler factor."
+ Because the observed radio clussion is produced bv clectrous with σαν lavger than 2c ouly by a factor =10. both assumptions would imply a rapid facing of the| radio briehtuess with distauce as well.," Because the observed radio emission is produced by electrons with $\gamma_{syn}$ larger than $\gamma_{\rm C}$ only by a factor $\lesssim 10$, both assumptions would imply a rapid fading of the radio brightness with distance as well."
+ Meauchile. exactly the opposite radio brightuess profiles are observed.," Meanwhile, exactly the opposite radio brightness profiles are observed."
+ Secondly. because the cooling time is shorter for radio than for Corupton ταν cutting electrons. oue should also expect that the N-ray knots in the EC model wemud be more diffuse than the radio or optica knots.," Secondly, because the cooling time is shorter for radio than for Compton X-ray emitting electrons, one should also expect that the X-ray knots in the EC model would be more diffuse than the radio or optical knots."
+ Again. this predictiou Is in disagreenieliowith observations.," Again, this prediction is in disagreement with observations."
+ Note in this regard that the adiabatic cooling of N-rav cluitting clectrous due to ast expansion of X-ray knots. as proposed receutlv (Taveccliio.Ctdis-ellinui.aud.Celotti2003) for the iuterpretation of facing of the N-vav brielitucss outside ιο knots. would also equally iupact the radio-ocuittiug clectrous (Stawarzctal.2001).," Note in this regard that the adiabatic cooling of X-ray emitting electrons due to fast expansion of X-ray knots, as proposed recently \citep{tgc03} for the interpretation of fading of the X-ray brightness outside the knots, would also equally impact the radio-emitting electrons \citep{sta04}."
+. Therefore this scenario would uot resolve either of these two discrepancies between the N-rav aud the radio brightuess patterns., Therefore this scenario would not resolve either of these two discrepancies between the X-ray and the radio brightness patterns.
+" One more diffiailty for the Compton interpretation arises for those jets where the spectral indices & (for the spectral flux £F,xvr "") are significantly different at racio and N-rav frequencies.", One more difficulty for the Compton interpretation arises for those jets where the spectral indices $\alpha$ (for the spectral flux $F_{\nu} \propto \nu^{-\alpha} $ ) are significantly different at radio and X-ray frequencies.
+" As also cliscussecl earlier. ooree, by IurisandNKrawezvu-ski (2002).. spectral profiles steeper in N-ravs thui in radio.withn ayD0, aud especially jets/knot:J4. in a nuniber of FRI galaxies where ayz 1. represent strong evidence for the svuchrotron origin of the X-rav flux. nuplviug acceleration of electrons with VeFouqnrLOfων,-2)e120= in the kuots of τικο jets."," As also discussed earlier, g., by \citet{hk02}, spectral profiles steeper in X-rays than in radio,with $\alpha_X > \alpha_r$, and especially jets/knots in a number of FR1 galaxies where $\alpha_{X} \gtrsim 1$ , represent strong evidence for the synchrotron origin of the X-ray flux, implying acceleration of electrons with $\gamma^\prime \gg
+10^7\sqrt{E_{keV}(1+z)/\delta_{10}B_{30}}$ in the knots of radio jets."
+" This also is the case for many terminal hot spots. such as the western hot spot of Pictor DÀ with a,zO7L aud ay=1.07x0.11 (Wilsonetal.2001)."," This also is the case for many terminal hot spots, such as the western hot spot of Pictor A with $\alpha_{r} \approx 0.74 $ and $\alpha_X = 1.07 \pm 0.11$ \citep{wys01}."
+. It should be pointoc out here that these difüculties with he interpretation of iffereut spectral aud opposite spatial profiles oei radio and Nay patterus of jets relate not ouly to EC models. but equalv to SSC models.," It should be pointed out here that these difficulties with the interpretation of different spectral and opposite spatial profiles in radio and X-ray patterns of jets relate not only to EC models, but equally to SSC models."
+ The N-rav flux in SSC inodeIs is Likewise produced bv low- electrons fiat cuit the svuchrotron racio photous forsubseqnent Compton interactions. aud the density of this arect docs not decline. aud may," The X-ray flux in SSC models is likewise produced by low-energy electrons that emit the synchrotron radio photons forsubsequent Compton interactions, and the density of this target does not decline, and may"
+ το.b 2—1 ἐν1). , $z\sim1$ $z\sim1$ $z\ga 1$ 
+Noting that the 0.80.9 keV feature could be due to an II-like oxveen VIII Ix edge (at resi. [rame energy of 0.87 keV. seen in many Sevlert 1: sources (seee.g.Revnolds1997) due lo presence of ionized gas in (he line of sight). we next attempt to test if this feature (present onlv during the low phase) could indeed be due to the oxvgen absorption edge.,"Noting that the 0.8–0.9 keV feature could be due to an H-like oxygen VIII K edge (at rest frame energy of 0.87 keV, seen in many Seyfert 1 sources \citep[see e.g.][]{reynolds1997} due to presence of ionized gas in the line of sight), we next attempt to test if this feature (present only during the low phase) could indeed be due to the oxygen absorption edge."
+ For this we included a multiplicative component to our models., For this we included a multiplicative component to our models.
+ The best-fit values after including an edge in both Mhi_bbbhi and AMhi256poi are given in Table I.., The best-fit values after including an edge in both $Mhi\_bb\_bkn$ and $Mhi\_2bb\_pow$ are given in Table \ref{tab:fits}. .
+ The improvement in X? Π statistic [or model Mo.bbblnedge between 7=0 (no edge) and 7=0.29 (best fit) isNA?=19.4. implving a single parameter confidence limitof 4.40.," The improvement in $\chi^2$ fit statistic for model $Mlo\_bb\_bkn\_edge$ between $\tau=0$ (no edge) and $\tau=0.29$ (best fit) is$\Delta\chi^2=19.4$, implying a single parameter confidence limitof $4.4\sigma$."
+ The improvement in 4? fit statistic [or Mo.2bbpowedge between 7=0 (no edge) and 7=0.23 (best fit) is NA?=16.6. implving a single parameter confidence limit of 4.076.," The improvement in $\chi^2$ fit statistic for $Mlo\_2bb\_pow\_edge$ between $\tau=0$ (no edge) and $\tau=0.23$ (best fit) is $\Delta\chi^2=16.6$, implying a single parameter confidence limit of $4.07\sigma$."
+ Thus both low phase models need an edee al 4c level. strongly. suggesting that the edee is present in the data.," Thus both low phase models need an edge at $>$ $\sigma$ level, strongly suggesting that the edge is present in the data."
+ While fits to both models. A/fo_bb_bhn ancl Affo2hb_pouw are formally acceptable. there remain hints of some svstematic trend in the residuals shown in Fig.," While fits to both models $Mlo\_bb\_bkn$ and $Mlo\_2bb\_pow$ are formally acceptable, there remain hints of some systematic trend in the residuals shown in Fig."
+ 4 at lower energies., \ref{f:lo_renorm} at lower energies.
+ This is discussed in greater detail in refconchision., This is discussed in greater detail in \\ref{conclusion}.
+ The 0.255 keV [lux in the low phase spectrum is (1.14€0.1)x10H Cnr> /s., The 0.25–5 keV flux in the low phase spectrum is $(1.14\pm0.1)\times10^{-11}$ $^2$ /s.
+ Next we attempted (o estimate the strength of the same OVI absorption edge in the hieh phase data. by adding an edge to both the models M.bbbli and Mhi200pow ancl recloing the fits.," Next we attempted to estimate the strength of the same OVIII absorption edge in the high phase data, by adding an edge to both the models $Mhi\_bb\_bkn$ and $Mhi\_2bb\_pow$ and redoing the fits."
+ We found (hat inclusion of the edee did not improve the statistical quality of the fits., We found that inclusion of the edge did not improve the statistical quality of the fits.
+ The OVIII edge could not be detected (at >36 level) in the data using either of the continuum models. and this confirms (that the optical depth of the edge varied significantly between the high ancl low phases.," The OVIII edge could not be detected (at $>3\sigma$ level) in the data using either of the continuum models, and this confirms that the optical depth of the edge varied significantly between the high and low phases."
+ Next we modeled the edge in the low phase somewhat more physically using the NSTAR photoionization code (Ixallman&Battista2001)., Next we modeled the edge in the low phase somewhat more physically using the XSTAR photoionization code \citep{xstar}.
+".. A grid of XSTAR models was created from the following variables: (1) the column density CVj. allowed range LO?"" 107 7). (2) ionization parameter (/og(£)=log(L/(n1)). allowed range 24). (3) oxvgen abundance (relative to solar. allowed range 0.14). (4) neon abundance (relative to solar. allowed range 0.11)."," A grid of XSTAR models was created from the following variables: (1) the column density $N_H$, allowed range $10^{20}$ – $10^{23}$ $^{-2}$ ), (2) ionization parameter $log(\xi)=log(L/(nR^2))$, allowed range 2–4), (3) oxygen abundance (relative to solar, allowed range 0.1–4), (4) neon abundance (relative to solar, allowed range 0.1–1)."
+ A multiplicative table model was ereated from this multidimensional grid., A multiplicative table model was created from this multidimensional grid.
+ Turbulent velocities of 300 km/s ancl above cause absorption lines to become more prominent. and lower turbulent velocities are required to fit the data well.," Turbulent velocities of 300 km/s and above cause absorption lines to become more prominent, and lower turbulent velocities are required to fit the data well."
+ The best-fit parameters to the low phase data using this model in conjunction with a continuum model composed of law suggest - :Vj/(107!E EN7) = 44 1. UN log(£)—3.2a0.2 (jj. and a relatively. high oxygen abundance (2.1 (2).," The best-fit parameters to the low phase data using this model in conjunction with a continuum model composed of double-blackbody+power law suggest $N_H$ $^{21}$$^{-2}$ ) = $4\pm1$ , $\xi$ $ 3.2 _{-0.1}^{+0.2}$ , and a relatively high oxygen abundance $ 2.1 _{-0.5}^{+0.8} $ )."
+ A sub-solar neon abundance (0.13) is preferred by the best-fit solution but its error-bar could not be constrained., A sub-solar neon abundance (0.13) is preferred by the best-fit solution but its error-bar could not be constrained.
+ See Table 1. for the detailed [it, See Table \ref{tab:fits} for the detailed fit
+part of the spectrum gets significantly. modified due to the general rclativisitic effects.,part of the spectrum gets significantly modified due to the general relativisitic effects.
+ A computation similar to that of Ebisawa et al. (, A computation similar to that of Ebisawa et al. (
+1991). for Galactic black hole cancliclates. Bas been done by Asaoka (1989) using the Ixerr metric.,"1991), for Galactic black hole candidates, has been done by Asaoka (1989) using the Kerr metric."
+ However. our work incorporating both: the full general. relativistic ellects. of rapid. rotation. as well as realistic equations of state describing neutron star interiors. using an appropriate metric is the first calculation for rotating neutron stars.," However, our work incorporating both: the full general relativistic effects of rapid rotation, as well as realistic equations of state describing neutron star interiors, using an appropriate metric is the first calculation for rotating neutron stars."
+ In the present work we ignore the elfects of the stellar magnetic field. so our results are applicable to weakly magnetised neutron stars.," In the present work we ignore the effects of the stellar magnetic field, so our results are applicable to weakly magnetised neutron stars."
+ The structure of the paper is as follows., The structure of the paper is as follows.
+ In section 2. we describe a framework for the caleulation of neutron star structure. disc temperature profile and observed spectrum.," In section 2, we describe a framework for the calculation of neutron star structure, disc temperature profile and observed spectrum."
+ We also comment here on the the chosen equations of state., We also comment here on the the chosen equations of state.
+ We describe the numerical procedure for spectrum calculation in section 3., We describe the numerical procedure for spectrum calculation in section 3.
+ The results. of the spectrum calculation are clisplaved in section d., The results of the spectrum calculation are displayed in section 4.
+ We summarise the conclusions in section. 5 and. enlist some relevant mathematical expressions in the Appendix., We summarise the conclusions in section 5 and enlist some relevant mathematical expressions in the Appendix.
+ We calculate the structure of a rapidly rotating neutron star for realistic IOS mocdels. using the same procedure as Cook et al. (," We calculate the structure of a rapidly rotating neutron star for realistic EOS models, using the same procedure as Cook et al. ("
+1994) (see also Datta et al 1998).,1994) (see also Datta et al 1998).
+" With the assumption that the star is rigidly rotating and a perfect Ελα, we take a stationary. axisvnunetric. asvmptoticallv [at and rellection-symmetric (about the equatorial plane) metric. given by where the metric potentials 5.9.0 and the angular speed. (a) of zero-angular-momoentum-observer. (ZAALO) with respect to infinity. are all functions of the quasi-isotropic radial coordinate (r) ancl polar angle (80)."," With the assumption that the star is rigidly rotating and a perfect fluid, we take a stationary, axisymmetric, asymptotically flat and reflection-symmetric (about the equatorial plane) metric, given by where the metric potentials $\gamma, \rho, \alpha$ and the angular speed $\omega$ ) of zero-angular-momentum-observer (ZAMO) with respect to infinity, are all functions of the quasi-isotropic radial coordinate $\bar r$ ) and polar angle $\theta$ )."
+ Thequantity. 7 ds. related to the Sehiwarzschild-like racial coordinate (r) by the equation rr—rePES AVG , Thequantity $\bar r$ is related to the Schwarzschild-like radial coordinate $r$ ) by the equation $r = \bar r e^{\rm {(\gamma - \rho)/2}}$.
+use the geometric units e=61 in this paper., We use the geometric units $c = G = 1$ in this paper.
+ We solve Einstein's field equations and the equation of hydrostatic equilibrium self-consistentlyancl numerically from the centre of the star upto infinity to obtain +.p.a. and © (angular speed of neutron star with respect to infinity) as functions of r and 8.," We solve Einstein's field equations and the equation of hydrostatic equilibrium self-consistentlyand numerically from the centre of the star upto infinity to obtain $\gamma, \rho, \alpha, \omega$ and $\Omega$ (angular speed of neutron star with respect to infinity) as functions of $\bar r$ and $\theta$."
+ This is done for a particular LOS and assumed. values of central density ancl ratio of polar to equatorial radii., This is done for a particular EOS and assumed values of central density and ratio of polar to equatorial radii.
+ The obtained numerical equilibrium solutions of the metric enable us to caleulate bulk structure parameters. such as gravitational mass (AL). cquatorial radius (2). angular momentum (./) etc.," The obtained numerical equilibrium solutions of the metric enable us to calculate bulk structure parameters, such as gravitational mass $M$ ), equatorial radius $R$ ), angular momentum $J$ ) etc."
+ of the neutron star., of the neutron star.
+" We also calculate the radius (75,405) of the innermost stable circular orbit (ISCO). specific energy GE) and specific angular momentum (/) of a test. particle in a Ixeplerian orbit and the Ixeplerian angular speed. (Qi) (see Phampan Datta 1998 for a detailed. description of the method of calculation)"," We also calculate the radius $r_{\rm orb}$ ) of the innermost stable circular orbit (ISCO), specific energy $\tilde E$ ) and specific angular momentum $\tilde l$ ) of a test particle in a Keplerian orbit and the Keplerian angular speed $\Omega_{\rm K}$ ) (see Thampan Datta 1998 for a detailed description of the method of calculation)."
+ The neutron star structure parameters are quite sensitive to the chosen EOS., The neutron star structure parameters are quite sensitive to the chosen EOS.
+ In the literature. there exist a large number of EOS ranging from very soft to very stil.," In the literature, there exist a large number of EOS ranging from very soft to very stiff."
+ For the purpose of a general studs. we have chosen four EOS. of which one is soft. (EOS: X. Panclharipande (1979). (bvperons)). one is intermediate (EOS: D. Baldo. Bombaci Durgio (1997) (AV14|8b0)) in stilffness and two are still (EOS: €. Walecka (1974) and EOS: D. Sahu. Dasu Datta (1993)).," For the purpose of a general study, we have chosen four EOS, of which one is soft (EOS: A, Pandharipande (1979) (hyperons)), one is intermediate (EOS: B, Baldo, Bombaci Burgio (1997) (AV14+3bf)) in stiffness and two are stiff (EOS: C, Walecka (1974) and EOS: D, Sahu, Basu Datta (1993))."
+ Of these EOS: D is the stillest., Of these EOS: D is the stiffest.
+ The structure of a neutron star for a given. EOS is described: uniquely by two parameters he gravitational mass (AJ) and the angular speed (Q)., The structure of a neutron star for a given EOS is described uniquely by two parameters: the gravitational mass $M$ ) and the angular speed $\Omega$ ).
+ For cach chosen EOS. we construct constant AL equilibrium sequences with O varving from the non-rotating case upto the centrifugal mass shed limit (rotation rate at which inwardly directed eravitational forces are balanced by outwardly directed centrifugal forces).," For each chosen EOS, we construct constant $M$ equilibrium sequences with $\Omega$ varying from the non-rotating case upto the centrifugal mass shed limit (rotation rate at which inwardly directed gravitational forces are balanced by outwardly directed centrifugal forces)."
+" Depending on O. AZ and the EOS model. neutron stars may have 2roa, or? mrugas"," Depending on $\Omega$ , $M$ and the EOS model, neutron stars may have $R>r_{\rm orb}$ or $R<r_{\rm orb}$."
+ The effective temperature of a thin blackhocky disc is given by where σ is the Stefan-Doltzmann constant and f° is the X-ray energy lux per unit surface areca., The effective temperature of a thin blackbody disc is given by where $\sigma$ is the Stefan-Boltzmann constant and $F$ is the X-ray energy flux per unit surface area.
+" We caleulate fusing the expression. of Page ""Thorne (1974). (valid for gcometrically thin non-scll-eravitating disc cmbecedect in a general axisvmmetric spacetime) for the accretion disc around a rotating blackhole: where llere nu is the disc inner edge radius and a comma followed by a variable as subscript to a quantity. represents a derivative of the quantity with respect to the variable."," We calculate $F$ using the expression of Page Thorne (1974) (valid for geometrically thin non-self-gravitating disc embedded in a general axisymmetric space–time) for the accretion disc around a rotating blackhole: where Here $r_{\rm in}$ is the disc inner edge radius and a comma followed by a variable as subscript to a quantity, represents a derivative of the quantity with respect to the variable."
+ Bhattacharyya ct al. (, Bhattacharyya et al. (
+2000) argue that. though formulated for the case of black holes. these expressions hold quite well also for the case of neutron stars.,"2000) argue that, though formulated for the case of black holes, these expressions hold quite well also for the case of neutron stars."
+ LERran. then the clise will touch the star anc we take my= A.," If $R > r_{\rm orb}$, then the disc will touch the star and we take $r_{\rm in} = R$ ."
+ Otherwise. Mu= Morn.," Otherwise, $r_{\rm in} = r_{\rm orb}$ ."
+ The temperature profile. of the accretion disc calculated as described. in this and the previous subsections is à function ofAL and © of the central star for any. adopted EOS.," The temperature profile, of the accretion disc calculated as described in this and the previous subsections is a function of$M$ and $\Omega$ of the central star for any adopted EOS."
+whether they also merge with the primary protostar. or whether they survive as secondary companions.,"whether they also merge with the primary protostar, or whether they survive as secondary companions."
+ Fig., Fig.
+ S shows velocity vectors superimposed on false-colour images of the density field in 1H0XU«140AU slices through the first fragment to form.," 8 shows velocity vectors superimposed on false-colour images of the density field in $140\,{\rm AU}\times 140\,{\rm AU}$ slices through the first fragment to form."
+ Densities and. velocities are computed as in Fig., Densities and velocities are computed as in Fig.
+ 2. but using co-ordinates measured relative to the centre of mass of the fragment.," 2, but using co-ordinates measured relative to the centre of mass of the fragment."
+ We see that the fragment is surrounded by a well-defined. rotationally supported: cireum-fragmentarvy disc (CED).," We see that the fragment is surrounded by a well-defined, rotationally supported circum-fragmentary disc (CFD)."
+ Very similar cdises are found around the other fragments., Very similar discs are found around the other fragments.
+ Typically they have radiibetween ~40AU and ~GOAL.," Typically they have radiibetween $\sim 40\,{\rm AU}$ and $\sim 60\,{\rm AU}$."
+ Although most of the core mass infalls directly onto the disc. the spiral arms in the disc are effective. at transporting angular momentum outwards. by means of gravitational torques. thereby allowing material to spiral inwarcs ancl accrete onto he primary protostar.," Although most of the core mass infalls directly onto the disc, the spiral arms in the disc are effective at transporting angular momentum outwards, by means of gravitational torques, thereby allowing material to spiral inwards and accrete onto the primary protostar."
+ Fragmentation further enhances the ransport of angular momentum., Fragmentation further enhances the transport of angular momentum.
+ This is because CEDs rotate in the same sense as the mother cise., This is because CFDs rotate in the same sense as the mother disc.
+ Consequently. on the inside of a CED (i.e. the side closest to the primary ootostar) the shear between the material in the CFD and he material in the mother disc acts to slow down the material in the mother disc. thereby. reducing its angular momentum.," Consequently, on the inside of a CFD (i.e. the side closest to the primary protostar) the shear between the material in the CFD and the material in the mother disc acts to slow down the material in the mother disc, thereby reducing its angular momentum."
+ Conversely. on the outside of the CEL (ie. urther from the primary protostar) the shear between the material in the CED and the material in the mother disc acts to speed up the material in the mother disc. thereby increasing its angular momentun.," Conversely, on the outside of the CFD (i.e. further from the primary protostar) the shear between the material in the CFD and the material in the mother disc acts to speed up the material in the mother disc, thereby increasing its angular momentum."
+ Efficient transport of angular momentum by the CED may cause a fragment to spiral inwards and merge with the primary protostar., Efficient transport of angular momentum by the CFD may cause a fragment to spiral inwards and merge with the primary protostar.
+ However. it is possible that the inspiral is being artificially accelerated. by our imposition of adiabaticitv at. densities pc10Lgcm since this forces a fragment. anc its CED to remain hot. extended and therefore. very dissipative. rather than allowing them to continue contracting towards protostellar densities.," However, it is possible that the inspiral is being artificially accelerated by our imposition of adiabaticity at densities $\rho>10^{-13}\,{\rm g}\,{\rm cm}^{-3}$, since this forces a fragment and its CFD to remain hot, extended and therefore very dissipative, rather than allowing them to continue contracting towards protostellar densities."
+ Whereas in Run 7 the azimuthallv averaged disc temperature never rises above ~200k. in Run 1 it rises above ~2001. and locally i0. reaches SOOM. sullicienthy high to produce crystalline. silicates.," Whereas in Run 7 the azimuthally averaged disc temperature never rises above $\sim 200\,{\rm K}$, in Run 1 it rises above $\sim 500\,{\rm K}$, and locally it reaches $\sim 800\,{\rm K}$, sufficiently high to produce crystalline silicates."
+" Thus we might expect ervstalline silicates to be present in discs forming from low-(3, cores. and hence to give rise to characteristic features in the SEDs of single protostars."," Thus we might expect crystalline silicates to be present in discs forming from $\Omega_{_{\rm O}}$ cores, and hence to give rise to characteristic features in the SEDs of single protostars."
+" Conversely. we might expect crystalline silicates to be rare in the cold. extended. disces forming from high-,, cores. and hence their characteristic features should be weak in or even absent from — the SEDs of multiple protostars."," Conversely, we might expect crystalline silicates to be rare in the cold, extended discs forming from $\Omega_{_{\rm O}}$ cores, and hence their characteristic features should be weak in – or even absent from – the SEDs of multiple protostars."
+ The existence of a link between the amount of crvstalline dust. which can be found in protoplanctary disces ancl the initial conditions within the parental molecular cloud core has already been suggested by 7.., The existence of a link between the amount of crystalline dust which can be found in protoplanetary discs and the initial conditions within the parental molecular cloud core has already been suggested by \citet{Dullemond2006}.
+" Referring now to the complete ensemble of Runs (1 to 7). we note that. as £2, increases. there is an essentially monotonic shift in the outcome: small departures from monotonicity at high values of €,TUO) can be attributed to the stochastic nature of fragmentation."," Referring now to the complete ensemble of Runs (1 to 7), we note that, as $\Omega_{_{\rm O}}$ increases, there is an essentially monotonic shift in the outcome; small departures from monotonicity at high values of $\Omega_{_{\rm O}}$ can be attributed to the stochastic nature of fragmentation."
+ Wey parameters are presented in Table 1., Key parameters are presented in Table 1.
+" For low O,,.350. a substantial primary protostar is formed. surrounded by a compact. thick. hot disc: the disc is unable to fragment. and simply accretes onto the primary protostar."," For low $\Omega_{_{\rm O}}$, a substantial primary protostar is formed, surrounded by a compact, thick, hot disc; the disc is unable to fragment, and simply accretes onto the primary protostar."
+" As ο, is Increased. the growth rate of the primary protostar decreases. not least because it acquires an increasing fraction of its mass via the disc (rather than by cürect infall)."," As $\Omega_{_{\rm O}}$ is increased, the growth rate of the primary protostar decreases, not least because it acquires an increasing fraction of its mass via the disc (rather than by direct infall)."
+ At the same time the dise becomes more massive. more extended and. cooler.," At the same time the disc becomes more massive, more extended and cooler."
+" Eventually. for ,,m135510Ds* (Runs 5. 6 and 7). the disc becomes unstable against fragmentation. ie. it satisfies both the Toomre and. Canimie conditions (?2).."," Eventually, for $\Omega_{_{\rm O}}>1.35\times 10^{-13}\,{\rm s}^{-1}$ (Runs 5, 6 and 7), the disc becomes unstable against fragmentation, i.e. it satisfies both the Toomre and Gammie conditions \citep{1964Toomre,Gammie2001}."
+ We are unable to ascertain whether all the resulting fragments meree with the primary protostar. or whether some of them survive as secondary companions.," We are unable to ascertain whether all the resulting fragments merge with the primary protostar, or whether some of them survive as secondary companions."
+" As ©, increases. the time. /,,. at which the primary formis increases. and its subsequent growth rate. AL. protostardecreases."," As $\Omega_{_{\rm O}}$ increases, the time, $t_{_{\rm O}}$, at which the primary protostar forms increases, and its subsequent growth rate, $\dot{M}_\star$ , decreases."
+ Consequently. the overall pattern of accretion shifts from one in whieh the primary protostar formis first.," Consequently, the overall pattern of accretion shifts from one in which the primary protostar forms first,"
+U1.28 and ULL respectively.,U1.28 and U1.11 respectively.
+ Clearly these are very. large sizes. placing these two LOCs among the largest features so [ar seen in the carly universe.," Clearly these are very large sizes, placing these two LQGs among the largest features so far seen in the early universe."
+ For comparison. Yacayctal.(2010) eive an idealised limit to the scale of homogeneity in the concordance Cosmology as  870 Alpe: it should. not be possible to find departures from homogeneity above this scale.," For comparison, \citet{Yadav2010} give an idealised limit to the scale of homogeneity in the concordance cosmology as $\sim$ 370 Mpc: it should not be possible to find departures from homogeneity above this scale."
+ The LOCs appear to be only marginally consistent with this scale of homogeneity., The LQGs appear to be only marginally consistent with this scale of homogeneity.
+ Calculation of the inertia tensor for these two LOCs shows a ratio 2.5 for the longest and shortest. principal axes in both cases., Calculation of the inertia tensor for these two LQGs shows a ratio $\sim 2.5$ for the longest and shortest principal axes in both cases.
+ They are therefore. substantially elongated., They are therefore substantially elongated.
+ By this measure. the longest axes are ~ 630 anc 78O Alpe for UL2s8 anc ULL respectively., By this measure the longest axes are $\sim$ 630 and 780 Mpc for U1.28 and U1.11 respectively.
+ Their morphologics appear to be markedly oblate. like a thick lens. each with two comparably laree long axes ancl a short axis that is smaller by the factor ~2.5.," Their morphologies appear to be markedly oblate, like a thick lens, each with two comparably large long axes and a short axis that is smaller by the factor $\sim 2.5$."
+ Clearly the long axes do exceed the expected scale of homogencity., Clearly the long axes do exceed the expected scale of homogeneity.
+" The estimated overdensitics are 6,=dp,/py0.83. 0.55 lor U1.28 and ULL respectively."," The estimated overdensities are $\delta_q = \delta \rho_q / \rho_q = 0.83$, $0.55$ for U1.28 and U1.11 respectively."
+ These overdensities are substantially lower than those found by Pilipenko(2007) and a little higher than those found by Milleretal.(2004)., These overdensities are substantially lower than those found by \citet{Pilipenko2007} and a little higher than those found by \citet{Miller2004}.
+. The occurrence of structure on à. particular scale is naturally taken to mean that the universe is. not homogeneous on that scale., The occurrence of structure on a particular scale is naturally taken to mean that the universe is not homogeneous on that scale.
+ These two LOQGs. UL.28 and UI.ll. as overdensities of the amplitudes ancl scales indicated. thus raise a question of compatibility with the scale of homogeneity in the concordance cosmology. if the Yadayvctal.(2010) fractal caleulations are adopted as reference.," These two LQGs, U1.28 and U1.11, as overdensities of the amplitudes and scales indicated, thus raise a question of compatibility with the scale of homogeneity in the concordance cosmology, if the \citet{Yadav2010} fractal calculations are adopted as reference."
+ A counter argument. could be made that these LOCs are chance associations of groups on sub-homogencity, A counter argument could be made that these LQGs are chance associations of groups on sub-homogeneity
+with measured masses ranging [rom a few 107M. to a [ow 10 (Magorrian. 1998). ancl predicted. X-ray luminosities (Gf the accretion. occurs at. the Dondi rate with a standard raciative ellieiencv) four to five orders of magnitude greater than is observed (Di Matteo 19992).,with measured masses ranging from a few $10^8 \Msun$ to a few $10^9$ (Magorrian 1998) and predicted X-ray luminosities (if the accretion occurs at the Bondi rate with a standard radiative efficiency) four to five orders of magnitude greater than is observed (Di Matteo 1999a).
+ The central cluster galaxies provide an extreme illustration of the phenomenon of quiescent black holes., The central cluster galaxies provide an extreme illustration of the phenomenon of quiescent black holes.
+ Beside possessing FR-Ltivpe radio sources anc very large lack hole masses. these galaxies exist in extremely. gas-rich environments in cooling Hows at the centres of clusters. ancl are therefore ideal sources in whieh to study he physics of low-raciative efficiency. accretion.," Beside possessing FR-I-type radio sources and very large black hole masses, these galaxies exist in extremely gas-rich environments in cooling flows at the centres of clusters, and are therefore ideal sources in which to study the physics of low-radiative efficiency accretion."
+ However. hese same reasons also imply that these galaxies are by no means tvpical.," However, these same reasons also imply that these galaxies are by no means typical."
+ Despite exhibiting lower luminosities. the »operties of the other three ellipticals in our sample. which do not exist in such preferential locations. may be more easily generalized. to other systems.," Despite exhibiting lower luminosities, the properties of the other three ellipticals in our sample, which do not exist in such preferential locations, may be more easily generalized to other systems."
+ These three galaxies nave also been studied at high radio frequencies anc sub-mm wavelengths (Di Matteo 19998) and. when coupled with the reduced. jet-racio-activity in these svstems (and he ASCA data. presented. in this paper). provide further crucial constraints on the primary accretion mechanisms in he nuclei of elliptical ealaxies.," These three galaxies have also been studied at high radio frequencies and sub-mm wavelengths (Di Matteo 1999a) and, when coupled with the reduced jet-radio-activity in these systems (and the ASCA data presented in this paper), provide further crucial constraints on the primary accretion mechanisms in the nuclei of elliptical galaxies."
+ The ASCA (CTanaka. Inoue Holt. 1994) observations were mace over a two-and-a-half vear period between 1993 June and 1995. December.," The ASCA (Tanaka, Inoue Holt 1994) observations were made over a two-and-a-half year period between 1993 June and 1995 December."
+ The ASCA X-ray telescope array (NIE) consists of four nested-foil telescopes. cach focussed onto one of four detectors: two X-ray CCD cameras. the Solid-state Imaging Spectrometers (SIS: S0 and SL) and two Gas scintillation Imaging Spectrometers (Cbs: G2 and C3).," The ASCA X-ray telescope array (XRT) consists of four nested-foil telescopes, each focussed onto one of four detectors; two X-ray CCD cameras, the Solid-state Imaging Spectrometers (SIS; S0 and S1) and two Gas scintillation Imaging Spectrometers (GIS; G2 and G3)."
+ The XIXE. provides a spatial resolution of 3 arcmin (Llalf Power Diameter) in the energy range 0.510 keV. The SIS detectors provide excellent spectral resolution A/E=.02(£/5.9keV)TRE ] over a −−−−22 aremirSo field of Znview.," The XRT provides a spatial resolution of $\sim 3$ arcmin (Half Power Diameter) in the energy range $0.5 - 10$ keV. The SIS detectors provide excellent spectral resolution $\Delta E/E = 0.02(E/5.9 {\rm
+keV})^{-0.5}$ ] over a $22 \times 22$ $^2$ field of view."
+ yThe GIS detectors provide poorer energy resolution AL/E=OSCE/5.9keV) ud but cover a larger circular field of view of ~50 arcmin diameter.," The GIS detectors provide poorer energy resolution $\Delta E/E = 0.08(E/5.9
+{\rm keV})^{-0.5}$ ] but cover a larger circular field of view of $\sim
+50$ arcmin diameter."
+ For our analysis of the ος. cata we have used he screened. event Lists from. the τον». processing. of he data sets available on the Coddard-Space Flight Center (GSFC) ASCA archive (for details sce the ASCA Data Recluetion Guide. published. by GSEC.)," For our analysis of the ASCA data we have used the screened event lists from the rev2 processing of the data sets available on the Goddard-Space Flight Center (GSFC) ASCA archive (for details see the ASCA Data Reduction Guide, published by GSFC.)"
+ The data were reduced using the E'POOLS software. (version 4.1) rom within the NSELEC'T environment. (version 1.4)., The data were reduced using the FTOOLS software (version 4.1) from within the XSELECT environment (version 1.4).
+ Further data-cleaning procedures as recommended. in the ASCA Data Reduction Guide. including appropriate grade selection. gain corrections and. manual screening based. on the individual instrument light curves. were followed.," Further data-cleaning procedures as recommended in the ASCA Data Reduction Guide, including appropriate grade selection, gain corrections and manual screening based on the individual instrument light curves, were followed."
+ A full summary of the observations is given in Table 1., A full summary of the observations is given in Table 1.
+ Spectra were extracted. [rom all four SCA detectors. except for NGC 4696 for which the S0 and G8 data exhibited residual calibration errors and so were excluded. from the analvsis.," Spectra were extracted from all four ASCA detectors, except for NGC 4696 for which the S0 and G3 data exhibited residual calibration errors and so were excluded from the analysis."
+ Phe spectra were extracted from. cireular regions. centred on the peaks of the X-ray emission from the galaxies.," The spectra were extracted from circular regions, centred on the peaks of the X-ray emission from the galaxies."
+ For the SIS data. the radii of the regions. studied: were selected to minimize the number of chip boundaries crossed (thereby minimizing the systematic uncertainties introduced bv such crossings) whilst covering as large a region of the ealaxies as possible.," For the SIS data, the radii of the regions studied were selected to minimize the number of chip boundaries crossed (thereby minimizing the systematic uncertainties introduced by such crossings) whilst covering as large a region of the galaxies as possible."
+ Data from the regions between the chips were masked out and excluded., Data from the regions between the chips were masked out and excluded.
+ The final extraction raclit for the SIS data are summarized in Fable 2., The final extraction radii for the SIS data are summarized in Table 2.
+ Also noted are the chip modes used for the observations and the number of chips from which the extracted data were drawn., Also noted are the chip modes used for the observations and the number of chips from which the extracted data were drawn.
+ For the CIS data. a fixed extraction radius of 6 aremin was adopted.," For the GIS data, a fixed extraction radius of 6 arcmin was adopted."
+ Background subtraction was carried out using the, Background subtraction was carried out using the
+"of the protoplanctary atmosphere) and metal-free (4.6. containing ouly IT aud We) atinosphlieres aud fouud values of Mj lower by up to an order of maguitude (as low as 2 M in the metal-free case for M=109 AL, vo) than in the dusty case.",of the protoplanetary atmosphere) and metal-free (i.e. containing only H and He) atmospheres and found values of $M_{crit}$ lower by up to an order of magnitude (as low as $2$ $_\odot$ in the metal-free case for $\dot M=10^{-6}$ $_\odot$ $^{-1}$ ) than in the dusty case.
+ It would certainly be iuterestiug to repeat calculations done in 5 aud Appendix À for the case of dust- atmosphere to see how this extreme reduction of opacity would extend the radial rauge available for CT.," It would certainly be interesting to repeat calculations done in \\ref{sect:mcrit}, \ref{sect:constr} and Appendix \ref{app:zeta} for the case of dust-free atmosphere to see how this extreme reduction of opacity would extend the radial range available for CI."
+" Iu practice. we cannot do such calculation at present since it is not possible to calibrate M, against the results of Tori Dsoma (2010) as we did in equations (L1)}-(17)) and Appendix A.."," In practice, we cannot do such calculation at present since it is not possible to calibrate $M_{atm}$ against the results of Hori Ikoma (2010) as we did in equations \ref{eq:Matm}) \ref{eq:zeta}) ) and Appendix \ref{app:zeta}."
+ This is because in the isence of dust & is a function uot only of gas temperature but also of gas density., This is because in the absence of dust $\kappa$ is a function not only of gas temperature but also of gas density.
+ As demonstrated iu ROG in this case AL. is no louger -independent 6: the ambicut temperature Ty aud density py of the uchula as the original analyses of Mizuno (1980) and Stevenson (1982) sugeest., As demonstrated in R06 in this case $M_{crit}$ is no longer independent of the ambient temperature $T_0$ and density $\rho_0$ of the nebula as the original analyses of Mizuno (1980) and Stevenson (1982) suggest.
+" Iustead oue finds that M4,X EEO where q=ὃνοΠερ (ROG)."," Instead one finds that $M_{atm}\propto (\rho_0/T_0^3)^{q/(1+q)}$ , where $q\equiv \partial\ln\kappa/\partial\ln\rho$ (R06)."
+ Obviously. ALi then also depend ou py aud Ty. and the knowledge 6 this dependence is verv nportaut for obtainiug Xii and η iu the dust-free case.," Obviously, $M_{crit}$ then also depend on $\rho_0$ and $T_0$, and the knowledge of this dependence is very important for obtaining $\Sigma_{lim}$ and $a_{lim}$ in the dust-free case."
+ Unfortunately. we do not possess this knowledge frou. first priuciples as opacity calculations are rather complicated. aud m anv case we cannot currently calibrate Mg as functions of AM. Ty and py against uuuercal results because calculations of Wo DTsoma (2010) were done for a single value of the planetary semui-1najor axis Gucaning fixed values of Ty and py). while AL was varied.," Unfortunately, we do not possess this knowledge from first principles as opacity calculations are rather complicated, and in any case we cannot currently calibrate $M_{crit}$ as functions of $M_c$, $T_0$ and $\rho_0$ against numerical results because calculations of Hori Ikoma (2010) were done for a single value of the planetary semi-major axis (meaning fixed values of $T_0$ and $\rho_0$ ), while $\dot M$ was varied."
+" The scaling of Ms, and £4; with Ty and py and its implication for the possibility of the CI thus remain worthwhile issues for future investigation.", The scaling of $M_{atm}$ and $M_{crit}$ with $T_0$ and $\rho_0$ and its implication for the possibility of the CI thus remain worthwhile issues for future investigation.
+ We can still ect a qualitative idea of how αν changes in the clust-free case by setting opacity in equation (30)) at the verv low level cousisteut with pure gas opacity. ee we=10 Pons] gt ," We can still get a qualitative idea of how $a_{lim}$ changes in the dust-free case by setting opacity in equation \ref{eq:a_lim}) ) at the very low level consistent with pure gas opacity, e.g. $\kappa=10^{-4}$ $^2$ $^{-1}$."
+"Wo then find aiezosü AU compared to LL AU that equation (30)) predicts for ky=OL en? Ἡ,", We then find $a_{lim}^{MMSN}\approx 80$ AU compared to $44$ AU that equation \ref{eq:a_lim}) ) predicts for $\kappa=0.1$ $^2$ $^{-1}$.
+ Thus. opacity reduction due to sedimentation and coagulation of dust erams iu the protoplanctary atinosphere nav help in exteudiug the range of distances in which the CI is possible.," Thus, opacity reduction due to sedimentation and coagulation of dust grains in the protoplanetary atmosphere may help in extending the range of distances in which the CI is possible."
+ Despite the robustuess of our argunients it is nof inconceivable that some additional factors can weaken them aucdimake eiat planet formation by the CT possible even bevond the limits represeuted bv equations (23]) and (30))., Despite the robustness of our arguments it is not inconceivable that some additional factors can weaken them and make giant planet formation by the CI possible even beyond the limits represented by equations \ref{eq:Sig_lim}) ) and \ref{eq:a_lim}) ).
+ Alternatively. if is quite possible that some of the asstuuptions used iu deriviug these results are too extreme and one can get even better constraints by focussing ou less dramatic assuniptious.," Alternatively, it is quite possible that some of the assumptions used in deriving these results are too extreme and one can get even better constraints by focussing on less dramatic assumptions."
+ Below we review factors that can work one wav or another., Below we review factors that can work one way or another.
+ Oue possible way to facilitate CT and increase η is to consider possibility of planetesinial accretion at ratescceccdimg Monae., One possible way to facilitate CI and increase $a_{lim}$ is to consider possibility of planetesimal accretion at rates $\dot M_{max}$.
+ This is very difficult (since there are many factors that tend to reduce Ar conrpared to HAM SEC refsubsect:strenethen)) but may be possible if ee. one takes into account the increase of planetesimal capture cross-section bv the core caused by its extended. dense atinosplere.," This is very difficult (since there are many factors that tend to reduce $\dot M$ compared to $\dot M_{max}$, see \\ref{subsect:strengthen}) ) but may be possible if e.g. one takes into account the increase of planetesimal capture cross-section by the core caused by its extended, dense atmosphere."
+ This effect has Όσοι previously investigated by Inaba Tkoma (2003) who demonstrated that au increase of AL by a factor of ~10 compared to the value computed without atmosphere is possible., This effect has been previously investigated by Inaba Ikoma (2003) who demonstrated that an increase of $\dot M$ by a factor of $\sim 10$ compared to the value computed without atmosphere is possible.
+ According to equation (30)) such an cuhancement of AL (incorporated v increasing X) would boost a?72457 by a factor of ~LL , According to equation \ref{eq:a_lim}) ) such an enhancement of $\dot M$ (incorporated by increasing $\chi$ ) would boost $a_{lim}^{MMSN}$ by a factor of $\sim 4$.
+Our preseut caleulatious asstme that the core is accreting planetesinals coutinnously until the xotoplanetary nebula dissipates this is nuportant at arge a since Massive core requires long time to be built., Our present calculations assume that the core is accreting planetesimals continuously until the protoplanetary nebula dissipates — this is important at large $a$ since massive core requires long time to be built.
+ But one may wonder if building smaller core im shorter iue aud then cutting off subsequent plauctesimal accretion (and energy release at the core surface. which supports atinosphere against eoiug uustable) completely uav still lead to the CI and potentially extend it to argor seni-niajor axes.," But one may wonder if building smaller core in shorter time and then cutting off subsequent planetesimal accretion (and energy release at the core surface, which supports atmosphere against going unstable) completely may still lead to the CI and potentially extend it to larger semi-major axes."
+ Such accretion scenario has been adopted bv e.g. Pollack (1996)., Such accretion scenario has been adopted by e.g. Pollack (1996).
+ The problem iu js case is that even if M=0 dt still takes loug time or the atinosphere around the core to grow to the mass comparable to AZ..., The problem in this case is that even if $\dot M=0$ it still takes long time for the atmosphere around the core to grow to the mass comparable to $M_c$.
+ INEO show that this process occurs on rermmal timescale of the atmosphere aud typically takes uillious of vears., INE0 show that this process occurs on thermal timescale of the atmosphere and typically takes millions of years.
+ Sinilaur problemi is also encouutered in a scenario where re core erows rapidly bw planetesimal accretion iu ie juner regions of protoplanetary disk aud then sets scattered out to large radii by some massive perturber., Similar problem is also encountered in a scenario where the core grows rapidly by planetesimal accretion in the inner regions of protoplanetary disk and then gets scattered out to large radii by some massive perturber.
+ One might expect that after the orbit of the scattered core cireularizes bv dynamical friction the core would eradually accrete massive atuosphere and undergo CT at sole poit., One might expect that after the orbit of the scattered core circularizes by dynamical friction the core would gradually accrete massive atmosphere and undergo CI at some point.
+ Caven that both the orbit circularization aud euvelope accretion are likely to take long time it is nof at all obvious whether the CT could be achieved in this scenario witlin several Myrs., Given that both the orbit circularization and envelope accretion are likely to take long time it is not at all obvious whether the CI could be achieved in this scenario within several Myrs.
+ There are nau factors that can potentially reduce (tim compared to LL AU estimated in equation (30))., There are many factors that can potentially reduce $a_{lim}$ compared to 44 AU estimated in equation \ref{eq:a_lim}) ).
+" For example. there are several reasons why it may nof (e possible for M. to reach the maximuun rate M,,,,."," For example, there are several reasons why it may not be possible for $\dot M$ to reach the maximum rate $\dot M_{max}$."
+ First. the erowine core may clear out a gap 1u dlanctesimal disk around its orbit. thus siguificautlv reducing AL (Tanaka Ida 1997: Rafikov 2001: 2003a).," First, the growing core may clear out a gap in planetesimal disk around its orbit, thus significantly reducing $\dot M$ (Tanaka Ida 1997; Rafikov 2001; 2003a)."
+ Iu our previous calculations we implicitly assumed. this [unot to happen e.g. because of the core migration through ie disk. which allows fresh planctesimal material to »© Constantly supplied for core accretion (Alibert Ww051.," In our previous calculations we implicitly assumed this not to happen e.g. because of the core migration through the disk, which allows fresh planetesimal material to be constantly supplied for core accretion (Alibert 2005)."
+ Second. as we mentioned ii retsectiacer.. a known pathway to AL&ων is via the erowth of the core to the size at which it starts dominating dvuamical evolution of nearby planuctesimals and trigecrs thei efiicicut collisional fragmentation (Rafikoy 2001).," Second, as we mentioned in \\ref{sect:accr}, a known pathway to $\dot M\approx\dot M_{max}$ is via the growth of the core to the size at which it starts dominating dynamical evolution of nearby planetesimals and triggers their efficient collisional fragmentation (Rafikov 2004)."
+ However. there is a «ποιο implicit assunption in this scenario that the core cau reach this critical size within the nebula lifetime.," However, there is a strong implicit assumption in this scenario — that the core can reach this critical size within the nebula lifetime."
+ Ra&ükov (2003)) has shown that ata30010 AU advnamically dominant core would need to have mass of order 1074 ο and would require on the order of 10100 Myr to grow in the MAISN., Rafikov (2003b) has shown that at $a\sim 30-40$ AU adynamically dominant core would need to have mass of order $10^{24}$ g and would require on the order of $10-100$ Myr to grow in the MMSN.
+ This time scale is apparently in conflict witli the typical dissipation times of protoplauetary disks., This time scale is apparently in conflict with the typical dissipation times of protoplanetary disks.
+ Thus. oue may need to either require a more massive disk at these radii or to iud other patlavays for accretion," Thus, one may need to either require a more massive disk at these radii or to find other pathways for accretion"
+magnitude selected sample. evolving the magnitude cut-olf based. on accurate measurements of galaxy evolution.,"magnitude selected sample, evolving the magnitude cut-off based on accurate measurements of galaxy evolution."
+ We have presented results of the fitting of analytical halo models. incorporating simple LOD mocels with minimal number of free parameters. to cata (in this case the projected 2-point correlation. function) from the VVDS survey.," We have presented results of the fitting of analytical halo models, incorporating simple HOD models with minimal number of free parameters, to data (in this case the projected 2-point correlation function) from the VVDS survey."
+ This allowed us to study the evolution of the average number weighted halo mass and satellite fraction., This allowed us to study the evolution of the average number weighted halo mass and satellite fraction.
+ On cillerent scales there are contributions from central-satellite. satellite-satellite. anc central-central pairs of galaxies to the correlation function. thereby provicling constraints on the evolution of the galaxy satellite fraction.," On different scales there are contributions from central-satellite, satellite-satellite, and central-central pairs of galaxies to the correlation function thereby providing constraints on the evolution of the galaxy satellite fraction."
+ The evolution was obtained from data observed in the same restframe band. and. provides for simpler interpretations as compared to previous studies using data from clillerent restframe bands.," The evolution was obtained from data observed in the same restframe band, and provides for simpler interpretations as compared to previous studies using data from different restframe bands."
+ Various luminosity threshold. samples. at different redshifts were selected and the corresponding best-fit HOD parameters for two similar HOD mocdels obtained., Various luminosity threshold samples at different redshifts were selected and the corresponding best-fit HOD parameters for two similar HOD models obtained.
+ This is done in order to single out. possible degeneracies and inconsistencies with the fitting procedure at high redshifts., This is done in order to single out possible degeneracies and inconsistencies with the fitting procedure at high redshifts.
+ On the whole. both models are in agreement with cach other and show similar trends in evolution.," On the whole, both models are in agreement with each other and show similar trends in evolution."
+ The impact of our selection on the average halo mass is addressed: using the Millennium simulation., The impact of our selection on the average halo mass is addressed using the Millennium simulation.
+" We find that a erowth in halo mass as seen in the data could rather be an underestimation of 10% to what is seen in an ""ideal sample containing all the descendants.", We find that a growth in halo mass as seen in the data could rather be an underestimation of $\sim 10\%$ to what is seen in an 'ideal' sample containing all the descendants.
+ Pherefore a measure in the growth of mass of a halo can be mainly attributed to the hierarchical formation of structure ancl not due to the typology of the selection., Therefore a measure in the growth of mass of a halo can be mainly attributed to the hierarchical formation of structure and not due to the typology of the selection.
+ We find that the number-weightec average halo mass erows by ~90% [rom redshift 1.0 to 0.5., We find that the number-weighted average halo mass grows by $\sim 90 \%$ from redshift 1.0 to 0.5.
+ This. is the first time à growth in the underlying halo mass has been measured within a single data survey. and provides evidence for the rapid. accretion phase of massive halos.," This is the first time a growth in the underlying halo mass has been measured within a single data survey, and provides evidence for the rapid accretion phase of massive halos."
+ The mass accretion history follows the form. given in Wechsler et al. (, The mass accretion history follows the form given in Wechsler et al. (
+2002) with AZ(2)=Moe7. where joLOT+0.57(1.540.13) when only the VVDS points were used and i~1.04EO.10(2.090.04) after including the SDSS data as reference points at low recdshift and depending on the model used to obtain the best fits.,"2002) with $M(z) = M_0 e^{-\beta z}$, where $\beta \sim 1.07 \pm 0.57 (1.54 \pm 0.13)$ when only the VVDS points were used and $\beta \sim 1.94 \pm 0.10 (2.09 \pm 0.04)$ after including the SDSS data as reference points at low redshift and depending on the model used to obtain the best fits."
+ Phe addition of the low redshift SDSS points adds complications due to the addition of possible svstematics hy comparing data from two cillerent rest-frame bands. even after conversion to à common fiducial band.," The addition of the low redshift SDSS points adds complications due to the addition of possible systematics by comparing data from two different rest-frame bands, even after conversion to a common fiducial band."
+ We adopt the average value of jo1340.30 from the VVDS 1points when discussing8 a erowth in halo mass. and found to be slightly higher than the results from N-body simulations.," We adopt the average value of $\beta \sim 1.3 \pm 0.30$ from the VVDS points when discussing a growth in halo mass, and found to be slightly higher than the results from N-body simulations."
+ If we express this result in terms of the expected halo mass at present times. Alyc 100575134.. such halos appear to acerete m 0.25 Aly between redshifts of 0.5 and 1.0.," If we express this result in terms of the expected halo mass at present times, $M_0 \simeq$ $10^{13.5} h^{-1}M_\odot$, such halos appear to accrete $m \sim$ 0.25 $M_0$ between redshifts of 0.5 and 1.0."
+ Stewart et al. (, Stewart et al. (
+2007) have shown that  (S0%)) of Alb=1075!AL. halos experienced an mc0.3Λο (mz 0.LMo) merger event in the last LO Cor. this would translate into a m20.LMo merger event over the redshift range z=0.5-1.0] for the high mass halos here.,"2007) have shown that $\sim$ ) of $M_0=10^{13} h^{-1}M_\odot$ halos experienced an $m > 0.3 M_0$ $m > 0.1 M_0$ ) merger event in the last 10 Gyr, this would translate into a $m > 0.1 M_0$ merger event over the redshift range z=[0.5-1.0] for the high mass halos here."
+ From merger rate studies one finds that of the stellar mass of massive galaxies with LOMAL.<Al1014341. has been assembled through mergers since z=! (de Ravel et al., From merger rate studies one finds that of the stellar mass of massive galaxies with $10^{10} M_\odot < M < 10^{11} M_\odot$ has been assembled through mergers since z=1 (de Ravel et al.
+ 2000. and references therein).," 2009, and references therein)."
+ The integrated stellar mass growth obtained can then be compared to the halo mass growth obtained here., The integrated stellar mass growth obtained can then be compared to the halo mass growth obtained here.
+ For samples at similar redshifts we see that the average halo mass. «Al>. generally increases with the luminosity threshold of the sample. with a very mild hint of a decreasing galaxy satellite fraction.," For samples at similar redshifts we see that the average halo mass, $<M>$, generally increases with the luminosity threshold of the sample, with a very mild hint of a decreasing galaxy satellite fraction."
+ This implies that galaxies in the faint sample show a stronger probability of being satellites, This implies that galaxies in the faint sample show a stronger probability of being satellites
+starburst ealaxics Meurer ct al.,starburst galaxies Meurer et al.
+ use a galaxw spectrum obtained from a coustaut star formation rate for at most 105 years whereas we use eunpirical broadband spectra: the contribution of the old evolved stars to dust heating js certainly lavecr in our approach leading to a OWCL UV extinction. for the same amount of dust emission., use a galaxy spectrum obtained from a constant star formation rate for at most $10^8$ years whereas we use empirical broadband spectra: the contribution of the old evolved stars to dust heating is certainly larger in our approach leading to a lower UV extinction for the same amount of dust emission.
+ Another major difference is the treatineut of ecometrical effects., Another major difference is the treatment of geometrical effects.
+ Moeurer et al;, Meurer et al.
+ use a screen model aud we caleulate the extinction with a radiation fransfer model im au infinite plane parallel ecometry where dust aud stars are uniformly distributed and which accounts for scattering effects aud disk inclination (Nu Buat 1995))., use a screen model and we calculate the extinction with a radiation transfer model in an infinite plane parallel geometry where dust and stars are uniformly distributed and which accounts for scattering effects and disk inclination (Xu Buat \cite{xubu}) ).
+ Finally we asstuue a Milkv Way extinction curve whereas Meurer et al., Finally we assume a Milky Way extinction curve whereas Meurer et al.
+ adopt a uniform UV extinction for the entire spectrum of the starburst., adopt a uniform UV extinction for the entire spectrum of the starburst.
+ Therefore. our model i$ probably nore appropriate for uormal star-forming sSalaxies and the entire disk of starburst ealaxies whereas the calculations of Meurer et al.," Therefore, our model is probably more appropriate for normal star-forming galaxies and the entire disk of starburst galaxies whereas the calculations of Meurer et al."
+ are made for starburst regious., are made for starburst regions.
+ Civen these fuudiaimoeuta differences aud the rather large uncertainties on the corrections for extinction the two methods are in reasonable aerecimneut., Given these fundamental differences and the rather large uncertainties on the corrections for extinction the two methods are in reasonable agreement.
+ The FIR to UV flux ratio appears relatively iuscusitive to the dust characteristics (type. distribution) and the stars/dust ecometry.," The FIR to UV flux ratio appears relatively insensitive to the dust characteristics (type, distribution) and the stars/dust geometry."
+ This has been confirmed by the recent study of Witt Gordon (1999)) who explore various dust distributions (homogencous or clumpy). extinction properties (Milky Wav or Small. Magellanic Cloud) and stars/dust distributions ( uuifoxuuu mixture or shells).," This has been confirmed by the recent study of Witt Gordon \cite{witt}) ) who explore various dust distributions (homogeneous or clumpy), extinction properties (Milky Way or Small Magellanic Cloud) and stars/dust distributions ( uniform mixture or shells)."
+ Such a robustness makes the FIR to UV. flux ratio a reliable quantitative tracer of the dust attenuation im star forming ealaxies., Such a robustness makes the FIR to UV flux ratio a reliable quantitative tracer of the dust attenuation in star forming galaxies.
+ One basic difficultv of these studies based on individual ealaxies is that the samples used are all biased and soletimes in a very complicated seuse., One basic difficulty of these studies based on individual galaxies is that the samples used are all biased and sometimes in a very complicated sense.
+ The diagnostics ou the UV slope of nearby galaxies all derive from the compilation of Iimuey et al. (1993)), The diagnostics on the UV slope of nearby galaxies all derive from the compilation of Kinney et al. \cite{kinney}) )
+ of IUE observations of starburst galaxies which is not complete m any seuse., of IUE observations of starburst galaxies which is not complete in any sense.
+ Duat aud Xu (1996)) have used samples of star forming ealaxies selected on their UVand FIR emissions leadiug ο very complicated biases., Buat and Xu \cite{buxu}) ) have used samples of star forming galaxies selected on their UV FIR emissions leading to very complicated biases.
+ Whereas the use of sample of ealaxies which may be strougly biased is probably not a Huitation to calibrate the plivsical link between the FIR ο UV flux ratio and the extinction. he presence of these jdases must be accounted for when generic properties of oealaxies are deduced from these samples.," Whereas the use of sample of galaxies which may be strongly biased is probably not a limitation to calibrate the physical link between the FIR to UV flux ratio and the extinction, the presence of these biases must be accounted for when generic properties of galaxies are deduced from these samples."
+ Ou purpose is to use our FIR selected galaxy sample o test the influence of such a selection on the deduced value of the FIR to UV ratio., Our purpose is to use our FIR selected galaxy sample to test the influence of such a selection on the deduced value of the FIR to UV ratio.
+ We will consider both fluxcs at GO jan and in the range. 10-120 4422. the so-called FIR Hux.," We will consider both fluxes at 60 $\mu$ m and in the range 40-120 $\mu$ m, the so-called FIR flux."
+ Each oue has its own advantages: ou one haud more ealaxies have a measured flux at GO jan than at 100421. and the luninosity fuuction has been derived at 60. un. onan other haud the FIR cussion over the range 10.120 jaa ds iore easily related to the total emission of the dust rence to the amount of extinction than a single baud Iu," Each one has its own advantages: on one hand more galaxies have a measured flux at 60 $\mu$ m than at $\mu$ m and the luminosity function has been derived at 60 $\mu$ m, on an other hand the FIR emission over the range 40–120 $\mu$ m is more easily related to the total emission of the dust and hence to the amount of extinction than a single band flux."
+ this section. we ouly discuss the observational biases herefore use the data at 60 gn. In figure 2 is reported the ratio of fluxes at GO jan and 0.2 pau. Foy/Fy2 as a Muction of the flux aud luuinosity of the galaxies at 60 Foy and Fy.» are of the form Fy=Acἓν where fq is à flux per uni waveleusth.," In this section, we only discuss the observational biases and therefore use the data at 60 $\mu$ m. In figure 2 is reported the ratio of fluxes at 60 $\mu$ m and 0.2 $\mu$ m, $\rm F_{60}/F_{0.2}$ as a function of the flux and luminosity of the galaxies at 60 $\mu$ m. $\rm F_{60}$ and $\rm F_{0.2}$ are of the form $\rm
+F_\lambda = \lambda \cdot f_\lambda $ where $\rm f_{\lambda}$ is a flux per unit wavelength."
+ The figure 2a with the fiux of the galaxies can be used to study the selection bias iu liuüted flux samples., The figure 2a with the flux of the galaxies can be used to study the selection bias in limited flux samples.
+ The figureOo 2b where are reported the —ninosities ofthe galaxies is useful to discuss the intrinsic properties of the galaxies., The figure 2b where are reported the luminosities of the galaxies is useful to discuss the intrinsic properties of the galaxies.
+ There is a clear trend in both figures in the seuse of a larger Fou/Fy.2 ratio for brighter galaxies at GO san. The tail found in figure 2a at large 60 yan flux toward low Fou/Fy.2 ratios is due to very nearby galaxies.," There is a clear trend in both figures in the sense of a larger $\rm 
+F_{60}/F_{0.2}$ ratio for brighter galaxies at 60 $\mu$ m. The tail found in figure 2a at large 60 $\mu$ m flux toward low $\rm F_{60}/F_{0.2}$ ratios is due to very nearby galaxies."
+ This effect of distance disappears when the huuinositv is considered (fieure 2b)., This effect of distance disappears when the luminosity is considered (figure 2b).
+ In order to lighheht the ecneral trend we have calculated a moving median on the sample., In order to highlight the general trend we have calculated a moving median on the sample.
+ The data are sorted according to the 60 422 luminosity. then a mecdian is calculated for bius of 11 objects each time shifted by 5 objects.," The data are sorted according to the 60 $\mu$ m luminosity, then a median is calculated for bins of 11 objects each time shifted by 5 objects."
+ The result is shown in figure 3., The result is shown in figure 3.
+ As expected the moving median has reduced the dispersion of the data aud flattened the dispersed trend of the figure 2h., As expected the moving median has reduced the dispersion of the data and flattened the dispersed trend of the figure 2b.
+ A linear fit elves These figures illustrate the bias introduced by a FIR selection., A linear fit gives These figures illustrate the bias introduced by a FIR selection.
+ As we consider galaxies with an increasing 60 jan flux or luninosity. their FoyΕνω ratio also increases andl is less aud less representative of the mean properties of the local Universe as we will see below.," As we consider galaxies with an increasing 60 $\mu$ m flux or luminosity, their $\rm 
+F_{60}/F_{0.2}$ ratio also increases and is less and less representative of the mean properties of the local Universe as we will see below."
+ The case of these galaxies is especially interesting since they are good candidates for very obscured galaxies., The case of these galaxies is especially interesting since they are good candidates for very obscured galaxies.
+ Nevertheless their low number (5 cases. section 2.2) makes them having5 no or little influence on the statistical properfies discussed in this paper.," Nevertheless their low number (5 cases, section 2.2) makes them having no or little influence on the statistical properties discussed in this paper."
+ Moreover. little information is known about these objects.," Moreover, little information is known about these objects."
+ Onlv one ealaxv (F12212)0919) has a known redshift in the NED database., Only one galaxy (F12242+0919) has a known redshift in the NED database.
+ The ΈσΕνω ratio of cach object is reported iu able 2 and plotted in figure 2a., The $\rm F_{60}/F_{0.2}$ ratio of each object is reported in table 2 and plotted in figure 2a.
+" The ""upper linuits ound for these galaxies are compatible with the values ond for some galaxies of the TRAS/FOCA sample but heir location in the figure is surprising since they do rot follow the general (although dispersed) trend of a arecr Foy fux for a larger Fouξυυ ratio."," The upper limits found for these galaxies are compatible with the values found for some galaxies of the IRAS/FOCA sample but their location in the figure is surprising since they do not follow the general (although dispersed) trend of a larger $\rm F_{60}$ flux for a larger $\rm 
+F_{60}/F_{0.2}$ ratio."
+ However. oulv he fleure 2b where the huninosity of the galaxies are reported has a pliysical meaning and unfortunately only one object (F12212|0919) has a measured redshift.," However, only the figure 2b where the luminosity of the galaxies are reported has a physical meaning and unfortunately only one object (F12242+0919) has a measured redshift."
+ Since F12235|0911 aud FI2259|11tL are classified by. Yuan, Since F12235+0914 and F12259+1141 are classified by Yuan
+Classification with Bayes’ rule niiunuizes the total umber of uuisclassificatious. if the true distribution of class- probabilities p(x[Q;) is used (?7)..,"Classification with Bayes' rule minimizes the total number of misclassifications, if the true distribution of class-conditional probabilities $p(\vec{x}|\Omega_i)$ is used \citep{Hand:1981,Anderson:1984}."
+ Using Bayes’ theorem. posterior probabilities P(Q;|a) can be calculated.," Using Bayes' theorem, posterior probabilities $P(\Omega_i|\vec{x})$ can be calculated."
+ A spectrum of unknown class. with given feature vector a. can then be classified using Baves” rule: Tn inost of the classification problems arising iu the IES it ix desired. to eather a sample of objects of a specific class. or a specific set of classes.," A spectrum of unknown class, with given feature vector $\vec{x}$, can then be classified using Bayes' rule: In most of the classification problems arising in the HES it is desired to gather a sample of objects of a specific class, or a specific set of classes."
+ In these cases. Bayes’ rule is not appropriate. because we do not want to miuinize he total uunber of misclassificatious. but the musclassificatious between the desired classtes) ofobjects. aud the remiaiuiug classes.," In these cases, Bayes' rule is not appropriate, because we do not want to minimize the total number of misclassifications, but the misclassifications between the desired class(es) of objects, and the remaining classes."
+ Suppose we have three classes. A-. F-. aud C-type stars. and we want to eather a complete sample of A-tvpe stars.," Suppose we have three classes, A-, F-, and G-type stars, and we want to gather a complete sample of A-type stars."
+ Then oulv iisclassificatious between A-type stars and F- aud C-type stars (aud vice versa) are of interest., Then only misclassifications between A-type stars and F- and G-type stars (and vice versa) are of interest.
+ More specifically. misclassificatious of A-type stars to F- and C-type stars (leadiug to incompleteness) are least desirable when a complete sample shall be gathered. aud erroneous classification of F- and C-type stars as A-type stars (resulting in sample contanunation) cau be accepted at a amocderate rate.," More specifically, misclassifications of A-type stars to F- and G-type stars (leading to incompleteness) are least desirable when a complete sample shall be gathered, and erroneous classification of F- and G-type stars as A-type stars (resulting in sample contamination) can be accepted at a moderate rate."
+ Misclassificatious between F- aud C-ype stars can be totally ignored. because the target object ype is not involved.," Misclassifications between F- and G-type stars can be totally ignored, because the target object type is not involved."
+ Classification aims like this can be realized by using a Μπα cost rule., Classification aims like this can be realized by using a minimum cost rule.
+ Cost factors rj4. with allow to assign relative weights o individual types of wisclassifications.," Cost factors $r_{hk}$, with allow to assign relative weights to individual types of misclassifications."
+" The cost factor ry, is the relative weight of a uuisclassification from class 95 to class ο.", The cost factor $r_{hk}$ is the relative weight of a misclassification from class $\Omega_h$ to class $\Omega_k$.
+ Sppose we have au object of uuknown class. with cature vector à.," Suppose we have an object of unknown class, with feature vector $\vec{x}$."
+" We ask how large the cost is if it belougs ο class O,. and would be assigned to class ο. hzkb."," We ask how large the cost is if it belongs to class $\Omega_h$, and would be assigned to class $\Omega_k$, $h\not= k$."
+ The cost CjGe) dx: Tn the last step we have used the abbreviations Γον)=ay and pla|Qn)=pla)., The cost $C_{h\to k}(\vec{x})$ is: In the last step we have used the abbreviations $P(\Omega_h)=a_h$ and $p(\vec{x}|\Omega_h)=p_h(\vec{x})$.
+ We do uot kuow to which of the possible classes O5. /=1.....Πρι the object actually belongs.," We do not know to which of the possible classes $\Omega_h$ , $h = 1,\dots,n_c$, the object actually belongs."
+" Therefore. we estimate the expected cost κα)[ for assigningeuim an object with feature vector x to the class O,. by computing the following suuni of costs: Now we can formulate the imuunuuni cost rule. which nudinizes the total cost (?2).."," Therefore, we estimate the expected cost $C_k(\vec{x})$ for assigning an object with feature vector $\vec{x}$ to the class $\Omega_k$ by computing the following sum of costs: Now we can formulate the minimum cost rule, which minimizes the total cost \citep{Hand:1981}."
+" Tf the cost factors are chosen such that ry),= the iain. cost rule classification is identical to classification using Dawes rule."," If the cost factors are chosen such that $r_{hk}\equiv\delta_{hk}$ , the minimum cost rule classification is identical to classification using Bayes' rule."
+" Tn this case hne cost for assigning the class O, to a spectrum with feature vector x is the probability that the object belongs to one of the other classes fizKk.", In this case the cost for assigning the class $\Omega_k$ to a spectrum with feature vector $\vec{x}$ is the probability that the object belongs to one of the other classes $h\not= k$.
+ This follows inunediatelv from Eq. C)., This follows immediately from Eq. \ref{V_k}) ).
+" Tf rn,mὃν. the total πο of musclassifications is miuimuizel. so that he quality of a nini cost rule classification has to he evaluated by other criteria."," If $r_{hk}\not=\delta_{hk}$, the total number of misclassifications is minimized, so that the quality of a minimum cost rule classification has to be evaluated by other criteria."
+ For auv given classification aim. one can divide the cost factors to be chosen iuto three sets: Since seauple completeness and contamination are interdependent. in. practice only the relative value has to be adjusted.," For any given classification aim, one can divide the cost factors to be chosen into three sets: Since sample completeness and contamination are interdependent, in practice only the relative value has to be adjusted."
+ For this purpose. the classification results as ai function of are evaluated.," For this purpose, the classification results as a function of are evaluated."
+ The expected error rates; estimated e.g. with the ‘leaving one out method (see Sect.," The expected error rates, estimated e.g. with the “leaving one out” method (see Sect."
+ ?? below). tell which level of completeness aud sample contamination will be achieved.," \ref{Sect:Evalu} below), tell which level of completeness and sample contamination will be achieved."
+ ?. preseuted a software tool for a convenieut choice of cost factors., \cite{ncADASSVII} presented a software tool for a convenient choice of cost factors.
+ Nowinathematically speaking. Bayes” rule assigus the class with the lighest relative resemblance to cach spectitun to be classified.," Non-mathematically speaking, Bayes' rule assigns the class with the highest relative resemblance to each spectrum to be classified."
+ However. it is ignorant of the absolute rescmblance: A spectrum with feature vector x inav be assigned to a cass with very low posterior wobability. p(9;aw). if p(Q;α) Is even lower for all other classes.," However, it is ignorant of the absolute resemblance: A spectrum with feature vector $\vec{x}$ may be assigned to a class with very low posterior probability $p(\Omega_i|\vec{x})$, if $p(\Omega_i|\vec{x})$ is even lower for all other classes."
+" This mecaus that a class is assigned to all spectra. even to ""earbase spectra” which are disturbed. or instance. bv plate artifacts."," This means that a class is assigned to all spectra, even to “garbage spectra” which are disturbed, for instance, by plate artifacts."
+ Therefore. it is useful to apply a rejection rule in addition to either the Daves rule or theminium cost rule.," Therefore, it is useful to apply a rejection rule in addition to either the Bayes rule or theminimum cost rule."
+ The rejection rule cau also ο used “stand alone for he identification of peculiar objects. e.g. quasars.," The rejection rule can also be used “stand alone” for the identification of peculiar objects, e.g., quasars."
+ALACPOP REN for financial support.,MAGPOP RTN for financial support.
+where VoSl. NLumucquivfrac2x(35 Recent high resolution numerical studies of pre-galactie objects indicate that massive stars (~LOO AL.) Dorm at the centers of progenitor dark matter halos atredshifts :=15—20 (Abel. Drvan. Norman 2000. 2001).,"where v_s^2 _s; 	Recent high resolution numerical studies of pre-galactic objects indicate that massive stars $\sim 100 \ M_{\odot}$ ) form at the centers of progenitor dark matter halos atredshifts $z=15-20$ (Abel, Bryan, Norman 2000, 2001)."
+ Subsequent evolution will end in supernovae on Myr timescales. leaving a populationof remnant (50 M.) seed black holes. long before the first galactic mass dark halos have formed. (Alaclan Rees 2001).," Subsequent evolution will end in supernovae on Myr timescales, leaving a populationof remnant $\sim 50 \ M_{\odot}$ ) seed black holes, long before the first galactic mass dark halos have formed (Madau Rees 2001)."
+ Fherelore. the material that accretes and merges to form galactic mass dark halos will be well seeded with black holes already. accreting dark matter on dvnamical timescales shorter than the timescale for SIDA collisional evolution.," Therefore, the material that accretes and merges to form galactic mass dark halos will be well seeded with black holes already accreting dark matter on dynamical timescales shorter than the timescale for SIDM collisional evolution."
+ Following Ostriker (2000). we consider the quasi-spherical accretion ol dark matter onto a single seed black hole at the center of a galactic dark matter halo.," Following Ostriker (2000), we consider the quasi-spherical accretion of dark matter onto a single seed black hole at the center of a galactic dark matter halo."
+ Clearly. (he assumption of a single seed black hole al the halo center is an over-simplification.," Clearly, the assumption of a single seed black hole at the halo center is an over-simplification."
+ llowever. given the rapid phase of initial growth (see below). one black hole is likely to dominate and will eat or eject the others. so that this complication should not alter our estimates of the final black hole mass.," However, given the rapid phase of initial growth (see below), one black hole is likely to dominate and will eat or eject the others, so that this complication should not alter our estimates of the final black hole mass."
+" The dark matter is treated as an adiabatic gas. which is valid provided the optical depth diverges lor r««r, (an assumption we check below)."," 	The dark matter is treated as an adiabatic gas, which is valid provided the optical depth diverges for $r<<r_s$ (an assumption we check below)."
+ The density. and velocity dispersion can be written rs," The density and velocity dispersion can be written . ,"
+ and, and
+start of the simulation.,start of the simulation.
+" As soon as material from the core is mixed into the envelope, small convective cells develop."," As soon as material from the core is mixed into the envelope, small convective cells develop."
+" At this early stage, the fluid has a large Reynolds number, with a characteristic eddy length of 50 km, fluid velocities ranging between v = 10?—10° cm s-!, and a dynamic of 2x10 P The fluid velocity v remains below the speed of sound Cs (that is, the Mach number Mazν/ος is always less than unity, see Fig."," At this early stage, the fluid has a large Reynolds number, with a characteristic eddy length of 50 km, fluid velocities ranging between $v$ = $10^5 - 10^6$ cm $^{-1}$, and a dynamic of $2 \times 10^4$ P The fluid velocity $v$ remains below the speed of sound $c_{\rm s}$ (that is, the Mach number $=v/c_{\rm s}$ is always less than unity, see Fig."
+" 2), hence, the fluid displays a deflagration rather than a detonation — see Williams (1985), for a thorough analysis of the differences between flame propagation under detonation and deflagration conditions."," 2), hence, the fluid displays a deflagration rather than a detonation — see Williams (1985), for a thorough analysis of the differences between flame propagation under detonation and deflagration conditions."
+" At t=235 s, The characteristic"," At $t = 235$ s, The characteristic"
+eeneralion of the strong toroidal field by the stretching of field lines must be taking place there.,generation of the strong toroidal field by the stretching of field lines must be taking place there.
+ However. it seems unlikely (hat the whole diamo process (as envisaged in (he PSAR approach) occurs at the bottom of the convection zone.," However, it seems unlikely that the whole dynamo process (as envisaged in the PSKR approach) occurs at the bottom of the convection zone."
+ The studies of buovant rise of the toroidal [lix from there suggest (hat the toroidal field at the bottom of the convection zone should be of the order 10? G. substantially stronger than the equipartition value (Chondhuri Gilman 1937: Choudhuri 1939: D'Silva Choudhuri 1993: Fan. Fisher. DeLuca 1993: Caligari. Moreno-Insertis. Schüsssler 1995).," The studies of buoyant rise of the toroidal flux from there suggest that the toroidal field at the bottom of the convection zone should be of the order $10^5$ G, substantially stronger than the equipartition value (Choudhuri Gilman 1987; Choudhuri 1989; D'Silva Choudhuri 1993; Fan, Fisher, DeLuca 1993; Caligari, Moreno-Insertis, Schüsssler 1995)."
+ Such a strong field would completely quench the o-elfect of the PSIXR. approach., Such a strong field would completely quench the $\alpha$ -effect of the PSKR approach.
+ To explain the generation of the poloidal field. the most natural wav is to invoke the DL idea of the decay. of tilled active regions. though there are still some allempts to work within the PSAR approach by considering an interface dvnamo (Parker 1993: Charbonneau MacGregor 1997; Markiel Thomas 1999).," To explain the generation of the poloidal field, the most natural way is to invoke the BL idea of the decay of tilted active regions, though there are still some attempts to work within the PSKR approach by considering an interface dynamo (Parker 1993; Charbonneau MacGregor 1997; Markiel Thomas 1999)."
+ In this paper. we assune that the poloidal field is produced by the decay of tilted active regions near the solar surface.," In this paper, we assume that the poloidal field is produced by the decay of tilted active regions near the solar surface."
+ Although the sunspots migrate equatorward with the solar evele. the weak cilfuse magnetic field on the solar surface migrates poleward (Dumba Howard 1965: Howard LaBonte 1981: Makarov. Fatianov. Sivaraman 1983: Makarov Sivaraman 1989).," Although the sunspots migrate equatorward with the solar cycle, the weak diffuse magnetic field on the solar surface migrates poleward (Bumba Howard 1965; Howard LaBonte 1981; Makarov, Fatianov, Sivaraman 1983; Makarov Sivaraman 1989)."
+ Mos of the dvnamo models based on the PSAR approach (starting Ivom Steenbeck Ixrause 1969) mainly concentrated on the sunspols and ignored (the poleward migration of the weak diffuse field., Most of the dynamo models based on the PSKR approach (starting from Steenbeck Krause 1969) mainly concentrated on the sunspots and ignored the poleward migration of the weak diffuse field.
+ The poleward migration has been explained bv assuming that the weak diffuse field (which is essentially the poloidal field) is carried by the meridional circulation (Wang. Nash. Sheeley 1989a. 1989b: Dikpati Choudhuri 1994. 1995: Choudhuri Dikpati 1999).," The poleward migration has been explained by assuming that the weak diffuse field (which is essentially the poloidal field) is carried by the meridional circulation (Wang, Nash, Sheeley 1989a, 1989b; Dikpati Choudhuri 1994, 1995; Choudhuri Dikpati 1999)."
+ II we now accept the DL idea that the poloidal field is produced by the decay of tilted bipolar active regions. then the meridional circulation should play an important role in (he dvnamo problem by bringing the poloidal field from the surface to the bottom οἱ the convection zone. where the poloidal field is stretched out to produce the toroidal field.," If we now accept the BL idea that the poloidal field is produced by the decay of tilted bipolar active regions, then the meridional circulation should play an important role in the dynamo problem by bringing the poloidal field from the surface to the bottom of the convection zone, where the poloidal field is stretched out to produce the toroidal field."
+ The challenge before us now is to develop a new (vpe of clvnamo model. in which the," The challenge before us now is to develop a new type of dynamo model, in which the"
+Tihank Julio Chanamé.. Subo Dong. and David Weinberg for valuable discussions.,"I thank Julio Chanamé,, Subo Dong, and David Weinberg for valuable discussions."
+ This work was supported by grant AST 02-01266 [rom the NSF., This work was supported by grant AST 02-01266 from the NSF.
+Massive. early-type binaries with powerful stelay winds can generate a complex. region of shock-heated plasma with temperatures in excess of 10* I whei these wiids collide (Prilutxkiü Usov 1976: Cherepasheclak 1976)).,"Massive, early-type binaries with powerful stellar winds can generate a complex region of shock-heated plasma with temperatures in excess of $10^{7}$ K when these winds collide (Prilutskii Usov \cite{PU1976}; ; Cherepashchuk \cite{C1976}) )."
+ A-rav observations provide a direct probe of the conditions within the coliNon Zone ux also of the uushocked attenuatiue niacrial along the ine of sieht in the svsteim Stevens 1992:: Pollock 1999: Pitti Stevens 19973)., X-ray observations provide a direct probe of the conditions within the collision zone and also of the unshocked attenuating material along the line of sight in the system Stevens \cite{SBP1992}; Pollock \cite{Po1999}; Pittard Stevens \cite{PS1997}) ).
+ Single carly-ty2ο and Wol-Roavet (WR) stars are also shown to be X-rav enuütters, Single early-type and Wolf-Rayet (WR) stars are also known to be X-ray emitters.
+ The most viable inodel arenly proposed is that the 1010* KK A-ravs are generated by shocss du the wind resulting from the wustade nature of the line-dyviven acceleration Owocki ]O88: Feldiucier 19973)., The most viable model currently proposed is that the $10^{6} - 10^{7}$ K X-rays are generated by shocks in the wind resulting from the unstable nature of the line-driven acceleration Owocki \cite{OCR1988}; Feldmeier \cite{FPP1997}) ).
+ Direct spectroscopic evidence of ehuiips/shocks in WR winds was first noted by Moffa (1988))., Direct spectroscopic evidence of clumps/shocks in WR winds was first noted by Moffat \cite{MDLR1988}) ).
+ Duissonu frou colliding winds ids often observable as a strong lieh-temperature N-rav excess above that expected from the iidividual stars. together with plase-locked orbital variabiliv resultiug from t16 changing line of sight iuto the svsteu Pollock 1987: Chlebowskli 19589: Williams 1990: Chlebowsü Carmany 1991:: Corcoran 19963).," Emission from colliding winds is often observable as a strong high-temperature X-ray excess above that expected from the individual stars, together with phase-locked orbital variability resulting from the changing line of sight into the system Pollock \cite{Po1987}; Chlebowski \cite{C1989}; Williams \cite{W1990}; Chlebowski Garmany \cite{CG1991}; Corcoran \cite{C1996}) )."
+ Receut Chandraoservations Lave revealed spectral featires which are noticeably differeut to those from sinele carly-tyvpe stars. dicludiug strong forbidden Lae ↸∖↕⋯↴∖∷∖↴↕∪∐⋖≼∪↥⋅↸⊳≺ ennerran .2001:: Pollock 2002:: Shi 20013).," Recent observations have revealed spectral features which are noticeably different to those from single early-type stars, including strong forbidden line emission (Corcoran \cite{CISP2001}; Pollock \cite{Po2002}; Skinner \cite{SGSS2001}) )."
+ Tn all previous N-ray observations. the colliling wiud emissoji has remained spatially unresolved. and contaminated w the emission from the individual wiids.," In all previous X-ray observations, the colliding wind emission has remained spatially unresolved, and contaminated by the emission from the individual winds."
+ Nou-thermal radio enuissionu frou WR stars is also a eood indicator of wiud-wiud interaction Dougherty Williams 2000))., Non-thermal radio emission from WR stars is also a good indicator of wind-wind interaction Dougherty Williams \cite{DW2000}) ).
+ Such «auission ds thought to be svuchrotron radiation arising from electrous accelerated to relativistic velocities in the wind-wind collisiou zo1ο (Eichler Usov 1993))., Such emission is thought to be synchrotron radiation arising from electrons accelerated to relativistic velocities in the wind-wind collision zone (Eichler Usov \cite{EU1993}) ).
+ The ]xeseuce of nouthenual radio ewission from WR 117 (AS 13lL) was reported by CaillatIt (1985)) and Abbott (1986))., The presence of nonthermal radio emission from WR 147 (AS 431) was reported by Caillault \cite{C1985}) ) and Abbott \cite{ABCT1986}) ).
+" Later observatioIs withAZERLIN at 5 CIIz resolved two radio Sources (WR L178 aud WR L17N) separated by. ~O.6""(Moran 19893).", Later observations with at 5 GHz resolved two radio sources (WR 147S and WR 147N) separated by $\sim 0.6$ (Moran \cite{M1989}) ).
+ The southeru hermal source (WR 1178) was coincident with the WwRh star. and the northern Xorce (AVR 117N) was proposed to be responsible for the OServed nouthermal ciissic1.," The southern thermal source (WR 147S) was coincident with the WR star, and the northern source (WR 147N) was proposed to be responsible for the observed nonthermal emission."
+" Independent observations with the confirmed this lvpothesis (Churchwell .1992:: Skimmer 10901}, ", Independent observations with the confirmed this hypothesis (Churchwell \cite{C1992}; Skinner \cite{SINZ1999}) ).
+Moran (1989)) further sugeested that WR 117 couk have an unuiseeu companion associated with the northern racio source., Moran \cite{M1989}) ) further suggested that WR 147 could have an unseen companion associated with the northern radio source.
+ WR LIT attracted the attention of many astronomers when radio. infrared and optical observatioi atf lieh spatial resoltion provided unaibiguous evience of its Dinarityv (Wiljus 1997. . hereafter W97: Nieinela 1998.. hereafter N98).," WR 147 attracted the attention of many astronomers when radio, infrared and optical observations at high spatial resolution provided unambiguous evidence of its binarity (Williams \cite{W1997}, , hereafter W97; Niemela \cite{N1998}, hereafter N98)."
+ New. higher resolutioiMERLIN observations w WOT revealed that at 5 GIIz both sources are elongated with sizes of —170«253 mas (WB 11758) and ~267&79 mas (WR 1LGN). and that they had a separation of 575+15 mas.," New, higher resolution observations by W97 revealed that at 5 GHz both sources are elongated with sizes of $\sim 170 \times 253$ mas (WR 147S) and $\sim 267 \times 79$ mas (WR 147N), and that they had a separation of $575 \pm 15$ mas."
+ TheCATRT iutrared audIIST optical nuages revealed for the first time a faint conrpanion sar located to the north of the routhermal source associated with WR LITN. Crucially. he relative positions of tje sources revealed thatthe companion star. althoughclose to the nouthermal radio cussion. was located sliehtlv. (2 60 mas) more distant from the WR star.," The infrared and optical images revealed for the first time a faint companion star located to the north of the nonthermal source associated with WR 147N. Crucially, the relative positions of the sources revealed thatthe companion star, althoughclose to the nonthermal radio emission, was located slightly $\approx$ 60 mas) more distant from the WR star."
+ The resulting picture is thus consistent with the, The resulting picture is thus consistent with the
+We acquired optical spectra for 155 galaxies in the region of Abell 1644 curing 2000 January-June with FAST. a high throughput slit spectrometer mounted on the 1.5 m telescope abt FLWO.,"We acquired optical spectra for 155 galaxies in the region of Abell 1644 during 2000 January-June with FAST, a high throughput slit spectrometer mounted on the 1.5 m telescope at FLWO."
+ The FAST setup includes a 300 groove ! erating and 3” slit. and vields spectra which cover at a resolution of ~GA..," The FAST setup includes a 300 groove $^{-1}$ grating and $^{\prime \prime}$ slit, and yields spectra which cover at a resolution of $\sim6$."
+ Although DresslerhereafterDS) measured 103 redshifts in the region of Abell 1644. we re-acquired spectra of 91 of their objects to assure uniformity throughout our data set ancl to assess the spectral ivpe of the galaxies.," Although \citet{dre8a} measured 103 redshifts in the region of Abell 1644, we re-acquired spectra of 91 of their objects to assure uniformity throughout our data set and to assess the spectral type of the galaxies."
+ The remaining twelve DS galaxies are either too faint for FAST redshilt measurements or are outside our survey region., The remaining twelve DS galaxies are either too faint for FAST redshift measurements or are outside our survey region.
+ We selected 64 aclditional targets on the basis of the DS catalog. digitized skv survey images. ancl our own inlrared survey.," We selected 64 additional targets on the basis of the DS catalog, digitized sky survey images, and our own infrared survey."
+ We obtained a reduced spectrum lor each object using IRAF., We obtained a reduced spectrum for each object using IRAF.
+ We made bias. and ilhunination corrections to each frame and then used the IRAF ACSAO task to obtain redshifts via cross-correlation with stancard eniission ancl absorption. templates.," We made bias, flat-field, and illumination corrections to each frame and then used the IRAF XCSAO task to obtain redshifts via cross-correlation with standard emission and absorption templates."
+ The resulting spectroscopic sample is complete to 7/7~13 in (he region where we have photometry., The resulting spectroscopic sample is complete to $H\sim13$ in the region where we have }-band photometry.
+ We tested the differences between the FAST and DS observations lor the 91 common galaxies., We tested the differences between the FAST and DS observations for the 91 common galaxies.
+" The FAST velocities are svstematically lower by ~39 kms ο, and the ris velocity cilference is NC...—80 kms |."," The FAST velocities are systematically lower by $\sim39$ km $^{-1}$, and the rms velocity difference is $\Delta v_{rms}=80$ km $^{-1}$ ."
+ These offsets ave consistent with the estimated velocity errors of roughly 30 kms ! (FAST) and 45 kins ! (DS)., These offsets are consistent with the estimated velocity errors of roughly 30 km $^{-1}$ (FAST) and 45 km $^{-1}$ (DS).
+ Table 2. lists properties of (he 155 galaxies with measured redshifts., Table \ref{table2} lists properties of the 155 galaxies with measured redshifts.
+ Columns (1) aud (2) give the right ascension and declination: column (3). the racial velocity: column (4). the error on the radial velocity: column (5). the spectral tvpe (either enission or absorption as discussed in Section 3.1): column (6). DDBEST Gf available): and column (7). the reference number from the DS catalog for galaxies in the DS list.," Columns (1) and (2) give the right ascension and declination; column (3), the radial velocity; column (4), the error on the radial velocity; column (5), the spectral type (either emission or absorption as discussed in Section \ref{spectral}) ); column (6), BEST (if available); and column (7), the reference number from the DS catalog for galaxies in the DS list."
+ Figure 2aa shows the measured. velocity distribution., Figure \ref{fig-vhist}a a shows the measured velocity distribution.
+ ALG4 is remarkably isolated in redshift space., A1644 is remarkably isolated in redshift space.
+ A lew background galaxies are visible at redshifts up to ez~65.000 km 1.," A few background galaxies are visible at redshifts up to $cz\sim65,000$ km $^{-1}$."
+ There is also evidence of foreground structure in the range 2500 «ο< 10.000 kin sf.," There is also evidence of foreground structure in the range 2500 $< cz
+<$ 10,000 km $^{-1}$."
+ Figure 2bb shows the ΓΗ] cluster members with redshifts between 10.000 and 20.000 kms !.," Figure \ref{fig-vhist}b b shows the 141 cluster members with redshifts between 10,000 and 20,000 km $^{-1}$."
+ The mean radial velocity of the cluster is ez=14.295493 kms ! with a dispersion of σ=1108zi km IF. computed at the confidence limit as in (1980).," The mean radial velocity of the cluster is $c\overline{z}=14,295\pm93$ km $^{-1}$ with a dispersion of $\sigma=1108^{+73}_{-61}$ km $^{-1}$, computed at the confidence limit as in \citet{dan80}."
+. The radial velocity of the «D (14.233433 lan 1!) is coincident with the overall cluster mean.," The radial velocity of the cD $14,233\pm33$ km $^{-1}$ ) is coincident with the overall cluster mean."
+ The cluster sample is isolated fromloreeround and background galaxies by velocity gaps of 3244 km ! on the low end and 4589 km ! on the high end., The cluster sample is isolated fromforeground and background galaxies by velocity gaps of 3244 km $^{-1}$ on the low end and 4589 km $^{-1}$ on the high end.
+ The sizes of, The sizes of
+hese two assumptions are valid. for isolated. haloes tvpical of the local Universe. at high. redshift when the Universe was denser. these assumptions are incorrect.,"these two assumptions are valid for isolated haloes typical of the local Universe, at high redshift when the Universe was denser, these assumptions are incorrect."
+ In addition to hese two assumptions. it is possible that systematic errors rom the halo finding algorithm may bias measurements of ido virialization.," In addition to these two assumptions, it is possible that systematic errors from the halo finding algorithm may bias measurements of halo virialization."
+ Finally. we explore the possibility that he departure from virialization is correlated with kev halo operties.," Finally, we explore the possibility that the departure from virialization is correlated with key halo properties."
+ The organization of this paper is as follows: we derive the virial equation from the momentum equation in section 2.. explicith retaining the surface terms that are usually neglected.," The organization of this paper is as follows: we derive the virial equation from the momentum equation in section \ref{sec:ve}, explicitly retaining the surface terms that are usually neglected."
+ In. section. 3 present simulation results and the correlations between halo energeties and. kev halo structural properties. namely. halo spin. concentration. and local environment.," In section \ref{sec:sim} present simulation results and the correlations between halo energetics and key halo structural properties, namely, halo spin, concentration, and local environment."
+ We conclude in section 4 with a discussion of our findings and their implications., We conclude in section \ref{sec:discussion} with a discussion of our findings and their implications.
+ llere we derive the virial theorem for dark matter haloes including boundary effects. ancl an external gravitational potential., Here we derive the virial theorem for dark matter haloes including boundary effects and an external gravitational potential.
+ An isolated. system: of collisionless particles in equilibrium satisfies the scalar virialtheorem (see Equation 1)., An isolated system of collisionless particles in equilibrium satisfies the scalar virialtheorem (see Equation \ref{EQ:scalarVE}) ).
+ This result. is derived: under the assumption that the mass density approaches zero at large radii for a finite mass clistribution., This result is derived under the assumption that the mass density approaches zero at large radii for a finite mass distribution.
+ In this work. our svstenis of interest are high redshift dark matter haloes. which form within a cosmic web of sheets. voids. and filaments.," In this work, our systems of interest are high redshift dark matter haloes, which form within a cosmic web of sheets, voids, and filaments."
+ We therefore cannot assume that the mass density approaches zero at the boundary of these systems., We therefore cannot assume that the mass density approaches zero at the boundary of these systems.
+ We also consider the potential due to nearby particles exterior to the The virial theorem is obtained from the conservation of momentum for a dark matter halo modeled. as a collisionless Duid., We also consider the potential due to nearby particles exterior to the The virial theorem is obtained from the conservation of momentum for a dark matter halo modeled as a collisionless fluid.
+ Following Binney&‘Tremaine(1987).. the corresponding [uid equation is multiplied. by. ary and integrated over the volume V. with the caveat that the surface terms resulting from. integration by parts do not vanish.," Following \citet{BTbook}, the corresponding fluid equation is multiplied by $x_k$ and integrated over the volume $V$, with the caveat that the surface terms resulting from integration by parts do not vanish."
+" In this case. the tensor virial theorem corresponding to a collisionless Εις is given by. where dj)=npror is the moment of inertia tensor. dvi;—sf,ρου,X is the kinetic energy. tensor. We,=liprjOdder is the potential energy. tensor."," In this case, the tensor virial theorem corresponding to a collisionless fluid is given by, where $I_{ij} \equiv \int_V{\rho x_i x_j \dd^3 \vec{x}}$ is the moment of inertia tensor, $K_{ij} \equiv \frac{1}{2}\int_V{\rho
+ v_iv_j \dd^3\vec{x}}$ is the kinetic energy tensor, $W_{ij} \equiv -
+ \int_V{\rho x_i\partial \Phi/\partial x_j \dd^3 \vec{x}}$ is the potential energy tensor."
+ We denote the mass fordensity. position. average velocity. ancl eravitational potential by p. αἲ £F. and & respectively.," We denote the mass density, position, average velocity, and gravitational potential by $\rho$, $\vec{x}$, $\vec{v}$ , and $\Phi$ respectively."
+ The vector 5H ds the outward-pointing unit normal to the surface S that encloses the volume V., The vector $\vec{n}$ is the outward-pointing unit normal to the surface $S$ that encloses the volume $V$.
+ The third term on the right of equation (2)) corresponds to the kinetic stresses on the halo boundary., The third term on the right of equation \ref{EQ:tensorvir}) ) corresponds to the kinetic stresses on the halo boundary.
+ The last term is the time derivative of the moment of inertia Hux through the surface. (Ballesteros-Paredes 2006)., The last term is the time derivative of the moment of inertia flux through the surface \citep{BP06}.
+. The scalar virial theorem is obtained by taking the trace of equation (2)):Following Binney&‘Tremaine(1987) anc Ballesteros-Paredes(2006).. the potential energv term may be broken up into contributions from particles inside and outside of the halo. ®=diu|dou. so that Note that the first term is the gravitational potential energy of the halo. which we denote as C from here on.," The scalar virial theorem is obtained by taking the trace of equation \ref{EQ:tensorvir}) ):Following \citet{BTbook} and \citet{BP06}, the potential energy term may be broken up into contributions from particles inside and outside of the halo, $\Phi = \Phi_{\mathrm{int}} + \Phi_{\mathrm{ext}}$, so that Note that the first term is the gravitational potential energy of the halo, which we denote as $U$ from here on."
+ In the case of a spherical volume. the last term in equation (3)) is I7ML/2. where 2 is the radius and A is the mass enclosed within the volume.," In the case of a spherical volume, the last term in equation \ref{EQ:scalarvir}) ) is $R^2 \ddot{M}/2$, where $R$ is the radius and $M$ is the mass enclosed within the volume."
+ Assuming that this term is negligible for ἃ steady-state system. we obtain [or a halo in virial equilibrium.," Assuming that this term is negligible for a steady-state system, we obtain for a halo in virial equilibrium."
+" (ox, ancl fy represent corrections to the standard. virial relation anc can be explicitly calculated for simulated dark matter halocs.", $U_{\mathrm{ext}}$ and $E_s$ represent corrections to the standard virial relation and can be explicitly calculated for simulated dark matter haloes.
+ Using CLADGIZI-2 (Springel2005).. we ran a dark matter only simulation with the WALADRS cosmological parameters (Qu.Oy.O8-feneas} = (00.258. 0.742. 0.044. 0.719. 0.963. from ςz100 t0 2=6 (sce Davis&Natarajan(2010). for full details).," Using GADGET-2 \citep{GADGET}, we ran a dark matter only simulation with the WMAP5 cosmological parameters \citep[\{$\Omega_M,
+\Omega_\Lambda, \Omega_b, h, n, \sigma_8$\} = 0.258, 0.742, 0.044, 0.719, 0.963, from $z \approx 100$ to $z = 6$ (see \citet{Davis10} for full details)."
+ Our clark matter particle mass is m=1.0103AL./h. which sets a comoving box size at 2.46Mpc/h. with 512? dark matter particles.," Our dark matter particle mass is $m = 1.0 \times 10^4 \Msunh$, which sets a comoving box size at $2.46 \Mpch$, with $512^3$ dark matter particles."
+ We use the LOD algorithm to identify. collapsed dark matter baloes (Eisenstein&Llut1998)., We use the HOP algorithm to identify collapsed dark matter haloes \citep{HOP}.
+. LIOP first calculates a density for each. particle bx smoothing over its nearest 64 neighbors using a cubic spline kernel., HOP first calculates a density for each particle by smoothing over its nearest 64 neighbors using a cubic spline kernel.
+ I0 then eroups particles with their clensest neighbor. and. density thresholds are used. to ensure that. haloes are. not. being over-counted as sub-haloes within a larger halo.," It then groups particles with their densest neighbor, and density thresholds are used to ensure that haloes are not being over-counted as sub-haloes within a larger halo."
+ We choose density thresholds to match the high redshift mass function described in Reectetal.(2007)., We choose density thresholds to match the high redshift mass function described in \citet{reed07}.
+. Lastly. we remove unbound particles from each halo.," Lastly, we remove unbound particles from each halo."
+ For cach halo. we calculate the total potential energy. U. using a direct summation over the particles assigned to the halo: where ry is the separation between particles / and j. €? is Newton's gravitational constant. and m the particle mass.," For each halo, we calculate the total potential energy, $U$, using a direct summation over the particles assigned to the halo: where $r_{ij}$ is the separation between particles $i$ and $j$, $G$ is Newton's gravitational constant, and $m$ the particle mass."
+ The total kinetic energy. Z/ is found by summing over the particles individual kinetic energies: 1n measuring { and C. we must ensure that we have enough particles in our halo sample to accurately measure the energies.," The total kinetic energy, $T$ is found by summing over the particle's individual kinetic energies: In measuring $T$ and$U$ , we must ensure that we have enough particles in our halo sample to accurately measure the energies."
+ We addressed this question in Davis&Natara- (2010).. and summarize the relevant. points here.," We addressed this question in \citet{Davis10}, , and summarize the relevant points here."
+ We re- our simulation with one eighth as many particles (256° , We re-ran our simulation with one eighth as many particles $256^3$ 
+the Evershed flow.,the Evershed flow.
+ Because of the rapid decline of deusity with height. this flow must be close to horizoutal. covering a vertical extent not exceeding a couple of pressure scale heights.," Because of the rapid decline of density with height, this flow must be close to horizontal, covering a vertical extent not exceeding a couple of pressure scale heights."
+ This iu turnheight. at least inits outer penumbra where the Evershed flow is observed.," This in turn, at least in its outer penumbra where the Evershed flow is observed."
+ Tn our niodel. this is guaranteed because it identifies he surface of the flament with the boundary between he peuunbral magnetic Geld and the convection zone )clow.," In our model, this is guaranteed because it identifies the surface of the filament with the boundary between the penumbral magnetic field and the convection zone below."
+ Pressure balance across this interface determines it» location iu depth., Pressure balance across this interface determines its location in depth.
+ The imaguetic pressure of the ovumanibra causes this depth o be depressed below the ional photosphere by not more than about one scale weight (somewhat more iu the ο peuunbra where 10 field strength increases to mubral values)., The magnetic pressure of the penumbra causes this depth to be depressed below the normal photosphere by not more than about one scale height (somewhat more in the inner penumbra where the field strength increases to umbral values).
+ A possible -1odel for the force driving he Evcrshe¢ flow is the ‘siphon effect described by Schimidt aud. Meyer., A possible model for the force driving the Evershed flow is the `siphon effect' described by Schmidt and Meyer.
+ The field becomes exactly horizontal just at the top of the eap. where the field lines followine the sides of the gap mect.," The field becomes exactly horizontal just at the top of the gap, where the field lines following the sides of the gap meet."
+ To be consistent with observations. this must happen close to the 7=1 laver.," To be consistent with observations, this must happen close to the $\tau=1$ layer."
+ This is also consistent with the heat fix requirement mentioned iu the Iutroduction: The observed leat flux cau be carried by radiation oulv if the surface of the gap is close to the coutinuuni 7=1 level (withiu a few teus of kin)., This is also consistent with the heat flux requirement mentioned in the Introduction: The observed heat flux can be carried by radiation only if the surface of the gap is close to the continuum $\tau=1$ level (within a few tens of km).
+lines. This is an intriguing possibility for explaining the ‘stray Πο contributions deduced. from observations. as well as the very low field streugths found in the deepest lavers in the inversions of Westendorp ct al..," This is an intriguing possibility for explaining the `stray light' contributions deduced from observations, as well as the very low field strengths found in the deepest layers in the inversions of Westendorp et al.,"
+ Bellot Rubio et al. (, Bellot Rubio et al. (
+200L). and Borrero et al. (,"2004), and Borrero et al. ("
+2005) (cf.,2005) (cf.
+ sect 2))., sect \ref{interp}) ).
+ Tn our analytic model. the shape of the gap is nof realistic: its top is approximately circular mstead of a cusp as sketched im rofeaps..," In our analytic model, the shape of the gap is not realistic: its top is approximately circular instead of a cusp as sketched in \\ref{gaps}."
+" This is because the analytic model accounts for the pressure balance between the maguctic field ane the eap ouly in au approximate wav,", This is because the analytic model accounts for the pressure balance between the magnetic field and the gap only in an approximate way.
+ Accurate calculations of field configurations in pressure balance are possible with the methods used for coustructing sunspot equilibria (Jahn and Schnudt 1991). but these are much uore demanding.," Accurate calculations of field configurations in pressure balance are possible with the methods used for constructing sunspot equilibria (Jahn and Schmidt 1994), but these are much more demanding."
+ We finally note that the nüagnetic field) iu the ΟΠΗ will of course not ο exactly a potential field: ou the one hand because it is not static. on the other jecause it ΙΣΤ contain material to provide the opacity seen in the lines.," We finally note that the magnetic field in the penumbra will of course not be exactly a potential field; on the one hand because it is not static, on the other because it must contain material to provide the opacity seen in the lines."
+ This eas exerts a pressure that varies with wieght. hence introducing some deviations from a poteutial field.," This gas exerts a pressure that varies with height, hence introducing some deviations from a potential field."
+ Given the relatively sinall amounts of gas needed to xovide line opacity. these deviations eau be sinall.," Given the relatively small amounts of gas needed to provide line opacity, these deviations can be small."
+ Due to the gap. the magnetic field streneth above it is reduced.," Due to the gap, the magnetic field strength above it is reduced."
+ This is evident from the shape of the field lines in rofpotfeld.. and shown explicitly in rofD-inc..," This is evident from the shape of the field lines in \\ref{potfield}, and shown explicitly in \\ref{B-inc}."
+ This is also in agrecment with the recent results of Langhans et al. (, This is also in agreement with the recent results of Langhans et al. (
+2005a).,2005a).
+ We have proposed that bright pemuubral filaments are caused by field.free convection just below the surface of the penunbra. and have discussed the theoretical aud observational evidence for this iuterpretation.," We have proposed that bright penumbral filaments are caused by field–free convection just below the surface of the penumbra, and have discussed the theoretical and observational evidence for this interpretation."
+ The field, The field
+In addition. the larger density in the compressed layer. by a actor of 4. tends to enhance recombination. and so for our models. we allow p5|1 in the shocked layer.,"In addition, the larger density in the compressed layer, by a factor of 4, tends to enhance recombination, and so for our models, we allow $p \rightarrow p + 1$ in the shocked layer."
+ Comparisons o detailed radiative transfer calculations. such as those of the Potsdam group or CMFGEN (Hillier Miller 1999). are needed o assess the plausibility of the selected p values.," Comparisons to detailed radiative transfer calculations, such as those of the Potsdam group or CMFGEN (Hillier Miller 1999), are needed to assess the plausibility of the selected $p$ values."
+ Although motivated by the ionization considerations. Q(r) mainly acts as an additional weighting function to shift the predominant region of ine formation in radius.," Although motivated by the ionization considerations, $Q(r)$ mainly acts as an additional weighting function to shift the predominant region of line formation in radius."
+ Ultimately. our profile modeling still falls short of accurately reproducing the observed emission lines.," Ultimately, our profile modeling still falls short of accurately reproducing the observed emission lines."
+ The [Calv] line does become more extreme in its lopsided appearance for some phases., The ] line does become more extreme in its lopsided appearance for some phases.
+ The [Ner] line can have a triple-horned appearance at some phases. which is a good sign: unfortunately. the triple-peaks are never as equal looking in appearance as observed.," The ] line can have a triple-horned appearance at some phases, which is a good sign; unfortunately, the triple-peaks are never as equal looking in appearance as observed."
+ With (2 a constant. the sulfur line continues to display mildly sloped or a mainly topped appearance for restricted phases.," With $Q$ a constant, the sulfur line continues to display mildly sloped or a mainly flat-topped appearance for restricted phases."
+ There are two main problems with the modeling for ~ Vel., There are two main problems with the modeling for $\gamma$ Vel.
+ First. [Nen] always shows a more or less double-horned morphology. never triple-horned as observed.," First, ] always shows a more or less double-horned morphology, never triple-horned as observed."
+ Second. using the ephemeris of North (2007). we have determined the orbital phase of ~ Vel for the line data.," Second, using the ephemeris of North (2007), we have determined the orbital phase of $\gamma$ Vel for the line data."
+ According to Northal... the O star orbits counter-clockwise on the sky. and so the CWIR spirals outward in a clockwise fashion.," According to North, the O star orbits counter-clockwise on the sky, and so the CWIR spirals outward in a clockwise fashion."
+ According to the archive. the data were obtained on 1996 May 17 at 23:36:24. or MID 50220.9836.," According to the archive, the data were obtained on 1996 May 17 at 23:36:24, or MJD 50220.9836."
+ Using the observation of North oon MID 50120.4 and a period of 2?=78.53 days. the O star is of an orbit past periastron in time. which corresponds to an orbital phase of 143°.," Using the observation of North on MJD 50120.4 and a period of $P=78.53$ days, the O star is of an orbit past periastron in time, which corresponds to an orbital phase of $143^\circ$."
+ In our simulations this is close to a phase of 150 that is plotted as a dashed line., In our simulations this is close to a phase of $150^\circ$ that is plotted as a dashed line.
+ Thus. the second problem is that we cannot simulaneously get all four profiles with the overall grossly correct morphologies at quite the same phase.," Thus, the second problem is that we cannot simulaneously get all four profiles with the overall grossly correct morphologies at quite the same phase."
+ Note that at a nearby phase of 120°. the [Neri] line does show a triple peak.," Note that at a nearby phase of $120^\circ$, the ] line does show a triple peak."
+ Our study concerns model calculations for forbidden emission line profile shapes that form in a colliding wind system., Our study concerns model calculations for forbidden emission line profile shapes that form in a colliding wind system.
+ The approach recognizes that spherically symmetric winds produce optically thin and flat-topped emission line shapes produced in the constant expansion flow., The approach recognizes that spherically symmetric winds produce optically thin and flat-topped emission line shapes produced in the constant expansion flow.
+ Deviations from that flat-top morphology represents an opportunity to infer information about the orbital properties of the binary and the colliding wind bow shock., Deviations from that flat-top morphology represents an opportunity to infer information about the orbital properties of the binary and the colliding wind bow shock.
+ Our approach makes a number of signiticant simplitications to make the problem tractable: the OB companion wind makes no contribution to the line emission. flow in the compressed layer is radial and with the same speed as the WR wind. the extent of the shocked laver is only approximate. the geometry of the bow shock head is treated poorly. the bow shock is assumed axisymmetric about the line of centers for the two stars. and more input regarding the ionization balance in the large scale wind is needed to model line shapes properly.," Our approach makes a number of significant simplifications to make the problem tractable: the OB companion wind makes no contribution to the line emission, flow in the compressed layer is radial and with the same speed as the WR wind, the extent of the shocked layer is only approximate, the geometry of the bow shock head is treated poorly, the bow shock is assumed axisymmetric about the line of centers for the two stars, and more input regarding the ionization balance in the large scale wind is needed to model line shapes properly."
+ In light of all of these shortcomings. the models do produce a remarkable degree of diverse profile morphologies.," In light of all of these shortcomings, the models do produce a remarkable degree of diverse profile morphologies."
+ Although even qualitative matches to observed forbidden lines from WR 147 and ~ Vel are far from satisfactory. the attempt does appear to capture a number of trends and shows potential as a tool for deducing or limiting orbital and wind parameters.," Although even qualitative matches to observed forbidden lines from WR 147 and $\gamma$ Vel are far from satisfactory, the attempt does appear to capture a number of trends and shows potential as a tool for deducing or limiting orbital and wind parameters."
+ Although of lower spectral resolution than/SO. we determined that theSpizer IRS could probably detect modulations of IR forbidden line shapes with orbital phase.," Although of lower spectral resolution than, we determined that the IRS could probably detect modulations of IR forbidden line shapes with orbital phase."
+" Unfortunately. the IRS will not be available during ""warm"" cycle."," Unfortunately, the IRS will not be available during “warm” cycle."
+ It is expected that the James Webb Space Telescope will have a infrared with a resolving power similar to/SO., It is expected that the James Webb Space Telescope will have a mid-infrared with a resolving power similar to.
+ In addition. the [Nel] emission line is observable from ground-based observatories. as for example in the study of Smith Houck (2001).," In addition, the ] emission line is observable from ground-based observatories, as for example in the study of Smith Houck (2001)."
+ In particular. their sample included WR 146 which has quite broad lines owing to a wind terminal speed of nearly 3000 km +.," In particular, their sample included WR 146 which has quite broad lines owing to a wind terminal speed of nearly 3000 km $^{-1}$."
+ So there is the possibility that future observations will provide new high quality data of these and other systems for which our diagnostics will be relevant., So there is the possibility that future observations will provide new high quality data of these and other systems for which our diagnostics will be relevant.
+ We wish to thank Ken Gayley for helpful discussions in the early stages of this project., We wish to thank Ken Gayley for helpful discussions in the early stages of this project.
+ We are also grateful a number of useful suggestions made by an anoymous referee., We are also grateful a number of useful suggestions made by an anoymous referee.
+" This project was funded by a partnership between the National Science Foundation (NSF AST-0552798). Research Experiences for Undergraduates (REU). and the Department of Defense (DoD) ASSURE (Awards to Stimulate and Support Undergraduate Research Experiences)programs,"," This project was funded by a partnership between the National Science Foundation (NSF AST-0552798), Research Experiences for Undergraduates (REU), and the Department of Defense (DoD) ASSURE (Awards to Stimulate and Support Undergraduate Research Experiences)programs."
+"Recent hydrodynamic studies of core-collapse supernovae have shown that the neutrino-driven outflows develop to be supersonic, which abruptly decelerate by the reverse shock from the outer layers(e.g.,Janka&Müller1995,1996;Burrowsetal.1995; 2006).","Recent hydrodynamic studies of core-collapse supernovae have shown that the neutrino-driven outflows develop to be supersonic, which abruptly decelerate by the reverse shock from the outer layers\citep[e.g.,][]{Jank1995, Jank1996, Burr1995, Bura2006}."
+". Arconesetal.(2007) have explored the effects of the reverse shock on the properties of neutrino-driven winds by one-dimensional, long-term hydrodynamic simulations of core-collapse supernovae."," \citet{Arco2007} have explored the effects of the reverse shock on the properties of neutrino-driven winds by one-dimensional, long-term hydrodynamic simulations of core-collapse supernovae."
+" Their result shows that, in all of their models (10—25Mo the outflows become supersonic and form the progenitors),termination shock when colliding with the slower preceding supernova ejecta."," Their result shows that, in all of their models $10-25\, M_\odot$ progenitors), the outflows become supersonic and form the termination shock when colliding with the slower preceding supernova ejecta."
+ This condition continues until the end of their computations (10 seconds after core in their all of “standard” models with reasonable bounce)parameter choices., This condition continues until the end of their computations (10 seconds after core bounce) in their all of “standard” models with reasonable parameter choices.
+ A recent self-consistently exploding model of a 9Mo star also shows qualitatively the same result (Hüdepohletal.," A recent self-consistently exploding model of a $9\, M_\odot$ star also shows qualitatively the same result \citep{Hued2010}."
+" In this subsection, we explore the 2010)..effect of the wind-termination on the vp-process."," In this subsection, we explore the effect of the wind-termination on the $\nu$ p-process."
+" The termination point is located at ry.=100, 200, 231, 300 (standard model), 400, 500, and 1000 km on the transonic wind trajectory (black line)shown in Figure 1 (top panel)."," The termination point is located at $r_\mathrm{wt} =
+100$, 200, 231, 300 (standard model), 400, 500, and 1000 km on the transonic wind trajectory (black line)shown in Figure 1 (top panel)."
+" The other parameters L,, Mss, and are kept to be the fiducial values (Table 1; 2nd to 9th Υοοlines)."," The other parameters $L_\nu$, $M_\mathrm{ns}$, and $Y_\mathrm{e,
+9}$ are kept to be the fiducial values (Table 1; 2nd to 9th lines)."
+" In Figure 1 (middle and bottom panels), we find shock-jumps of density and temperature by wind termination only for the Τι=1000 km case, since the termination points are placed below the sonic radius (rs=515 km; Figure 1, top panel) for the other 'The result of nucleosynthesis calculations is shown in Figure 2."," In Figure 1 (middle and bottom panels), we find shock-jumps of density and temperature by wind termination only for the $r_\mathrm{wt} = 1000$ km case, since the termination points are placed below the sonic radius $r_\mathrm{s}= 515$ km; Figure 1, top panel) for the other The result of nucleosynthesis calculations is shown in Figure 2."
+" The top panel shows the mass fractions, X4, of nuclei as a function of atomic mass number, A."," The top panel shows the mass fractions, $X_A$ , of nuclei as a function of atomic mass number, $A$ ."
+ We, We
+slope of 1 to one of |3.,slope of $-1$ to one of $-3$.
+" For a halo of a given mass AMooo within the virial radius σου (which. following common practice. we define as the radius of a sphere within which the mean density is 200p). 9, is related to the ""concentration parameter"" e—reo£r; by Since in à CDM cosmogony halocs of a given mass are preferentially assembled in a fairly narrow range of redshifts (with smaller haloes forming first). there is a correlation between 3, (or c) and. Aogo: less massive haloes are οπου and more centrally concentrated."," For a halo of a given mass $M_{200}$ within the virial radius $r_{200}$ (which, following common practice, we define as the radius of a sphere within which the mean density is $200\rho_c$ ), $\delta_c$ is related to the “concentration parameter” $c \equiv r_{200}/r_s$ by Since in a CDM cosmogony haloes of a given mass are preferentially assembled in a fairly narrow range of redshifts (with smaller haloes forming first), there is a correlation between $\delta_c$ (or $c$ ) and $M_{200}$: less massive haloes are denser and more centrally concentrated."
+" In the case of154... a very. isolated cwarl disc ealaxy whose rotation curve has been measured to a considerable distance from its centre (Carignan&Freeman1988:Loll-manetal. 1993).. the ""universal. profile” of NEW cannot simultaneously account forthe observations at small and at laree racii."," In the case of, a very isolated dwarf disc galaxy whose rotation curve has been measured to a considerable distance from its centre \cite{CF88,H.93}, the “universal profile” of NFW cannot simultaneously account forthe observations at small and at large radii."
+ I£ the parameters (A. rs) are chosen to fit the observations at small racii (after allowance has been made for the formation of a uniform-censity core by mass outllow). the NEW profile predicts that the rotation curve should continue to rise well ονομα the racius at which it is observed. to peak.," If the parameters $\delta_c$, $r_s$ ) are chosen to fit the observations at small radii (after allowance has been made for the formation of a uniform-density core by mass outflow), the NFW profile predicts that the rotation curve should continue to rise well beyond the radius at which it is observed to peak."
+ If. on the other hand the parameters are adjusted to fit the declining part of the rotation curve. 16 NEW profile predicts an unreasonably large mass excess in 1e central regions.," If on the other hand the parameters are adjusted to fit the declining part of the rotation curve, the NFW profile predicts an unreasonably large mass excess in the central regions."
+ Burkert Silk (1997.hereafterBS) rave invoked this difficulty as an argument for the presence ofan additional dark component formed through cdissipative orocesses. analogous to the putative NLACTIIO component of our own Galaxy.," Burkert Silk \shortcite[hereafter BS]{BS97} have invoked this difficulty as an argument for the presence of an additional dark component formed through dissipative processes, analogous to the putative MACHO component of our own Galaxy."
+ In their view. this new component should »e spheroidal rather than clise-like since an overly. massive disc would be unstable to fragmentation and. expected. to orm large quantities of stars. contrary to what is observed in his galaxy.," In their view, this new component should be spheroidal rather than disc-like since an overly massive disc would be unstable to fragmentation and expected to form large quantities of stars, contrary to what is observed in this galaxy."
+ Γον still assume that before this dissipational collapse process the halo was well described. by the NEW »ofile., They still assume that before this dissipational collapse process the halo was well described by the NFW profile.
+ is not the only clwarl galaxy. lor which the observational data are at. variance with the cuspy profiles oedieted bv most W-bocly simulations., is not the only dwarf galaxy for which the observational data are at variance with the cuspy profiles predicted by most $N$ -body simulations.
+ Burkert (1905). aux Ixravtsov et al (1998). have argued that the observationa data can be deseribed by a different one-parameter Family of profiles with at most a very weak cusp., Burkert \shortcite{B95} and Kravtsov et al \shortcite{K.98} have argued that the observational data can be described by a different one-parameter family of profiles with at most a very weak cusp.
+ Of the galaxies examined in these two works. however. only arc have the very extended rotation Curves require for a reallv stringent test of secular evolution moclels for the halo density profile.," Of the galaxies examined in these two works, however, only and have the very extended rotation curves required for a really stringent test of secular evolution models for the halo density profile."
+ (Aleurerctal.1996). is in many wavs a peculiar object. making it much more cillicul to model.," \cite{M.96} is in many ways a peculiar object, making it much more difficult to model."
+ We therefore follow the example of Burkert Silk (1907). and concentrate on alone.," We therefore follow the example of Burkert Silk \shortcite{BS97}
+ and concentrate on alone."
+ In this paper we explore the possibility of achieving a reasonable fit to the observed rotation curve of without invoking such a clissipative spheroidal component., In this paper we explore the possibility of achieving a reasonable fit to the observed rotation curve of without invoking such a dissipative spheroidal component.
+ Πω we experiment with cdillerent mass ejection fractions and disc density distributions.," Instead, we experiment with different mass ejection fractions and disc density distributions."
+ We improve on the work of NEF by using halo distribution functions drawn fron ab initio cosmological simulations: the Llernquist (1990) distributions unrealistically fast clrop-olf at. large: raclii müeht vield a spuriously better fit to the falling part. of the rotation curve., We improve on the work of NEF by using halo distribution functions drawn from ab initio cosmological simulations: the Hernquist \shortcite{H90} distribution's unrealistically fast drop-off at large radii might yield a spuriously better fit to the falling part of the rotation curve.
+" Also. NISEs preferred fits for involve characteristic scale lengths (e, in their notation. corresponding roughly to or, in this paper) an order of magnitude larger than expected. for the halo of if one assumes a total mass. of order LOMAL. or less. commensurate with the observed radial extent and rotation velocity of this galaxy. ("," Also, NEF's preferred fits for involve characteristic scale lengths $a_h$ in their notation, corresponding roughly to $r_s$ in this paper) an order of magnitude larger than expected for the halo of if one assumes a total mass of order $10^{10}\Msun$ or less, commensurate with the observed radial extent and rotation velocity of this galaxy. ("
+A larger total mass would. Likely only exacerbate the difficulty in reproducing the observed ecline in the rotation curve.),A larger total mass would likely only exacerbate the difficulty in reproducing the observed decline in the rotation curve.)
+ Alternatively. NELs cise scale —engths are unreasonably small 454) by comparison to 1e halo scale length.," Alternatively, NEF's disc scale lengths are unreasonably small ) by comparison to the halo scale length."
+ In the present work. by contrast. the haracteristic halo racius à is determined by our ab initio cosmological simulations. and is in good agreement with the results of NEW: our disc scale lengths are also based on observations.," In the present work, by contrast, the characteristic halo radius $r_s$ is determined by our ab initio cosmological simulations, and is in good agreement with the results of NFW; our disc scale lengths are also based on observations."
+ We study both models in which the present-cay disc mass is as implied by optical and observations of (Carignan&Freeman1988) and models with additional barvonic mass.," We study both models in which the present-day disc mass is as implied by optical and observations of \cite{CF88}
+ and models with additional baryonic mass."
+ In the latter case. the additional mass is assumed to be hidden in some hitherto undetected form and distributed. like the observedHEEL. (," In the latter case, the additional mass is assumed to be hidden in some hitherto undetected form and distributed like the observed. ("
+Lhe mean metallicity of154.. as measured by the OIL abundance ratio. is only O.05Z. (vanZee.Havnes&Salzer1997).. and at such a low metallicity there are both theoretical (Maloney&Black1988) and observational (Verter&Lodge1995) indications that the CO-to-Lls conversion ratio should be very large. implying that significant amounts of ll» could be present without conspicuous associated CO emission.),"The mean metallicity of, as measured by the O/H abundance ratio, is only $0.05 \Zsun$ \cite{vZ.97}, and at such a low metallicity there are both theoretical \cite{MB88} and observational \cite{VH95} indications that the $_2$ conversion ratio should be very large, implying that significant amounts of $_2$ could be present without conspicuous associated CO emission.)"
+ We examine the argument that a more massive disc would be Toomre-unstable and. formi stars at a higher rate than observed. and find that we can accommodate a isc up to three times more massive than reported. from observations without having to assume an implausibly strong stabilizing pressure within the disc.," We examine the argument that a more massive disc would be Toomre-unstable and form stars at a higher rate than observed, and find that we can accommodate a disc up to three times more massive than reported from observations without having to assume an implausibly strong stabilizing pressure within the disc."
+ Section 2. describes in detail the parameters of our A’- body simulations., Section \ref{s:simulations} describes in detail the parameters of our $N$ -body simulations.
+ Results are presented in section 3.. and conclusions drawn in section 4..," Results are presented in section \ref{s:results}, and conclusions drawn in section \ref{s:conclusions}."
+ We selected: initial conditions out of a realization of a standard. CDM. power spectrum (ο= 1A=0. h=Hy/A00kms+Alpe1j2 0.5) normalized to ex=0.6.," We selected initial conditions out of a realization of a standard CDM power spectrum $\Omega=1$, $\Lambda=0$, $h\equiv H_0/(100\,\rm{km}\,\rm{s}^{-1}\,\rm{Mpc}^{-1})=0.5$ ) normalized to $\sigma_8=0.6$."
+ This realization was computed on a 128% grid in a periodic box of only 4h.+ Mpe on a side: this would be too small for stucies of clustering. or even of larger single galaxies. but turned out to be adequate for our purpose of obtaining haloes matching the profile of equation 1..," This realization was computed on a $128^3$ grid in a periodic box of only $4 h^{-1}$ Mpc on a side; this would be too small for studies of clustering, or even of larger single galaxies, but turned out to be adequate for our purpose of obtaining haloes matching the profile of equation \ref{q:nfw1}. ."
+ The box was evolved to the present. epoch with a standard. particle-mesh code. and haloes within the desired range of virial masses were identified in the final state.," The box was evolved to the present epoch with a standard particle-mesh code, and haloes within the desired range of virial masses were identified in the final state."
+ ‘The corresponding regions within the initial conditions were, The corresponding regions within the initial conditions were
+We also extracted spectra [rom the observation of 2000 October 17.,We also extracted spectra from the observation of 2000 October 17.
+ The observation lasted for 12116 sec but high backgrouud significantly coustrained the extracted spectra to shorter times., The observation lasted for 42416 sec but high background significantly constrained the extracted spectra to shorter times.
+ The EPIC pu spectrum totaled 15910 s with an average rate of 0.263 counts s.I: the MOS-1 aud NIOS-2 spectra totalecl 23838 and 21121 sec. respectively. at an average count rate of 0.085 and 0.080 counts +.," The EPIC pn spectrum totaled 15940 s with an average rate of 0.263 counts $^{-1}$; the MOS-1 and MOS-2 spectra totaled 23838 and 24124 sec, respectively, at an average count rate of 0.088 and 0.089 counts $^{-1}$."
+ Just under of the total extracted counts formed the pu spectrum., Just under of the total extracted counts formed the pn spectrum.
+ We fit the pu aud MOS spectra using the best-lit model. namely the dual variable Mekal model.," We fit the pn and MOS spectra using the best-fit model, namely the dual variable Mekal model."
+ Table 3. lists the resulting model parameters., Table \ref{ascafit} lists the resulting model parameters.
+ The Πίος spectrum is shown in Figure 5.., The fitted spectrum is shown in Figure \ref{fit_xmm}.
+ The best-fit temperatures are 0.67 aud 2.87 keV. Rather than plot the parameter coutours separately. we included them ou the ACIS plots of Figure 3:) the contours are thelower set of contours in each figure.," The best-fit temperatures are 0.67 and 2.87 keV. Rather than plot the parameter contours separately, we included them on the ACIS plots of Figure \ref{cont_dvmek}; the contours are the set of contours in each figure."
+ Note that the temperatures measured from the two instruments are essentially identical within the errors: the coutours differ only in Nyy by a factor of 1.65., Note that the temperatures measured from the two instruments are essentially identical within the errors; the contours differ only in $_{\rm H}$ by a factor of $\sim$ 1.65.
+ One possible explanation is an increasii[n]ο column density as the material cools aud 'ecombiues., One possible explanation is an increasing column density as the material cools and recombines.
+ The inferred Iumiuosity in the soft band depends. very seusitively on the adopted value for wg because of that sensitivity. we adopt for Nj a value deteriniued from the measured Ejy or theROSAT HRI aud PSPC spectra as well as the spectra.," The inferred luminosity in the soft band depends very sensitively on the adopted value for $_{\rm H}$ ; because of that sensitivity, we adopt for $_{\rm H}$ a value determined from the measured $_{\rm
+B-V}$ for the HRI and PSPC spectra as well as the spectra."
+ This approach assigns a 'easonable value to Ny for the HRI data points for which no spectral information would otherwise je available.," This approach assigns a reasonable value to $_{\rm
+H}$ for the HRI data points for which no spectral information would otherwise be available."
+ In addition. adoption of an Ey yv-determiued value for Nyy removes the artificially ueh value determined [rom fits to the spectra for which the calibration below ~0.7 keV is uicertain (Sambrunaetal.2000).," In addition, adoption of an $_{\rm
+B-V}$ -determined value for $_{\rm H}$ removes the artificially high value determined from fits to the spectra for which the calibration below $\sim$ 0.7 keV is uncertain \citep{SCU00}."
+. We use the Chandra- and NALAL-New/on-determinecd values of 4 to probe the seusitivity of the unabsorbed fluxes to the adopted values.," We use the - and -determined values of $_{\rm
+H}$ to probe the sensitivity of the unabsorbed fluxes to the adopted values."
+ For direct comparison of the and results with the previously-described spectra (Petreetal.199L. $99). we re-fit the spectra using the dual variable Mekal mocel as well as a siugle variable Mekal model.," For direct comparison of the and results with the previously-described spectra \citealt{Petre94}, S99), we re-fit the spectra using the dual variable Mekal model as well as a single variable Mekal model."
+ We included a single variable Mekal model to test whether the dual moclel is statistically required to fit the SIS spectra., We included a single variable Mekal model to test whether the dual model is statistically required to fit the SIS spectra.
+ Necessarily. the lower count rate yields larger error bars for each parameter of the dual Mekal model.," Necessarily, the lower count rate yields larger error bars for each parameter of the dual Mekal model."
+ We restrict our attention solely to the SIS spectra to ensure the best seusitivity to line emission., We restrict our attention solely to the SIS spectra to ensure the best sensitivity to line emission.
+ The fitted parameters values are listed in Table 3. while the fluxes are listed in Table 2.., The fitted parameters values are listed in Table \ref{ascafit} while the fluxes are listed in Table \ref{fluxval}.
+ From the \7/v values. the dual variable model provides a statistically better fit.," From the ${\chi}^2/{\nu}$ values, the dual variable model provides a statistically better fit."
+ For theROSAT PSPC aud HRI fluxes. we re-extracted and re-fit the PSPC spectra using a sinele variable Mekal mocdel.," For the PSPC and HRI fluxes, we re-extracted and re-fit the PSPC spectra using a single variable Mekal model."
+ A dual model is not statistically required to provide a good Lit largely because of the narrow bandpass., A dual model is not statistically required to provide a good fit largely because of the narrow bandpass.
+ All of the abundance values were consistent with solar. within the errors. so we froze all at solar values.," All of the abundance values were consistent with solar, within the errors, so we froze all at solar values."
+ Unabsorbed I[Iuxes were calculated in the 0.5-2 keV band., Unabsorbed fluxes were calculated in the 0.5-2 keV band.
+ As HRI does not. possess spectral resolution. we must adopt a fixed spectral inodel.," As HRI does not possess spectral resolution, we must adopt a fixed spectral model."
+ We used a best lit model of a single variable Mekal spectrum at kT = 0.6 keV and Ny = L8x107! 7., We used a `best fit' model of a single variable Mekal spectrum at kT = 0.6 keV and $_{\rm H}$ = $\times$ $^{21}$ $^{-2}$.
+ This spectrum is essentially au average of the PSPC spectra audadopting the Ep y-cdeterminecd Ny value.," This spectrum is essentially an average of the PSPC spectra andadopting the $_{\rm
+B-V}$ -determined $_{\rm H}$ value."
+ Only. the mocel normalization was adjusted for each HRI point., Only the model normalization was adjusted for each HRI point.
+ Unabsorbed fluxes in the, Unabsorbed fluxes in the
+motion but has another component which allects the line of sight. velocity.,motion but has another component which affects the line of sight velocity.
+ The angular size of the velocity. perturbations seen in the renzograms (as predicted by eqn 14)) is plotted against x in Fig., The angular size of the velocity perturbations seen in the renzograms (as predicted by eqn \ref{eqn:angsize}) ) is plotted against $x$ in Fig.
+ Lifa) for four inclination angles: 0 degrees (solid line). 30 degrees (dashed line). 60 degrees (dot-dashed line) and 90 degrees (dotted line).," \ref{fig:angsize} for four inclination angles; 0 degrees (solid line), 30 degrees (dashed line), 60 degrees (dot-dashed line) and 90 degrees (dotted line)."
+ ���Phis plot assumes the galaxy is at a distance of 785 kpe (i.e the distance to M31) witha shock 7 kpe from the galactic centre (representative of our model galaxy)., This plot assumes the galaxy is at a distance of 785 kpc (i.e the distance to M31) with a shock 7 kpc from the galactic centre (representative of our model galaxy).
+ The shock is assumed to reduce the circular velocity by a factor f=0.9 (i.e. à perturbation of 20 km/s on a 200 km/s circular velocity)., The shock is assumed to reduce the circular velocity by a factor $f=0.9$ (i.e. a perturbation of 20 km/s on a 200 km/s circular velocity).
+ A similar plot showing the size of the velocity features in kpe is shown in Fig. Lib)., A similar plot showing the size of the velocity features in kpc is shown in Fig. \ref{fig:linsize}.
+ The predicted size of the features is in agreement with the eatures seen in the renzograms of the svnthetic observations (i.c. a tvpical size of approximately a kpe)., The predicted size of the features is in agreement with the features seen in the renzograms of the synthetic observations (i.e. a typical size of approximately a kpc).
+ In all cases the size of the velocity perturbation tends owards zero as e tends towards zero. due to the assumption hat the perturbation only allects the circular velocity (as discussed. above).," In all cases the size of the velocity perturbation tends towards zero as $x$ tends towards zero, due to the assumption that the perturbation only affects the circular velocity (as discussed above)."
+ As ur increases towards f the size of he feature increases because the circular velocity has a arger component in the line of sight., As $x$ increases towards $f$ the size of the feature increases because the circular velocity has a larger component in the line of sight.
+ The renzograms for our svnthetic observations also show Larger velocity kinks in channels further [rom the svstemic velocity., The renzograms for our synthetic observations also show larger velocity kinks in channels further from the systemic velocity.
+ There is a change in the sign of Aa when .r becomes negative (eqn is an odd function of μη) due to the assumption that the shock always reduces the circular. velocity., There is a change in the sign of $\Delta \alpha$ when $x$ becomes negative (eqn \ref{eqn:angsize} is an odd function of $x$ ) due to the assumption that the shock always reduces the circular velocity.
+ This results in post-shock gas always moving towards the velocity channel of the svstemic velocity (or zero i£ the svstemic velocity has been subtractecd)., This results in post-shock gas always moving towards the velocity channel of the systemic velocity (or zero if the systemic velocity has been subtracted).
+ The size of the velocity kinks increases as the inclination angle increases. due to projection effects. however in practice it. will be cillicult to see these features for nearly facc-on cases because very high velocity resolution would. be requirecl (for nearly. face-on cases the range of .r corresponds to a small velocity range)," The size of the velocity kinks increases as the inclination angle increases, due to projection effects, however in practice it will be difficult to see these features for nearly face-on cases because very high velocity resolution would be required (for nearly face-on cases the range of $x$ corresponds to a small velocity range)."
+ We have presented. a method for generating. synthetic spectral cubes of the 21 cm hydrogen. line from SPILL simulations of spiral galaxies., We have presented a method for generating synthetic spectral cubes of the 21 cm hydrogen line from SPH simulations of spiral galaxies.
+ Phe methocl successfully maps the density. temperature and velocity from the SPLE particles onto an AMI grid. while preserving important structures (e.g. spiral arms ancl spurs) and accurately representing the total mass.," The method successfully maps the density, temperature and velocity from the SPH particles onto an AMR grid, while preserving important structures (e.g. spiral arms and spurs) and accurately representing the total mass."
+ Synthetic cata cubes are generated: using a ταν tracing method., Synthetic data cubes are generated using a ray tracing method.
+ The svathetic cata show good agreement with observations of M31. anc M33. whereby increasing velocity channels trace out the main cise of the galaxy., The synthetic data show good agreement with observations of M31 and M33 whereby increasing velocity channels trace out the main disc of the galaxy.
+ Velocity contours of the synthetic data show perturbations due to non-circular motions. similar to those already: observed in M81 (222).. whieh are not seen in the observations of M31 and M33.," Velocity contours of the synthetic data show perturbations due to non-circular motions, similar to those already observed in M81 \citep{Adler96,Visser1980,rots_1975}, which are not seen in the observations of M31 and M33."
+ Our model galaxy is a grand. design spiral galaxy ancl shows velocity structure associated with the spiral perturbation., Our model galaxy is a grand design spiral galaxy and shows velocity structure associated with the spiral perturbation.
+ “Phe method of generating synthetic observations can also be applied to simulations in. which velocity. structure is. generated. by other mechanisms., The method of generating synthetic observations can also be applied to simulations in which velocity structure is generated by other mechanisms.
+ Internal mechanisms. such as self-eravity. ancl stellar fecdback. can generate velocity structure. ancl indeed: are dominant in Hocculent spiral galaxies.," Internal mechanisms, such as self-gravity and stellar feedback, can generate velocity structure, and indeed are dominant in flocculent spiral galaxies."
+ External influences can also alleet the morphology of a spiral galaxy e.g. interactions with a companion galaxy. high velocity clouds or the intracluster medium.," External influences can also affect the morphology of a spiral galaxy e.g. interactions with a companion galaxy, high velocity clouds or the intracluster medium."
+ As the majority. of. galaxies reside in groups or clusters. it ds likely that external environmental effects will influence the structure of most real galaxies: our model galaxy is perfectly isolated and our results show idealised behaviour in the absence of external inlluences.," As the majority of galaxies reside in groups or clusters, it is likely that external environmental effects will influence the structure of most real galaxies; our model galaxy is perfectly isolated and our results show idealised behaviour in the absence of external influences."
+ A [auge parameter space could. potentially be studied. involving galaxies with dillerent. internally ancl externally eencrated velocity. structure.," A large parameter space could potentially be studied, involving galaxies with different internally and externally generated velocity structure."
+ Some of. these aspects can already be modelled (c.g. a galaxy undergoing an interaction. ?)) and some require further model development (e.g inclusion of stellar feedback. which is ongoing).," Some of these aspects can already be modelled (e.g. a galaxy undergoing an interaction, \cite{Dobbs09}) ) and some require further model development (e.g inclusion of stellar feedback, which is ongoing)."
+ A large number of models must be be run in order to generate the SPL input for the radiative transfer caleulation., A large number of models must be be run in order to generate the SPH input for the radiative transfer calculation.
+ Once such a library of simulations was available. one could compare these with real observations in order to understand the velocity structure of spiral galaxies.," Once such a library of simulations was available, one could compare these with real observations in order to understand the velocity structure of spiral galaxies."
+ However given the potentially complex nature of the velocity structures generated there would. need to be a reliable way of matching an observed galaxy with its counterpart in the synthetic observation library., However given the potentially complex nature of the velocity structures generated there would need to be a reliable way of matching an observed galaxy with its counterpart in the synthetic observation library.
+ In our present. models. we can relate the strength of the spiral shock. ancl the inclination of the galaxy. to the size of the perturbations in the renzograms.," In our present models, we can relate the strength of the spiral shock, and the inclination of the galaxy, to the size of the perturbations in the renzograms."
+ When we inclucle additional physics. c.g. self gravity and. feedback. we can investigate whether these perturbations are still observable (for a given shock strength) or whether they are overwhelmed by other motions in the σας.," When we include additional physics, e.g. self gravity and feedback, we can investigate whether these perturbations are still observable (for a given shock strength) or whether they are overwhelmed by other motions in the gas."
+ Our method has also recently been applied to. an observer placed inside the galaxy (2) to generate a synthetic galactic plane survey., Our method has also recently been applied to an observer placed inside the galaxy \citep{douglas_2010} to generate a synthetic galactic plane survey.
+ sell absorption features ancl the conversion of cold to molecular clouds. as seen in a real ealactic plane survey (e.g. 27)). are seen in the synthetic data.," self absorption features and the conversion of cold to molecular clouds, as seen in a real galactic plane survey (e.g. \cite{Taylor03,Stil06}) ), are seen in the synthetic data."
+ A powerful future application of using a simulation for eenerating a survey is that given a series of time frames we will be able to trace the evolution of molecular clouds. ancl related observed features. to the physical properties of the material in the cloud.," A powerful future application of using a simulation for generating a survey is that given a series of time frames we will be able to trace the evolution of molecular clouds, and related observed features, to the physical properties of the material in the cloud."
+ The method can also be extended to eenerate svnthetie observations of other tracer species. such as CO. or dust. for simulated surveys or simulated external ealaxies.," The method can also be extended to generate synthetic observations of other tracer species, such as CO, or dust, for simulated surveys or simulated external galaxies."
+" We would like to thank ""Tyler. Foster for. supplying the ADSL data used. in this paper.", We would like to thank Tyler Foster for supplying the M31 data used in this paper.
+ We would. like to thank Laurent Chemin for providing the M31 observed line profile data., We would like to thank Laurent Chemin for providing the M31 observed line profile data.
+ Caleulations presented here were performed using the University of Exeter Supercomputer., Calculations presented here were performed using the University of Exeter Supercomputer.
+" The research. leading to these results has received. funding from the European Communitys Seventh Framework Programme under grant agreement n"" PLHE-CA-2008-221289.", The research leading to these results has received funding from the European Community's Seventh Framework Programme under grant agreement $^o$ PIIF-GA-2008-221289.
+ €. L. Do's research ab Exeter was conducted as part of the award. 77Ehe formation of stars ancl planets: Iacdiation hvedrodyvnamical and magnetohvdrodynamical simulations.” mace under the European Heads of Research Councils and European Science Foundation EURY (european Young Investigator) Awarels scheme ancl supported. by funds from the Participating Organisations of EUIRYI and the EC Sixth Framework Progranine.," C. L. D.'s research at Exeter was conducted as part of the award “'The formation of stars and planets: Radiation hydrodynamical and magnetohydrodynamical simulations,” made under the European Heads of Research Councils and European Science Foundation EURYI (European Young Investigator) Awards scheme and supported by funds from the Participating Organisations of EURYI and the EC Sixth Framework Programme."
+No frequency dependence was detected in the 2002 September 30 bursts. ancl none is detected across the 60-channel. 15-MIIz banclpass for the 2003 September 28 burst.,"No frequency dependence was detected in the 2002 September 30 bursts, and none is detected across the 60-channel, 15-MHz bandpass for the 2003 September 28 burst."
+ A fit across the 15-MlIz bandpass of the 2003 September 28 observation vields a wavelength dependence of SxA=.," A power-law fit across the 15-MHz bandpass of the 2003 September 28 observation yields a wavelength dependence of $S \propto \lambda^{4 \pm
+5}$."
+ No circular polarization is detected in the bursts with an upper limit of obtained for both epochs., No circular polarization is detected in the bursts with an upper limit of obtained for both epochs.
+ Linear polarization measurements are not available for either the discovery or recovery observations., Linear polarization measurements are not available for either the discovery or recovery observations.
+ No emission is detected from wwhen imaging the 2003 September 28 observation al times when the burst is not occurring., No emission is detected from when imaging the 2003 September 28 observation at times when the burst is not occurring.
+ We are able to improve the (50) upper limit for 330 MlIz interburst emission trom 75 mJ. for the discovery observations. to 25 mJy. for the recovery observation.," We are able to improve the $5\sigma$ ) upper limit for 330 MHz interburst emission from 75 mJy, for the discovery observations, to 25 mJy, for the recovery observation."
+ The upper limit on «quiescent emission during periods of no burst activity is 15 mJv at 330 MIIz 2005)., The upper limit on quiescent emission during periods of no burst activity is 15 mJy at 330 MHz \citep{hlkrmy-z05}.
+". Nondetections on 2005 March 25 at both 330 and 1400 MIIZ also vield an upper limit of 15 mJy al 330 MlIIz. but a significantly reduced upper limit of 0.4 mJy at 1400 MIIZ. as compared to a 35 mJy upper limit obtained from a 2003 January observation αἱ that frequency,"," Nondetections on 2005 March 25 at both 330 and 1400 MHz also yield an upper limit of 15 mJy at 330 MHz, but a significantly reduced upper limit of 0.4 mJy at 1400 MHz, as compared to a 35 mJy upper limit obtained from a 2003 January observation at that frequency."
+ We have also learned that hhas been observed in early and mid-2005 with the Westerbork Synthesis Radio Telescope (WSRT) at both 330 and 1400 MIIz., We have also learned that has been observed in early and mid-2005 with the Westerbork Synthesis Radio Telescope (WSRT) at both 330 and 1400 MHz.
+ Upper limits on any emission are approximately a lew milliJanskvs (R. Brawn 2005. private communication).," Upper limits on any emission are approximately a few milliJanskys (R. Braun 2005, private communication)."
+ As Ivanοἱal.(2005). reported. during the discovery epoch (2002 September October 1) the bursts from the source had an approximate 77 pperiodicitv.," As \cite{hlkrmy-z05} reported, during the discovery epoch (2002 September 30--October 1) the bursts from the source had an approximate 77 periodicity."
+ As noted above. the recovery on 2003 September 28 occurs in the midst of a set of sscans. spread over several hours.," As noted above, the recovery on 2003 September 28 occurs in the midst of a set of scans, spread over several hours."
+ Assuming that the source was emitting periodic bursts at this epoch. each of 10 dduration. we have determined (he periods al which bursts could occur while being consistent with the gaps and non-detections during the 2003 September 28 epoch.," Assuming that the source was emitting periodic bursts at this epoch, each of 10 duration, we have determined the periods at which bursts could occur while being consistent with the gaps and non-detections during the 2003 September 28 epoch."
+ There are 212 bbetween the detected burst and the end of the observation., There are 212 between the detected burst and the end of the observation.
+ Thus. we can place no constraints on periods longer than 212 min.," Thus, we can place no constraints on periods longer than 212 min."
+ For shorter periods. only the following ranges of periods are allowed: 3637 min.," For shorter periods, only the following ranges of periods are allowed: 36–37 min.,"
+ 7178 min.," 71–78 min.,"
+ LOT112 min..," 107–112 min.,"
+ 131.160 min..," 131–160 min.,"
+ and 179195 min., and 179–195 min.
+ We estimate (hat (he uncertainty in making these determinations is perhaps l aand results from slightly varying noise levels within the scans and (the assumption that the duration of the bursts remains fixed at 10 min., We estimate that the uncertainty in making these determinations is perhaps 1 and results from slightly varying noise levels within the scans and the assumption that the duration of the bursts remains fixed at 10 min.
+ Thus. while consistent with the gaps and non-detections. the possible 3637 pperioclicity is perhaps only mareinally so.," Thus, while consistent with the gaps and non-detections, the possible 36–37 periodicity is perhaps only marginally so."
+ A T1 pperiodidtv remains consistent with the 2003 September 28 observations., A 77 periodicity remains consistent with the 2003 September 28 observations.
+ We cannot use the interval between. 2002 September 30 and 2003 September 28 {ο constrain the burst activity because (he uncertainty on (he 77 pperiod determined from the 2002 September 30 observations is sulliciently large (~15 s) that we cannol connect the phase between the (wo observations., We cannot use the interval between 2002 September 30 and 2003 September 28 to constrain the burst activity because the uncertainty on the 77 period determined from the 2002 September 30 observations is sufficiently large $\sim 15$ s) that we cannot connect the phase between the two observations.
+ Indeed. the nearest observation to the 2002," Indeed, the nearest observation to the 2002"
+We have found the electron fraction with thermal ionisation and cosmic rav ionisation with recombination.,We have found the electron fraction with thermal ionisation and cosmic ray ionisation with recombination.
+ La order to determine the dead zone height we take the electron fraction to be the maximum. of equation(8)) and. the solution to equation (17)., In order to determine the dead zone height we take the electron fraction to be the maximum of equation\ref{xet}) ) and the solution to equation \ref{cubic}) ).
+" Now we can find the magnetic ltevnolds number in equation. (3)) and for each radius. H. tine the eritieal height. in. the disc. co, by solving Rey,=Regpan."," Now we can find the magnetic Reynolds number in equation \ref{rey}) ) and for each radius, $R$, find the critical height in the disc, $z_{\rm
+ crit}$ by solving $Re_{\rm M}=Re_{\rm M,crit}$."
+ Where thermal ionisation is dominant. this is straightforward.," Where thermal ionisation is dominant, this is straightforward."
+ However. it is slightly more complicated where cosmic rav ionisation dominates.," However, it is slightly more complicated where cosmic ray ionisation dominates."
+ For cxample. with the high metallicity limit for the electron. fraction (equation 19)) we solve In the limit of low metallicity we replace 3) with ή.," For example, with the high metallicity limit for the electron fraction (equation \ref{high}) ) we solve In the limit of low metallicity we replace $\beta_{\rm r}$ with $\beta_{\rm
+ d}$."
+" Given a total surface density and temperature. this equation can be solved numerically to find 23, and hence the active laver surface density. given in equation (21)) and the dead zone surface density NM=XM."," Given a total surface density and temperature, this equation can be solved numerically to find $z_{\rm crit}$ and hence the active layer surface density given in equation \ref{sa}) ) and the dead zone surface density $\Sigma_{\rm d}=\Sigma-\Sigma_{\rm m}$."
+ We illustrate this with an example in the next section., We illustrate this with an example in the next section.
+ We choose a fidiucial mocel for the disc structure around and star of mass Al=LAL. and consider properties of the dead zone.," We choose a fiducial model for the disc structure around and star of mass $M=1\,\rm M_\odot$ and consider properties of the dead zone."
+ We take the total surface density to be where No=I0*gem.> and the midplane temperature is where ο=100Ix.," We take the total surface density to be where $\Sigma_0=10^3\,\rm g\,cm^{-2}$ and the midplane temperature is where $T_0=100\,\rm K$."
+ We choose the cise such that without the dead zone. it has à constant acerction rate at all ασ.," We choose the disc such that without the dead zone, it has a constant accretion rate at all radii."
+ This requires MxeS=const and so we take s|15.," This requires $\dot M \propto \nu \Sigma =\,\rm const$ and so we take $s+t=1.5$."
+ Unless otherwise stated. we choose s=/0.75.," Unless otherwise stated, we choose $s=t=0.75$."
+ This disc is in a steady state if it is supplied with a constant source of mass in the outer parts and the disc is fully NIU active., This disc is in a steady state if it is supplied with a constant source of mass in the outer parts and the disc is fully MRI active.
+ Llowever. we now show that the disc may have a dead zone.," However, we now show that the disc may have a dead zone."
+ μα Fig., In Fig.
+ 1 we show. at a fixed radius in the disc. 7?=LAU. the surface density in the active laver as a function of the total surface density.," \ref{thick} we show, at a fixed radius in the disc, $R=1\,\rm
+AU$, the surface density in the active layer as a function of the total surface density."
+ We vary the metallicity ancl critical magnetic Itevnolds number in the disc., We vary the metallicity and critical magnetic Reynolds number in the disc.
+ The dotted. lines show the hieh metallicity limit. for the electron. [raction given in equation (19))., The dotted lines show the high metallicity limit for the electron fraction given in equation \ref{high}) ).
+ “Phe solid lines show intermediate metallicities found. by solving the cubic equation (17)) for the electron. fraction., The solid lines show intermediate metallicities found by solving the cubic equation \ref{cubic}) ) for the electron fraction.
+ For fixed. radius. the active laver reaches a constant surface cdensitv as the total surface density increases.," For fixed radius, the active layer reaches a constant surface density as the total surface density increases."
+ The active laver is smallest with a low metallicity and a high critical magnetic Ievnolds number., The active layer is smallest with a low metallicity and a high critical magnetic Reynolds number.
+ Llowever. the finite metallicities considered in Fig.," However, the finite metallicities considered in Fig."
+ 1 are tiny., \ref{thick} are tiny.
+" Even the highest finite metal fraction we consider here. ry,=107. corresponds to à very small metallicity of about Z=3O° (see equation 10)) that is several orders of magnitude smaller than that in the SAIC."," Even the highest finite metal fraction we consider here, $x_{\rm m}=10^{-7}$, corresponds to a very small metallicity of about $Z=3\times 10^{-6}$ (see equation \ref{mf}) ) that is several orders of magnitude smaller than that in the SMC."
+ The metallicity required for the dead zone to be significantly different from the high metallicity limit is very small. unless the critical magnetic ltevnolds number is found to be significantlv larger than 107.," The metallicity required for the dead zone to be significantly different from the high metallicity limit is very small, unless the critical magnetic Reynolds number is found to be significantly larger than $10^4$."
+ Hence. for reasonable metallicities. those in the observable universe. the dead zone structure cloes not depend on metallicity and the high metallitiev limit is appropriate.," Hence, for reasonable metallicities, those in the observable universe, the dead zone structure does not depend on metallicity and the high metalliticy limit is appropriate."
+ In the rest of this work we consider only the hieh metallicity limit., In the rest of this work we consider only the high metallicity limit.
+ 1n Fig., In Fig.
+ 2. we show the height of the dead. zone as a function of radius in the disc for varving critical magnetic ]tevnolds number in the high metallicity limit., \ref{zones} we show the height of the dead zone as a function of radius in the disc for varying critical magnetic Reynolds number in the high metallicity limit.
+ As the critical magnetic Itevnolds number increases. the size of the dead zone increases.," As the critical magnetic Reynolds number increases, the size of the dead zone increases."
+ The results are similar to those of Matsumura&Pudritz(2003) who find the dead zone in a disc with a different surface density and temperature distribution., The results are similar to those of \cite{matsumura03} who find the dead zone in a disc with a different surface density and temperature distribution.
+ In lig., In Fig.
+ 3. we also plot the inner most parts of the dead. zone and see that the elfect of varving the constant A7] (the Potassium abundance relative to hydrogen. see Section 2.1)) is not very significant.," \ref{zones2} we also plot the inner most parts of the dead zone and see that the effect of varying the constant $[K/H]$ (the Potassium abundance relative to hydrogen, see Section \ref{thermal}) ) is not very significant."
+ Lt only alfects the innermost parts of the disc where thermal ionisation is the dominant. source., It only affects the innermost parts of the disc where thermal ionisation is the dominant source.
+ Further out where cosmic rav ionisation becomes dominant the dead zone does not depend on A]., Further out where cosmic ray ionisation becomes dominant the dead zone does not depend on $[K/H]$ .
+ In the rest of the work we take the solar abundance value. A/1/]= 0.," In the rest of the work we take the solar abundance value, $[K/H]=0$ ."
+forcing by Neptune causes p and q to evolve as (to first order in i and iy): where The constants igee and y are determined from the initial conditions.,forcing by Neptune causes $p$ and $q$ to evolve as (to first order in $i$ and $i_N$ ): where The constants $\ifree$ and $\gamma$ are determined from the initial conditions.
+" Here, Qy is the longitude of ascending node of Neptune and iy is the inclination of Neptune."," Here, $\Omega_N$ is the longitude of ascending node of Neptune and $i_N$ is the inclination of Neptune."
+" While eforcea depends on a (Eqn. 3)),"," While $\eforced$ depends on $\alpha$ (Eqn. \ref{eqn:simples}) ),"
+" to first order éforcea is equal to in, regardless of the particle’s semi-major axis.", to first order $\iforced$ is equal to $i_N$ regardless of the particle's semi-major axis.
+ Three quantities are plotted in Fig., Three quantities are plotted in Fig.
+" 4 as a function of a (for two different αν): 1) eforcea/@n, 2) tforcea/tn, and the secular period ma."," \ref{fig:freqs} as a function of $a$ (for two different $a_N$ ): 1) $\eforced/e_N$, 2) $\iforced/i_N$, and the secular period $\frac{2\pi}{\gkbo}$."
+ These values describe the amplitude and timescale of the secular excitation of the planetesimal population., These values describe the amplitude and timescale of the secular excitation of the planetesimal population.
+" For example, when Neptune is at 20 AU, a planetesimal at 45 AU has a forced eccentricity of 0.55eyw and forced inclination of in."," For example, when Neptune is at 20 AU, a planetesimal at 45 AU has a forced eccentricity of $0.55 e_N$ and a forced inclination of $i_N$."
+" If the planetesimal begins with ae=i0, it will reach its maximum eccentricity e=2x0.55ey1.1εΝ and maximum inclination i=2iy on a timescale of $5 = 13 Myr (i.e. half a secular oscillation period)."," If the planetesimal begins with $e = i = 0$, it will reach its maximum eccentricity $e = 2\times0.55 e_N = 1.1 e_N$ and maximum inclination $i = 2 i_N$ on a timescale of $\frac{1}{2} \frac{2\pi}{\gkbo}$ = 13 Myr (i.e. half a secular oscillation period)."
+ The secular excitation of a planetesimal disk can be modeled using the expressions above., The secular excitation of a planetesimal disk can be modeled using the expressions above.
+ Here we consider that Neptune is temporarily on an inclined and/or eccentric orbit after undergoing scattering on an effectively instantaneous timescale (i.e. much less than the secular timescale)., Here we consider that Neptune is temporarily on an inclined and/or eccentric orbit after undergoing scattering on an effectively instantaneous timescale (i.e. much less than the secular timescale).
+ Then Neptune imparts a forced eccentricity €forceq and inclination igorceq on the initially cold planetesimal disk., Then Neptune imparts a forced eccentricity $\eforced$ and inclination $ \iforced$ on the initially cold planetesimal disk.
+ Thus each planetesimal has €freev€forced., Thus each planetesimal has $\efree \sim \eforced$.
+ AN 2&eevtforced-, and $\ifree \sim \iforced$ .
+" The planetesimal's total eccentricity and inclination, each a vector sum of the free and forced, now oscillates from the initial values of e(0)~0 and {(0)~0, reaching maximum values of e=€free+Cforcedv2€forced ANd =treetiforcedv2torced; on a timescale set by the secular{ evolution rate gxgo."," The planetesimal's total eccentricity and inclination, each a vector sum of the free and forced, now oscillates from the initial values of $e(0) \sim 0$ and $i(0) \sim 0$, reaching maximum values of $e = \efree + \eforced \sim 2 \eforced$ and $i = \ifree + \iforced \sim 2 \iforced$, on a timescale set by the secular evolution rate $\gkbo$."
+ Thus an initially cold planetesimal disk will become excited., Thus an initially cold planetesimal disk will become excited.
+ We limit our parameter study to include only the planet Neptune., We limit our parameter study to include only the planet Neptune.
+ We have several reasons for choosing this approach., We have several reasons for choosing this approach.
+" First, unlike previous approaches, we are not attempting to create a single model that reproduces the entire Kuiper belt in detail, but to constrain which histories of Neptune are consistent with a major qualitative feature: the unexcited orbits of the cold classicals."," First, unlike previous approaches, we are not attempting to create a single model that reproduces the entire Kuiper belt in detail, but to constrain which histories of Neptune are consistent with a major qualitative feature: the unexcited orbits of the cold classicals."
+ Our constraints will feed into more detailed models., Our constraints will feed into more detailed models.
+" Second, restricting the parameter study to just Neptune drastically reduces the number of parameters, allowing us to thoroughly explore the remaining parameter space."," Second, restricting the parameter study to just Neptune drastically reduces the number of parameters, allowing us to thoroughly explore the remaining parameter space."
+" Third, we find (?) that the primary effects of the other planets on the Kuiper belt are indirect: they alter the orbit of Neptune, which in turn affects the Kuiper belt."," Third, we find \citep{2012D} that the primary effects of the other planets on the Kuiper belt are indirect: they alter the orbit of Neptune, which in turn affects the Kuiper belt."
+ Our modifications to the integrator (Appendix A) allow us to model any orbital evolution of Neptune without needing to include the other planets., Our modifications to the integrator (Appendix A) allow us to model any orbital evolution of Neptune without needing to include the other planets.
+ Fig., Fig.
+" 5 provides one example pair of integrations, demonstrating that cold planetesimals evolve similarly in the presence of only Neptune and in the presence of all four giant planets."," \ref{fig:otherplanets} provides one example pair of integrations, demonstrating that cold planetesimals evolve similarly in the presence of only Neptune and in the presence of all four giant planets."
+ One main effect of other planets is to cause Neptune's longitude of periapse to precess., One main effect of other planets is to cause Neptune's longitude of periapse to precess.
+" However, the disruption of the cold classicals is not significantly affected if Neptune's precession period is comparable to or longer than the secular excitation time (Fig. 5))."," However, the disruption of the cold classicals is not significantly affected if Neptune's precession period is comparable to or longer than the secular excitation time (Fig. \ref{fig:otherplanets}) )."
+" If Neptune precesses very quickly,the cold classicals couldbe"," If Neptune precesses very quickly,the cold classicals couldbe"
+contribute a fairly lavee Traction of the EGB (Thompsonetal.2007:Bhattacharvamar2009:Dhattacharvaetal. 2009).,"contribute a fairly large fraction of the EGB \citep{t07,bs09,b09}. ."
+. Recently. Abdoetal.(2010b) built a source count distribution al GeV οποιον and vielded that point sources which had not been detected by the LAT can contribute EGBL.," Recently, \citet{abdo10b} built a source count distribution at GeV energy and yielded that point sources which had not been detected by the LAT can contribute EGB."
+ At the fluxes currently reached by the LAT. they ruled out the hypothesis that sources bblazars) produce a laree Traction of the EGD.," At the fluxes currently reached by the LAT, they ruled out the hypothesis that point-like sources blazars) produce a large fraction of the EGB."
+ llowever. if the property of undetected sources is not similar to the detected sources. these conclusions maybe not correct.," However, if the property of undetected sources is not similar to the detected sources, these conclusions maybe not correct."
+ Therefore. we apply an image stacking method to directly. study. the undetected point sources.," Therefore, we apply an image stacking method to directly study the undetected point sources."
+ For a sample of possible y-ray point sources which have not been detected by theFermi due to their faint [Inxes or soft spectra 20106)... we can stack a large number of them to improve the statistics 2010).," For a sample of possible $\gamma$ -ray point sources which have not been detected by the due to their faint fluxes or soft spectra \citep{abdo10c}, we can stack a large number of them to improve the statistics \citep{ando10}."
+. IL their fluxes are not too faint. we can derive (heir mean flux aud photon index by Alaximum Likelihood (ML) method.," If their fluxes are not too faint, we can derive their mean flux and photon index by Maximum Likelihood (ML) method."
+ The (Murphyetal.2010) is the largest catalog of hieh frequency radio sources ancl contains 5890 sources with the Πακ at GGlIIz exceeding 40 mJy in the whole skv south of declination Q. In the south sky. about are associated with the AT20G sources (Ghirlandaοἱal.2010a).," The \citep{murphy10} is the largest catalog of high frequency radio sources and contains 5890 sources with the flux at GHz exceeding 40 mJy in the whole sky south of declination $^\circ$ In the south sky, about are associated with the AT20G sources \citep{G10a}."
+. Through studying the correlation between the +-rav and radio (lux density. of AT20G sources. vielded that AT20G sources not detected by the LAT can contribute We exclude the sources identified as Galactic Hj regions. Galactic Planetary Nebulas and parts of the Magellanie Clouds in AT20G. then the majority of sources (5803) we obtain are quasi-stellar objects (seeMurphyetal.2010.buttheopticalpropertiesof(hesehavenotbeen published)..," Through studying the correlation between the $\gamma$ -ray and radio flux density of AT20G sources, \citet{G10b} yielded that AT20G sources not detected by the LAT can contribute We exclude the sources identified as Galactic $_{\rm II}$ regions, Galactic Planetary Nebulas and parts of the Magellanic Clouds in AT20G, then the majority of sources (5808) we obtain are quasi-stellar objects \citep[see][but the optical properties of these objects have not been
+published]{murphy10}. ."
+ In order to minimize the influence of strong sources. we only use the sources that are at least 2° away [rom the nearest First Fermi-LAT catalog (LFCL) source and locate at highGalactic latitudes.|b]> 15°.Finally. we obtain 2900 sources to analvze their contribution to the EGD.," In order to minimize the influence of strong sources, we only use the sources that are at least $2^\circ$ away from the nearest First -LAT catalog (1FGL) source and locate at highGalactic latitudes,$|b|>15^\circ$ .Finally, we obtain 2900 sources to analyze their contribution to the EGB."
+The study of molecule formation in the post-recombination epoch has grown cousiderablv in recent vears.,The study of molecule formation in the post-recombination epoch has grown considerably in recent years.
+ Saslaw Zipov (1967) aud Peebles Dicke (1968) wore the fust to realize the importance of gas phase reactions for the formation of the simplest molecule. Il.," Saslaw Zipoy (1967) and Peebles Dicke (1968) were the first to realize the importance of gas phase reactions for the formation of the simplest molecule, $_2$."
+ Thev showed that trace amounts of molecular ivdrogen. of order 9 10°. could indeed form via the intermediaries species HL] and once the radiation Ποια 10 louger contained a high deusity of photous with chergics above the threshold of dissociation (2.61 aud V75 eV. respectively).," They showed that trace amounts of molecular hydrogen, of order $^{-6}$ $^{-5}$, could indeed form via the intermediaries species $_2^+$ and $^-$ once the radiation field no longer contained a high density of photons with energies above the threshold of dissociation (2.64 and 0.75 eV, respectively)."
+" The presence of even a trace abundance of IL, is of direct relevance for the cooling properties of the mimordial eas which. im its absence. would be au extremely poor radiator: cooling by Ly-a photons is in fact ineffective at temperatures XSOOO Is. well above the matter and raciation temperature iu the post-recombination cra."," The presence of even a trace abundance of $_2$ is of direct relevance for the cooling properties of the primordial gas which, in its absence, would be an extremely poor radiator: cooling by $\alpha$ photons is in fact ineffective at temperatures $\la 8000$ K, well above the matter and radiation temperature in the post-recombination era."
+ Since the evolution of primorcial density fuctuations is controlled by the ability of the eas to cool down to low temperatures. if is very müportaut to obtain a firm picture of the chemistry of the dust-free gas nmüxture. not linited to the formation of IL. but also to other molecules of potential interest.," Since the evolution of primordial density fluctuations is controlled by the ability of the gas to cool down to low temperatures, it is very important to obtain a firm picture of the chemistry of the dust-free gas mixture, not limited to the formation of $_2$, but also to other molecules of potential interest."
+ In this LORAM. Lepp Shull (1981) and Puy et al. (," In this regard, Lepp Shull (1984) and Puy et al. ("
+1993) have computed the abundances of Wo. WD and Lill as a function of redshift for various cosmological models.,"1993) have computed the abundances of $_2$, HD and LiH as a function of redshift for various cosmological models."
+ Although the final abundances of IT) and WD agree in the two calculations. their evolution with redshift is markedly different. since the epoch of formation varies by a factor of 2.," Although the final abundances of $_2$ and HD agree in the two calculations, their evolution with redshift is markedly different, since the epoch of formation varies by a factor of $\sim 2$."
+ Also. the ΤΠ abnucdauce shows a large discrepancy of about two orders of magnitude.," Also, the LiH abundance shows a large discrepancy of about two orders of magnitude."
+ More recently. Palla et al. (," More recently, Palla et al. ("
+1995) have analyzed the effects on the chemistry of the pregalactie eas of a high primordial D abundance in the light of the coutroversia results obtained towards high redshift quasars (see e.g. Tytler Burles 1997).,1995) have analyzed the effects on the chemistry of the pregalactic gas of a high primordial D abundance in the light of the controversial results obtained towards high redshift quasars (see e.g. Tytler Burles 1997).
+" They fouxd that the abundance of IL, is rather iuseusitive to variaions in the cosmologica parameters προς, by a factor of ~10 cuhanccemen of primordial |D/TI]. while ΠΟ aud Lill abundances vary by larecr amounts."," They found that the abundance of $_2$ is rather insensitive to variations in the cosmological parameters implied by a factor of $\sim 10$ enhancement of primordial [D/H], while HD and LiH abundances vary by larger amounts."
+ However. the abundance of," However, the abundance of"
+and is believed to have important implications.,and is believed to have important implications.
+ The upscattered photons are capable of producing the clectron-positron pairs and may compete with the curvature emission of the primary particles in controlling the pair production cascade in the polar eap of a pulsar (Sturner&Dermer1994:Sturner1995:Luo1996).," The upscattered photons are capable of producing the electron-positron pairs and may compete with the curvature emission of the primary particles in controlling the pair production cascade in the polar gap of a pulsar \citep{sd94,s95,l96}."
+. FPhus. the resonant Compton scattering may substantially alfect the characteristics of the secondary pulsar plasma (Llibschman&AronsAluslimoy 2002).," Thus, the resonant Compton scattering may substantially affect the characteristics of the secondary pulsar plasma \citep{ha01a,ha01b,ae02,hm02}."
+. his process can also account for the high-enerey spectra of magnetars 2007)..," This process can also account for the high-energy spectra of magnetars \citep*{tlk02,lg06,rzlt07,ft07,bt07,bh07}."
+ For a review of other applications of the resonant Compton scattering to pulsars see. e... (2006).," For a review of other applications of the resonant Compton scattering to pulsars see, e.g., \citet{hl06}."
+. The scattering of the non-thermal racio emission of pulsars olf the secondary plasma particles is an essentially distinct process and is also of interest., The scattering of the non-thermal radio emission of pulsars off the secondary plasma particles is an essentially distinct process and is also of interest.
+ Phe secondary. clectron-positron plasma streams ultrarelativisticallv. with 107. along the open magnetic lines of a pulsar ancl ultimately leaves the magnetosphere as a pulsar wind.," The secondary electron-positron plasma streams ultrarelativistically, with $\gamma\sim 10^2$, along the open magnetic lines of a pulsar and ultimately leaves the magnetosphere as a pulsar wind."
+ The radio emission is believed to originate inside the plasma How deep in the magnetosphere., The radio emission is believed to originate inside the plasma flow deep in the magnetosphere.
+ Hence. on its way in the magnetosphere and. bevond. pulsar raciation passes through the plasma and is subject to scattering olf the plasma particles.," Hence, on its way in the magnetosphere and beyond, pulsar radiation passes through the plasma and is subject to scattering off the plasma particles."
+ Since for radio frequencies 1. the quantum elfects on the scattering are neeligible and the classical treatment is appropriate.," Since for radio frequencies $h\nu/mc^2\ll 1$ , the quantum effects on the scattering are negligible and the classical treatment is appropriate."
+ In the vicinity of the emission region. the photon frequency in the particle rest frame is much less than the electron evrolrequency. o(1/—cos)óc; (here 3 is the particle velocity in units of e and @ is the angle between the photon wavevector and the particle velocity). ic. the magnetic field is strong enough to allect the scattering process.," In the vicinity of the emission region, the photon frequency in the particle rest frame is much less than the electron gyrofrequency, $\omega\gamma
+(1-\beta\cos\theta)\ll\omega_G$ (here $\beta$ is the particle velocity in units of $c$ and $\theta$ is the angle between the photon wavevector and the particle velocity), i.e. the magnetic field is strong enough to affect the scattering process."
+ As the magnetic field strength rapidly. decreases with distance from the neutron star. Dx/?57. in the outer magnetosphere the waves sulfer evclotron resonance. and in the pulsar wind the scattering is non-magnetic.," As the magnetic field strength rapidly decreases with distance from the neutron star, $B\propto R^{-3}$, in the outer magnetosphere the waves suffer cyclotron resonance, and in the pulsar wind the scattering is non-magnetic."
+ The scattering of pulsar radio emission in the magnetic and non-magnetic regimeshas first been considered in Dlandford (1978).," The scattering of pulsar radio emission in the magnetic and non-magnetic regimeshas first been considered in \citet{bs76,wr78}."
+. Because of extremely high brightness temperatures of pulsar radiation. Ix.qo the induced. scattering. strongly dominates. the spontaneous one and can be ellicient- in. both regimes.," Because of extremely high brightness temperatures of pulsar radiation, $T_B\sim 10^{25}-10^{30}$ K, the induced scattering strongly dominates the spontaneous one and can be efficient in both regimes."
+. Further. gauclics (Lvubarskii&Petrova1996:2004a.b) have demonstrated that the magnetized induced scattering can lead o substantial redistribution of intensity in frequency and pulse longitude and thus can account for. various phenomena haracteristie of the observed radio pulses.," Further studies \citep{lp96,p04a,p04b} have demonstrated that the magnetized induced scattering can lead to substantial redistribution of intensity in frequency and pulse longitude and thus can account for various phenomena characteristic of the observed radio pulses."
+ Close to the neutron star surface. the magnetic field is strong enough for any transverse momentum of the electrons o be almost immediately lost. via svnchrotron re-emission.," Close to the neutron star surface, the magnetic field is strong enough for any transverse momentum of the electrons to be almost immediately lost via synchrotron re-emission."
+ However. it is not the case in the outer magnetosphere. where 10 radio waves meet the condition of evclotron resonance.," However, it is not the case in the outer magnetosphere, where the radio waves meet the condition of cyclotron resonance."
+ Correspondingly. in the resonance region the waves are subject o evelotron absorption rather than resonant scattering.," Correspondingly, in the resonance region the waves are subject to cyclotron absorption rather than resonant scattering."
+ At the conditions relevant to pulsar magnetosphere. the process of cvclotron absorption not only allects the radio wave intensity (Blandford&Scharlemann1976:LvubarskiiPetrovaetrova.2002:Luo&AlelrosePOOL:FussellMelrose 2003).. but also leads to a substantial increase of the transverse momenta of the absorbing particles (Blandford&Scharlemann1976:LyubarskiiPetrova1998:2002.2003).," At the conditions relevant to pulsar magnetosphere, the process of cyclotron absorption not only affects the radio wave intensity \citep*{bs76,lp98,p02,melr_a,melr_b}, but also leads to a substantial increase of the transverse momenta of the absorbing particles \citep{bs76,lp98,p02,p03}."
+. Since he pulsar radiation. is. broadband. οn-10Tαν107Lu Hz. the resonance region. is. sullicientlyM extended.," Since the pulsar radiation is broadband, $\nu\sim n\cdot 10^7-n\cdot 10^{10}$ Hz, the resonance region is sufficiently extended."
+ EThe particles. entering. he resonance region acquire relativistic evration energies straight near its lower boundary. in the course of absorption of the waves with izLOM C (Petrova2002).," The particles entering the resonance region acquire relativistic gyration energies straight near its lower boundary, in the course of absorption of the waves with $\nu\geq 10^{10}$ G \citep{p02}."
+.. Then the lower-Irequeney waves. à«LOM Lz. which are still below the resonance. we(leos?)«ve; are subject to the magnetized scattering olf the relativistically evrating particles.," Then the lower-frequency waves, $\nu\ll 10^{10}$ Hz, which are still below the resonance, $\nu\gamma (1-\beta\cos\theta)\ll\nu_G$, are subject to the magnetized scattering off the relativistically gyrating particles."
+ This process will be examined in detail in the present paper., This process will be examined in detail in the present paper.
+ The scattering by a evrating electron cdillers substantially from that by an electron at rest., The scattering by a gyrating electron differs substantially from that by an electron at rest.
+ For the electron at rest. the scattering to high. harmonics holds only within the framework of the relativistic treatment. in the magnetic fields close to 1 critical value and at high enough frequencies (Herold1979:Melrose&Parle1983:DaughertyHarding1986).. (," For the electron at rest, the scattering to high harmonics holds only within the framework of the relativistic treatment, in the magnetic fields close to the critical value and at high enough frequencies \citep{h79,mp83,dh86}. ("
+Note wt the relativistic formalism. of the scattering has been developed. only for the case when the electron is initially at the eround Landau orbital.),Note that the relativistic formalism of the scattering has been developed only for the case when the electron is initially at the ground Landau orbital.)
+ For the evrating electron. the high-harmonic scattering is a purely classical effect and it may be ellicient in arbitrary magnetic fields for the incident. frequencies below the electron evrolrequency.," For the gyrating electron, the high-harmonic scattering is a purely classical effect and it may be efficient in arbitrary magnetic fields for the incident frequencies below the electron gyrofrequency."
+ In application to pulsars. je scattering by the electrons at high Landau levels transfers the radio photons into the optical and X-ray ranges and thus ον]λος a physical connection between the radio and high-energy emissions of a pulsar.," In application to pulsars, the scattering by the electrons at high Landau levels transfers the radio photons into the optical and X-ray ranges and thus provides a physical connection between the radio and high-energy emissions of a pulsar."
+ Other astrophysical applications of 10 process. e.g.. to svnchrotron sources. are not excluded as well.," Other astrophysical applications of the process, e.g., to synchrotron sources, are not excluded as well."
+ Lt should be noted that the pulsar radio emission is believed to be generated at frequencies of order of the local proper plasma frequency. ονwy. where ωναπλο/m and IN ds the number density of the plasma particles.," It should be noted that the pulsar radio emission is believed to be generated at frequencies of order of the local Lorentz-shifted proper plasma frequency, $\omega\sim\omega_p\sqrt{\gamma}$, where $\omega_p\equiv\sqrt{4\pi Ne^2/m}$ and $N$ is the number density of the plasma particles."
+ Therefore in the radio emission region and its close vicinity the scattering is a collective plasma process., Therefore in the radio emission region and its close vicinity the scattering is a collective plasma process.
+ The induced scatterings in the plasma are suggested as an important ingredient of the pulsar racio emission mechanism (see.e.g.Lominadze1996:Lyutikoy1998). and as a significant propagation ellect of the generated radiation (Gangadhara 2006).," The induced scatterings in the plasma are suggested as an important ingredient of the pulsar radio emission mechanism \citep[see, e.g.,][]{l79,l92,l93,lyub96,l98} and as a significant propagation effect of the generated radiation \citep{gk93,l98,lm06}."
+. However. as the plasma number density rapidly decreases with distance from the neutron star. INxA”. far enough from the emission region wπα and the plasma effects on the scattering become negligible.," However, as the plasma number density rapidly decreases with distance from the neutron star, $N\propto
+R^{-3}$, far enough from the emission region $\omega\gg\omega_p\sqrt{\gamma}$ and the plasma effects on the scattering become negligible."
+ Thus. the scattering in the resonance region can be considered as a single-particle process.," Thus, the scattering in the resonance region can be considered as a single-particle process."
+ Besides that. in this region the wave dispersion can also be ignored. so that the incident radiation presents transverse electromagnetic waves.," Besides that, in this region the wave dispersion can also be ignored, so that the incident radiation presents transverse electromagnetic waves."
+ A general formalism of the scattering by evrating electrons in the magnetoactive plasma has been developed in (1972).., A general formalism of the scattering by gyrating electrons in the magnetoactive plasma has been developed in \citet{ms72}. .
+ In the present paper. we introduce some corrections and simplifications to that treatment and obtain the scattering," In the present paper, we introduce some corrections and simplifications to that treatment and obtain the scattering"
+"Stellar parameters in this subsection are set as L,—1L...M, =LAL... and Tig=40001Ix for comparison wilh the solar svstem.","Stellar parameters in this subsection are set as $L_{*} = 1L_{\odot}$ , $M_{*} = 1M_{\odot}$ , and $T_{\mathrm{eff}} = 4000\mathrm{K}$ for comparison with the solar system."
+ Figure G shows the heliocentric distance of snow line as a [function of mass accretion rate., Figure \ref{solar_parameter_snowline} shows the heliocentric distance of snow line as a function of mass accretion rate.
+ Several curves correspond to dillerent dust grain sizes assumed in each calculation., Several curves correspond to different dust grain sizes assumed in each calculation.
+ Calculations are stopped when the optical depth in the region around (he snow line for the stellar radiation becomes less than unitv., Calculations are stopped when the optical depth in the region around the snow line for the stellar radiation becomes less than unity.
+ When the size of erains is >LOO yan. the scattering coefficient is set to be zero.," When the size of grains is $\geq 100$ $\mathrm{\mu m}$, the scattering coefficient is set to be zero."
+ During the inward migration of the snow line. the snow line location is almost independent of the cust erain size if the grain size is smaller than 10 jan. The wavelength of radiation from the disk interior at the condensation temperature (7LTO) is longer than 10 jan. so the opacity of the dust particle is in the Ravleigh regime. dependent only on the total mass of the dust particles. and dominated by the absorption.," During the inward migration of the snow line, the snow line location is almost independent of the dust grain size if the grain size is smaller than 10 $\mu$ m. The wavelength of radiation from the disk interior at the condensation temperature $T \simeq 170\mathrm{K}$ ) is longer than 10 $\mu$ m, so the opacity of the dust particle is in the Rayleigh regime, dependent only on the total mass of the dust particles, and dominated by the absorption."
+ On the other hand. when the size is larger than LO jan. ie. in (he geometrical optics regime. the opacity depends on the total cross section of the dust particles and decreases with the dust erain size. and hence the temperature enhancement by (he viscous heating becomes inefficient. although both absorption and scattering contribute in (his case.," On the other hand, when the size is larger than 10 $\mathrm{\mu m}$, i.e., in the geometrical optics regime, the opacity depends on the total cross section of the dust particles and decreases with the dust grain size, and hence the temperature enhancement by the viscous heating becomes inefficient, although both absorption and scattering contribute in this case."
+ When the dust grain size is larger than 10 jan. the evolutional track of the snow line levels off at the minimum value. because the opacity lor large erains is almost independent of the wavelength emitted [rom the disk.," When the dust grain size is larger than 10 $\mu$ m, the evolutional track of the snow line levels off at the minimum value, because the opacity for large grains is almost independent of the wavelength emitted from the disk."
+ This implies that the snow line necessarily passes the heliocentric distance of 1 AU. unless the dust grain size exceeds a certain large value.," This implies that the snow line necessarily passes the heliocentric distance of 1 AU, unless the dust grain size exceeds a certain large value."
+ This point will be considered in 84.4., This point will be considered in 4.4.
+" The shift ratio of the snow line fa, is qualitatively the same as in Figure 5..", The shift ratio of the snow line $f_{\rm SL}$ is qualitatively the same as in Figure \ref{snowline_evolution}.
+ The snow line shifts to a larger heliocentric cistance in ils inwarcl migration phase. whereas itdoesnol shift substantially in its outward migration and leveling-off phases.," The snow line shifts to a larger heliocentric distance in its inward migration phase, whereas itdoesnot shift substantially in its outward migration and leveling-off phases."
+Table 2 for features with fixes up o LO% of the brightest niaser spots;,Table \ref{results} for features with fluxes up to $10\%$ of the brightest maser spots.
+ Some selected spectra are shown with fits of the LTE Zecai models and the uou-LTE models., Some selected spectra are shown with fits of the LTE Zeeman models and the non-LTE models.
+ When using the non-LTE models. fits with Ομ=1.0 kan/s Oogenerally are |vost.," When using the non-LTE models, fits with $v_{\rm thermal} = 1.0$ km/s generally are best."
+ As expected the splitting is small aud the circulay polarization is ecucrally not more than 0.5.LOX., As expected the splitting is small and the circular polarization is generally not more than $0.5 - 1.0\%$.
+ Magnetic field strenetis obtained using the LTE Zeeman models are given in οmun 7 while the results obtained uxiug the nou-LTE models are shown iu column 8., Magnetic field strengths obtained using the LTE Zeeman models are given in column 7 while the results obtained using the non-LTE models are shown in column 8.
+ The featires that show less than a 236 circular polarization signal are considered nou-detections., The features that show less than a $3\sigma$ circular polarization signal are considered non-detections.
+" For these we have determined (ahsolu, pper hmits using the Appg cocfficieuts obtained. using he NW radiative transfer models.", For these we have determined (absolute) upper limits using the $A_{\rm F-F'}$ coefficients obtained using the NW radiative transfer models.
+" When using the LTE Zeeman analysis the limits iucrease by a factor of z1.1,", When using the LTE Zeeman analysis the limits increase by a factor of $\approx 1.4$.
+ Features labeled je V-spectra that did not allow accurate fitting. possibv due to blending effects.," Features labeled $^*$ have V-spectra that did not allow accurate fitting, possibly due to blending effects."
+ They do show cimeulaur polarization above the 23e level. although eenerallv ouly slightly more.," They do show circular polarization above the $3\sigma$ level, although generally only slightly more."
+" For these we have used Viu, and Vias to estimate the magnetic field streugth. aud we asstumed the best estiuate for the coefficient 0.015."," For these we have used $V_{\rm min}$ and $V_{\rm max}$ to estimate the magnetic field strength, and we assumed the best estimate for the coefficient $A_{\rm F-F'} = 0.018$ ."
+ The exrors are t1ο formal lo errors., The errors are the formal $1\sigma$ errors.
+ We have not been able to detect linear poluizatiou Π anv of the 1| sources., We have not been able to detect linear polarization in any of the 4 sources.
+ This gives upper limits to the fractional linear polarization of z1€ for the weakest features down to zc0.01% for the strongest., This gives upper limits to the fractional linear polarization of $\approx 1\%$ for the weakest features down to $\approx 0.01\%$ for the strongest.
+ Fie., Fig.
+ 9 shows the total intensity map of the IT2O. maser features around ο Per. in which we cau ideutifv most of the features detected m earlier observations(Diuuond ct al.," \ref{sper} shows the total intensity map of the $_2$ O maser features around S Per, in which we can identify most of the features detected in earlier observations(Diamond et al."
+ 1987: Marvel 1997)., 1987; Marvel 1997).
+ Positions are shown relaive to the feature laheledb. discussed iu more detail iu. VOI.," Positions are shown relative to the feature labeled, discussed in more detail in V01."
+ Circular polarization between 0.20.74 has becu etected in 2ofthe brightest features after a careful recxamunation of he data.," Circular polarization between $0.2 -
+0.7\%$ has been detected in 3 of the brightest features after a careful reexamination of the data."
+ This has also resulted in a slightlv lower magnetic feld streneth value for than he one presented in VOL after optimizing the fting rouines., This has also resulted in a slightly lower magnetic field strength value for than the one presented in V01 after optimizing the fitting routines.
+ The fim YO:so shows the total power and circular volarization spectra for 2 of the features. with fits for both tιο LTE alc unon-LTE Zeeman models.," The figure also shows the total power and circular polarization spectra for 2 of the features, with fits for both the LTE and non-LTE Zeeman models."
+ The velocity abeliug has chauge from VOL to correspond to the couvenion used by NW., The velocity labeling has changed from V01 to correspond to the convention used by NW.
+ We fiid the magnetic field poiutiugo away fro11 is Oll all featiPCs., We find the magnetic field pointing away from us on all features.
+ From these results we estimate the magnetic field in the II2O maser region around S Per o be zc 200 mC using the LTE Zeeman analysis or z 15! iu using the nou-LTE motels., From these results we estimate the magnetic field in the $_2$ O maser region around S Per to be $\approx$ 200 mG using the LTE Zeeman analysis or $\approx$ 150 mG using the non-LTE models.
+ The iaxiuunu radius of the TeO maser region was determined to be z60 mas by Diuuond et al(987). correspondiug to zLOO AU.," The maximum radius of the $_2$ O maser region was determined to be $\approx 60$ mas by Diamond et al.(1987), corresponding to $\approx
+100$ AU."
+ They conclude for a thick shell model that the ner radius of the water maser shell Is ~--30 amas. while the outer radius is ~--TU as.," They conclude for a thick shell model that the inner radius of the water maser shell is $\approx 30$ mas, while the outer radius is $\approx 70$ mas."
+ For comparison. the ΟΠ maser shell extent was determined by them to be 80.30 11as. using MERLEN observations.," For comparison, the OH maser shell extent was determined by them to be $80 \times 30$ mas, using MERLIN observations."
+Gamma-ray Bursts (GRB) are the most luminous flashes of y-rays known to humankind.,Gamma-ray Bursts (GRB) are the most luminous flashes of $\gamma$ -rays known to humankind.
+ It is generally believed that they originate from a compact source with highly relativistic collimated outflows (TF> 100)., It is generally believed that they originate from a compact source with highly relativistic collimated outflows $\Gamma > 100$ ).
+ A large fraction of our knowledge of the prompt emission comes from the Burst and Transient Explorer (BATSE.?) onboard the Compton Gamma-Ray Observatory(CGRO.. 1991-2000).," A large fraction of our knowledge of the prompt emission comes from the Burst and Transient Explorer \citep[\batse,][]{meegan92} onboard the Compton Gamma-Ray Observatory, 1991-2000)."
+ Unfortunately. only a handful of bursts had a measured redshift because the error boxes were too large and thus. follow-up observations with X-ray and optical instrumentation was very limited.," Unfortunately, only a handful of bursts had a measured redshift because the error boxes were too large and thus, follow-up observations with X-ray and optical instrumentation was very limited."
+ Moreover. the first afterglow was only detected in 1997. already near the end of the mission.," Moreover, the first afterglow was only detected in 1997, already near the end of the mission."
+ The lack of distance measurements led to a focus of GRB studies in the observer frame without redshift corrections., The lack of distance measurements led to a focus of GRB studies in the observer frame without redshift corrections.
+ Due to the cosmological origin of GRBs. such a correction is likely to be necessary to understand the intrinsic nature of these events.," Due to the cosmological origin of GRBs, such a correction is likely to be necessary to understand the intrinsic nature of these events."
+ With the two dedicated satellites. (?) and (?).. the situation has changed and afterglow and host galaxy spectroscopy has provided redshifts for more than 200 events by now.," With the two dedicated satellites, \citep{boella97} and \citep{gehrels04}, the situation has changed and afterglow and host galaxy spectroscopy has provided redshifts for more than 200 events by now."
+" Unfortunately. the relatively narrow energy band of (0.1 keV - 300 keV) and Swift//BAT (15 keV - 150 keV) limits the constraints on the prompt emission spectrum because the peak energy.£).. in the v7, spectrum of GRBs. can only be determined for low values and is often unconstrained (2?).."," Unfortunately, the relatively narrow energy band of (0.1 keV - 300 keV) and /BAT (15 keV - 150 keV) limits the constraints on the prompt emission spectrum because the peak energy, in the $\nu F_{\nu}$ spectrum of GRBs, can only be determined for low values and is often unconstrained \citep[][]{butler07,saka09}."
+ In this work we will take advantage of the broad energy coverage of the Fermi//GBM (8 keV - 40 MeV) to study the primary spectral and temporal properties. such asEp..Του. the time interval in which of the burst fluence has been observed. and£i... the isotropic equivalent bolometric energy. in the rest-frame of the progenitors of 32 GRBs with measured redshift.," In this work we will take advantage of the broad energy coverage of the /GBM (8 keV - 40 MeV) to study the primary spectral and temporal properties, such as, the time interval in which of the burst fluence has been observed, and, the isotropic equivalent bolometric energy, in the rest-frame of the progenitors of 32 GRBs with measured redshift."
+ The Gamma-Ray Burst Monitor (GBM) ts one of the instruments onboard the Gamma-Ray Space Telescope (?) launched on June 11. 2008.," The Gamma-Ray Burst Monitor (GBM) is one of the instruments onboard the Gamma-Ray Space Telescope \citep{atwood09} launched on June 11, 2008."
+ Specifically designed for GRB studies. GBM observes the whole unocculted sky with a total of 12 thallium-activated sodium iodide (NalCTD) scintillation detectors covering the energy range from 8 keV to 1 MeV and two bismuth germanate scintillation detectors (BGO) sensitive to energies between 190 keV and 40 MeV (?)..," Specifically designed for GRB studies, GBM observes the whole unocculted sky with a total of 12 thallium-activated sodium iodide (NaI(Tl)) scintillation detectors covering the energy range from 8 keV to 1 MeV and two bismuth germanate scintillation detectors (BGO) sensitive to energies between 150 keV and 40 MeV \citep{meegan09}."
+ Thus. GBM offers a unprecedented view of GRBs. covering more then 3 decades in energy.," Thus, GBM offers a unprecedented view of GRBs, covering more then 3 decades in energy."
+ The selection criterion. for our sample is solely based on the redshift determination., The selection criterion for our sample is solely based on the redshift determination.
+ We form a sample of 32 bursts detected by GBM up to October l6th. 2010. with known redshift (determined either spectroscopically or photometrically).," We form a sample of 32 bursts detected by GBM up to October 16th, 2010, with known redshift (determined either spectroscopically or photometrically)."
+ Our sample contains 4+ short and 28 long GRBs., Our sample contains 4 short and 28 long GRBs.
+ The redshift distribution of the GBM GRBs is shown in Fig., The redshift distribution of the GBM GRBs is shown in Fig.
+ | together with a histogram of all 239 GRBs with redshift determinations todate!., \ref{fig:histoz} together with a histogram of all 239 GRBs with redshift determinations to.
+. A two-tailed Kolmogorov-Smirnov (KS) test between the full sample and the GBM-only sample shows that the two distributions are very similar (P=84 %))., A two-tailed Kolmogorov-Smirnov (KS) test between the full sample and the GBM-only sample shows that the two distributions are very similar $P=84$ ).
+ In conclusion. the GBM-only sample is representative of the full GRB sample with redshift.," In conclusion, the GBM-only sample is representative of the full GRB sample with redshift."
+stronely clustered. but. distributed. fairly randomly across the region except for a central swathe of exceptionally. high completeness introduced hy virtue of DRT data.,strongly clustered but distributed fairly randomly across the region except for a central swathe of exceptionally high completeness introduced by virtue of DR7 data.
+ llavingD established. a sub-region5 anc a mechanism of correcting for the spectroscopic incompleteness as a function of apparent magnitude ancl spatial location we now require a. πο of test particles to. start exploring the cosmic variance within the test region., Having established a sub-region and a mechanism of correcting for the spectroscopic incompleteness as a function of apparent magnitude and spatial location we now require a set of test particles to start exploring the cosmic variance within the test region.
+ We adopt as test xuwiicles the most common galaxy. in a locally observed sample. Le. AZ”+LO mag galaxies.," We adopt as test particles the most common galaxy in a locally observed sample, i.e., $M^*\pm1.0$ mag galaxies."
+" ""These are both numerous and readily detectable to large distances.", These are both numerous and readily detectable to large distances.
+ Selecting xieghter svstemis potentially introduces additional variance as brighter sources are known to cluster more. strongly (ic.. atvpical: ef.," Selecting brighter systems potentially introduces additional variance as brighter sources are known to cluster more strongly (i.e., atypical; c.f.,"
+ Norberg et al., Norberg et al.
+ 2002b). whereas low unminosity sources would restrict. the depth of our volume and their intrinsic clustering properties are less well known.," 2002b), whereas low luminosity sources would restrict the depth of our volume and their intrinsic clustering properties are less well known."
+ We adopt. AJSlog(toc)=21.58 mag taken [rom he recent. ugrizYJHIN LE estimates— of Hill et al. (, We adopt $M_r^*-5\log_{10}(h_{0.7})=-21.58$ mag taken from the recent $ugrizYJHK$ LF estimates of Hill et al. (
+2010).,2010).
+ The αμ is used as this is the filter in which the SDSS main spectroscopic sample is selected., The r-band is used as this is the filter in which the SDSS main spectroscopic sample is selected.
+ Fig., Fig.
+ 3. shows the clistribution of absolute magnitude versus redshift with the systems within our defined volume shown in green., \ref{mz} shows the distribution of absolute magnitude versus redshift with the systems within our defined volume shown in green.
+" In deriving the absolute magnitudes we adopted universal &(z) and e(z) corrections (&(z)=1.2632|O.S95270.566? and e(z)=2.5logy,(112) ΜΕ. ", In deriving the absolute magnitudes we adopted universal $k(z)$ and $e(z)$ corrections $k(z)=1.263z+0.895z^2-0.566z^3$ and $e(z)=2.5\log_{10}[(1+z)^{-0.5}]$ ).
+The choice of &(z) and c(z) corrections are not critical nor is the assumption of universal rather than individual corrections as firstly the corrections are small at low-z (2«0.1. see Lill οἱ al.," The choice of $k(z)$ and $e(z)$ corrections are not critical nor is the assumption of universal rather than individual corrections as firstly the corrections are small at low-z $z<0.1$, see Hill et al."
+ 2010). and secondly we will always compare cells constructed over identical redshift ranges.," 2010), and secondly we will always compare cells constructed over identical redshift ranges."
+ From Fig., From Fig.
+ 3. we can see that he maximum redshift we're capable of sampling. due to he SDSS spectroscopic limit. is 2= 0.1.," \ref{mz} we can see that the maximum redshift we're capable of sampling, due to the SDSS spectroscopic limit, is $z=0.1$ ."
+ Fig., Fig.
+ 4. shows he variance along the line of sight bx counting the space-density of test particles (ie. AdE1.0 galaxies) as a function of redshift. both cumulative (mauve curve) ancl cülferential (ercen data points with errorbars) distributions are shown.," \ref{cosvar} shows the variance along the line of sight by counting the space-density of test particles (i.e., $M^*\pm1.0$ galaxies) as a function of redshift, both cumulative (mauve curve) and differential (green data points with errorbars) distributions are shown."
+ The cilferential counts show that the SDSS appears to sulfer rom an extreme uncderdensity locally (2« 0.02) followed by a strong overdensity at (2= 0.022) with the cumulative clensity well behaved from z—0.03 to z—0.1 (within of the mean density in this range. see red lines on Fig. 4))," The differential counts show that the SDSS appears to suffer from an extreme underdensity locally $z<0.02$ ) followed by a strong overdensity at $z=0.022$ ) with the cumulative density well behaved from $z=0.03$ to $z=0.1$ (within of the mean density in this range, see red lines on Fig. \ref{cosvar}) )"
+ after which incompleteness (i... traditional Malmequist bias) starts to alfect the sample (rf," after which incompleteness (i.e., traditional Malmquist bias) starts to affect the sample (r.f."
+ Fig. 3))., Fig. \ref{mz}) ).
+" We therefore adopt az range of 0.03 to 0.10 (minimum/maximunm transverse co-moving scale of 2.2/7.3 Mpce/deg) which contains 100117 test galaxies distributed over an area of 5150 sq.deg with a co-moving radial length of 291, a lookback intervalof 0.96vr. and a volume of 3.7£M«pc.10"" A, 2Mpce."," We therefore adopt a z range of 0.03 to 0.10 (minimum/maximum transverse co-moving scale of 2.2/7.3 Mpc/deg) which contains 100117 test galaxies distributed over an area of 5150 sq.deg with a co-moving radial length of $291h^{-1}_{0.7}$ Mpc, a lookback intervalof $\sim$ 0.9Gyr, and a volume of $3.7 \times 10^7$ $h^{-3}_{0.7}$ $^{3}$."
+" In the results that follow we will therefore only be sensitive to cosmic variance on scales below 10, 2Mpe. however as cosmic variance is generally seen to be decreasing the loss of sensitivity to variance on scales &reater than this volume is not expected to be a significant issue."," In the results that follow we will therefore only be sensitive to cosmic variance on scales below $10^7h^{-3}_{0.7}$ $^3$, however as cosmic variance is generally seen to be decreasing the loss of sensitivity to variance on scales greater than this volume is not expected to be a significant issue."
+" In what follows the racial co-moving distance is typically comparable or larger (~290h, ENIpe) than the transverse lengths (< 2504EMpe) and is therefore not the dominating dimension contributing to the cosmic variance.", In what follows the radial co-moving distance is typically comparable or larger $\sim 290h^{-1}_{0.7}$ Mpc) than the transverse lengths $<250h^{-1}_{0.7}$ Mpc) and is therefore not the dominating dimension contributing to the cosmic variance.
+" This would not be the case if the depth was quantisecl more discreet. hence for what follows we have two caveats: (1) We are not sensitive to the cosmic variance on volumes larger than 10'5,2Mpc* (2) We require that the shortest. co-moving lengths defining the volume are tangential to the line-of-5ight."," This would not be the case if the depth was quantised more discreetly, hence for what follows we have two caveats: (1) We are not sensitive to the cosmic variance on volumes larger than $10^7h^{-3}_{0.7}$ $^{3}$ (2) We require that the shortest co-moving lengths defining the volume are tangential to the line-of-sight."
+ In general this will be the case if the depth interval is >O.9Civr in lookback time., In general this will be the case if the depth interval is $\geq 0.9$ Gyr in lookback time.
+ Finally Fig., Finally Fig.
+ ο shows the histogram of the counts in 1 sq., \ref{gauss} shows the histogram of the counts in 1 sq.
+ deg cells for our test. particles indicating a well behavecl distribution somewhere between an ideal Normal and logNormal distribution typically in. previous studies a Normal distribution is assumed. and we therefore. follow convention., deg cells for our test particles indicating a well behaved distribution somewhere between an ideal Normal and logNormal distribution — typically in previous studies a Normal distribution is assumed and we therefore follow convention.
+" We elect to quantify cosmic variance using the simplest measure of standard: deviation. defined. as: GCosVarusCAyuON100 where Cay.—NNredsSTNT] N> is the mean ofks the counts in cells for that particular cell size. and N; the counts in the "" cell"," We elect to quantify cosmic variance using the simplest measure of standard deviation defined as: $\zeta_{\rm Cos. Var.}(\%)=\frac{\sigma_{\rm Var.}}{<N>}
+ \times 100$ where $\sigma_{\rm Var.}^2=\frac{\Sigma[<N> -
+ N_i]^2}{n}$, $<N>$ is the mean of the counts in cells for that particular cell size, and $N_i$ the counts in the $i^{th}$ cell."
+ Note that this statistic will by definition include the intrinsic Poisson component which in all cases is the minor component (see blue line on Vig. 7))., Note that this statistic will by definition include the intrinsic Poisson component which in all cases is the minor component (see blue line on Fig. \ref{fig:sample1}) ).
+ llaving defined an extensive test. region (see Fig. 5)), Having defined an extensive test region (see Fig. \ref{testvol}) )
+ containing test particles which reflect the underlying local structures contained in the volume. we can now sample the variance by repetitively extracting counts in fixed sized cells (truncated square-based pyramicds) at random locations and repeat for increasing cell sizes.," containing test particles which reflect the underlying local structures contained in the volume, we can now sample the variance by repetitively extracting counts in fixed sized cells (truncated square-based pyramids) at random locations and repeat for increasing cell sizes."
+ In the next two subsections we firstly. explore the basic variance in square cells. [rom l sq«deg το 2048 sq.deg. ancl secondly the variance. in rectangular cells where the aspect ratio is varied [rom 1:1 to 1:128.," In the next two subsections we firstly explore the basic variance in square cells from 1 sq.deg to 2048 sq.deg, and secondly the variance in rectangular cells where the aspect ratio is varied from 1:1 to 1:128."
+ In the sample extractions which follow we only sample individual test. particles once. resulting in a Lully uncorrelated measure of the cosmic variance but a reduction in the statistical accuracy for the Larger samples and aspect ratios where fewer independent samples can be constructed.," In the sample extractions which follow we only sample individual test particles once, resulting in a fully uncorrelated measure of the cosmic variance but a reduction in the statistical accuracy for the larger samples and aspect ratios where fewer independent samples can be constructed."
+ Fig., Fig.
+ 7 shows the (cosmic) variance (data points) derived [rom our test region., \ref{fig:sample1} shows the (cosmic) variance (data points) derived from our test region.
+ This is shown as à percentage versus volume sampled., This is shown as a percentage versus volume sampled.
+ The range of individual values measured is shown bv the grey. bars and the accuracy to which the mean variance is measured. is shown as errorbars., The range of individual values measured is shown by the grey bars and the accuracy to which the mean variance is measured is shown as errorbars.
+" We can see that the variance decreases steadily from GO per cent for volumes of 1075, 2Mipe? to 10 per cent at our sampling limit of 10h,pe.", We can see that the variance decreases steadily from 60 per cent for volumes of $10^4h^{-3}_{0.7}$ $^3$ to 10 per cent at our sampling limit of $10^7h^{-3}_{0.7}$ $^{3}$.
+ We fit à simple second order polynomial to the data 2Mand find a good (<x1 percent) fit given by: where Ccosvae represents the cosmic variance for a volume. V.M as a percentage (1.0..M S>LouAN ).," We fit a simple second order polynomial to the data and find a good $< \pm 1$ percent) fit given by: where $\zeta_{\rm Cos. Var}$ represents the cosmic variance for a volume, $V$ , as a percentage (i.e., $\frac{100\Delta N}{N}$ )."
+ This expression can be used to provide a robust estimate of the cosmic variance for any given square shaped z«0.1 survey given thetotal volume sanipled., This expression can be used to provide a robust estimate of the cosmic variance for any given square shaped $z<0.1$ survey given thetotal volume sampled.
+exhibiting lour peak components.,exhibiting four peak components.
+ IH the hole was produced by star formation activity. it would seem that (he gas is still somewhat unsettlecd.," If the hole was produced by star formation activity, it would seem that the gas is still somewhat unsettled."
+ Overall. DDO 43 is a fairly. (vpical eas-rich. optically faint. small dwarl galaxy.," Overall, DDO 43 is a fairly typical gas-rich, optically faint, small dwarf galaxy."
+ I6 has an extensive gas disk and a high. Mj; /Lj ratio., It has an extensive gas disk and a high $_{HI}$ $_B$ ratio.
+ The hhas been sculpted into a high density ridge forming a parGal ring that contains knots and holes., The has been sculpted into a high density ridge forming a partial ring that contains knots and holes.
+ This is most likely the result of multiple episodes of star formation al various times., This is most likely the result of multiple episodes of star formation at various times.
+ The mechanical energy. deposited in the gas has carved out regions of lower gas clensily as eas was swept up into expanding shells., The mechanical energy deposited in the gas has carved out regions of lower gas density as gas was swept up into expanding shells.
+ High density regions occur where the shells have overlapped., High density regions occur where the shells have overlapped.
+ As a result. star formation was triggered in these areas as current star formation aclivily is located (here.," As a result, star formation was triggered in these areas as current star formation activity is located there."
+ The shells are not currently expanding. so the events that formed (hem must have happened long ago.," The shells are not currently expanding, so the events that formed them must have happened long ago."
+ The current star formation rate is modest: coupled with the large gas reservoir and modest sizes of the older clusters seen in V. it seems that DDO 43 has evolved slowly. and will continue to do so for the next several (αντ.," The current star formation rate is modest; coupled with the large gas reservoir and modest sizes of the older clusters seen in V, it seems that DDO 43 has evolved slowly, and will continue to do so for the next several Gyr."
+ We wish to thank Emily Bowsher for examining the color ratio images of DDO 43 as part of the 2003 Research Experiences for Undergraduates program of Northern Arizona University., We wish to thank Emily Bowsher for examining the color ratio images of DDO 43 as part of the 2003 Research Experiences for Undergraduates program of Northern Arizona University.
+ DAI obtained partial support for this research from the Lowell Research Fund., DAH obtained partial support for this research from the Lowell Research Fund.
+ Additional support to DAIL. support for TEN. and support for travel to Flagstalf for CES came from grant AST-0204922 from the National Science Foundation.," Additional support to DAH, support for TEN, and support for travel to Flagstaff for CES came from grant AST-0204922 from the National Science Foundation."
+ This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Acministration.," This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
+an optical depth of 110.,an optical depth of $1-10$.
+ Το investigate even lower density systems a pixel. correlation technique has been emploved where. instead of fitting. an associated. absorption line. a pixel-by-pixel search. for excess absorption at the corresponding wavelength is. performed. (2)...," To investigate even lower density systems a pixel correlation technique has been employed where, instead of fitting an associated absorption line, a pixel-by-pixel search for excess absorption at the corresponding wavelength is performed \citep{1998Natur.394...44C}."
+ Τις method was hoped to be. more robust than the use of line ‘stacking’ (7: 7)). but there is some ambiguity in the interpretation of the results.," This method was hoped to be more robust than the use of line `stacking' \citealt{1995qal..conf..289T}; \citealt{2000AJ....120.1175E}) ), but there is some ambiguity in the interpretation of the results."
+ At the relevant redshifts. the species is the most sensitive tracer of metals in. low density regions of the IGM. provided that the spectrum of the UV background is sulliciently hard (? and ?)).," At the relevant redshifts, the species is the most sensitive tracer of metals in low density regions of the IGM, provided that the spectrum of the UV background is sufficiently hard \citealt*{1997ApJ...481..601R} and \citealt{1998ApJ...499..172H}) )."
+" However. the relevant absorption lines (Ao,=1032... 1038A)) occur in the same wavelength range as the forest and. higher order lines of the Lyman series. which makes the detection of weak absorption challenging."," However, the relevant absorption lines $\lambda_{O\textsc{vi}a} = 1032$, $\lambda_{O\textsc{vi}b} = 1038$ ) occur in the same wavelength range as the forest and higher order lines of the Lyman series, which makes the detection of weak absorption challenging."
+ ? claim to have detected at optical depths significantly. smaller than unity in some QSO absorption spectra ancl used. the optical depth-density relation of numerical hyelroclvnamic simulations to argue that this indicates a detection. of metals in underdense regions of the IM., \citet{2000ApJ...541L...1S} claim to have detected at optical depths significantly smaller than unity in some QSO absorption spectra and used the optical depth-density relation of numerical hydrodynamic simulations to argue that this indicates a detection of metals in underdense regions of the IGM.
+ In this work. we apply the pixel-bv-pixel search to synthetic spectra and thereby. investigate instrumenta and physical ellects on this technique (see also 2)).," In this work, we apply the pixel-by-pixel search to synthetic spectra and thereby investigate instrumental and physical effects on this technique (see also \citealt*{2002ApJ...576....1A}) )."
+ We then assess the statistical significance of the detection aux the implied density threshold above which the presence of metals is confidently detected., We then assess the statistical significance of the detection and the implied density threshold above which the presence of metals is confidently detected.
+ The speetra (both simulated αμα observed) are described. in Section 2.. In., The spectra (both simulated and observed) are described in Section \ref{Simulated and Observed Spectra}.
+ Section 3 we outline the pixel. correlation method used. to detect. the presence of absorption.," In Section \ref{Pixel-by-pixel search}
+ we outline the pixel correlation method used to detect the presence of absorption."
+ The results of the search for are shown and. discussed in Section 4.., The results of the search for are shown and discussed in Section \ref{Results From Simulated Spectra}.
+ In Section 5 Νο explore the consequences for the metal enrichment of the low density 1M., In Section \ref{Search for OVI in the low density IGM} we explore the consequences for the metal enrichment of the low density IGM.
+ We use the method developed by Bi (?:: 72)) to analytically calculate svnthetic spectra in the fluctuating iunn-DPeterson approximation for a warm photoionised LGAL with a density distribution that Is expected in a ACDAL structure formation model., We use the method developed by Bi \citealt*{1992A&A...266....1B}; \citealt{1997ApJ...479..523B}) ) to analytically calculate synthetic spectra in the fluctuating Gunn-Peterson approximation for a warm photoionised IGM with a density distribution that is expected in a $\Lambda$ CDM structure formation model.
+ Phe basic assumptions are: ‘Table 1 shows the model parameters., The basic assumptions are: Table \ref{standardparam} shows the model parameters.
+ A brief summary of how we calculate the svnthetic spectra follows., A brief summary of how we calculate the synthetic spectra follows.
+ According to standard. structure. formation. models the matter distribution was initially a Gaussian. random Ποιά with small density DHuctuations that evolve. uncer gravity into the pattern of sheets ancl filaments characteristic of CDI models., According to standard structure formation models the matter distribution was initially a Gaussian random field with small density fluctuations that evolve under gravity into the pattern of sheets and filaments characteristic of CDM models.
+ Many. features of such a distribution can be reproduced by a rank-ordered mapping of a linear Ciaussian random field onto a lognormal PDE of the density., Many features of such a distribution can be reproduced by a rank-ordered mapping of a linear Gaussian random field onto a lognormal PDF of the density.
+ This idea is at the heart of the method used to calculate svnthetic absorption spectra (2)..., This idea is at the heart of the method used to calculate synthetic absorption spectra \citep{1992A&A...266....1B}.
+ Phis method is an ellicient way of »oducing large numbers of realistic synthetic spectra., This method is an efficient way of producing large numbers of realistic synthetic spectra.
+ Note hat we are interested in producing synthetic spectra with a realistic density PDE and realistic instrumental properties., Note that we are interested in producing synthetic spectra with a realistic density PDF and realistic instrumental properties.
+ teproducing the detailed clustering properties of the density 101 is less important., Reproducing the detailed clustering properties of the density field is less important.
+" We beein by creating a linear Gaussian random field which is Lully determined by its power spectrum. Loa(he). or which we have used the form given hy ?.. where αΞξο. b23.0/F. e—LT/U. P=Q,h (δη, h have their normal definitions) and vy=1.13."," We begin by creating a linear Gaussian random field which is fully determined by its power spectrum, $P_{DM}(k)$, for which we have used the form given by \citet*{1992MNRAS.258P...1E}, where $a=6.4/\Gamma$, $b=3.0/\Gamma$, $c=1.7/\Gamma$, $\Gamma=\Omega_mh$ $\Omega_m$ $h$ have their normal definitions) and $\nu=1.13$."
+ The power spectrum. was normalised by the Ductuation amplitude on a Sh!Mpe scale. ox. as given in Table 1.," The power spectrum was normalised by the fluctuation amplitude on a $\mathrm{8h^{-1}Mpc}$ scale, $\sigma_8$, as given in Table 1."
+ In order to take into account pressure effects that suppress I[uctuations on scales smaller than a certain filtering scale (7) the power spectrum of the barvon density in the linear regime is assumed to have the form. where aj is à comoving scale related to the Jeans length (?: ?)). 5 is the ratio of specific heats. Zr is the effective mean temperature and fis the molecular weight of the ICM.," In order to take into account pressure effects that suppress fluctuations on scales smaller than a certain filtering scale \citep{1998MNRAS.296...44G} the power spectrum of the baryon density in the linear regime is assumed to have the form, where $x_{\rmn{b}}$ is a comoving scale related to the Jeans length \citealt{1980lssu.book.....P}; ; \citealt{1997ApJ...479..523B}) ), $\gamma$ is the ratio of specific heats, $\bar{T}_{\rm{eff}}$ is the effective mean temperature and $\mu$ is the molecular weight of the IGM."
+ An appropriate value for Zi has been determined. [from comparison with hvdrodynamic simulations by 2.., An appropriate value for $\bar{T}_{\rm{eff}}$ has been determined from comparison with hydrodynamic simulations by \citet{1997ApJ...479..523B}.
+ This is the barvonic matter power spectrum at >=0 ancl linear theory is used to evolve it backward to high recshillt., This is the baryonic matter power spectrum at $z=0$ and linear theory is used to evolve it backward to high redshift.
+ Equations (1)) and (2)) provide the 3D power spectrum., Equations \ref{powspec}) ) and \ref{jeans}) ) provide the 3D power spectrum.
+ A 1D power spectrum needed for the description of line-of-sight (LOS) IIuctuations is obtained by an integration of the 3D power spectrum (2).., A 1D power spectrum needed for the description of line-of-sight (LOS) fluctuations is obtained by an integration of the 3D power spectrum \citep{1991ApJ...379..482K}.
+ “Phe matter density and the peculiar. velocity field. are related ancl can be cleseribecl as coupled Gaussian random fields (see 2? for more details)., The matter density and the peculiar velocity field are related and can be described as coupled Gaussian random fields (see \citealt{1992A&A...266....1B} for more details).
+ Random realisations of the density. ancl peculiar velocity in real space (co-moving coordinates) are obtained using a Fourier transform routine., Random realisations of the density and peculiar velocity in real space (co-moving coordinates) are obtained using a Fourier transform routine.
+ Phe non-linear. evolution is modelled by a local rank-ordercd mappingto a lognormal density probability. distribution.," The non-linear evolution is modelled by a local rank-ordered mappingto a lognormal density probability distribution,"
+exponential disks.,exponential disks.
+ We displax snapshots of the disk's particles in Figures 5. and 6.., We display snapshots of the disk's particles in Figures \ref{fig5} and \ref{fig6}.
+. Examining these slapshots. we see that at 0.08. 0.16. 0.24. and 0.32 Gyr times the svstem presents alvpical substructures. in particular in its central region (2?S10.0 kpe).," Examining these snapshots, we see that at 0.08, 0.16, 0.24, and 0.32 Gyr times the system presents atypical substructures, in particular in its central region $R \lesssim 10.0$ kpc)."
+ As the disk evolves. an annular substructure appears. and (he central region becomes almost empty of particles. a void.," As the disk evolves, an annular substructure appears, and the central region becomes almost empty of particles, a void."
+ Subsequently. the big ring collapses and fragments into three spiral unbound arms.," Subsequently, the big ring collapses and fragments into three spiral unbound arms."
+ Al {his stage. there is no core.," At this stage, there is no core."
+ At /=0.66 Gvr. the arms merge themselves and become a bar-like structure surrounded bv two weak spiral arms.," At $t = 0.66$ Gyr, the arms merge themselves and become a bar-like structure surrounded by two weak spiral arms."
+ This bar-like structure is small when compared to the Newtonian case., This bar-like structure is small when compared to the Newtonian case.
+ This anomalous behavior is due to the fact that the particles are not in perfect initial equilibrinm with the Moffatian potential., This anomalous behavior is due to the fact that the particles are not in perfect initial equilibrium with the Moffatian potential.
+ llowever. it is important to bear in mind that we are testing the DM claim. which considers (hat the Moffatian potential can explain many observed rotation curves of galactic disks without dark matter halos.," However, it is important to bear in mind that we are testing the BM claim, which considers that the Moffatian potential can explain many observed rotation curves of galactic disks without dark matter halos."
+ We recall (hat our simulated initial disk reproduces the observational features of real galaxies: exponential-5pitzer disk. velocity dispersions. rotation curves and barvonic mass. and has secular equilibrium in (he Newtonian case.," We recall that our simulated initial disk reproduces the observational features of real galaxies: exponential-Spitzer disk, velocity dispersions, rotation curves and baryonic mass, and has secular equilibrium in the Newtonian case."
+ From (his point of view. we would expect (that. if DM hypotheses were reliable. our disk would be stable all the time in the simulation.," From this point of view, we would expect that, if BM hypotheses were reliable, our disk would be stable all the time in the simulation."
+ We obtain that an equilibrium configuration is reached only after /Z0.66 Gvr., We obtain that an equilibrium configuration is reached only after $t \gtrsim 0.66$ Gyr.
+ After this time. the svstem is similar. in some dynamical aspects. to the Newtonian case. but we will see below that this similarity is only apparent.," After this time, the system is similar, in some dynamical aspects, to the Newtonian case, but we will see below that this similarity is only apparent."
+ The results of model Ls simulations bring some interesting questions., The results of model II's simulations bring some interesting questions.
+ One could claim (hat the simulations of model II need a finetuning of the parameters Mo and ry to maintain the entire disk in equilibrium., One could claim that the simulations of model II need a finetuning of the parameters $\rm{M_0}$ and $\rm{r_0}$ to maintain the entire disk in equilibrium.
+ These parameters could. therefore. be changed to avoid alvpical rings ancl central voids during the first 0.33 Gyr.," These parameters could, therefore, be changed to avoid atypical rings and central voids during the first 0.33 Gyr."
+ To study these sub-morphology features. we have mace some tests changing (he parameters Mo and rp. bul we have not," To study these sub-morphology features, we have made some tests changing the parameters $M_0$ and $r_0$ , but we have not"
+The second mieasuremoent is made analysing an eutire xvsteni evolution.,The second measurement is made analysing an entire system evolution.
+ We start with the initial condition of 2 million particles with redshift Z= 50. At=0.001. 0=0.5. and particle mass about 1.655-1019 solar masses.," We start with the initial condition of 2 million particles with redshift $Z = 50$ , $\Delta t = 0.001$, $\theta = 0.8$, and particle mass about $1.655 \cdot 10^{10}$ solar masses."
+ The system evolution is carried out up to redshift Z=0., The system evolution is carried out up to redshift $Z = 0$.
+ The evolution is executed with the original BIT aleorithim aud with the WD99 code., The evolution is executed with the original BH algorithm and with the WD99 code.
+ As refercuce case. we adopt the critical level equal to 6.," As reference case, we adopt the critical level equal to 6."
+ We measure the absolute error e iu the position of particles aud in the velocities of particles in the mean square seus. aud The study of the &ual evolution is described in the next sub-section.," We measure the absolute error $\epsilon$ in the position of particles and in the velocities of particles in the mean square sense, and The study of the final evolution is described in the next sub-section."
+" Were we eive the measured value at the end of the simulation: ej,4=0.002 aud ερ0.01.", Here we give the measured value at the end of the simulation: $\epsilon_{pos} = 0.003$ and $\epsilon_{vel} = 0.01$.
+ Similar values are measured when runing simulations with more than 2 milliou particles aud with a critical level equal to 6., Similar values are measured when running simulations with more than 2 million particles and with a critical level equal to 6.
+ The obtained € values lead us to conchide that the WD99 procedure does not introduce significaut errors in comparison with the BIT algorithiu., The obtained $\epsilon$ values lead us to conclude that the WD99 procedure does not introduce significant errors in comparison with the BH algorithm.
+ The final stages obtained rumuing simulations with the DIT algoritlian and the WD99 code are very similar., The final stages obtained running simulations with the BH algorithm and the WD99 code are very similar.
+The two-point correlation function is defiued as,The two-point correlation function is defined as
+and optical are at very. low luminosity ancl it suggests they correlate from the LS to coll” state.,and optical are at very low luminosity and it suggests they correlate from the LS to `off' state.
+ More recently. Corbel ct al. (," More recently, Corbel et al. ("
+"1999) show that the N-ray and radio fluxes during the ""ol state and the LS are well correlated.",1999) show that the X-ray and radio fluxes during the `off' state and the LS are well correlated.
+ 3339.4 was detected at a level of 0.27£0.06 niv (at 8640 Mllz) with the Australia Telescope Compact Array (ATCA) on 1999 Xugust 17 (NLJD 51407) and at a comparable level two weeks later (Corbel et al., 339–4 was detected at a level of $0.27\pm0.06$ mJy (at 8640 MHz) with the Australia Telescope Compact Array (ATCA) on 1999 August 17 (MJD 51407) and at a comparable level two weeks later (Corbel et al.
+ 1999)., 1999).
+ Ehis is one of the lowest levels ever detected in the radio. but still considerably above the upper limits to the radio emission during the LIS (see Fender et al.," This is one of the lowest levels ever detected in the radio, but still considerably above the upper limits to the radio emission during the HS (see Fender et al."
+ 1999)., 1999).
+ It is very similar to 22023|338 which still showed ongoing radio emission after the source was olf (see Llan IHjellming 1992)., It is very similar to 2023+338 which still showed ongoing radio emission after the source was off (see Han Hjellming 1992).
+" Acting all the results from the reported. X-ray. coll, state ancl LS. and the behaviour at other wavelengths. it can be concluded that the ""ol state is indeed an extension of the LS. but at lower intensity."," Adding all the results from the reported X-ray `off' state and LS, and the behaviour at other wavelengths, it can be concluded that the `off' state is indeed an extension of the LS, but at lower intensity."
+ One might therefore expect rapid (< 1008) temporal variability in the “oll” state but we were unable to verily this due to the low counting statistics., One might therefore expect rapid $<$ 100s) temporal variability in the `off' state but we were unable to verify this due to the low counting statistics.
+ We note that Alénndez van der Ixlis (1997) derived a 30 upper limit on the variability in the 0.00210 Lz range of rms in their observations in the “oll” state., We note that Ménndez van der Klis (1997) derived a $\sigma$ upper limit on the variability in the 0.002–10 Hz range of rms in their observations in the `off' state.
+ Our observed luminosity in the 0.510 keV. band. is comparable to the BIISNT. €6822023|338 and 1163047 in quiescence (Asai et al., Our observed luminosity in the 0.5–10 keV band is comparable to the BHSXT 2023+338 and 1630--47 in quiescence (Asai et al.
+ 1998: Alenou ct al., 1998; Menou et al.
+ 1999: Parmar et al., 1999; Parmar et al.
+ 1997)., 1997).
+ The quiescent X-ray luminosities and X-rav/optical luminosity ratios of several DIISNEs are given in Table 3 for comparison., The quiescent X-ray luminosities and X-ray/optical luminosity ratios of several BHSXTs are given in Table 3 for comparison.
+ Note that the quiescent spectrum of A0G2000 can be fitted either by a power law or black-body model. presumably due to the narrow energy range ofR," Note that the quiescent spectrum of A0620–00 can be fitted either by a power law or black-body model, presumably due to the narrow energy range of."
+OSAT. MI the objects except 11630.47 in Table 3 have firm detections such that the spectra. can be determined., All the objects except 1630–47 in Table 3 have firm detections such that the spectra can be determined.
+ We note that DIISNTs can be fitted. with a power law spectrum in general. while neutron star SNTs. (NSSNTS) can be fitted by power law or black-hocky mocels (see Asai et al.," We note that BHSXTs can be fitted with a power law spectrum in general, while neutron star SXTs (NSSXTs) can be fitted by power law or black-body models (see Asai et al."
+ 1998)., 1998).
+ Recently. Rutlecege ct al. (," Recently, Rutledge et al. ("
+1999) fitted the spectra from NSSNTs with a hvdrogen atmosphere model and found that the derived. parameters (radius. and KE) of 0620)0 ancl 22023|338 were dillerent. from those found [or NSSNTs.,1999) fitted the spectra from NSSXTs with a hydrogen atmosphere model and found that the derived parameters (radius and kT) of A0620--00 and 2023+338 were different from those found for NSSXTs.
+ Although the results are based on data. Asai ct al. (," Although the results are based on data, Asai et al. ("
+1998) show similar findings by re-analvsing data.,1998) show similar findings by re-analysing data.
+ Ht suggests that the quiescent N-rav spectrum can provide additional information to distinguish. between black holes and neutron stars., It suggests that the quiescent X-ray spectrum can provide additional information to distinguish between black holes and neutron stars.
+ Significant. N-rav. variability in quiescence was observed in. 22023|338 (Wagner e al., Significant X-ray variability in quiescence was observed in 2023+338 (Wagner et al.
+ 1994). 1163047 (Parmar et al.," 1994), 1630–47 (Parmar et al."
+ 1997) and AOGZO00 Asai et al., 1997) and A0620–00 (Asai et al.
+ 1998: Moenou et al., 1998; Menou et al.
+ 1999). while we have οίaine similar result. for 33394 (i.e. by comparing with Asai et al.," 1999), while we have obtained a similar result for 339–4 (i.e. by comparing with Asai et al."
+" 1998 and other ""oll state observations).", 1998 and other `off' state observations).
+ This is a strong indication that the DIISNTs in quiescence are. no totally turned olf aad that the soll” state of 3339.4 is an extended: LS. as discussed above.," This is a strong indication that the BHSXTs in quiescence are not totally turned off and that the `off' state of 339–4 is an extended LS, as discussed above."
+ 33394 is similar to the quiescent state of DIISNEs. as will also be cliseussce below.," 339–4 is similar to the quiescent state of BHSXTs, as will also be discussed below."
+ We convert the optical magnitude into an optical (300-TOO nm) luminosity; of ⋅−24.10731 ergs 1 (assuming: ly=3.6J. and feey=1.2: Zdziarski et al., We convert the optical magnitude into an optical (300--700 nm) luminosity of $2.4\times10^{34}$ erg $^{-1}$ (assuming $A_V=3.6$ and $E_{B-V}=1.2$; Zdziarski et al.
+ 1998)., 1998).
+ The ratio of the soft X-ray (0.510 keV) and optical (300.700 nm) luminosities. Lyflow is 0.27. which is higher than other DIISNTS (sce Table 3).," The ratio of the soft X-ray (0.5–10 keV) and optical (300–700 nm) luminosities, $L_X/L_{opt}$ is $\sim 0.27$, which is higher than other BHSXTs (see Table 3)."
+ This could. be due to a somewhat higher X-ray luminositw for 33394 (see Table 3)., This could be due to a somewhat higher X-ray luminosity for 339–4 (see Table 3).
+ ALL these results resemble the quicscent state. spectrum. predicted x acdvection-clominatecl accretion How models (ADA: see varavan Yi 1995: Esin et al., All these results resemble the quiescent state spectrum predicted by advection-dominated accretion flow models (ADAF; see Narayan Yi 1995; Esin et al.
+ 1997)., 1997).
+ In the current ADA model for DIISNTs in quiescence. the viscously dissipated energv ds stored in the gas and advects into the black idle so as to account for the low luminosity o£. DIISNTS in quiescence.," In the current ADAF model for BHSXTs in quiescence, the viscously dissipated energy is stored in the gas and advects into the black hole so as to account for the low luminosity of BHSXTs in quiescence."
+ This model is in good agreement with the observed optical. UV. and X-ray emission of systems such as 22023|338 and 0620:00 (Naravan et al.," This model is in good agreement with the observed optical, UV and X-ray emission of systems such as 2023+338 and A0620–00 (Narayan et al."
+ 1996: Naravan et al., 1996; Narayan et al.
+ 1997a)., 1997a).
+ More recently. Fender et al. (," More recently, Fender et al. ("
+1999) and Wilms et al. (,1999) and Wilms et al. (
+1999) have applied ADAL mocels to explain. the observed. X-ray emission of 33394 in the LS.,1999) have applied ADAF models to explain the observed X-ray emission of 339–4 in the LS.
+ Narayan et al. (, Narayan et al. (
+1997b) point out that the ADAP luminosity cepends significantly on the accretion rate.,1997b) point out that the ADAF luminosity depends significantly on the accretion rate.
+ Therefore GX33394 maintains a certain accretion rate in the LS for most of the time. but the soll” state reported here represents a sudden decrease in the accretion rate.," Therefore 339–4 maintains a certain accretion rate in the LS for most of the time, but the `off' state reported here represents a sudden decrease in the accretion rate."
+" Lt therefore suggests that the source is not totally turned ο, as is observed."," It therefore suggests that the source is not totally turned off, as is observed."
+ Following Naravan et al. (, Following Narayan et al. (
+"1997). the mass accretion rate of 33394 in the coll” state corresponds to ~107 & st,","1997), the mass accretion rate of 339–4 in the `off' state corresponds to $\sim 10^{16}$ g $^{-1}$."
+ Lis also consistent with the AM~Lol101 es + predicted by binary-evolution models (Ixing ct al., It is also consistent with the $\dot M \sim 10^{15}-10^{16}$ g $^{-1}$ predicted by binary-evolution models (King et al.
+ 1996: Alenou et al., 1996; Menou et al.
+ 1999). assuming an orbital period of 14.8 h (Callanan ct al.," 1999), assuming an orbital period of 14.8 h (Callanan et al."
+ 1992) ancl black bole mass of 5 Al. (Zeliarski ct al., 1992) and black hole mass of 5 $M_{\odot}$ (Zdiarski et al.
+ 1998)., 1998).
+ This aceretion rate is very similar to that of 22023|338 (Naravan et al., This accretion rate is very similar to that of 2023+338 (Narayan et al.
+ 1997a) in quiescence., 1997a) in quiescence.
+ This relatively large accretion rate could be due to an evolved companion star like 22023|338 which supports the possibility of an evolved subgiant in 3339.4 (Callanan et al., This relatively large accretion rate could be due to an evolved companion star like 2023+338 which supports the possibility of an evolved subgiant in 339–4 (Callanan et al.
+ 1992)., 1992).
+" More detailed ADAE modelling may be needed in order to study the nature of the ""ol state of 3339.4 ancl make a direct comparison with DIISNXTSs in quiescence.", More detailed ADAF modelling may be needed in order to study the nature of the `off' state of 339–4 and make a direct comparison with BHSXTs in quiescence.
+ Note that apart from the above. the observed. range in V magnitude of 3339.4 is similar to the expected outburst. amplitude for an SNT with an orbital period. of 1458 hr (Shahbaz lIxuulkers 1908).," Note that apart from the above, the observed range in $V$ magnitude of 339–4 is similar to the expected outburst amplitude for an SXT with an orbital period of 14.8 hr (Shahbaz Kuulkers 1998)."
+ Together with the evidence for the similarity of the spectrum and luminosity. we sugeest the “oll state of 3339.4 to be equivalent to DIISNTs in. quiescence.," Together with the evidence for the similarity of the spectrum and luminosity, we suggest the `off' state of 339–4 to be equivalent to BHSXTs in quiescence."
+ The result. is important. since 33394 undergoes state transitions on much shorter timescales. while most of the BLIISN’PS exhibit. outbursts every 10.20 vears.," The result is important since 339–4 undergoes state transitions on much shorter timescales, while most of the BHSXTs exhibit outbursts every 10–20 years."
+" Finally. we note that we will need a σους radial velocity measurement of 33394 in its ""oll state when the contamination of the X-ray irracdated disc is at a minimum. so as to measure or constrain the black hole mass."," Finally, we note that we will need a good radial velocity measurement of 339–4 in its `off' state when the contamination of the X-ray irradiated disc is at a minimum, so as to measure or constrain the black hole mass."
+ We are grateful to the SDC science team for their assistance in the data reduction., We are grateful to the SDC science team for their assistance in the data reduction.
+ We thank Robert Rutledge ancl Lars Bilelsten for constructive commoents., We thank Robert Rutledge and Lars Bildsten for constructive comments.
+ EK thanks Rob Fender for stimulating discussions., EK thanks Rob Fender for stimulating discussions.
+ This paper utilizies, This paper utilizies
+in the Appendix.,in the Appendix.
+ Using the values of mp and my from SpanglerandSpitler(2004) in Equation (9). we have e=0.84.," Using the values of $m_{B \parallel}$ and $m_{B \perp}$ from \cite{Spangler04} in Equation (9), we have $\epsilon = 0.84$."
+ The anisotropy. of solar wind fluctuations had been consklered prior to (2004) bv Bavassanoetal(1982) and Wleinοἱal(1993)., The anisotropy of solar wind fluctuations had been considered prior to \cite{Spangler04} by \cite{Bavassano82} and \cite{Klein93}.
+. Davassano also studied magnetic field fIuctuations. in conditions of high speed solar wind at several heliocentric distances.," \cite{Bavassano82} also studied magnetic field fluctuations, in conditions of high speed solar wind at several heliocentric distances."
+ The results of Bavassanoetal(1982). (see data in their Figure 2) vield values lor e in the range 0.8 - 0.9 and ereater. ie. verv similar (o that quoted above.," The results of \cite{Bavassano82} (see data in their Figure 2) yield values for $\epsilon$ in the range 0.8 - 0.9 and greater, i.e. very similar to that quoted above."
+ It should also be noted that a highlv anisotropic anc Alfvénnic nature is a characteristic of turbulence in the inner solar svstem. (hat might not be valid (hroughout interplanetary space.," It should also be noted that a highly anisotropic and Alfvénnic nature is a characteristic of turbulence in the inner solar system, that might not be valid throughout interplanetary space."
+ Davassanoetal(1982) and kleinetal(1993) show thal properties of solar wind turbulence depend on heliocentric distance. specifically that the degree of anisotropy decreases wilh increasing heliocentric distance.," \cite{Bavassano82} and \cite{Klein93} show that properties of solar wind turbulence depend on heliocentric distance, specifically that the degree of anisotropy decreases with increasing heliocentric distance."
+ Ixleinetal(1993). further report Chat the anisolropy is less in slow speed (han high speed solar wind., \cite{Klein93} further report that the anisotropy is less in slow speed than high speed solar wind.
+ The upper limit we can place {ο anisotropy. of velocity [Inctuations in the Local Clouds is less than. though comparable to. that of solar wind turbulence at a heliocentric distance of 1 AU.," The upper limit we can place to anisotropy of velocity fluctuations in the Local Clouds is less than, though comparable to, that of solar wind turbulence at a heliocentric distance of 1 AU."
+ As discussed in Section 3. the analysis of line widths by Redfield)andLinskv(2004) also vields (he ion temperature 7.," As discussed in Section 3, the analysis of line widths by \cite{Redfield04} also yields the ion temperature $T$."
+ Strictly speaking. this is a line-ol-sight temperature: il is a measure of the line-ol-sight component of thermal motion of atoms and ions.," Strictly speaking, this is a line-of-sight temperature; it is a measure of the line-of-sight component of thermal motion of atoms and ions."
+" If the Local Cloud turbulence is similar to heliospheric turbulence. the ion temperature Z"" might also depend on the angle A between the Iime-of-sight and the local interstellar magnetic field. as a consequence of TσεTy."," If the Local Cloud turbulence is similar to heliospheric turbulence, the ion temperature $T$ might also depend on the angle $A$ between the line-of-sight and the local interstellar magnetic field, as a consequence of $T_{\perp} \neq T_{\parallel}$."
+" In regions of the solar wind in which ion cvclotron resonance heating appears to be active. and in particular in the solar corona. T>>T, (Cranmer 2002).. IxXasper"," In regions of the solar wind in which ion cyclotron resonance heating appears to be active, and in particular in the solar corona, $T_{\perp} \gg T_{\parallel}$ \citep{Cranmer02}."
+"etal(2009). have made an extensive study of the perpendicular-to-parallel temperature ratio At,=T for protons in the solar wind.", \cite{Kasper09} have made an extensive study of the perpendicular-to-parallel temperature ratio $R_p = \frac{T_{\perp}}{T_{\parallel}}$ for protons in the solar wind.
+" Although it is not the case that R,> Lin all cases. itis iue that 1 15 larger than it would be in an adiabaticallv-expanding solar wind."," Although it is not the case that $R_p > 1$ in all cases, it is true that $T_{\perp}$ is larger than it would be in an adiabatically-expanding solar wind."
+" Furthermore. INasperetal(2009) identify boundaries in a (/2,..9) plane. where dis the conventional plasma phlivsies quantity of the ratio of thermal (ο magnetic pressure."," Furthermore, \cite{Kasper09} identify boundaries in a $(R_p, \beta)$ plane, where $\beta$ is the conventional plasma physics quantity of the ratio of thermal to magnetic pressure."
+" These boundaries are clearly seen in the distribution of measured values of 722, and 2. and correspond to instability thresholds for generation of plasma waves."," These boundaries are clearly seen in the distribution of measured values of $R_p$ and $\beta$, and correspond to instability thresholds for generation of plasma waves."
+" Apparently these plasma waves heat the protons in a way which keeps (hem ""in bounds” in a subset of the (Z0.3) plane."," Apparently these plasma waves heat the protons in a way which keeps them “in bounds” in a subset of the $(R_p, \beta)$ plane."
+ The important point here is (hat in the corona ancl solar wind. the perpendicular and parallel temperatures are usually not equal. and 7>T where ion evclotron resonant effects are important.," The important point here is that in the corona and solar wind, the perpendicular and parallel temperatures are usually not equal, and $T_{\perp} > T_{\parallel}$ where ion cyclotron resonant effects are important."
+the compact source. Mejouu. from the luminosity.,"the compact source, $M_{\rm cloud}$, from the luminosity."
+ where ον=20kms7! based on the width of the narrow component in the CRIRES spectrum (FWHM=S50kms7!).," where $\sigma_{\rm v} = 20 \, {\rm km \, s^{-1}}$ based on the width of the narrow component in the CRIRES spectrum $\rm FWHM = 50 \, km \, s^{-1}$ )."
+ We find This mass estimate depends on the ratio e/fu..," We find This mass estimate depends on the ratio $\epsilon / f_{\rm
+H_2}$."
+ Assuming that this ratio is ~] as estimated for the CO complex. we find a mass slightly larger than the virial mass Mai;=1.3x107 dderived for the same velocity dispersion in Sect. 3.3...," Assuming that this ratio is $\sim 1$ as estimated for the CO complex, we find a mass slightly larger than the virial mass $M_{\rm
+vir} = 1.3 \times 10^7\,$ derived for the same velocity dispersion in Sect. \ref{sec:cs}."
+ We conclude from this comparison that the eemission from the compact source may be accounted for by the dissipation of turbulent kinetic energy if the source Is a gravitationally bound cloud with a mass in the range | to a few 10M.," We conclude from this comparison that the emission from the compact source may be accounted for by the dissipation of turbulent kinetic energy if the source is a gravitationally bound cloud with a mass in the range 1 to a few $10^7\,M_\odot $."
+ With this interpretation. the compact ssource Is a gravitationally bound cloud. massive enough to form a super star cluster.," With this interpretation, the compact source is a gravitationally bound cloud, massive enough to form a super star cluster."
+ Like what we have done in the previous section for the CO complex. we can use equation (4)) to estimate the mass accretion rate and thereby the cloud formation time scale.," Like what we have done in the previous section for the CO complex, we can use equation \ref{eqA}) ) to estimate the mass accretion rate and thereby the cloud formation time scale."
+ For Cu. We use the width of the broad component 160 kms!) in the line profile of the 1-0 H» S(1) measured with CRIRES.," For $\sigma_{\rm H_2}$, we use the width of the broad component $\rm FWHM \, = \, 160 \, km \,
+s^{-1}$ ) in the line profile of the $-$ 0 $_2$ S(1) measured with CRIRES."
+" We find: The ratio between Mí and M yields a cloud formation time scale by accretion of t,=2.9x10°(e/0.25) yr."," We find: The ratio between $M_{\rm cloud}$ and $\dot{M}_{\rm acc} $ yields a cloud formation time scale by accretion of $t_{\rm acc} = 2.9 \times
+10^6 \times (\epsilon/0.25) \,$ yr."
+ This value is about twice the cloud dynamical time Κον., This value is about twice the cloud dynamical time $R/\sigma_{\rm v}$.
+ We would thus be observing a pre-cluster cloud early in its evolution., We would thus be observing a pre-cluster cloud early in its evolution.
+ Even at this early stage it is remarkable that the fraction of the cloud mass that has been transformed into stars is very small (~0.3%. Sect.," Even at this early stage it is remarkable that the fraction of the cloud mass that has been transformed into stars is very small $\sim 0.3\%$, Sect."
+ 3.3 ).," \ref{sec:cs}
+ )."
+ We observe that the cloud has formed and become gravitationally bound without having formed massive stars., We observe that the cloud has formed and become gravitationally bound without having formed massive stars.
+ If our interpretation is correct. by observing the gas cooling through H2 lines. we may have discovered a massive cloud on its way to form a SSC within the next few Myr.," If our interpretation is correct, by observing the gas cooling through $_2$ lines, we may have discovered a massive cloud on its way to form a SSC within the next few Myr."
+ However. this conclusion is only tentative. and will remain incomplete. until we obtain the missing information about the mass. the density structure and the kinematics of the bulk of the gas.," However, this conclusion is only tentative, and will remain incomplete, until we obtain the missing information about the mass, the density structure and the kinematics of the bulk of the gas."
+ The H» line emission provides this information only for the fraction of warm. shock-heated gas.," The $_2$ line emission provides this information only for the fraction of warm, shock-heated gas."
+ The missing information can be obtained with ALMA., The missing information can be obtained with ALMA.
+ Our observations are related to two main questions in the field of star formation., Our observations are related to two main questions in the field of star formation.
+ How do super star clusters form?, How do super star clusters form?
+ What are the roles of turbulence and stellar feedback in regulating the star formation efficiency in galaxy mergers?, What are the roles of turbulence and stellar feedback in regulating the star formation efficiency in galaxy mergers?
+ We discuss how our observations and interpretation may be of general relevance for these questions., We discuss how our observations and interpretation may be of general relevance for these questions.
+ Star formation is the result of the hierarchical structure of the molecular interstellar medium established by turbulent compression and gravitational contraction (?).., Star formation is the result of the hierarchical structure of the molecular interstellar medium established by turbulent compression and gravitational contraction \citep{maclow04}.
+ The available gas fragments over a range of masses and time scales., The available gas fragments over a range of masses and time scales.
+ Massive clusters form at different times from the densest and most massive gas concentrations., Massive clusters form at different times from the densest and most massive gas concentrations.
+ For SGMC 2. this view is supported by the presence of two massive clusters (the near-IR SSC and the compact ionizec source). Which have formed before the compact ssource.," For SGMC 2, this view is supported by the presence of two massive clusters (the near-IR SSC and the compact ionized source), which have formed before the compact source."
+ More clusters. too small to be identified individually. are likely to have formed or to be in the process of being formed.," More clusters, too small to be identified individually, are likely to have formed or to be in the process of being formed."
+ Thus. the formation of super star clusters in SGMC 2 can be understood within the same framework as the formation of smaller clusters in disk galaxies like the Milky Way.," Thus, the formation of super star clusters in SGMC 2 can be understood within the same framework as the formation of smaller clusters in disk galaxies like the Milky Way."
+ Our work highlights three physical parameters which may contribute to account for the formation of exceptionally massive clusters in the overlap region of the Antennae., Our work highlights three physical parameters which may contribute to account for the formation of exceptionally massive clusters in the overlap region of the Antennae.
+ These may be of general relevance for the formation of super star clusters im other extragalactic objects: other interacting/merging galaxies. and also starbursts and irregular dwarf galaxies (22).. (," These may be of general relevance for the formation of super star clusters in other extragalactic objects: other interacting/merging galaxies, and also starbursts and irregular dwarf galaxies \citep{keto05,weidner10}. ("
+1) The presence of compressive motion on large scaleswhich collects the gas.,1) The presence of compressive motion on large scaleswhich collects the gas.
+ The formation of the SGMC 2 complex and that of the compact ssource is driven by the galaxy interaction., The formation of the SGMC 2 complex and that of the compact source is driven by the galaxy interaction.
+ The large scale dynamics triggers the gas compression necessary for their formation and their subsequent growth by gas aceretion. (, The large scale dynamics triggers the gas compression necessary for their formation and their subsequent growth by gas accretion. (
+2) As discussed 1n earlier studies (??).. shear is the second key parameter in the formation of massive clusters because it sets the maximum mass that the gas self-gravity can collect.,"2) As discussed in earlier studies \citep{escala08,weidner10}, shear is the second key parameter in the formation of massive clusters because it sets the maximum mass that the gas self-gravity can collect."
+ In the Antennae. the present geometry of the interaction. which temporarily makes the tidal forces compressive (?).. Is favorable to the formation of massive clusters.," In the Antennae, the present geometry of the interaction, which temporarily makes the tidal forces compressive \citep{renaud08}, is favorable to the formation of massive clusters."
+ As discussed by ? such conditions occur repeatedly in galaxy mergers. (, As discussed by \citet{renaud09} such conditions occur repeatedly in galaxy mergers. (
+3) High turbulence ts an additional key parameter.,3) High turbulence is an additional key parameter.
+ Since the Jeans mass is proportional to o (see Sect. ??)).," Since the Jeans mass is proportional to $\sigma^4 $ (see Sect. \ref{sec:gf}) ),"
+ a high value of the velocity dispersion c scales up the masses of the clouds formed by gravitational fragmentation., a high value of the velocity dispersion $\sigma$ scales up the masses of the clouds formed by gravitational fragmentation.
+ To form a super star cluster. it is not sufficient to bind a large mass of gas.," To form a super star cluster, it is not sufficient to bind a large mass of gas."
+ It is also necessary that the star formation efficiency becomes high when the gas mass is highly concentrated., It is also necessary that the star formation efficiency becomes high when the gas mass is highly concentrated.
+ It ts likely that high turbulence is a key factor which prevents star formation to be efficient before the cloud mass has been concentrated by gravity., It is likely that high turbulence is a key factor which prevents star formation to be efficient before the cloud mass has been concentrated by gravity.
+ This is ἃ prerequisite to form a dense. potentially bound. cluster rather than a loose OB association (?)..," This is a prerequisite to form a dense, potentially bound, cluster rather than a loose OB association \citep{elmegreen08}."
+ The impact of turbulence on the star formation efficiency is along standing topic of research (2).., The impact of turbulence on the star formation efficiency is a long standing topic of research \citep{maclow04}.
+ ? and ? used numerical simulations to quantify the dependence of the star-formation rate per free-fall time on the Mach number of turbulence., \citet{krumholz05} and \citet{padoan11} used numerical simulations to quantify the dependence of the star-formation rate per free-fall time on the Mach number of turbulence.
+ Both agree that. independent of the mean cloud density. for à cloud near virial equilibrium. the fraction of the gas mass that is unstable to collapse is small.," Both agree that, independent of the mean cloud density, for a cloud near virial equilibrium, the fraction of the gas mass that is unstable to collapse is small."
+ If this is correct. then," If this is correct, then"
+the mereed 7 aud. 131uu data. respectively.,"the merged 7 and 13mm data, respectively."
+ The cross-correlation analysis use the ZDC fuuctiou, The cross-correlation analysis use the ZDC function
+of non fragmenting self-gravitating accretion dises are going to be significantly affected by poor resolution. in the sense expressed by Eq. 10).,"of non fragmenting self-gravitating accretion discs are going to be significantly affected by poor resolution, in the sense expressed by Eq. \ref{eq:fit}) )."
+ Probably the most important property. apart from fragmentation. is related to the stress induced by the instabiliy. and in particular its magnitude and tqe locality of the associated dissipation (which are in turn related ο the power spectrum of the modes excited by the instability).," Probably the most important property, apart from fragmentation, is related to the stress induced by the instability, and in particular its magnitude and the locality of the associated dissipation (which are in turn related to the power spectrum of the modes excited by the instability)."
+ As mentioned in the Introduction. several studies {90909) have considered the issue of locality of the transport induced by the instability.," As mentioned in the Introduction, several studies \citep{LR04,LR05,CLC09} have considered the issue of locality of the transport induced by the instability."
+" Here the question is whether the typical wavelength ""the disturbances is a significant fraction of the radius /7.", Here the question is whether the typical wavelength of the disturbances is a significant fraction of the radius $R$.
+" The above studies (and in particular 2) have shown that the spectrum ""the disturbances peaks at a wavelength Azz// and becomes negligible for A3/1.", The above studies (and in particular \citealt{CLC09}) ) have shown that the spectrum of the disturbances peaks at a wavelength $\lambda \approx H$ and becomes negligible for $\lambda\gtrsim 3H$.
+ In this respect. thus. the relevant resolution requirement is the generally less restrictive condition ///1 ready discussed in 2..," In this respect, thus, the relevant resolution requirement is the generally less restrictive condition $h/H\lesssim 1$ already discussed in \citet{nelson06}."
+ Furthermore. not resolving the smallest Pdcales would only preven the disc from developing small scale P2ructures. which would not οςontribute to global transport anyway.," Furthermore, not resolving the smallest scales would only prevent the disc from developing small scale structures, which would not contribute to global transport anyway."
+ Let us now discuss whetyer the evaluation of the stress from such simulations is affecec by Eq. (10). Clearly.di, Let us now discuss whether the evaluation of the stress from such simulations is affected by Eq. \ref{eq:fit}) ).
+sturbance.. regardless of the way it has been excited. a correct evaluation of the resulting stress only requires that the typical wavelengths of the disturbance are resolved.," Clearly, regardless of the way it has been excited, a correct evaluation of the resulting stress only requires that the typical wavelengths of the disturbance are resolved."
+ However. here the cuestion is whether the dise response to an external boundary condition (in this ease. to an externally imposed cooling rate) depends on resolution.," However, here the question is whether the disc response to an external boundary condition (in this case, to an externally imposed cooling rate) depends on resolution."
+ The condition of thermal equilibrium relates in a unique way the external cooling to the magnitude of the stress induced by the instability (eq. C143).," The condition of thermal equilibrium relates in a unique way the external cooling to the magnitude of the stress induced by the instability (eq. \ref{eq:alpha}) )),"
+ under the hypothesis that ye only heating source in the dise is provided by the instability self and is dissipated locally., under the hypothesis that the only heating source in the disc is provided by the instability itself and is dissipated locally.
+ Here. we have shown that the new empirically determined resolution condition corresponds to qe case Where artificial viscosity only provides 5 per cent of the jeating necessary for thermal balance.," Here, we have shown that the new empirically determined resolution condition corresponds to the case where artificial viscosity only provides 5 per cent of the heating necessary for thermal balance."
+ Therefore. unless the effects of artificial viscosity are much larger than expected. one would not expect a change in the stress above the 5 per cent level.," Therefore, unless the effects of artificial viscosity are much larger than expected, one would not expect a change in the stress above the 5 per cent level."
+ In his respect. the effect of poor resolution is just to suppress the development of small scale disturbances. and thus reduce the peak value of the density perturbation. while keeping the overall average value of the stress essentially unchanged.," In this respect, the effect of poor resolution is just to suppress the development of small scale disturbances, and thus reduce the peak value of the density perturbation, while keeping the overall average value of the stress essentially unchanged."
+ Finally. it is worth notingyat the discussion above is focussed on SPH simulations.," Finally, it is worth noting that the discussion above is focussed on SPH simulations."
+ Many t»revious results (999y. including several determinations of the fragmentaion boundary (22??2222). have been obtained using grid-based codes.," Many previous results \citep{gammie01,mejia05,boley06}, including several determinations of the fragmentation boundary \citep{gammie01,mejia05,boss04,boss06}, have been obtained using grid-based codes."
+ Artificial viscosity Is also present to various degrees is grid-based codes. and one might thus conclude that an analogous resoluion requirement should be formulated also in thus ease.," Artificial viscosity is also present to various degrees is grid-based codes, and one might thus conclude that an analogous resolution requirement should be formulated also in thus case."
+" Whether also grid-based simulations do not converge in the determination o ""the critical time scale for fragmentation and the exact form of the resolution requirement in this case are questions that still need to be worked out.", Whether also grid-based simulations do not converge in the determination of the critical time scale for fragmentation and the exact form of the resolution requirement in this case are questions that still need to be worked out.
+ We thank Phil Armitage. Matthew Bate. Farzana Meru and Ken Rice for interesting discussions.," We thank Phil Armitage, Matthew Bate, Farzana Meru and Ken Rice for interesting discussions."
+ We also thank Charles Gammie. for an insightful referee's report.," We also thank Charles Gammie, for an insightful referee's report."
+Negueruela&Marco(2006).. while the EFOSC observations are unpublished: programs [IDs 079.D-0371(A) and 083.D-Ο11Ο(Α).,"\citet{negueruela06}, while the EFOSC observations are unpublished: programs IDs 079.D-0371(A) and 083.D-0110(A)."
+ Data reduction was performed using the ΠΑΕ packages. following the standard procedure.," Data reduction was performed using the IRAF packages, following the standard procedure."
+ We analysed data obtained by all three instruments aboard (Bradtetal..1993): the All sky Monitor (ASM) data consist of daily flux averages in the energy range 1.3-12.] keV. The Proportional Counter Array (PCA) covers the lower part of the energy range 2-60 keV. and consists of five identical coaligned gas-filled proportional units giving a total collecting area of 6500 οπή” and provides an energy resolution of at 6 keV. The High Energy Timing Experiment (HEXTE) is constituted by 2 clusters of 4 Nal/CsI scintillation counters. with a total collecting area of 2 x 800 cm’. sensitive in the 15-250 keV band with à nominal energy resolution of at 60 keV. Data taken during satellite slews. passage through the South Atlantic Anomaly and Earth occultation were removed.," We analysed data obtained by all three instruments aboard \citep{bradt93}: the All sky Monitor (ASM) data consist of daily flux averages in the energy range 1.3-12.1 keV. The Proportional Counter Array (PCA) covers the lower part of the energy range 2–60 keV, and consists of five identical coaligned gas-filled proportional units giving a total collecting area of 6500 $^{-2}$ and provides an energy resolution of at 6 keV. The High Energy Timing Experiment (HEXTE) is constituted by 2 clusters of 4 NaI/CsI scintillation counters, with a total collecting area of 2 $\times$ 800 $^2$, sensitive in the 15–250 keV band with a nominal energy resolution of at 60 keV. Data taken during satellite slews, passage through the South Atlantic Anomaly and Earth occultation were removed."
+ The overall on-source time analysed amounts to 496.3 ks., The overall on-source time analysed amounts to 496.3 ks.
+ The mission-specific packages of Heasarc FTOOLS (version 6.6.3) were employed to perform data reduction. while XSPEC was used for spectral analysis.," The mission-specific packages of Heasarc FTOOLS (version 6.6.3) were employed to perform data reduction, while XSPEC was used for spectral analysis."
+ For each observation we obtained an average energy spectrum and a power spectrum., For each observation we obtained an average energy spectrum and a power spectrum.
+" The energy spectra were generated from standard mode data. namely. the PCA of PCU2 and HEXTE “PSS8"" cluster B. data. which have a time resolution of 16 s and cover the 2-60 keV range with 129 channels and the 15-250 keV with 64 channels. respectively."," The energy spectra were generated from standard mode data, namely, the PCA of PCU2 and HEXTE ""FS58"" cluster B data, which have a time resolution of 16 s and cover the 2-60 keV range with 129 channels and the 15–250 keV with 64 channels, respectively."
+ All spectra were background-subtracted., All spectra were background-subtracted.
+ A systematic error of was added in quadrature to the PCA spectra to account for systematic errors., A systematic error of was added in quadrature to the PCA spectra to account for systematic errors.
+ The power spectra were generated by computing Fourier transforms of 128-s segments of the 2~°-s binned light curves in the energy range 2-15 keV (PCA channels 0-35) and averaging them together., The power spectra were generated by computing Fourier transforms of 128-s segments of the $2^{-6}$ -s binned light curves in the energy range 2–15 keV (PCA channels 0–35) and averaging them together.
+ The final power spectra were logarithmically rebinned in frequency and corrected for dead time effects according to the prescription given in. Nowaketal.(1999)., The final power spectra were logarithmically rebinned in frequency and corrected for dead time effects according to the prescription given in \citet{nowak99}.
+. Power spectra were normalized such that the integral over a given frequency range equals to the squared fractional rats amplitude (Belloni&Hasinger.1990;Miyamotoetal..1991).," Power spectra were normalized such that the integral over a given frequency range equals to the squared fractional $rms$ amplitude \citep{belloni90,miyamoto91}."
+. To fit the energy spectra. we used a model composed by a combination of photoelectric absorption. a power law with high-energy exponential cutoff. a Gaussian line profile at 6.5 keV and an absorption edge at around 9 keV to account for Fe K fluorescence.," To fit the energy spectra, we used a model composed by a combination of photoelectric absorption, a power law with high-energy exponential cutoff, a Gaussian line profile at 6.5 keV and an absorption edge at around 9 keV to account for Fe K fluorescence."
+ We did not find evidence for a cyclotron absorption feature. commonly seen in other acereting X-ray pulsars.," We did not find evidence for a cyclotron absorption feature, commonly seen in other accreting X-ray pulsars."
+ The power spectra were fitted using Lorentzian profiles only., The power spectra were fitted using Lorentzian profiles only.
+ The main peak of the X-ray pulsations and two to three of its harmonics are clearly seen in the power spectra., The main peak of the X-ray pulsations and two to three of its harmonics are clearly seen in the power spectra.
+ These peaks were fitted with narrow Lorentzians (the width was normally fixed at ~0.001 Hz)., These peaks were fitted with narrow Lorentzians (the width was normally fixed at $\sim 0.001$ Hz).
+ The broad-band noise was fitted with two zero-centred Lorentzians., The broad-band noise was fitted with two zero-centred Lorentzians.
+ The brighter observations required one extra Lorentzian to account for the noise above ~7 Hz., The brighter observations required one extra Lorentzian to account for the noise above $\sim 7$ Hz.
+ The information on the optical counterpart to Hs very scarce., The information on the optical counterpart to is very scarce.
+ To the authors’ knowledge no optical photometric observations are published., To the authors' knowledge no optical photometric observations are published.
+" The only dedicated observations are those reported in the Astronomer’s Telegrams ATel#7713 (Reaetal.2006) who obtained infrarecphotometric magnitudes and ATel#7739 (Negueruela&Marco. 2006).. who performed spectroscopic observations and reported the presence of H« strongly in emission,"," The only dedicated observations are those reported in the Astronomer's Telegrams 713 \citep{rea06} who obtained infraredphotometric magnitudes and 739 \citep{negueruela06}, who performed spectroscopic observations and reported the presence of $\alpha$ strongly in emission."
+ Information on photometric magnitudes and colours cai be found in astronomical catalogues., Information on photometric magnitudes and colours can be found in astronomical catalogues.
+ We searched for possible optical/infrared counterparts. in. the catalogue around the best-fit X-ray position (Campanaetal.. 2006).., We searched for possible optical/infrared counterparts in the catalogue around the best-fit X-ray position \citep{campana06}. .
+ The optical counterpart to Πδ the star named as OMADI-0380-07035569, The optical counterpart to is the star named as NOMAD1-0380-0705569
+colour space. possibly making the optical morphologica preselection of quasar candidates obsolete.,"colour space, possibly making the optical morphological preselection of quasar candidates obsolete."
+ Alodels (e.g. Granato&Danese 1994: Nenkovaetal. 2002)) and observations (e.g. Elvisetal. 1904)) suggest tha quasar It spectra are more or less flat. (in 4L) from ~ down to at least 25pum.. even in the absence of starburs activity.," Models (e.g. \citealt{granato94}; ; \citealt{nenkova02}) ) and observations (e.g. \citealt{elvis94}) ) suggest that quasar IR spectra are more or less flat (in $\nu L_{\nu}$ ) from $\sim 1$ down to at least 25, even in the absence of starburst activity."
+ Therefore. the deeper and better quality SWIRL 5 and 24 observations will allow the detection of high numbers of quasars and the easy. adaptation of the tools oesented here.," Therefore, the deeper and better quality SWIRE 8 and 24 observations will allow the detection of high numbers of quasars and the easy adaptation of the tools presented here."
+ The stellar contamination will probably be ugher at S than at 15 but this problem will most ikelv be solved due to the multitude of LR bands. that will allow an easier separation of the galactic and extragalactic »opulations.," The stellar contamination will probably be higher at 8 than at 15 but this problem will most likely be solved due to the multitude of IR bands, that will allow an easier separation of the galactic and extragalactic populations."
+ As a last remark. we would like to stress that a robust candidate selection. technique and. subsequent photometric redshift estimates such as the ones. presented rere will be increasingly requirecl by all future largearea," As a last remark, we would like to stress that a robust candidate selection technique and subsequent photometric redshift estimates such as the ones presented here will be increasingly required by all future largearea"
+Furthermore. RX J0852.0—4622 is. like RX J1713.7—3946. an extended SNR. at X-ray energies.,"Furthermore, RX $-$ 4622 is, like RX $-$ 3946, an extended SNR at X-ray energies."
+ The relatively large angular size enables us to study the detailed morphology even with the existing angular resolution., The relatively large angular size enables us to study the detailed morphology even with the existing angular resolution.
+ We have observed RX J0352.0—4622 with the mun diameter CANGAROO-II Imaging Aimospherie Cherenkov Telescope (ACT)., We have observed RX $-$ 4622 with the m diameter CANGAROO-II Imaging Atmospheric Cherenkov Telescope (IACT).
+ The telescope (IXawachietal.2001). detects eanmnma-ravs above several hundred GeV by reconstructing optical Cherenkov images generated by relativistic secondary particles in the cascades produced when high-energv. gamma-rays (and background. cosmic rays) interact with the Earth's upper atiosphere.," The telescope \citep{kawachi01}
+ detects gamma-rays above several hundred GeV by reconstructing optical Cherenkov images generated by relativistic secondary particles in the cascades produced when high-energy gamma-rays (and background cosmic rays) interact with the Earth's upper atmosphere."
+ Parametrizing the optical images enables the direction and energv of gamma-ray events to be inlerred., Parametrizing the optical images enables the direction and energy of gamma-ray events to be inferred.
+ The telescope is located near Woomera. South Australia (136 47' EE. 31706 8). [rom which RX J0352.0—4622 culminates at ~15° from the zenith. enabling observations with a lower energv (threshold.," The telescope is located near Woomera, South Australia $^\circ$ $'$ E, $^\circ$ $'$ S), from which RX $-$ 4622 culminates at $\sim$ $^\circ$ from the zenith, enabling observations with a lower energy threshold."
+ The angular resolution of CANGAROO-II was estimated to be 0.397 ppc at a distance of O.5kkpc) with an enerev threshold of GGeV depending on the observing conditions and the spectrum., The angular resolution of CANGAROO-II was estimated to be $^\circ$ pc at a distance of kpc) with an energy threshold of GeV depending on the observing conditions and the spectrum.
+ The observations were carried out between 2001 December 12 and 2002 February 15 (19 niehts). and between 2003 January 5 and February 28 (25 nights).," The observations were carried out between 2001 December 12 and 2002 February 15 (19 nights), and between 2003 January 5 and February 28 (25 nights)."
+ The telescope tracked the peak of the X-ray enission. in (he north-west rim of the SNR (a=132.25. ὃ=—45.65°. J2000 coordinates).," The telescope tracked the peak of the X-ray emission, in the north-west rim of the SNR $\alpha =132.25^\circ$ , $\delta =-45.65^\circ$, J2000 coordinates)."
+ The field of view (FOV) of the camera is 2.76°*2.76°., The field of view (FOV) of the camera is $^\circ \times$ $^\circ$.
+" The brightest star in the FOV. SAO 220422. has a visual magnitude of 4.1. and this region is brighter than our typical on-source ΕΟΝ,"," The brightest star in the FOV, SAO 220422, has a visual magnitude of 4.1, and this region is brighter than our typical on-source FOV."
+ As a result. we used a higher (trigger threshold. requiring 5 triggered pixels rather than the usual 3 (Ποetal.2003).," As a result, we used a higher trigger threshold, requiring 5 triggered pixels rather than the usual 3 \citep{cito03}."
+. The recently installed lishts several kilometers from the telescope were expected to have little effect on these observations as the telescope was pointed close to the zenith (Itohetal.2003).. and (his was confirmed from the azimuthal angle dependence of the shower rate (see 33).," The recently installed lights several kilometers from the telescope were expected to have little effect on these observations as the telescope was pointed close to the zenith \citep{cito03}, and this was confirmed from the azimuthal angle dependence of the shower rate (see 3)."
+ Each night was divided into (wo or three periods. iie. ONOFF. OFFONOFF. or OFFON observations.," Each night was divided into two or three periods, i.e., ON–OFF, OFF–ON–OFF, or OFF–ON observations."
+ ON-source observations were timed (o contain the mericlian passage ol the target. as was done by Enomotoοἱal.(2002b).," ON-source observations were timed to contain the meridian passage of the target, as was done by \cite{enomoto02b}."
+. In total. 5900 oof ON- and 5300 oof OFF-source data were obtained.," In total, 5900 of ON- and 5300 of OFF-source data were obtained."
+ First. cleaning cuts on camera images wereapplied. requiring (0.115-square) pixel," First, `cleaning' cuts on camera images wereapplied, requiring $^\circ$ -square) pixel"
+The observations were obtained with the during 2008 January 17.,The observations were obtained with the during 2008 January 17.
+" The SMA was in its very extended configuration, which included 28 independent baselines ranging in projected length from 30 to 385 Kl."," The SMA was in its very extended configuration, which included 28 independent baselines ranging in projected length from 30 to 385 $\lambda$."
+" The phase reference center of the observations was R.A. = 19h23m40.05s, decl= 1431'05.0"" (J2000.0)."," The phase reference center of the observations was R.A. = 19h23m40.05s, decl.= $^\circ$ $'$ $''$ (J2000.0)."
+" The frequency was centered at 217.1049 GHz in the Lower Sideband (LSB), while the Upper Sideband (USB) was centered at 228.1049 GHz."," The frequency was centered at 217.1049 GHz in the Lower Sideband (LSB), while the Upper Sideband (USB) was centered at 228.1049 GHz."
+ The primary beam of the SMA at 230 GHz has a FWHM of about 60”., The primary beam of the SMA at 230 GHz has a FWHM of about $''$ .
+" We detected the lines C3H;SOH(3743 59-369 54) and HCOCH5OH(62,3 49-6215 30) in the LSB."," We detected the lines $_2$ $_5$ $_{8,29}$ $_{9,28}$ ) and $_2$ $_{13,49}$ $_{12,50}$ ) in the LSB."
+ See Table | for their rest frequencies and rotational temperatures above of the ground state., See Table 1 for their rest frequencies and rotational temperatures above of the ground state.
+ The full bandwidth of the SMA correlator was 4 GHz (2 GHz in each band)., The full bandwidth of the SMA correlator was 4 GHz (2 GHz in each band).
+" The SMA digital correlator was configured in 24 spectral windows (“chunks”) of 104 MHz each, with 256 channels distributed over each spectral window, providing a resolution of 0.40 MHz (0.56 km s!) per channel."," The SMA digital correlator was configured in 24 spectral windows (“chunks”) of 104 MHz each, with 256 channels distributed over each spectral window, providing a resolution of 0.40 MHz (0.56 km $^{-1}$ ) per channel."
+" However, in this study we smoothed our spectral resolution to about 1 km s."," However, in this study we smoothed our spectral resolution to about 1 km $^{-1}$."
+" The zenith opacily (7230652. measured with the NRAO upping radiometer located at the Caltech Submillimeter Observatory, was ~ 0.15, indicating good weather conditions during the observations."," The zenith opacity $\tau_{230 GHz}$ ), measured with the NRAO tipping radiometer located at the Caltech Submillimeter Observatory, was $\sim$ 0.15, indicating good weather conditions during the observations."
+ Observations of Uranus provided the absolute scale for the flux density calibration., Observations of Uranus provided the absolute scale for the flux density calibration.
+ Phase and amplitude calibrators were the quasars 19254211 and 20354109., Phase and amplitude calibrators were the quasars 1925+211 and 2035+109.
+" The uncertainty in the flux scale is esumated to be 15—20*56, based on the SMA monitoring of quasars."," The uncertainty in the flux scale is estimated to be $\%$, based on the SMA monitoring of quasars."
+ Observations of Uranus provided the absolute scale for the flux density calibration., Observations of Uranus provided the absolute scale for the flux density calibration.
+ Further technical descriptions of the SMA and its calibration schemes can found in ?.., Further technical descriptions of the SMA and its calibration schemes can found in \citet{Hoetal2004}. .
+" The data were calibrated using the IDL superset MIR, originally developed for the Owens Valley Radio Observatory (?) and adapted for the The calibrated data were imaged and analyzed in the standard manner using the MIRIAD and AIPS packages."," The data were calibrated using the IDL superset MIR, originally developed for the Owens Valley Radio Observatory \citep{Scovilleetal1993} and adapted for the The calibrated data were imaged and analyzed in the standard manner using the MIRIAD and AIPS packages."
+ We used the ROBUST parameter of the INVERT task of MIRIAD set to 2 to obtain a slightly better sensiüvity while losing some angular resolution., We used the ROBUST parameter of the INVERT task of MIRIAD set to 2 to obtain a slightly better sensitivity while losing some angular resolution.
+ The resulting image rms noise of line images was around 30 mJy beam”! foreach velocity channel (with a smoothed size of 1 km s!) at an angular resolution of 0757 x 0/42 with a BA. = 57.67., The resulting image rms noise of line images was around 30 mJy $^{-1}$ foreach velocity channel (with a smoothed size of 1 km $^{-1}$ ) at an angular resolution of $0\rlap.{''}57$ $\times$ $0\rlap.{''}42$ with a P.A. = $57.6^\circ$.
+ The observations were obtained with the SMA on 2008 July 13., The observations were obtained with the SMA on 2008 July 13.
+ At the time of these observations the SMA had seven antennas in its extended configuration with baselines ranging in projected length from 30 to 255 kt., At the time of these observations the SMA had seven antennas in its extended configuration with baselines ranging in projected length from 30 to 255 $\lambda$.
+" Theprimary beam of the SMA at 340 GHz has à FWHM of 37"".", Theprimary beam of the SMA at 340 GHz has a FWHM of $''$ .
+ The molecular emission from, The molecular emission from
+ 2E).,(Figure 2E).
+" This implies that if dissipation was to be (Figureincreased, no apsidally anti-aligned attractor would exist."," This implies that if dissipation was to be increased, no apsidally anti-aligned attractor would exist."
+" If the eccentricity of the perturber is enhanced beyond e,>0.12, the anti-aligned fixed point disappears completely from the portrait (Figure "," If the eccentricity of the perturber is enhanced beyond $e_p > 0.12$, the anti-aligned fixed point disappears completely from the portrait (Figure 2D)."
+"In both “end-member” scenarios we 2D).considered (ep=0 and ej= 0.2), in presence of dissipation, only a single attractor would exist."," In both “end-member"" scenarios we considered $e_p = 0$ and $e_p = 0.2$ ), in presence of dissipation, only a single attractor would exist."
+" However, the attractors in these two cases arise from fixed points (i.e. one can not be transformed into another by a change in ej), centered around different branches of the separatrix."," However, the attractors in these two cases arise from fixed points (i.e. one can not be transformed into another by a change in $e_p$ ), centered around different branches of the separatrix."
+" Recall that the sole fixed point that was present in the €y=0 case, disappeared, when e,=0.12."," Recall that the sole fixed point that was present in the $e_p = 0$ case, disappeared, when $e_p = 0.12$."
+" Similarly, The fixed point that is present in the ej=0.2 portrait is not present when the perturber’s orbit is circular."," Similarly, The fixed point that is present in the $e_p = 0.2$ portrait is not present when the perturber's orbit is circular."
+ This has important implications for the motion of the particle when eccentricity of the perturber is not maintained at a constant value., This has important implications for the motion of the particle when eccentricity of the perturber is not maintained at a constant value.
+" Consider a scenario where no dissipation is applied, but the eccentricity of the perturber is varied adiabatically between e,=ϐ and ej,=02."," Consider a scenario where no dissipation is applied, but the eccentricity of the perturber is varied adiabatically between $e_p = 0$ and $e_p = 0.2$."
+" Here, “adiabatically” means that the oscillation period of the perturber's eccentricity greatly exceeds the apsidal circulation/libration period of the test-particle."," Here, “adiabatically"" means that the oscillation period of the perturber's eccentricity greatly exceeds the apsidal circulation/libration period of the test-particle."
+" Regardless of the particle's starting condition, its orbit will eventually encounter the separatrix."," Regardless of the particle's starting condition, its orbit will eventually encounter the separatrix."
+" Since the separatrix is an orbit with an infinite preriod, its crossing necessarily leads to chaotic motion (BruhwilerandCary 1989).."," Since the separatrix is an orbit with an infinite preriod, its crossing necessarily leads to chaotic motion \citep{1989PhyD...40..265B}."
+ In fact the situation is analogous to the motion of an amplitude-modulated pendulum., In fact the situation is analogous to the motion of an amplitude-modulated pendulum.
+ It has been shown that the chaotic region of such a system occupies the phase-space area that is swept by the separatrix (HenrardandHenrard1991;Morbidelli 1993).," It has been shown that the chaotic region of such a system occupies the phase-space area that is swept by the separatrix \citep{1991PhyD...54..135H, 1993PhyD...68..187H}."
+". As a result, by ensuring that the eccentricity of the perturber reaches zero and extends above e,=0.12 at every oscillation, we enforce the entire phase-space within 0.6 to be swept by the critical curve, causing all particles within this e-limit to become chaotic."," As a result, by ensuring that the eccentricity of the perturber reaches zero and extends above $e_p = 0.12$ at every oscillation, we enforce the entire phase-space within $e \lesssim 0.6$ to be swept by the critical curve, causing all test-particles within this $e$ -limit to become chaotic."
+ Large variations in the perturber's orbital eccentricity can be induced by secular interactions with a distant pair of planets., Large variations in the perturber's orbital eccentricity can be induced by secular interactions with a distant pair of planets.
+" Consider placing the perturber described above at a—1.85AU, initially on a circular orbit, in presence of Jupiter and Saturn, whose initial conditions correspond to their actual orbits in 1983 (see Ch."," Consider placing the perturber described above at $a = 1.85$ AU, initially on a circular orbit, in presence of Jupiter and Saturn, whose initial conditions correspond to their actual orbits in 1983 (see Ch."
+ 7 of MurrayandDermott(1999)))., 7 of \cite{1999ssd..book.....M}) ).
+" The orbital evolution of the massive system can be computed using Laplace-Lagrange secular theory, and the resulting solution is presented in Figure 3."," The orbital evolution of the massive system can be computed using Laplace-Lagrange secular theory, and the resulting solution is presented in Figure 3."
+" Note that this solution is approximate at large eccentricity, since high order terms are neglected."," Note that this solution is approximate at large eccentricity, since high order terms are neglected."
+" Due to the perturber's proximity to the ve secular resonance, its orbital eccentricity undergoes excursions between ej,=(0 and e,®&0.2 on aT &2m/(uy—Us)1 Myr time-scale, where u's are the corresponding eigen-frequencies of the perturbation matrix and s refers to Saturn."," Due to the perturber's proximity to the $\nu_6$ secular resonance, its orbital eccentricity undergoes excursions between $e_p = 0$ and $e_p \approx 0.2$ on a $\tau \approx 2 \pi / (u_p - u_s) \approx 1$ Myr time-scale, where $u$ 's are the corresponding eigen-frequencies of the perturbation matrix and $s$ refers to Saturn."
+" Furthermore, the perturber’s longitude of perihelion in this solution precesses at a nearly constant rate of g=23” /year."," Furthermore, the perturber's longitude of perihelion in this solution precesses at a nearly constant rate of $g = 23""$ /year."
+" For our purposes we approximate the variation in the perturber’s eccentricity as Addition of Jupiter and Saturn to the system does not significantly modify the evolution of the test-particle because of the substantial orbital separation between them (a;~0.2 and a,-0.1).", For our purposes we approximate the variation in the perturber's eccentricity as Addition of Jupiter and Saturn to the system does not significantly modify the evolution of the test-particle because of the substantial orbital separation between them $\alpha_j \sim 0.2$ and $\alpha_s \sim 0.1$ ).
+" In fact, evaluation of the corresponding constants 7, 9 and * shows that the particle's interactions with Saturn can be neglected all together, as they only contribute at the ~1% level, while for Jupiter, it suffices to account only for the additional apsidal precession, to which it contributes at the ~30% level, compared to the effect of the considered planet (15mg)."," In fact, evaluation of the corresponding constants $\eta$, $\beta$ and $\gamma$ shows that the particle's interactions with Saturn can be neglected all together, as they only contribute at the $\sim 1\%$ level, while for Jupiter, it suffices to account only for the additional apsidal precession, to which it contributes at the $\sim 30\%$ level, compared to the effect of the considered planet $15 m_{\oplus}$ )."
+" Quantitatively, this corresponds to an enhancement of the coefficients η and 8, but not y."," Quantitatively, this corresponds to an enhancement of the coefficients $\eta$ and $\beta$, but not $\gamma$."
+" As stated in the Appendix, where the expressions for the constants are given, we have been implicitly retaining the apsidal contribution due to Jupiter since the beginning of the paper, for consistency of the phase-space portraits."," As stated in the Appendix, where the expressions for the constants are given, we have been implicitly retaining the apsidal contribution due to Jupiter since the beginning of the paper, for consistency of the phase-space portraits."
+" The difference in longitude of perihelia between the test-particle and Jupiter forms a comparatively fast angle, and thus can be averaged out."," The difference in longitude of perihelia between the test-particle and Jupiter forms a comparatively fast angle, and thus can be averaged out."
+" Consequently, we avoid its introduction into the Hamiltonian."," Consequently, we avoid its introduction into the Hamiltonian."
+" This is further warranted, as the magnitude of the interaction term between Jupiter and the test-particle (i.e. y;) is about an order of magnitude smaller than that of the test-particle and the m=perturber!.."," This is further warranted, as the magnitude of the interaction term between Jupiter and the test-particle (i.e. $\gamma_j$ ) is about an order of magnitude smaller than that of the test-particle and the $m = 15 m_{\oplus}$."
+" Since the 15maintroduction of the variation of the perturber's eccentricity, we are now faced with a one-and- degrees of freedom Hamiltonian."," Since the introduction of the variation of the perturber's eccentricity, we are now faced with a one-and-a-half degrees of freedom Hamiltonian."
+ The dynamics of such a system is best visualized by using a Poincaré surface of section., The dynamics of such a system is best visualized by using a Poincarè surface of section.
+" As in the phase-space portraits above, on a Poincaré surface of section, a periodic orbit will appear as a point, or a finite sequence of points."," As in the phase-space portraits above, on a Poincarè surface of section, a periodic orbit will appear as a point, or a finite sequence of points."
+" À quasi-periodic orbit will appear as a curve that closes upon itself, while a chaotic orbit will appear as a sea of points, which fill a portion of the phase-space."," A quasi-periodic orbit will appear as a curve that closes upon itself, while a chaotic orbit will appear as a sea of points, which fill a portion of the phase-space."
+" Here, we take a section through phase-space every time ey goes through zero i.e. with a period of approximately 1Myr."," Here, we take a section through phase-space every time $e_p$ goes through zero i.e. with a period of approximately $1$ Myr."
+ The Poincaré surface of section illustrating the chaotic dynamics of the particle is shown in Figure 4., The Poincarè surface of section illustrating the chaotic dynamics of the particle is shown in Figure 4.
+ The, The
+"reached, which means the run is not reliable.","reached, which means the run is not reliable."
+ The bottom panel shows the number of triggers per block for one hour of data of a regular observation period., The bottom panel shows the number of triggers per block for one hour of data of a regular observation period.
+ The observation runs for which the number of triggers per data block is exceptionally high for a long period of time have been excluded from the analysis., The observation runs for which the number of triggers per data block is exceptionally high for a long period of time have been excluded from the analysis.
+ Figure 4 shows the distribution of S for the triggered events., Figure \ref{distribution} shows the distribution of $S$ for the triggered events.
+ The top curve corresponds to the raw triggers., The top curve corresponds to the raw triggers.
+" Several additional cuts are applied: Figure 4 displays distributions of S after application of only the timer cut (T), the timer and width cut (TW), and a combination of all cuts (TWA)."," Several additional cuts are applied: Figure \ref{distribution} displays distributions of $S$ after application of only the timer cut (T), the timer and width cut (TW), and a combination of all cuts (TWA)."
+ The line enclosing the black area in Fig., The line enclosing the black area in Fig.
+ 4. corresponds to the distribution of triggers that are expected if the background is pure Gaussian noise (see Appendix AppendixB: for details of the calculation).," \ref{distribution}
+ corresponds to the distribution of triggers that are expected if the background is pure Gaussian noise (see Appendix \ref{app:stat} for details of the calculation)."
+" After all cuts have been appliedthe number of triggers for which S>23 is a factor of 3-4 higher than the amount of triggers expected for Gaussian noise, while the largest S value in the distribution is about three times as large as the highest S value for Gaussian noise."," After all cuts have been appliedthe number of triggers for which $S>23$ is a factor of 3-4 higher than the amount of triggers expected for Gaussian noise, while the largest $S$ value in the distribution is about three times as large as the highest $S$ value for Gaussian noise."
+" Apparently, the background includes pulsed noise that produces triggers and contains pulses that are narrow enough to survive the cut on width."," Apparently, the background includes pulsed noise that produces triggers and contains pulses that are narrow enough to survive the cut on width."
+ The properties of these pulses are further explored in refsec:peak.., The properties of these pulses are further explored in \\ref{sec:peak}.
+" Due to the pulsed noise, the limit that we derive for the neutrino flux is less stringent than estimated in ?),, where the existence of this background was not anticipated."," Due to the pulsed noise, the limit that we derive for the neutrino flux is less stringent than estimated in \citet{scholten}, , where the existence of this background was not anticipated."
+"pseudo-C, methods (c.¢.. Chon et 22001) can lead to aliasing of E-uode power into the much smaller B-ocde power spectrum.","$C_{\ell}$ methods (e.g., Chon et 2004) can lead to aliasing of E-mode power into the much smaller B-mode power spectrum."
+ Although it is possible to construct ways around lis problem (Simith2005)... exact methods such as full liselibood evaluations or Gibbs sampling are clearly preferable solutious.," Although it is possible to construct ways around this problem \citep{smith:2005}, exact methods such as full likelihood evaluations or Gibbs sampling are clearly preferable solutions."
+ Wo start bv discussing the algorithius used for polarized Cajs senrplius. extending the signal and power spectrn sampling steps from teniperature to polarization.," We start by discussing the algorithms used for polarized Gibbs sampling, extending the signal and power spectrum sampling steps from temperature to polarization."
+ Then we analyze simulated data to verify that t1ο algorithm works aud to determine the computation:d efficiency of the method., Then we analyze simulated data to verify that the algorithm works and to determine the computational efficiency of the method.
+ , 
+This fu[usction at /=0 ts disceretized onto a grid with 100 zones along each side. aud the evolution f«llowed using the radiation diffusion term aloue.,"This function at $t=0$ is discretized onto a grid with 100 zones along each side, and the evolution followed using the radiation diffusion term alone."
+ The timestep chosen is a factor οἱ a hundred longer than the time (Aσιn for⋅ radiation⋅⋅ to diffusem across a zone., The timestep chosen is a factor of a hundred longer than the time $(\Delta x)^2\over D$ for radiation to diffuse across a zone.
+ The largest zone-by-zone dillereuce between the numerical auc analytic solutions during the evolution towards equilibrium is5.356... and the largest root-Inean-square cilerence is«.," The largest zone-by-zone difference between the numerical and analytic solutions during the evolution towards equilibrium is, and the largest root-mean-square difference is."
+. When the timestep is shortened by a [actor ten. the corresponcdiug residuals are audC.," When the timestep is shortened by a factor ten, the corresponding residuals are and."
+ The implicit numerical solution relaxes towards equ[unibrium more slowly thau the aialvtic solution in each case., The implicit numerical solution relaxes towards equilibrium more slowly than the analytic solution in each case.
+ When the cirectionally-split advection substeps describedin section 1.5 are applied in the sale order each timestep. the splitting leads to a systematic delay in trausport aloug the first axis. aud an advance in. trausport along the second axis.," When the directionally-split advection substeps describedin section \ref{sec:transport} are applied in the same order each timestep, the splitting leads to a systematic delay in transport along the first axis, and an advance in transport along the second axis."
+ As a test of this effect. the source terms were switched off. aud a circular regiou of increased radiation energy deusity initially at the center of the domain was carried twice diagonally across a grid of 327 zones with periodic boundaries.," As a test of this effect, the source terms were switched off, and a circular region of increased radiation energy density initially at the center of the domain was carried twice diagonally across a grid of $32^2$ zones with periodic boundaries."
+ About 2000 timesteps were taken. and the advection substeps were applied always in the order 1. 2.," About 2000 timesteps were taken, and the advection substeps were applied always in the order 1, 2."
+ At the end of this calculation. the cent'okl of the region lay off the grid diagonal by 0.02 zones aloug each axis.," At the end of this calculation, the centroid of the region lay off the grid diagonal by 0.02 zones along each axis."
+ When the calculation was repeated with the order of the substeps reversed each timestep. the centroid remained within 3»10' zones of the diagonal.," When the calculation was repeated with the order of the substeps reversed each timestep, the centroid remained within $3\times 10^{-7}$ zones of the diagonal."
+ We have choseu to reverse the order of advection each timestep iu all turther caleulatious discussed iu this article. in order to minimize the effects of the directional splitting of the trausport terms.," We have chosen to reverse the order of advection each timestep in all further calculations discussed in this article, in order to minimize the effects of the directional splitting of the transport terms."
+ Waves in a radiatiug 1uid may be dau;»ed by escape of photous., Waves in a radiating fluid may be damped by escape of photons.
+ Damping is often rapid for oscillations with (requeicy near the rate at which material cools by radiatiug. auc also [or oscillations with wavelengtl uear the distauce radiation dilluses in a wave period.," Damping is often rapid for oscillations with frequency near the rate at which material cools by radiating, and also for oscillations with wavelength near the distance radiation diffuses in a wave period."
+ Mibalas Mihalas (1983) obtained a Lilear dispersion reatiou for radiation and acoustic plane waves. drivel by a boundary into a uniform medium with oj»acity due to absorption.," Mihalas Mihalas (1983) obtained a linear dispersion relation for radiation and acoustic plane waves, driven by a boundary into a uniform medium with opacity due to absorption."
+ They made the Exldiugtou approximation. auc assuuiec a uniform backe'ouud iu thermal equilibrium.," They made the Eddington approximation, and assumed a uniform background in thermal equilibrium."
+ The coefficients of the dispersion relation may be specifiel by two dimeusionless parameters., The coefficients of the dispersion relation may be specified by two dimensionless parameters.
+ The Boltzimmaun uumber Bo=Ises€/(cE) is a ratio of typical rates of οποιον transport due to material advection and raciation in the background state. aud r=E/(Ise) fixes the ratio ofthe energy deusities.," The Boltzmann number $\mbox{Bo}\equiv{4\gamma c_a e/(c E)}$ is a ratio of typical rates of energy transport due to material advection and radiation in the background state, and $r\equiv {E/(4\gamma e)}$ fixes the ratio of the energy densities."
+" In these definitions. 5 is the index in the equation of state aud c, is the adiabatic souud speed."," In these definitions, $\gamma$ is the index in the equation of state and $c_a$ is the adiabatic sound speed."
+ For the tests discussed below. Bo=10. and r= 0.1.," For the tests discussed below, $\mbox{Bo}=10^{-3}$ and $r=0.1$ ."
+ Three [requencies are considered. selected for optical depth per wavelength τν=10.7. unity. and LO00.," Three frequencies are considered, selected for optical depth per wavelength $\tau_\lambda=10^{-3}$, unity, and 1000."
+ The first of these frequencies is near the raciative cooling rate. and the last mode has wavelength near the radiative diffusion distance.," The first of these frequencies is near the radiative cooling rate, and the last mode has wavelength near the radiative diffusion distance."
+ In the test caleulatious.H absorptionH opacityH isH set to 0.1 cm?> ο| timesH the density.," In the test calculations, absorption opacity is set to 0.4 $^2$ $^{-1}$ times the density."
+H The domainH isH initiallyHm ⋅ ⋅⋅ ⊳∖↕⋜↕⋃∩∐⋜↕↓⋅⊽∖⊽⋜↕⋯⇂⋯∐↥∩⋅⋯⋅⋜↕∐≺⊔∐↩∢∙∐∩⊳∖≺↵∐⊟∩⋜⋯≺⇂∣⊲∢∙∩↓⋅↕⋅≺↵⊳∖↥↽≻∩∐≺⇂↕≺⊓⇂≺↵⊳∖∐⊽∖⊽∙⋝⋅−≻∐≽≍⋯⋝∑≟∢∙⋯ ⋅ ⋅⋅∣↽⋅↽≝ o," The domain is initially stationary and uniform, and the chosen $\mbox{Bo}$ and $r$correspond to density $3.216\times
+10^{-9}$ g $^{-3}$ ,"
+H The domainH isH initiallyHm ⋅ ⋅⋅ ⊳∖↕⋜↕⋃∩∐⋜↕↓⋅⊽∖⊽⋜↕⋯⇂⋯∐↥∩⋅⋯⋅⋜↕∐≺⊔∐↩∢∙∐∩⊳∖≺↵∐⊟∩⋜⋯≺⇂∣⊲∢∙∩↓⋅↕⋅≺↵⊳∖↥↽≻∩∐≺⇂↕≺⊓⇂≺↵⊳∖∐⊽∖⊽∙⋝⋅−≻∐≽≍⋯⋝∑≟∢∙⋯ ⋅ ⋅⋅∣↽⋅↽≝ oE," The domain is initially stationary and uniform, and the chosen $\mbox{Bo}$ and $r$correspond to density $3.216\times
+10^{-9}$ g $^{-3}$ ,"
+spectral parameters.,spectral parameters.
+ Notwithstanding the significant count-rate variability. the spectral parameters show little variation in time.," Notwithstanding the significant count-rate variability, the spectral parameters show little variation in time."
+ All the stars in the sample show significant variability. although with different characteristics.," All the stars in the sample show significant variability, although with different characteristics."
+ To quantify this variability we computed the normalized cumulative distributions of the amplitude variability., To quantify this variability we computed the normalized cumulative distributions of the amplitude variability.
+ These represent the fraction of time that a source spends in a state with the flux larger than a given value. eXpressedin terms of a given a normalization value. which can be the minimum count rate. the median count rate. etc.," These represent the fraction of time that a source spends in a state with the flux larger than a given value, expressed in terms of a given a normalization value, which can be the minimum count rate, the median count rate, etc."
+ For the present sample we took as normalization value the count rate above which a source spends of the time: this ts less sensitive to noise fluctuations than the real minimum., For the present sample we took as normalization value the count rate above which a source spends of the time; this is less sensitive to noise fluctuations than the real minimum.
+ Figure 16 shows the distribution for the stars in our sample (except V71O Tau for which the count rate uncertainties are too large)., Figure \ref{fig:intAll} shows the distribution for the stars in our sample (except V710 Tau for which the count rate uncertainties are too large).
+ Some sources (V826 Tau. HD285845 and XZ Tau) show mostly low-amplitude variability. in which less than of the time is spent in a state 1.3—1.5 above the minimum.," Some sources (V826 Tau, HD285845 and XZ Tau) show mostly low-amplitude variability, in which less than of the time is spent in a state $-$ 1.5 above the minimum."
+ On the other hand. V827 Tau and HL Tau spend more than of the time in such a state. while V1075 Tau shows intermediate behavior.," On the other hand, V827 Tau and HL Tau spend more than of the time in such a state, while V1075 Tau shows intermediate behavior."
+ The differences between V827 Tau and HL Tau are not due to the large flares present in their light curves. as the flare points affect only the 10 bins relative to the highest count rate. while the low-variability tails of the distributions are also significantly different.," The differences between V827 Tau and HL Tau are not due to the large flares present in their light curves, as the flare points affect only the 10 bins relative to the highest count rate, while the low-variability tails of the distributions are also significantly different."
+ Kolmogorov-Smirnov tests for the cumulative distributions of various pairs of sources also indicate the presence of the different behaviors., Kolmogorov-Smirnov tests for the cumulative distributions of various pairs of sources also indicate the presence of the different behaviors.
+ While this would suggest differences in the X-ray emission processes. the two groups are heterogeneous. and both contain CTTS and WTTS. so that a simple interpretation in terms. of e.g. accreting vs. non-aecreting sources is not possible.," While this would suggest differences in the X-ray emission processes, the two groups are heterogeneous, and both contain CTTS and WTTS, so that a simple interpretation in terms of e.g. accreting vs. non-accreting sources is not possible."
+ The X-ray monitoring observations of the L1551 star-forming complex discussed here. in conjunction with the 2000 oobservations (FGMO3) and the 2001 oobservations (Ballyetal..2003) ). has allowed us to probe the X-ray intensity and spectral variability from six. YSOs (both CTTS and WTTS). on a variety of time scales. from hours to days to years.," The X-ray monitoring observations of the L1551 star-forming complex discussed here, in conjunction with the 2000 observations (FGM03) and the 2001 observations \citealp{bfr2003}) ), has allowed us to probe the X-ray intensity and spectral variability from six YSOs (both CTTS and WTTS), on a variety of time scales, from hours to days to years."
+ Most stars in the sample show X-ray variability at an amplitude of a factor of =2 on time scales of a few days., Most stars in the sample show X-ray variability at an amplitude of a factor of $\simeq 2$ on time scales of a few days.
+ In the WTTS V1075 Tau. the source count rate varies on time scales similar to its optically derived rotational period. which would be compatible with rotational modulation (as observed in YSO in the ONC. Flaccomioetal..2005)). but the significant correlation with spectral variability would appear to point to intrinsic variations.," In the WTTS V1075 Tau, the source count rate varies on time scales similar to its optically derived rotational period, which would be compatible with rotational modulation (as observed in YSO in the ONC, \citealp{fms+2005b}) ), but the significant correlation with spectral variability would appear to point to intrinsic variations."
+ Plasma temperature variations of the order of on similar time scales are also present. and in some cases they appear to be correlated with the intensity ffor V1075 Tau).," Plasma temperature variations of the order of on similar time scales are also present, and in some cases they appear to be correlated with the intensity for V1075 Tau)."
+ Again XZ Tau emerges as a peculiar source., Again XZ Tau emerges as a peculiar source.
+ While the short-term (day seale) spectral variability of XZ Tau ts not linked. to variations in the absorbing column density. as originally speculated by ΕΟΜΟΣ - reprocessing and reanalysing the original 2000 ddata showing that the variation reported by FGMO3 Was probably spurious — this is the only source in the sample to show significant long-term changes in its spectrum.," While the short-term (day scale) spectral variability of XZ Tau is not linked to variations in the absorbing column density, as originally speculated by FGM03 – reprocessing and reanalysing the original 2000 data showing that the variation reported by FGM03 was probably spurious – this is the only source in the sample to show significant long-term changes in its spectrum."
+ The spectrum of XZ Tau observed with tin 2000 had a significant hard component (7=4 keV). which was not visible in the 2001 ddata and is not present in the 2004 ddata and both are compatible with a hard component of at most 1.5 keV. The difference between the two spectral states is very well visiblein Fig. 1H..," The spectrum of XZ Tau observed with in 2000 had a significant hard component $T \simeq 4$ keV), which was not visible in the 2001 data and is not present in the 2004 data and both are compatible with a hard component of at most 1.5 keV. The difference between the two spectral states is very well visible in Fig. \ref{fig:xz_comp}."
+ The X-ray spectral hardening takes place close in time to the strong optical outburst of the N component of the binary., The X-ray spectral hardening takes place close in time to the strong optical outburst of the N component of the binary.
+ While the data by Coffeyetal. (2004).. who reported the outburst. only cover its beginning. our own inspection of new HST data indicates that the outburst was not long-lived. with the source returning to its state fainter than XZ Tau S by early 2002. one year after the early 2001 peak.," While the data by \cite{cdr04}, who reported the outburst, only cover its beginning, our own inspection of new HST data indicates that the outburst was not long-lived, with the source returning to its state fainter than XZ Tau S by early 2002, one year after the early 2001 peak."
+ Finally. while (Fig. 10))," Finally, while (Fig. \ref{fig:xzall}) )"
+ the count rate varied smoothly between 0.04 and 0.22 ets/see during the 2000 observation. during the 2004 campaign it never decreased below 0.12 cts/sec (with a peak value of =0.35 cts/sec). again showing the source to be in a different state.," the count rate varied smoothly between 0.04 and 0.22 cts/sec during the 2000 observation, during the 2004 campaign it never decreased below 0.12 cts/sec (with a peak value of $\simeq 0.35$ cts/sec), again showing the source to be in a different state."
+ The binary system XZ Tau is not resolved in the X- therefore one cannot exclude that XZ Tau S may have been the dominant X-ray source. at one or all the epochs at which it was observed.," The binary system XZ Tau is not resolved in the X-ray, therefore one cannot exclude that XZ Tau S may have been the dominant X-ray source, at one or all the epochs at which it was observed."
+ Nevertheless the observed phenomena are very reminiscent of the spectral hardening and. variability observed in the outbursting source. V1647 Ori (Kastneretal.. 2004:; Grossoetal.. 2005)). so that 1t seems reasonable to interpret the X-ray spectral changes and variability of XZ Tau as connected with the outburst of XZ Tau N. The potential physical mechanism causing the significantly harder X-ray spectrum (with a somewhat lower emission level) during the," Nevertheless the observed phenomena are very reminiscent of the spectral hardening and variability observed in the outbursting source, V1647 Ori \citealp{krg+2004}; \citealp{gko+2005}) ), so that it seems reasonable to interpret the X-ray spectral changes and variability of XZ Tau as connected with the outburst of XZ Tau N. The potential physical mechanism causing the significantly harder X-ray spectrum (with a somewhat lower emission level) during the"
+hat these caleulatious are doue iu a regiou where the funnel outflow ts well established he uunerical ceiling on the LorenZ factor does not come iuto play.,that these calculations are done in a region where the funnel outflow is well established the numerical ceiling on the Lorentz factor does not come into play.
+ To account for mass loss froi the disk to the black hole. we compute the rate of mass accretjon hrough the inner radial boutdary.," To account for mass loss from the disk to the black hole, we compute the rate of mass accretion through the inner radial boundary."
+ The mass and energy €arried by the jets and coronal wiud increase with iucreasiug black Itole spin. while the mass accreted onto the black hole decreases with increasiug spin. consistent will the finclines of DOS.," The mass and energy carried by the jets and coronal wind increase with increasing black hole spin, while the mass accreted onto the black hole decreases with increasing spin, consistent with the findings of D05."
+ However. he energy in the jets is highe“i all in the analogous models i Dxο. in. part due to initlal couchtions. but also due to improvementslue o the GRMHD cocle’s abiliiy o handle the οναι]s of tie jets.," However, the energy in the jets is higher than in the analogous models in D05, in part due to initial conditions, but also due to improvements to the GRMHD code's ability to handle the dynamics of the jets."
+" The inass auc euergy carried by he jets aud coronal wiud i1 3D model Ryyp38D are coriparable o the analogous 2D model B, suggests that tlie other axisvüuimetric models are adequatey capturiug the fluxes of unbouix aud euergy.", The mass and energy carried by the jets and coronal wind in 3D model ${\rm R_{vf}}$ 3D are comparable to the analogous 2D model ${\rm R_{vf}}$; this suggests that the other axisymmetric models are adequately capturing the fluxes of unbound mass and energy.
+ The main dillerenc'e between the 3D mocel aud its 2D analog is that the jet is s more energetic than the coronub wind in 3D. while the sitiation is reversed in 2D. To extract physical units [rom the mass aud euergy f[Iuxes. we use tlie light-crossiug time and the gravitational leugth ry=GΓη... to rescale time aud distance measurements.," The main difference between the 3D model and its 2D analog is that the jet is slightly more energetic than the coronal wind in 3D, while the situation is reversed in 2D. To extract physical units from the mass and energy fluxes, we use the light-crossing time $t_{\rm lc}=G\,M_{\rm BH}/c^3$ and the gravitational length $r_g=G\,M_{\rm BH}/c^2$ to rescale time and distance measurements."
+ We take collapsar fiducial parameters comparable to those of W93. namely a black bole mass ob Adpy=3ÀL. aud au initial torus of inass 0.3.A/.. a value consistent with the initial torus parameters and a central density in the range of 101 —10Pgem?," We take collapsar fiducial parameters comparable to those of W93, namely a black hole mass of $M_{\rm BH} = 3\,M_\odot$ and an initial torus of mass $0.3\,M_\odot$, a value consistent with the initial torus parameters and a central density in the range of $10^{11}$ $10^{12}\,{\rm g}\,{\rm cm}^{-3}$."
+ Iu 3D süinulations. the lifetime of the jets is set by the reservoir of inass in tle 1ualu disk: as long as turbulence drives the trausport of angular momentum. accretion aud ejection will continue.," In 3D simulations, the lifetime of the jets is set by the reservoir of mass in the main disk; as long as turbulence drives the transport of angular momentum, accretion and ejection will continue."
+ In axisyiumnetry. the MRI which drives accretion is uot sustainable. so a preinaure end to the accretion/ejection process occurs when the MRI subsides (alter approximately 10 orbits of the nain disk).," In axisymmetry, the MRI which drives accretion is not sustainable, so a premature end to the accretion/ejection process occurs when the MRI subsides (after approximately 10 orbits of the main disk)."
+ Nevertheless. we use the mass fluxes (AL) of the axisyiumetric models to infer the ifetime of the jets.," Nevertheless, we use the mass fluxes $\dot{M}$ ) of the axisymmetric models to infer the lifetime of the jets."
+ This lifetime is determined by the depletion of disk mass. Le. nass carried by he jet aud coronal wind (through 7 /A/=100 here) and the mass accreted into the black hole.," This lifetime is determined by the depletion of disk mass, i.e. mass carried by the jet and coronal wind (through $r/M$ =100 here) and the mass accreted into the black hole."
+ Table 2. lists the jet Lifetimes inferred from the simulation data: the values range from 0.1 to Q.2 s lor he high-spin simulatious to 1.0 to 1.1 s for the zero-spin case., Table \ref{results} lists the jet lifetimes inferred from the simulation data; the values range from 0.1 to 0.2 s for the high-spin simulations to 1.0 to 1.4 s for the zero-spin case.
+ Since the lifetime for mocel Εν 21) is comparable to its axisviunmetric analog. we can infer that jet lifetimes of a few tentlis to a few seconds are predicted by these simulations. with louger lifetimes produced by the owel mass fluxes generated by low-spin black holes.," Since the lifetime for model ${\rm R_{vf}}$ 3D is comparable to its axisymmetric analog, we can infer that jet lifetimes of a few tenths to a few seconds are predicted by these simulations, with longer lifetimes produced by the lower mass fluxes generated by low-spin black holes."
+ The jet lifetime is used to calculate the energy. transported by the jets tl‘ough the simple relation Ev=lojoctEEv. where Ei=0:3M.0 is the rest mass of the iitial torus.," The jet lifetime is used to calculate the energy transported by the jets through the simple relation $E_{\rm total} = t_{\rm eject}\,\dot{E}\,E_{\rm torus}$, where $E_{\rm torus} = 0.3\,M_\odot\,c^2$ is the rest mass of the initial torus."
+ Iu all simulations. the energy. carried by the jets is between LO aud LOY erg and some fraction of this euergy should be converted to radiation ouce jet material reaches the optically tlin outer region of the collapsar.," In all simulations, the energy carried by the jets is between $10^{48}$ and $10^{49}$ erg and some fraction of this energy should be converted to radiation once jet material reaches the optically thin outer region of the collapsar."
+ An analogous calculation shows that the coronal wind carries a comparable amount ol euergy., An analogous calculation shows that the coronal wind carries a comparable amount of energy.
+ These energetic coronal winds flow through a much larger surface a‘ea than the jets.," These energetic coronal winds flow through a much larger surface area than the jets,"
+"Let us first consider the ideal noiseless case. when the sampling resolution in the o space is equal with the half maxinmnm of the main peak of the RAISF. ὁμ=δώ. and the computational window is Oyj,=Omer.","Let us first consider the ideal noiseless case, when the sampling resolution in the $\phi$ space is equal with the half maximum of the main peak of the RMSF, $\phi_{R}=\delta\phi$, and the computational window is $\phi_{win}=\phi_{max}$."
+ In this particular case. as shown in Figure 1. the algorithm provides an exact solution. since V=M126 and therefore no information is lost in the measurement.," In this particular case, as shown in Figure 1, the RM-MP algorithm provides an exact solution, since $N=M=126$ and therefore no information is lost in the measurement."
+ One can also notice that in this noiseless exact sampling case. the solution is independent of the scale used in the dictionary. as it can beseen on the multi-scale representation lor ο«5x25.," One can also notice that in this noiseless exact sampling case, the solution is independent of the scale used in the dictionary, as it can beseen on the multi-scale representation for $0<S\leq25$."
+ However. (he problem becomes ill-clefinecl in the following situations: (he noise is present: (he sampling resolution becomes liner than the half maximum of the main peak of the RAISE. ὁν<δώ: and the number of independent observed channels issmaller than the number of points in the ó space. Vo«AL.," However, the problem becomes ill-defined in the following situations: the noise is present; the sampling resolution becomes finer than the half maximum of the main peak of the RMSF, $\phi_{R}<\delta\phi$; and the number of independent observed channels issmaller than the number of points in the $\phi$ space, $N<M$."
+ In this case the svstem becomes under-deterimined. and (therefore some information is lost.," In this case the system becomes under-determined, and therefore some information is lost."
+ In order (o exemplify (his situation. we consider a scenario in which all these [actors are present.," In order to exemplify this situation, we consider a scenario in which all these factors are present."
+ We add noise with (he standard deviation o=YN11.22. to the Q and U values., We add noise with the standard deviation $\sigma=\sqrt{N}=11.22$ to the $Q$ and $U$ values.
+" We limit the computational windowto o4,=126radm?<6,,; and we increase the number of points on the ó grid to AF=252. whichis double of the number of observation channels N=126."," We limit the computational windowto $\phi_{win}=126\,\mathrm{rad}\,\mathrm{m}^{-2}<\phi_{max}$, and we increase the number of points on the $\phi$ grid to $M=252$, whichis double of the number of observation channels $N=126$."
+" This results in a sampling resolution of o,=lradim?«96."," This results in a sampling resolution of $\phi_{R}=1\,\mathrm{rad}\,\mathrm{m}^{-2}\ll\delta\phi$."
+ The obtained results (lor 9]=y2N. S= 25) are shown in Figure 2.," The obtained results (for $\beta=\sqrt{2N}$, $S=25$ ) are shown in Figure 2."
+ One can see that the phase of some components cannot be reliably recovered anyanore. since (here is not enough information in the sienal to be detected properly.," One can see that the phase of some components cannot be reliably recovered anymore, since there is not enough information in the signal to be detected properly."
+ We should note that the problem is correctly resolved in (he noiseless case (not shown here)., We should note that the problem is correctly resolved in the noiseless case (not shown here).
+ Thus. the effect of noise addition consists in a partial loss of information about the phase of ΕΦ). which is expected. since the number of solutions compatible with the data increases dramatically with the added noise.," Thus, the effect of noise addition consists in a partial loss of information about the phase of $F(\phi)$, which is expected, since the number of solutions compatible with the data increases dramatically with the added noise."
+ Also. we should point out that the solution improves by increasing the signal-to-noise ratio. as shown in Figure 3. where we have increased the amplitude of the components by a factor of 1.5. keeping (heir phase unchanged.," Also, we should point out that the solution improves by increasing the signal-to-noise ratio, as shown in Figure 3, where we have increased the amplitude of the components by a factor of 1.5, keeping their phase unchanged."
+ One can see (hat in this case. the IRM-MP. method resolves correctly all (he components.," One can see that in this case, the RM-MP method resolves correctly all the components."
+ This result suggests that an adequate signal to noise ratio should be taken into account. in order for the method to be successful.," This result suggests that an adequate signal to noise ratio should be taken into account, in order for the method to be successful."
+" The GALFACTS survev. carried ont with the Arecibo telescope. has the following parameters: Frequency range: 545,5,=1225 MlIIZ. tyne,=1525 MIIZ: Wave length range: A2,,,=0.0386 n.A2,,,=0.0598 m2. AA?=0.0212 n: Half maxinium of the main peak of the RMSF: οὐ=163.044radm. 7"," The GALFACTS survey, carried out with the Arecibo telescope, has the following parameters: Frequency range: $\nu_{min}=1225\,\mathrm{MHz}$ , $\nu_{max}=1525\,\mathrm{MHz}$ ; Wave length range: $\lambda_{min}^{2}=0.0386\,\mathrm{m}^{2}$ ,$\lambda_{max}^{2}=0.0598\,\mathrm{m}^{2}$ , $\triangle\lambda^{2}=0.0212\,\mathrm{m}^{2}$ ; Half maximum of the main peak of the RMSF: $\delta\phi=163.044\,\mathrm{rad}\,\mathrm{m}^{-2}$ ;"
+" The GALFACTS survev. carried ont with the Arecibo telescope. has the following parameters: Frequency range: 545,5,=1225 MlIIZ. tyne,=1525 MIIZ: Wave length range: A2,,,=0.0386 n.A2,,,=0.0598 m2. AA?=0.0212 n: Half maxinium of the main peak of the RMSF: οὐ=163.044radm. 7:"," The GALFACTS survey, carried out with the Arecibo telescope, has the following parameters: Frequency range: $\nu_{min}=1225\,\mathrm{MHz}$ , $\nu_{max}=1525\,\mathrm{MHz}$ ; Wave length range: $\lambda_{min}^{2}=0.0386\,\mathrm{m}^{2}$ ,$\lambda_{max}^{2}=0.0598\,\mathrm{m}^{2}$ , $\triangle\lambda^{2}=0.0212\,\mathrm{m}^{2}$ ; Half maximum of the main peak of the RMSF: $\delta\phi=163.044\,\mathrm{rad}\,\mathrm{m}^{-2}$ ;"
+ , 
+3 for cach cluster and for cach,\ref{table_neISO} for each cluster and for each.
+for. Figure 2 shows the electron temperature (top). τανο density (1iiddle). aud pressure (bottoni) profiles for only the cluster ZW2316. as an example. of both the original CC ICAL aud the recovered ISO model.," Figure \ref{fig_ne_Te_pe_profiles1} shows the electron temperature (top), number density (middle), and pressure (bottom) profiles for only the cluster ZW3146, as an example, of both the original CC ICM and the recovered ISO model."
+ The errors associated to the curves account for the lo uucertaimties on the parameters., The errors associated to the curves account for the $\sigma$ uncertainties on the parameters.
+ We want tosress the consequences of the assumptions ou the ICAL physjes when we iiss X-ray uforinatiou., We want to stress the consequences of the assumptions on the ICM physics when we miss X-ray information.
+ A quantity. such as the total mass Mj. cau be biased by a different plivsical state of the ICAL (1.0. 1iergers or cooling flow mechanisius).," A quantity, such as the total mass $M_{tot}$, can be biased by a different physical state of the ICM (i.e. mergers or cooling flow mechanisms)."
+ In particular we estimate the mass for both the ICM discussed templates (CC and £90). by using the followiie ifferent approaches: The masses are caleulated. iu particular. within a fixed iutegration radius rjj; (aperture radius). which we arbitrarily choose equal to the rsyy values as reported in Morandi et al. (," In particular we estimate the mass for both the ICM discussed templates $CC$ and $ISO$ ), by using the following different approaches: The masses are calculated, in particular, within a fixed integration radius $r_{int}$ (aperture radius), which we arbitrarily choose equal to the $r_{500}$ values as reported in Morandi et al. ("
+2007. see Table 2)).,"2007, see Table \ref{table_te}) )."
+ This choice is motivated bv the need to fix au aperture radius within which to estimate integrated quantities., This choice is motivated by the need to fix an aperture radius within which to estimate integrated quantities.
+ We poiut out that Ming does not always correspond to the same overdeusity. due to the different. asstmed ICAL templates.," We point out that $r_{int}$ does not always correspond to the same overdensity, due to the different assumed ICM templates."
+ It is clear that. hereafter. results associated to the clusters simulated in this work cumot be considered as describing the true ICM physics o ‘the observed objects.," It is clear that, hereafter, results associated to the clusters simulated in this work cannot be considered as describing the true ICM physics of the observed objects."
+ We choose. however. to maintain the link with the “native” cluster iu the name CNAME»CTv arr).," We choose, however, to maintain the link with the “native” cluster in the name $NAME \rightarrow
+Cl_{NAME}$ )."
+" The first approach. ITE. assumes a spherical svuumetry for the cluster. so that the bydrostatic equilibrium equation can be written (Sarazin 1988)) as where AZ&gg4(orn) is the total cluster lnass within radius & and under the ideal gas assumption. P4.=(PyashBTyas) Gun). with Pyasfllty poosud gn are the total aud. electron mean molecular weights (1.6. the mean particle mass per electron in units of the proton lass 5). Gis the exavitational constant and TyasTiT.. the jon and olectrou temperatures respectively, because the system is assuned to be in thermal equilibrimu."," The first approach, HE, assumes a spherical symmetry for the cluster, so that the hydrostatic equilibrium equation can be written \cite{Sarazin88}) ) as where $M_{tot}(<r)$ is the total cluster mass within radius $r$ and under the ideal gas assumption, $P_{gas}=(\rho_{gas}k_BT_{gas})/(\mu
+m_p)$ , with $\rho_{gas}=\mu_e m_p n_e$, $\mu$ and $\mu_e$ are the total and electron mean molecular weights (i.e. the mean particle mass per electron in units of the proton mass $m_p$ ), $G$ is the gravitational constant and $T_{gas}=T_i=T_e$, the ion and electron temperatures respectively, because the system is assumed to be in thermal equilibrium."
+ To adapt this equation to the SZ, To adapt this equation to the SZ
+along the ENE/WSW axis through the nucleus.,along the ENE/WSW axis through the nucleus.
+ The position angle of the observed extension 1s rotated relative to this axis by approximately 307., The position angle of the observed extension is rotated relative to this axis by approximately $^\circ$.
+" In addition, the nucleus contains only ~2000 counts, so 30 counts are expected over the entire ~8' length of the FTS."," In addition, the nucleus contains only $\sim$ 2000 counts, so $\sim$ 30 counts are expected over the entire $\sim$ $'$ length of the FTS."
+HST observations of the host ealaxy show a similar elongation in the stellar distribution suggesting that the extended X-ray emission is from thermal gas., observations of the host galaxy show a similar elongation in the stellar distribution suggesting that the extended X-ray emission is from thermal gas.
+" There is an additional spatially extended X-ray component ~5"" (5.7 kpc) to the northeast.", There is an additional spatially extended X-ray component $\sim$ $''$ (5.7 kpc) to the northeast.
+ It is unclear whether this is thermal emission from an extended gascous component or X-ray emission from a nuclear jet., It is unclear whether this is thermal emission from an extended gaseous component or X-ray emission from a nuclear jet.
+ There is no evidence of a kpe scale jet in any of our radio maps., There is no evidence of a kpc scale jet in any of our radio maps.
+" The spectrum of the nucleus was extracted from à source-centered circle of radius 10 pixels (4.92""), with background sampled from a surrounding annulus of inner radius 10 pixels and outer radius 20 pixels (9.543)."," The spectrum of the nucleus was extracted from a source-centered circle of radius 10 pixels $''$ ), with background sampled from a surrounding annulus of inner radius 10 pixels and outer radius 20 pixels $''$ )."
+ The count rate from the nucleus is ~0.05 cts !. indicating little pileup.," The count rate from the nucleus is $\sim$ 0.05 cts $^{-1}$, indicating little pileup."
+" Initially we fitted the spectrum of the nucleus with a three-component model: an absorbed power law, a Gaussian line representing emission from neutral Fe K,,, and a thermal (the XSPEC Astrophysical Plasma Emission Code model. or simply APEC) component."," Initially we fitted the spectrum of the nucleus with a three-component model: an absorbed power law, a Gaussian line representing emission from neutral Fe $_\alpha$, and a thermal (the XSPEC Astrophysical Plasma Emission Code model, or simply APEC) component."
+" The fit was poor (2219.9 for 26 degrees of freedom) with large residuals between | and 3 keV, so we investigated two alternative models."," The fit was poor $\chi^2_\nu$ =19.9 for 26 degrees of freedom) with large residuals between 1 and 3 keV, so we investigated two alternative models."
+" Model | consisted of a four component model, the three components listed in the previous paragraph plus a reflection component."," Model 1 consisted of a four component model, the three components listed in the previous paragraph plus a reflection component."
+" The reflection model has been used to fit the X-ray spectra of Seytert galaxies, but has rarely been found to provide an adequate description of the X-ray spectra of radio-galaxy nuclei."," The reflection model has been used to fit the X-ray spectra of Seyfert galaxies, but has rarely been found to provide an adequate description of the X-ray spectra of radio-galaxy nuclei."
+ The X-ray spectrum with the best-fit four-component reflection and two power-law component models and residuals are shown in the left and right plots of Figure 14.. respectively.," The X-ray spectrum with the best-fit four-component reflection and two power-law component models and residuals are shown in the left and right plots of Figure \ref{specfit}, respectively."
+" Model 2 also consisted of the three components listed in the previous paragraph plus a second, less absorbed power law representing emission from a parsec scale jet."," Model 2 also consisted of the three components listed in the previous paragraph plus a second, less absorbed power law representing emission from a parsec scale jet."
+ Such a model has been shown to be a good description of emission from the nuclei of ER II radio galaxies 2006)., Such a model has been shown to be a good description of emission from the nuclei of FR II radio galaxies \citep{evans06}.
+. A summary of best-fit parameters and uncertainties for both models is contained in Table 5.., A summary of best-fit parameters and uncertainties for both models is contained in Table \ref{spectab}.
+ The APEC component in the model that includes a reflection component required abundances of Mg and Si that were larger than the Solar value and a sub-Solar value for Fe., The APEC component in the model that includes a reflection component required abundances of Mg and Si that were larger than the Solar value and a sub-Solar value for Fe.
+ This is probably not physical and may be indicative of systematics in the model such as multi-temperature gas in the ISM or more complex reflection ecometry than a uniformly illuminated slab., This is probably not physical and may be indicative of systematics in the model such as multi-temperature gas in the ISM or more complex reflection geometry than a uniformly illuminated slab.
+" We note that the large column density toward the primary power law in both spectral fits is consistent with the expectations from unified models for anarrow-line radio galaxy, as is generally seen in X-ray observations of large samples of such sources (e.g. Sambrunaefaf1999;Belsoleetαἰ.Evansetaf.2006:Hardcastleefαἰ. 2006))."," We note that the large column density toward the primary power law in both spectral fits is consistent with the expectations from unified models for anarrow-line radio galaxy, as is generally seen in X-ray observations of large samples of such sources (e.g. \citealt{sam99,bel06,evans06,mjh06}) )."
+ This strongly supports the idea that the jet axis 1s close to the plane of the sky and that relativistic beaming is not important in modifying the appearance of the hot spots., This strongly supports the idea that the jet axis is close to the plane of the sky and that relativistic beaming is not important in modifying the appearance of the hot spots.
+" Finally, we mention that we fitted the spectra (both MOS cameras and PN) with the same models and find thatthe fits are consistent with the results."," Finally, we mention that we fitted the spectra (both MOS cameras and PN) with the same models and find thatthe fits are consistent with the results."
+one of the ALACΠΟ and OGLE data sets at the determined period for that data set.,one of the MACHO and OGLE data sets at the determined period for that data set.
+ “Phe result. of this is illustrated in Figure 5.., The result of this is illustrated in Figure \ref{lc}.
+ Lighteurves from four different filters are shown in this figure., Lightcurves from four different filters are shown in this figure.
+ In the case of the V band. the OGLE data are rather sparse since this is not their main filter. and so the data have been supplemented. by observations taken over several nights from the SAAQ 1.0m in January 1999 and 2000.," In the case of the V band, the OGLE data are rather sparse since this is not their main filter, and so the data have been supplemented by observations taken over several nights from the SAAO 1.0m in January 1999 and 2000."
+ Since this data set is not as uniform as the other 3 bands the individual results are presented. in the phase cliaegram rather than a folded light curve., Since this data set is not as uniform as the other 3 bands the individual results are presented in the phase diagram rather than a folded light curve.
+ Overall. from this ligure. the extremely sinusoidal nature of the modulation is very clear.," Overall, from this figure, the extremely sinusoidal nature of the modulation is very clear."
+ The topmost curve in Figure 5. is a measure of the colour of the object obtained by subtracting individual MACHO blue measurments from their red. measurements taken on the same night., The topmost curve in Figure \ref{lc} is a measure of the colour of the object obtained by subtracting individual MACHO blue measurments from their red measurements taken on the same night.
+ LE the colour data are also subjected to the same Lomb-Scargle analysis. then the same period emerges from the data. but the depth of modulation is clearly very small.," If the colour data are also subjected to the same Lomb-Scargle analysis, then the same period emerges from the data, but the depth of modulation is clearly very small."
+ Perhaps the most important result to emerge from the combined OCGLISZ/NLACTIO data set is the period. history., Perhaps the most important result to emerge from the combined OGLE/MACHO data set is the period history.
+ Figure 6 illustrates this bv showing the periods determined from the 12 individual data sets., Figure \ref{per} illustrates this by showing the periods determined from the 12 individual data sets.
+ From this figure one can clearly see that the period. is changing bv a significant amount over the 7 vears., From this figure one can clearly see that the period is changing by a significant amount over the 7 years.
+ The simplest interpretation of the period change is a linear one. and from such an assumption a value of 13.5s/vear is found.," The simplest interpretation of the period change is a linear one, and from such an assumption a value of 13.5s/year is found."
+ More complex changes. for example a sinusoidal modulation with a period of 73600d. may be speculated upon but are not required by the data as vel.," More complex changes, for example a sinusoidal modulation with a period of $\sim$ 3600d, may be speculated upon but are not required by the data as yet."
+ The upper curve in Figure 6 shows the result of fitting a sine wave to cach individual data set and determining its amplitude., The upper curve in Figure \ref{per} shows the result of fitting a sine wave to each individual data set and determining its amplitude.
+ Phough there is clear evidence of amplitude variation by ~40% from both this figure ancl Figure 3.. it is not obvious that there is any significant pattern to the change.," Though there is clear evidence of amplitude variation by $\sim$ from both this figure and Figure \ref{mo}, it is not obvious that there is any significant pattern to the change."
+ In trving to establish the optical counterpart to the ASCA pulsar one must keep in mind that the ROSAT source is too weak to have shown any detectable. pulsations., In trying to establish the optical counterpart to the ASCA pulsar one must keep in mind that the ROSAT source is too weak to have shown any detectable pulsations.
+ Thus putting the IROSAT source aside for the moment. the mos objective approach is to just look at the colour-magnitucle diagram (Figure ?22)).," Thus putting the ROSAT source aside for the moment, the most objective approach is to just look at the colour-magnitude diagram (Figure \ref{cmd}) )."
+ This diagram reveals just two objects inside the best ASCA positional error circle - nos., This diagram reveals just two objects inside the best ASCA positional error circle - nos.
+ 512 anc 499., 512 and 499.
+ Phe other two objects lie in. or very close το the original ASCA circle. but are now significantly loss attractive as counterparts.," The other two objects lie in, or very close to the original ASCA circle, but are now significantly less attractive as counterparts."
+ Object. 499 is very. faint compared. to al other known counterparts to LAINBs in the SAIC. which typically have V15.16.," Object 499 is very faint compared to all other known counterparts to HMXBs in the SMC, which typically have $V\sim15-16$."
+ Hs colours and the presence of llo in eniussion suggest a D3-4Ve star - a somewhat later spectral type than most other De/X-ray binary systems., Its colours and the presence of $\alpha$ in emission suggest a B3-4Ve star - a somewhat later spectral type than most other Be/X-ray binary systems.
+ Is, Its
+" The upper MS corrected in this way appears considerably tighter in the CMD than the raw data (see reffig2bb) and this convinced us that the method, albeit statistical in nature, can be applied to the entire data set."," The upper MS corrected in this way appears considerably tighter in the CMD than the raw data (see \\ref{fig2}b b) and this convinced us that the method, albeit statistical in nature, can be applied to the entire data set."
+" We also experimented with a different number of neighbours, namely 5, 10, 20 and 50."," We also experimented with a different number of neighbours, namely 5, 10, 20 and 50."
+" We found that, by using the 20 nearest upper MS stars, in the CMD at V~19 appeared the tightest."," We found that, by using the 20 nearest upper MS stars, in the CMD at $V\simeq 19$ appeared the tightest."
+" Since our goal is to correct for reddening all candidate PMS stars, which are red, this choice seemed the most appropriate."," Since our goal is to correct for reddening all candidate PMS stars, which are typically red, this choice seemed the most appropriate."
+ The de-reddenedtypically CMD obtained in this way is shown in reffig2bb. One of the characteristic signatures of the accretion process on to PMS stars is the presence of excess emission in Ha (see e.g. Calvet et al., The de-reddened CMD obtained in this way is shown in \\ref{fig2}b b. One of the characteristic signatures of the accretion process on to PMS stars is the presence of excess emission in $\alpha$ (see e.g. Calvet et al.
+ 2000)., 2000).
+" Recently, we showed in perIl that PMS stars undergoing active accretion can be reliably identified, and their Ha luminosity measured, by using a combination of broad-band (V,/) and narrow-band (Ho) photometry, with no need for spectroscopy."," Recently, we showed in I that PMS stars undergoing active accretion can be reliably identified, and their $\alpha$ luminosity measured, by using a combination of broad-band $V, I$ ) and narrow-band $\alpha$ ) photometry, with no need for spectroscopy."
+" This way of PMS stars is more accurate and reliable than the simple identifyingclassification based on the position of the objects in the Hertzsprung-Russell diagram, i.e. stars placed well above the main sequence (e.g. Gilmozzi et al."," This way of identifying PMS stars is more accurate and reliable than the simple classification based on the position of the objects in the Hertzsprung–Russell diagram, i.e. stars placed well above the main sequence (e.g. Gilmozzi et al."
+ 1994; Hunter et al., 1994; Hunter et al.
+ 1995; see in particular Nota et al., 1995; see in particular Nota et al.
+ 2006 and Gouliermis et al., 2006 and Gouliermis et al.
+ 2007 for the specific case of NGC3346)., 2007 for the specific case of 346).
+" The true advantages of the method presented in II are that: it allows us to discriminate between bona-fide PMS stars and interlopers that occupy the same region of the CMD; it provides a secure detection of relatively old PMS stars, already close to their MS, whose colours and magnitudes would otherwise be indistinguishable from those of MS stars; by measuring the Ha luminosity L(Ho) of these PMS stars we can obtain their mass accretion rates Mace from L(Ho) and the stellar parameters (mass and radius);(iv) finally, this approach is more practical, efficient and economical than slitless spectroscopy for studying PMS stars in dense star forming regions, although it cannot provide constraints for magnetospheric accretion models that require information on the profile of the Ho emission line (e.g. Muzerolle et al."," The true advantages of the method presented in I are that: it allows us to discriminate between bona-fide PMS stars and interlopers that occupy the same region of the CMD; it provides a secure detection of relatively old PMS stars, already close to their MS, whose colours and magnitudes would otherwise be indistinguishable from those of MS stars; by measuring the $H\alpha$ luminosity $L(H\alpha)$ of these PMS stars we can obtain their mass accretion rates $\dot M_{\rm acc}$ from $L(H\alpha)$ and the stellar parameters (mass and radius); finally, this approach is more practical, efficient and economical than slitless spectroscopy for studying PMS stars in dense star forming regions, although it cannot provide constraints for magnetospheric accretion models that require information on the profile of the $\alpha$ emission line (e.g. Muzerolle et al."
+ 2001)., 2001).
+" As explained in detail in II, the method hinges on the simple consideration that, at any time, the largest majority of stars of a given effective temperature Te in a stellar population will have no excess Ha emission."," As explained in detail in I, the method hinges on the simple consideration that, at any time, the largest majority of stars of a given effective temperature $T_{\rm eff}$ in a stellar population will have no excess $\alpha$ emission."
+" Therefore, for stars of that effective temperature, the value of any colour index involving Ha (e.g. V— Ho) effectively defines a reference template with respect to which the Ho colour spectralexcess should be sought."," Therefore, for stars of that effective temperature, the value of any colour index involving $\alpha$ (e.g. $V-H\alpha$ ) effectively defines a spectral reference template with respect to which the $\alpha$ colour excess should be sought."
+ In reffig3 we show with small dots the V—Ho colour as a function of V—1 for all 187764 stars that have a combined photometric uncertainty 63<0.1 mmag (see refeq1))., In \\ref{fig3} we show with small dots the $V-H\alpha$ colour as a function of $V-I$ for all 764 stars that have a combined photometric uncertainty $\delta_3 \le 0.1$ mag (see \\ref{eq1}) ).
+" The value of 63 is dominated by the uncertainty on the Ha magnitude, while the median value of the uncertainty in the other two bands is ὃν=0.006 and 6,= 0.008, respectively."," The value of $\delta_3$ is dominated by the uncertainty on the $\alpha$ magnitude, while the median value of the uncertainty in the other two bands is $\delta_V=0.006$ and $\delta_I=0.008$ , respectively."
+ The thick dashed line in represents the median V—Ho colour obtained as, The thick dashed line in \\ref{fig3} represents the $V-H\alpha$ colour obtained as
+Tuterestinely. the Balmer cussion liue profiles shown by Clhoruocketal.(2009) are noticeably different than those in Borosou&Lauer (2009)..,"Interestingly, the Balmer emission line profiles shown by \citet{cho09} are noticeably different than those in \citet{bor09}. ."
+ For IL) the the peak of the r component is about of the peak of b component. as measured above the “valley” between the » components. while in Borosou&Lauer(2009) routGOM%.," For $\beta$ the the peak of the r component is about of the peak of b component, as measured above the “valley” between the r and b components, while in \citet{bor09} it is about."
+. Simibulv. for Πα the r peak is about : the b peak in Chornocketal.(2009).. while in Dorosou&Lauer(2009) it is about95%.," Similarly, for $\alpha$ the r peak is about of the b peak in \citet{cho09}, while in \citet{bor09} it is about."
+. These could be due to profile variability. but it could also be duc to different contributions of two poiut sources to the total spectrum. due to different slit positions iu the Dorosou&Lauer(2009). aud Choruocketal.(2009). spectra.," These could be due to profile variability, but it could also be due to different contributions of two point sources to the total spectrum, due to different slit positions in the \citet{bor09} and \citet{cho09}
+ spectra."
+ If the 5-kpc binary interpretation is correct then eround-based adaptive optics imaging iu the visible band. with sub-arcseconcd resolution. should clearly show two point sources. and spectroscopy of those sources should exhibit two differeut spectra. each oue with a single broad cuussiou-line svsteii.," If the 5-kpc binary interpretation is correct then ground-based adaptive optics imaging in the visible band, with sub-arcsecond resolution, should clearly show two point sources, and spectroscopy of those sources should exhibit two different spectra, each one with a single broad emission-line system."
+ Deep optical imagine on aresecond scales may also reveal the nature of the host ealaxyv. aud whether strongly interacting sealaxies are imvolved.," Deep optical imaging on arcsecond scales may also reveal the nature of the host galaxy, and whether strongly interacting galaxies are involved."
+ Deep optical πασάς on arciuiuute scales can establish whether or not SDSS J1536|OLLI resides m a cluster., Deep optical imaging on arcminute scales can establish whether or not SDSS J1536+0441 resides in a cluster.
+ Doroson&Lauer(2009) noted two nearby galaxies with simular photometric redshifts. which may indicate the xeseuce of such a cluster.," \citet{bor09} noted two nearby galaxies with similar photometric redshifts, which may indicate the presence of such a cluster."
+ However. a cluster is likely in both the 0.1-pc aud 5-kpe binary interpretations. as a inarv black-hole also requires anu earlier merger.," However, a cluster is likely in both the 0.1-pc and 5-kpc binary interpretations, as a binary black-hole also requires an earlier merger."
+ The oulv difference between the two interpretation ids the nuescale since the mereer occurred. which can be a raction of the Itubble time for a 0.1-pe binary.," The only difference between the two interpretation is the timescale since the merger occurred, which can be a fraction of the Hubble time for a 0.1-pc binary."
+ Ou such a timescale the iunediate environmenut may relax. but the larecr scale cluster environment evolves iudependcutly of the merger process.," On such a timescale the immediate environment may relax, but the larger scale cluster environment evolves independently of the merger process."
+ Follow-up radio observations at other frequencies aud epochs. as well as with higher angular resolutions. cau help distinguish between the O.l-pe aud 5-kpe binary interpretations.," Follow-up radio observations at other frequencies and epochs, as well as with higher angular resolutions, can help distinguish between the 0.1-pc and 5-kpc binary interpretations."
+" Specifically, demonstrating that cach radio source has a flat spectruin. is time-variable. aud relmaius compact at higher resolution would stronely support the 5-kpe binary interpretation."," Specifically, demonstrating that each radio source has a flat spectrum, is time-variable, and remains compact at higher resolution would strongly support the 5-kpc binary interpretation."
+ Iu contrast. in the Q-pe binary case. oue (or both) of the radio sources would be expected to be steady in time. have a steep spectrum. and begin to show an outflow-like structure when examined with higher resolution: a flat-spectrum source paired with a steep-spectrum source would sugeest a core-jet morpholoey.," In contrast, in the 0.1-pc binary case, one (or both) of the radio sources would be expected to be steady in time, have a steep spectrum, and begin to show an outflow-like structure when examined with higher resolution; a flat-spectrum source paired with a steep-spectrum source would suggest a core-jet morphology."
+ It the 5-kpe binary interpretation is correct. then SDSS J1536|OLLI may form the first known conrpact binary quasar of two hnuuinous broad-line svstenis. uulike the earlier N-rav-discovered compact svstenis where one or both of the AGNs are partly obscured. or broad emission lines are missing altogether.," If the 5-kpc binary interpretation is correct, then SDSS J1536+0441 may form the first known compact binary quasar of two luminous broad-line systems, unlike the earlier X-ray-discovered compact systems where one or both of the AGNs are partly obscured, or broad emission lines are missing altogether."
+ We thauk Shai Waspi for his help with the X-rav analysis., We thank Shai Kaspi for his help with the X-ray analysis.
+" We acknowledge useful feedback frou two anonvinous referees,", We acknowledge useful feedback from two anonymous referees.
+ This research was supported by THE ISRAEL SCIENCE. FOUNDATION (eraut ll07/08).aud by a eraut from the Norman and IHcleu Asher Space Research Tustitute.," This research was supported by THE ISRAEL SCIENCE FOUNDATION (grant 407/08),and by a grant from the Norman and Helen Asher Space Research Institute."
+"'The next steps in our analysis include combining multi-configuration and multi-band data for the G16.74-0.1 and G19.6—0.2 fields to introduce more constraints during image-reconstruction, to complete the software integration required to combine MS-MFS and W-Projection with the A-Projection algorithm to correct for time-varying primary-beams, and further tests with mosaicking observations and how the additional information aids the image reconstruction.","The next steps in our analysis include combining multi-configuration and multi-band data for the $+$ 0.1 and $-$ 0.2 fields to introduce more constraints during image-reconstruction, to complete the software integration required to combine MS-MFS and W-Projection with the A-Projection algorithm to correct for time-varying primary-beams, and further tests with mosaicking observations and how the additional information aids the image reconstruction."
+ These steps are in accordance with our long-term goal of doing a wide-band mosaiced survey of the Galactic plane at low frequencies., These steps are in accordance with our long-term goal of doing a wide-band mosaiced survey of the Galactic plane at low frequencies.
+" We used the ATNF Pulsar Catalogue (Manchesteretal.2005), accessible via the URL for the co-ordinates of the pulsar in the G55.74-3.4 field."," We used the ATNF Pulsar Catalogue \citep{PSRCat2005}, accessible via the URL for the co-ordinates of the pulsar in the $+$ 3.4 field."
+these values do not correspond to the detection limits. which are better defined using the peak count rate of the PSF in comparison with the background scatter in the skv anuulus.,"these values do not correspond to the detection limits, which are better defined using the peak count rate of the PSF in comparison with the background scatter in the sky annulus."
+ All objects for which [Iuxes are given in the tables are detected with a signal-to-noise-ratio of at least 4., All objects for which fluxes are given in the tables are detected with a signal-to-noise-ratio of at least 4.
+ The total errors quoted in Tables 1 are comprised of photometric errors aud svstemaltics., The total errors quoted in Tables \ref{targets_so} are comprised of photometric errors and systematics.
+ The latter ones are the combination of uncertainties in calibration (556)) and aperture correction.(1-254... from the IRAC! data handbook).," The latter ones are the combination of uncertainties in calibration ) and aperture correction, from the IRAC data handbook)."
+ No colour correction was carried out: instead we added à small contribution (0-2%)) to the error budget., No colour correction was carried out; instead we added a small contribution ) to the error budget.
+ The upper limits quoted in Tables 1. correspond to έτος times the photometric errors., The upper limits quoted in Tables \ref{targets_so} correspond to three times the photometric errors.
+ In channels 1 and 2. our fluxes are in excellent agreement. (within mimae) with the values reported recently by οἱal.(2007). for a subsample of our targets.," In channels 1 and 2, our fluxes are in excellent agreement (within mag) with the values reported recently by \citet{2007A&A...472L...9Z} for a subsample of our targets."
+ In channels 3 and 4. the magnitude differences between our work amd ZapateroOsorioetal.(2007) twpically remain within mmag. still within the lo errorbars.," In channels 3 and 4, the magnitude differences between our work and \citet{2007A&A...472L...9Z} typically remain within mag, still within the $\sigma$ errorbars."
+ We note. however. that in these two bands our f[Iuxes are in most cases substantially lower than the values given by ZapateroOsorioetal.(2007)..," We note, however, that in these two bands our fluxes are in most cases substantially lower than the values given by \citet{2007A&A...472L...9Z}."
+ We search for mid-inlrared. emission [rom dusty disks in our sample by comparing the measured IRAC colours with published values for (disk-less) field objects in the same Zur range (Pattenetal.2006).., We search for mid-infrared emission from dusty disks in our sample by comparing the measured IRAC colours with published values for (disk-less) field objects in the same $T_\mathrm{eff}$ range \citep{2006ApJ...651..502P}.
+ The main ciscriminators to distinguish between the photospheric contribution and the disk excess are the fluxes in the IRAC channels 3 and 4: at shorter wavelength the contrast between the photosphere ancl a possible disk is difficult to detect., The main discriminators to distinguish between the photospheric contribution and the disk excess are the fluxes in the IRAC channels 3 and 4; at shorter wavelength the contrast between the photosphere and a possible disk is difficult to detect.
+ In Fig., In Fig.
+ 3. we plot IRAC colours (1-H and I-I3) for all targets (except 770. see Sect. 4)).," \ref{f1} we plot IRAC colours (I1-I4 and I1-I3) for all targets (except 70, see Sect. \ref{sori70}) )."
+ As abscissa in these plots we use the 7—J colour. whieh can be assumed to be purely photometric ancl (hus represents a proxy [or Teg. which in (urn correlates with mass [or coeval objects.," As abscissa in these plots we use the $I-J$ colour, which can be assumed to be purely photometric and thus represents a proxy for $T_\mathrm{eff}$, which in turn correlates with mass for coeval objects."
+ The plot covers a mass range from ~20 to δρ.," The plot covers a mass range from $\sim 20$ to $\sim 8\,M_{\mathrm{Jup}}$."
+ IPMOSs are shown as crosses: (he solid lines are linear fits to the colours of the field objects and thus delineate the photospheric level in (hese figures., IPMOs are shown as crosses; the solid lines are linear fits to the colours of the field objects and thus delineate the photospheric level in these figures.
+ As can he seen in (his plot. the majority of the IPMOs have IRAC colours indistinguishable from the photospheric values and can thus be assumed to have either no disk or only small amounts of dust in the inner disk.," As can be seen in this plot, the majority of the IPMOs have IRAC colours indistinguishable from the photospheric values and can thus be assumed to have either no disk or only small amounts of dust in the inner disk."
+ Sources that lack excess in our data include 554. 5011555. 556 for which ZapateroOsorioetal.(2007) reported disk detections based on the shallow GTO IRAC images.," Sources that lack excess in our data include 54, 55, 56 for which \citet{2007A&A...472L...9Z} reported disk detections based on the shallow GTO IRAC images."
+ If we use the IRACS colour as disk criterion (upperpanel). we find (hat four objects | 771. 660. 666. 558 — show excess at the 36 level. ie. 4 out of 17 or 244-116..," If we use the IRAC3 colour as disk criterion (upperpanel), we find that four objects – 71, 60, 66, 58 – show excess at the $\sigma$ level, i.e. 4 out of 17 or $24\pm ^{13}_{11}$."
+ Since disk excesses may appear only bevond jin. the result [rom IRACS should be considered to be a lower limit.," Since disk excesses may appear only beyond $\,\mu m$ , the result from IRAC3 should be considered to be a lower limit."
+in the 30 fainter galaxies of the sample.,in the 30 fainter galaxies of the sample.
+ Considering now the high end mass. the probability law saturates to at M.—I0! implying that all galaxies above this level are expected M...to be seen at the depth of our radio images.," Considering now the high end mass, the probability law saturates to at $_\star \sim
+10^{11}M_{\sun}$, implying that all galaxies above this level are expected to be seen at the depth of our radio images."
+ This is not the case since we have three radio undetected. optically bright. galaxies.," This is not the case since we have three radio undetected, optically bright, galaxies."
+ Previous studies. characterized by far higher radio luminosity thresholds concur that a saturation of the detection rate at a ~ 30 level occurs at the high mass end (e.g..Mauch&Sadler2007).," Previous studies, characterized by far higher radio luminosity thresholds concur that a saturation of the detection rate at a $\sim$ 30 level occurs at the high mass end \citep[e.g.,][]{mauch07}."
+. Not surprisingly. expanding the radio luminosity coverage downward by 4 orders of magnitude we reach a higher detection rate. ~ 70σοι.," Not surprisingly, expanding the radio luminosity coverage downward by 4 orders of magnitude we reach a higher detection rate, $\sim$ 70."
+ Nonetheless. we still have three bright galaxies without a radio source associated with them.," Nonetheless, we still have three bright galaxies without a radio source associated with them."
+ These objects are 30000 times fainter than the radio core of M 87 (the brightest, These objects are $\gtrsim 30000$ times fainter than the radio core of M 87 (the brightest
+ These objects are 30000 times fainter than the radio core of M 87 (the brightest=, These objects are $\gtrsim 30000$ times fainter than the radio core of M 87 (the brightest
+far from the observations since they predict the population of pulsars in binaries with spin periods below 20ms.,far from the observations since they predict the population of pulsars in binaries with spin periods below $20$ ms.
+ In ? the authors have calculated the expected masses of neutron stars observed in gravitational waves using the StarTrack population synthesis code (Belczyfsski et al., In \citet{2004MNRAS.352.1372B} the authors have calculated the expected masses of neutron stars observed in gravitational waves using the StarTrack population synthesis code (Belczyńsski et al.
+ 2002 for detailed description)., 2002 for detailed description).
+ The assumed value for the minimum mass and maximum mass of a neutron star was 1.2Me and 3.0Mo respectively.," The assumed value for the minimum mass and maximum mass of a neutron star was $1.2 \ M_{\odot}$ and $ 3.0
+\ M_{\odot}$ respectively."
+ They have verified the results taking the maximum mass to be 2Mo and 2.5Mo (?).., They have verified the results taking the maximum mass to be $2 \ M_{\odot}$ and $2.5 \ M_{\odot}$ \citep{2005MmSAI..76..632G}.
+" They have shown that the distributions of mass ratio q (defined as the ratio of less to more massive component) of neutron star binaries observed in gravitational waves, have two peaks: the first one for nearly equal mass systems with two masses close to 1.4 Mo; and the second one with small mass ratio q~0.6—0.7 with binaries consisting of a neutron star with gravitational mass close to the maximum mass with relatively smaller companion."," They have shown that the distributions of mass ratio $q$ (defined as the ratio of less to more massive component) of neutron star binaries observed in gravitational waves, have two peaks: the first one for nearly equal mass systems with two masses close to $ 1.4 \ M_{\odot}$ ; and the second one with small mass ratio $q\sim 0.6- 0.7$ with binaries consisting of a neutron star with gravitational mass close to the maximum mass with relatively smaller companion."
+ The strength of the second peak depends on the value of the maximum mass of a neutron star., The strength of the second peak depends on the value of the maximum mass of a neutron star.
+ For the models with the higher maximal mass the sample starts to be dominated by binariesbelonging to, For the models with the higher maximal mass the sample starts to be dominated by binariesbelonging to
+occur and play an important role iu gas dyvnauics.,occur and play an important role in gas dynamics.
+ The differential Euler equations are ill-defined at shock iscontiuuities where derivatives are infinite., The differential Euler equations are ill-defined at shock discontinuities where derivatives are infinite.
+ Much ffort has been devoted to solving this problem aud a field of work has resulted frou it., Much effort has been devoted to solving this problem and a field of work has resulted from it.
+ Computational fliid vuaiics (CED) is a powerful approach to simulating Hid flow with emphasis ou high resolution capturing Mf shocks aud prevention of nuuerical iustabilitics., Computational fluid dynamics (CFD) is a powerful approach to simulating fluid flow with emphasis on high resolution capturing of shocks and prevention of numerical instabilities.
+ Both Eulerian and Lagrangian methods have beeu eveloped., Both Eulerian and Lagrangian methods have been developed.
+ Lagrangian methods based on smoothed particle wdrodvuamics (SPI:Cangold& Monaghanu1977:Lucy1977) consider à \oute-Carlo approximation to solving the fluid equations. somewhat analogous to JV- )ody: inethods for the Vlasov equation.," Lagrangian methods based on smoothed particle hydrodynamics \citep[SPH;][]{gin77,lucy77} consider a Monte-Carlo approximation to solving the fluid equations, somewhat analogous to $N$ -body methods for the Vlasov equation."
+ SPIT schemes ollow the trajectories of particles of fixed mass which represent fluid elements., SPH schemes follow the trajectories of particles of fixed mass which represent fluid elements.
+ The Lagrangian forms of the Euler equations are solved to determine smoothed Huicl variables like density. velocity. auc temperature.," The Lagrangian forms of the Euler equations are solved to determine smoothed fluid variables like density, velocity, and temperature."
+ The particle formulation does uot naturally capture shocks and artificial viscosity is added to preven uuphnvsical oscillations., The particle formulation does not naturally capture shocks and artificial viscosity is added to prevent unphysical oscillations.
+ Wowever. the addition of artificial viscosity broadeus shocks over several πιοhing leugths aud deeraces the resolution.," However, the addition of artificial viscosity broadens shocks over several smoothing lengths and degrades the resolution."
+ The Laeraneian approach has a large dynamic range in leneth but uo in dass., The Lagrangian approach has a large dynamic range in length but not in mass.
+ [t achieves eood spatial resolution iu lieh density regions but does poorly in low deusity regions., It achieves good spatial resolution in high density regions but does poorly in low density regions.
+ SPII schemes must smooth over a large number of neighbouring particles. making it computationally expensive aud challenging to implement in parallel.," SPH schemes must smooth over a large number of neighbouring particles, making it computationally expensive and challenging to implement in parallel."
+ The standard approach to Eulerian methods is to discretize the problemi and solve the iutegral Euler equations on a Cartesian erid bw cxuputing the flux of nass. momentum. and cnerev across ericl cell bounaries.," The standard approach to Eulerian methods is to discretize the problem and solve the integral Euler equations on a Cartesian grid by computing the flux of mass, momentum, and energy across grid cell boundaries."
+ In couscrvative schemes. the fux taken out of one cell is added to the uciglibouriug cell and this eusures the correct shock propagation.," In conservative schemes, the flux taken out of one cell is added to the neighbouring cell and this ensures the correct shock propagation."
+ Flux assiguimnent schemes based on the (Marten1983) have been designed for lieh order accuracy aud high resolution capturing of shocks. while preveuting unphysical oscillations.," Flux assignment schemes based on the \citep{har83} have been designed for high order accuracy and high resolution capturing of shocks, while preventing unphysical oscillations."
+ The Eulerian approach has a laree dvnamic rauge iu mass but not in leneth. opposite to that of Lagraugian schemes.," The Eulerian approach has a large dynamic range in mass but not in length, opposite to that of Lagrangian schemes."
+ In general. Envdevian algorithms are computationally faster bv several orders of maguitude.," In general, Eulerian algorithms are computationally faster by several orders of magnitude."
+ They are also easv to nuplement aud to parallelize., They are also easy to implement and to parallelize.
+ The purpose of this paper is to preseut a pedagogical review of some of the methods employed im Eulerian conrputational fluid dyuzauics., The purpose of this paper is to present a pedagogical review of some of the methods employed in Eulerian computational fluid dynamics.
+ In 52 we briefly review he Euler equations and discuss the standard approach ο discretizing couscrvation laws., In $\S2$ we briefly review the Euler equations and discuss the standard approach to discretizing conservation laws.
+ We describe traditional ceutral differencing methods such as the Las-Wendroft scheme in 52 and more modern flux assiguineut uethods like the TWD scheme iu 51., We describe traditional central differencing methods such as the Lax-Wendroff scheme in $\S3$ and more modern flux assignment methods like the TVD scheme in $\S4$.
+ Iu 55 we review the relaxing TWD method for svstenis of conservation laws like the Euler equations. which las been successfully naplenmiented for simulating cosinological astrophysical fluids by Pen(1998).," In $\S5$ we review the relaxing TVD method for systems of conservation laws like the Euler equations, which has been successfully implemented for simulating cosmological astrophysical fluids by \citet{pen98}."
+. In 66 we apply a sclferavitating lvdro code to simulatiug he formation of blic strageler stars through stellar nerecrs., In $\S6$ we apply a self-gravitating hydro code to simulating the formation of blue straggler stars through stellar mergers.
+ A sample 3-D relaxing TVD code is provided in the appendix., A sample 3-D relaxing TVD code is provided in the appendix.
+ The Euler equatious which govern lvdrodwuamics are asystei of couservation laws for ass. momenta. and energv.," The Euler equations which govern hydrodynamics are a system of conservation laws for mass, momentum, and energy."
+ Iu differential conservation form. the continuity equation. momentum equation. and euergv equation are elven as: We have omitted eravitational and other source terns like heating aud cooliug.," In differential conservation form, the continuity equation, momentum equation, and energy equation are given as: We have omitted gravitational and other source terms like heating and cooling."
+ The plivsical state of the fluid is specified by its density p. velocity field v. and total energy density. Tn practice. the thermal enerev 2 is evaluated by subtracting the kinetic euergv from the total cucrey.," The physical state of the fluid is specified by its density $\rho$ , velocity field $\boldsymbol{v}$, and total energy density, In practice, the thermal energy $\varepsilon$ is evaluated by subtracting the kinetic energy from the total energy."
+ For an ideal gas. the pressure P(2) is related to the thermal enerev by the equation of state. where 5 is the ratio of specific heats.," For an ideal gas, the pressure $P(\varepsilon)$ is related to the thermal energy by the equation of state, where $\gamma$ is the ratio of specific heats."
+ Another thermodvuaimic variable which is of iuportanceis the sound speed ος Which is given by, Another thermodynamic variable which is of importanceis the sound speed $c_s$ which is given by
+(seeTurner2010).,\citep[see][]{tu10}.
+. The augulur brackets denote intensity iueans., The angular brackets denote intensity means.
+ Spectroscopic observations are essential for confirming the picture implied by the 2MLASS observations. since star counts for stars identified photometrically as Ὁ chwarts display ouly a iuareinal concentration towards the cluster core.," Spectroscopic observations are essential for confirming the picture implied by the 2MASS observations, since star counts for stars identified photometrically as B dwarfs display only a marginal concentration towards the cluster core."
+ It appears that the main concentration of stars in Turner 1 corresponds to the old. sparse. foreground cluster identified in Fig. 6..," It appears that the main concentration of stars in Turner 1 corresponds to the old, sparse, foreground cluster identified in Fig. \ref{fig6},"
+ with contaninatioun bv a vouneg group of backeround D-dwarfs located aloug the same line of sight., with contamination by a young group of background B-dwarfs located along the same line of sight.
+ Collinder 119 is à sparse group of stars (Fig. 8)), Collinder 419 is a sparse group of stars (Fig. \ref{fig8}) )
+ of sinall aueular diameter (17.5accordingtoCollinder1931) associated with the massive O9 Ví((n)) double star TID 193322 (Walborn1971)., of small angular diameter \citep[$4\arcmin$.5 according to][]{co31} associated with the massive O9 V((n)) double star HD 193322 \citep{wa71}.
+. Collinder estimated a distance of 1170 pe for the eroup from its aneular diameter. while Cüesetal.(2010) used astrometric and photometric data from the UCACS catalogue for bright stars in the field to derive a distance of ¢=711+436 pc. with a reddening ofV) =0.37£0.05.," Collinder estimated a distance of 1470 pc for the group from its angular diameter, while \citet{gi10} used astrometric and photometric data from the UCAC3 catalogue for bright stars in the field to derive a distance of $d = 741 \pm36$ pc, with a reddening of $= 0.37 \pm0.05$."
+ Iucluded as a possible member of the eroup was the 1ο TII star IRAS 20161|1035., Included as a possible member of the group was the M3 III star IRAS 20161+4035.
+ The cluster does not stand out well at optical wavelengths. but we were able to detect a sleht density euliaucemeut visually aud estimate a crude center of svinmetry at 2000 co-ordinates: 20:18:08. |LO:12:12. close to the coordinates estimated by Collinder.," The cluster does not stand out well at optical wavelengths, but we were able to detect a slight density enhancement visually and estimate a crude center of symmetry at 2000 co-ordinates: 20:18:08, +40:42:42, close to the coordinates estimated by Collinder."
+ The cluster also appears as a set of reddened B-type stars in available 2MÀSS ΠΛ photometry for stars ling within 5! of that counter of svunuetiy (Fie. 9)).," The cluster also appears as a set of reddened B-type stars in available 2MASS $_s$ photometry for stars lying within $5\arcmin$ of that center of symmetry (Fig. \ref{fig9}) ),"
+ reasonably consistent with Collinder’s inferred dimeusions., reasonably consistent with Collinder's inferred dimensions.
+" A best fit by eve to the JH, data vields a reddening of =0.35940.02 = L.3240.07). significantly lavecr than the value estimated by Cüesoetal. (2010).. a distance modulus of Aly=11.740.2 (Vy My=10.75 40.28). and an interred distance of d=1.12x0.18 kpc. which ds also larecr than the (οςetal.(2010) value. although consistent with a location within the voung complex of stars and IT IT regions GucludingBerkeley87.Turuceretassociated with the Cveuus005 XN region[m] at"," A best fit by eye to the $_s$ data yields a reddening of $= 0.39 \pm0.02$ $= 1.32 \pm0.07$ ), significantly larger than the value estimated by \citet{gi10}, , a distance modulus of $_J = 11.7 \pm0.2$ $_0$ $_V = 10.75 \pm0.28$ ), and an inferred distance of $d = 1.42 \pm0.18$ kpc, which is also larger than the \citet{gi10} value, although consistent with a location within the young complex of stars and H II regions \citep[including Berkeley 87,][]{te10} associated with the Cygnus X region at"
+of view aud passing in frout of the cluster as shown iu Fie. L.,of view and passing in front of the cluster as shown in Fig. \ref{gal_x_dist}.
+ Its morphology sugeestsOO that this structure may be a filament perpendicular to the line of sight. with a projected overdensity of galaxies on top of the northeru substructure of ((sce section 8 refealgsf)," Its morphology suggests that this structure may be a filament perpendicular to the line of sight, with a projected overdensity of galaxies on top of the northern substructure of (see section \\ref{gal_dist}) )."
+hThepresenecofthisforegroundstractureshouldbetaken regems," The presence of this foreground structure should be taken into account when estimating the mass of the cluster through the weak lensing technique, which is sensitive to the distribution of matter along the line of sight."
+sionscemstobeassociatedtothisstrueturc.iuidicatingthatth," No X-ray emission seems to be associated to this structure, indicating that the galaxies in this group do not have a hot gas component."
+the reproduction of the observed correlation between the velocity aud ionization potential of ultraviolet absorption features towards the central star (Suuth 198 1)).,the reproduction of the observed correlation between the velocity and ionization potential of ultraviolet absorption features towards the central star (Smith \cite{SPDH1984}) ).
+ Our solutions also often reproduce this correlation seeFie., Our solutions also often reproduce this correlation seeFig.
+ 5. which shows the velocity aud temperature decreasing tosether for the heavily mass-loaded solution). as do those of Tanami Sakashita (1987)).," \ref{fig:lam-2_MLcomp} which shows the velocity and temperature decreasing together for the heavily mass-loaded solution), as do those of Hanami Sakashita \cite{HS1987}) )."
+" Interestinely, solutions with mass-loading from conductive evaporation of the swept-up shell do not reproduce this correlation (see Weaver 1977))."," Interestingly, solutions with mass-loading from conductive evaporation of the swept-up shell do not reproduce this correlation (see Weaver \cite{WMCSM1977}) )."
+ We have investigated the evolution of a mass-loaded wind blown bubble with a coustaut rate of injection of mechanical energv from a central wind source., We have investigated the evolution of a mass-loaded wind blown bubble with a constant rate of injection of mechanical energy from a central wind source.
+ The mass loading occurs due to the lycwodvuamic ablation of distributed clumps. aud is “switched-ou at a specified radius. interior to which the wind expands freely.," The mass loading occurs due to the hydrodynamic ablation of distributed clumps, and is `switched-on' at a specified radius, interior to which the wind expands freely."
+" The requirement that the solution be sclf-similar imposes a lnk between the racial variation of the intercluuip deusity (px i7) and the rate of mass loading from the chuups (px ry which forces A2(24.5/3,", The requirement that the solution be self-similar imposes a link between the radial variation of the interclump density $\rho \propto r^{\beta}$ ) and the rate of mass loading from the clumps $\dot{\rho} \propto r^{\lambda}$ ) which forces $\lambda = (2\beta - 5)/3$.
+ We first produced solutions with negligiblelass loading and»=2 (which correspond to coustaut HA and μι Which we compared with results obtained x Dyson (1973)).," We first produced solutions with negligiblemass loading and $\beta=-2$ (which correspond to constant $\dot{M}_{wind}$ and $v_{wind}$ ), which we compared with results obtained by Dyson \cite{D1973}) )."
+ Excellent agreement for the structure οἳ the bubble was found., Excellent agreement for the structure of the bubble was found.
+We also confirmed that for neglieible mass loading. the value of αμ had no effect on the resulting solutions. as desired.,"We also confirmed that for negligible mass loading, the value of $x_{ml}$ had no effect on the resulting solutions, as desired."
+ We then investigated thechanges iu the structure of he bubble for differcut values of A. o. ος reas aud iyu.," We then investigated thechanges in the structure of the bubble for different values of $\lambda$ , $\phi$ , $x_{\rm is}/x_{\rm cd}$ , and $x_{\rm ml}$ ."
+ The ceutral conclusions are:, The central conclusions are:
+"the middle row we show an example of the same kind of substructure, but here the tidal tails are more pronounced and the stream passes within 8 kpc of halo centre.","the middle row we show an example of the same kind of substructure, but here the tidal tails are more pronounced and the stream passes within $8$ kpc of halo centre."
+ The bottom two rows show two typical stream-like structures which have no associated self-bound subhalo., The bottom two rows show two typical stream-like structures which have no associated self-bound subhalo.
+ Their masses are similar to those of the biggest subhaloes., Their masses are similar to those of the biggest subhaloes.
+kkpe.,$^{-1}$.
+ In addition to the large scale bar. Emsellem et al.," In addition to the large scale bar, Emsellem et al."
+ postulate a small nuclear bar (~ 37)., \nocite{2003MNRAS.345.1297E} postulate a small nuclear bar $\sim 3^{\prime \prime}$ ).
+ The eeas is more suitable for measuring the elongation of the »otential. since the cold gas has a small velocity dispersion (tvpical values <19 13) and is on nearly circular orbits.," The gas is more suitable for measuring the elongation of the potential, since the cold gas has a small velocity dispersion (typical values $< 10 $ ) and is on nearly circular orbits."
+ Taking the mean value of the elongation as obtained rom the field. we find οσα=0.016+0.022.," Taking the mean value of the elongation as obtained from the field, we find $\epsilon_{\mathrm{pot}}\sin 2 \varphi = 0.016 \pm 0.022$."
+ We conclude hat the potential of NGC 2974 is well approximated: by an axisvmnmetric one., We conclude that the potential of NGC 2974 is well approximated by an axisymmetric one.
+ To find the rotation curve of NGC 2974. we subtract the svstemic velocities [rom the ionised and neutral gas velocity [fields separately.," To find the rotation curve of NGC 2974, we subtract the systemic velocities from the ionised and neutral gas velocity fields separately."
+ Next. we fix P—47 and q=0.50 (or equivalently 7= 607) of the ellipses to the mean values obtained from the neutral gas. and rerun kinemetry on both the velocity maps. now forcing the position angle and Ilattening to be the same everywhere in the gas disc.," Next, we fix $\Gamma = 47^\circ$ and $q=0.50$ (or equivalently $i=60^\circ$ ) of the ellipses to the mean values obtained from the neutral gas, and rerun kinemetry on both the velocity maps, now forcing the position angle and flattening to be the same everywhere in the gas disc."
+ Also. because velocity is an odd moment. the even terms in the harmonic expansion should be zero. and are not taken into account during the fit (see INrajnovié et al. 2006)).," Also, because velocity is an odd moment, the even terms in the harmonic expansion should be zero, and are not taken into account during the fit (see Krajnović et al. \nocite{2006MNRAS.366..787K}) )."
+ The rotation curve of the ionised eas is shown in Figure 6 (open cliamonels)., The rotation curve of the ionised gas is shown in Figure \ref{fig:fit_vphi} (open diamonds).
+ The ionised gas has a high observed velocity. dispersion Tong. exceeding 250 iin the centre of the galaxy.," The ionised gas has a high observed velocity dispersion $\sigma_{\mathrm{obs}}$, exceeding 250 in the centre of the galaxy."
+ Three phenomena can contribute to the observed velocity dispersion of gas: thermal motions. turbulence and gravitational interactions:a The thermal velocity clispersion is always present. and caused by the thermal energy of the gas molecules: where & is Boltzmann's constant. Z/ the temperature of the eas and m the typical mass of a eas particle.," Three phenomena can contribute to the observed velocity dispersion of a gas: thermal motions, turbulence and gravitational interactions: The thermal velocity dispersion is always present, and caused by the thermal energy of the gas molecules: where $k$ is Boltzmann's constant, $T$ the temperature of the gas and $m$ the typical mass of a gas particle."
+ The contribution of σωμα to the total velocity distribution in ionised eas is small: a typical temperature for ionised. gas is 103 Ix. which implies Mera~LO kms. ‘Turbulence can be caused. by e.g. internal motions within the gas clouds orshocks induced. by a non-axisvmmetric perturbation to the potential. such as a bar.," The contribution of $\sigma_{\mathrm{thermal}}$ to the total velocity distribution in ionised gas is small: a typical temperature for ionised gas is $10^4$ K, which implies $ \sigma_{\mathrm{thermal}} \sim 10$ km/s. Turbulence can be caused by e.g. internal motions within the gas clouds orshocks induced by a non-axisymmetric perturbation to the potential, such as a bar."
+ ‘This increases the dispersion. but has a negligible ellect on the circular. velocity of the gas.," This increases the dispersion, but has a negligible effect on the circular velocity of the gas."
+ In. contrast. gravitational interactions of individual gas clouds not only increase random motions of the cloucs and therefore their dispersion. but also lower the observed. velocity.," In contrast, gravitational interactions of individual gas clouds not only increase random motions of the clouds and therefore their dispersion, but also lower the observed velocity."
+ To correct for this last ellect. we need to apply an asymmetric drift correction to recover the true circular velocity.," To correct for this last effect, we need to apply an asymmetric drift correction to recover the true circular velocity."
+ Unfortunately. it is not possible to determine which fraction of the high velocity dispersion in the ionised eas is caused by turbulence. and which by gravitational interactions.," Unfortunately, it is not possible to determine which fraction of the high velocity dispersion in the ionised gas is caused by turbulence and which by gravitational interactions."
+ We therefore now first investigate the effect of asvnnmetric drift on the rotation curve of the ionisecl gas., We therefore now first investigate the effect of asymmetric drift on the rotation curve of the ionised gas.
+ Due to gravitational interactions of gas clouds on circular orbits. the observed velocity is lower than the circular velocity connected to the gravitational potential.," Due to gravitational interactions of gas clouds on circular orbits, the observed velocity is lower than the circular velocity connected to the gravitational potential."
+ Since we are interested in the mass distribution of NGC 2974. we need to trace the potential. and therefore we have to Increase our observed. velocity with an asymmetric drift correction. to obtain the true circular velocity.," Since we are interested in the mass distribution of NGC 2974, we need to trace the potential, and therefore we have to increase our observed velocity with an asymmetric drift correction, to obtain the true circular velocity."
+ We follow the formalism described. in Appendix A. which is based on the Jeans equations and the higher order velocity. moments of the collisionless Boltzmann equation.," We follow the formalism described in Appendix \ref{sec:adc}, which is based on the Jeans equations and the higher order velocity moments of the collisionless Boltzmann equation."
+ We assume that the galaxy. is axisvmmetric. which is a valid approach given the low elongation of the potential that we derived in section 3.2..," We assume that the galaxy is axisymmetric, which is a valid approach given the low elongation of the potential that we derived in section \ref{sec:elong}."
+ Further we assume that the eas lies in a thin disc., Further we assume that the gas lies in a thin disc.
+ We fit the prescription that Evans de Zeeuw (1094)) used. for their power-law models. to the rotation curve extracted [rom the ionised. gas. where V4 is the rotation velocity at large rach. and we introduce with 2. the core radius of the model.," We fit the prescription that Evans de Zeeuw \nocite{1994MNRAS.271..202E}) ) used for their power-law models to the rotation curve extracted from the ionised gas, where $V_\infty$ is the rotation velocity at large radii, and we introduce with $R_c$ the core radius of the model."
+ This is Equation CX14)) evaluated in the plane of the dise (2= 0). with a Lat rotation curve at large radii (2= 0).," This is Equation \ref{eq:powerlaw}) ) evaluated in the plane of the disc $z =
+0$ ), with a flat rotation curve at large radii $\beta = 0$ )."
+ Since we observe the gas only in the equatorial plane of the galaxy. we cannot constrain the Dlattening of the potential ga.," Since we observe the gas only in the equatorial plane of the galaxy, we cannot constrain the flattening of the potential $q_\Phi$."
+ We therefore. assumed. a spherical potential qa=1. which is not a bad approximation even if the density distribution is Ilattened. since the dependence on dq is weak.," We therefore assumed a spherical potential $q_\Phi = 1$, which is not a bad approximation even if the density distribution is flattened, since the dependence on $q_\Phi$ is weak."
+ Moreover. even though the density distribution. of most galaxies. is clearly Dattened. the potential is in. general. significantly rounder than the densitv.," Moreover, even though the density distribution of most galaxies is clearly flattened, the potential is in general significantly rounder than the density."
+ For example. an axisvmmoetric logarithmic potential is only about a third as [attened as the corresponding density. distribution (e.g. £22.2.2 of Binney ‘Tremaine 19873).," For example, an axisymmetric logarithmic potential is only about a third as flattened as the corresponding density distribution (e.g. 2.2.2 of Binney Tremaine \nocite{1987BT}) )."
+ lo be able to fit the observed. velocity we need. to, To be able to fit the observed velocity we need to
+of the total (see reffig:tempec).,of the total (see \\ref{fig:temp}c c).
+" Cases of vanishing initial pressure (4,<1) will have a fraction (5+1)sin!(6/2) of the released energy thermalized.", Cases of vanishing initial pressure $\beta_e\ll1$ ) will have a fraction $(\gamma+1)\sin^4(\zeta/2)$ of the released energy thermalized.
+" Compression work done on larger initial pressure will raise this fraction slightly, but all cases with 24<1/3 thermalize less than 10*4. of the released energy."," Compression work done on larger initial pressure will raise this fraction slightly, but all cases with $\beta_3<1/3$ thermalize less than $10\%$ of the released energy."
+" The foregoing illustrates, through a simplified analytic model, a process we believe must be common in the flaring corona."," The foregoing illustrates, through a simplified analytic model, a process we believe must be common in the flaring corona."
+" Localized reconnection does not directly dissipate magnetic energy, but rather initiates its release through subsequent shortening of field lines."," Localized reconnection does not directly dissipate magnetic energy, but rather initiates its release through subsequent shortening of field lines."
+" This shortening propagates from the reconnection site at the Alfvénn speed, converting energy from magnetic into kinetic form."," This shortening propagates from the reconnection site at the Alfvénn speed, converting energy from magnetic into kinetic form."
+" The shortening flux tubes compress plasma within them, raising its temperature as they do so."," The shortening flux tubes compress plasma within them, raising its temperature as they do so."
+ Due to the supersonic (Alfvénnic) flows this compression occurs at strong shocks. which raise the temperature far beyond that from adiabatic compression.," Due to the supersonic (Alfvénnic) flows this compression occurs at strong shocks, which raise the temperature far beyond that from adiabatic compression."
+" Both the GDS in our model and the SMS in Petsheek's model result, ultimately, from shortening or weakening of magnetic field lines."," Both the GDS in our model and the SMS in Petshcek's model result, ultimately, from shortening or weakening of magnetic field lines."
+" The two-dimensional Petschek model permits shortening only perpendicular to the symmetry direction, accompanied by weakening."," The two-dimensional Petschek model permits shortening only perpendicular to the symmetry direction, accompanied by weakening."
+ The shock is therefore a SMS whose normal is mostly perpendicular to the sheet., The shock is therefore a SMS whose normal is mostly perpendicular to the sheet.
+" After three-dimensional patchy reconnection, on the other hand, field lines also shorten in the erstwhile symmetry direction which is the orientation of the shock normal (CX)."," After three-dimensional patchy reconnection, on the other hand, field lines also shorten in the erstwhile symmetry direction which is the orientation of the shock normal $\pm\xhat$ )."
+" The GDSs are disconnected from the diffusion region, and are instead features of the ideal relaxation reconnection."," The GDSs are disconnected from the diffusion region, and are instead features of the ideal relaxation reconnection."
+" This new scenario requires two-step energy conversion, from magnetic to kinetic to thermal energy, in contrast to the SMSs which thermalize magnetic energy directly where they weaken the field."," This new scenario requires two-step energy conversion, from magnetic to kinetic to thermal energy, in contrast to the SMSs which thermalize magnetic energy directly where they weaken the field."
+ This work was supported by NASA and the NSF., This work was supported by NASA and the NSF.
+the telescope positioning.,the telescope positioning.
+ We conclude that there is no significant periodicity between 0.05 and 0.5 Hz., We conclude that there is no significant periodicity between 0.05 and 0.5 Hz.
+" A Lomb-Scargle periodogram of the full dataset is computationally too expensive, instead we calculated the Fast Fourier transform (FFT) of the full lightcurve sampled at 100 Hz (Fig. 9))."," A Lomb-Scargle periodogram of the full dataset is computationally too expensive, instead we calculated the Fast Fourier transform (FFT) of the full lightcurve sampled at 100 Hz (Fig. \ref{AATau_periodloglog}) )."
+" The largest peak seen is again the systematic signal at 0.038 Hz (as already discussed) and we can exclude significant periodicity down to 50 Hz, which is the Nyquist frequency for a sampling rate of 100 Hz."," The largest peak seen is again the systematic signal at 0.038 Hz (as already discussed) and we can exclude significant periodicity down to 50 Hz, which is the Nyquist frequency for a sampling rate of 100 Hz."
+ If periodic oscillations exist in the accretion shocks of CTTS their frequency may be unstable., If periodic oscillations exist in the accretion shocks of CTTS their frequency may be unstable.
+ Therefore we performed a wavelet analysis of the lightcurves using the IDL functionwv_applet., Therefore we performed a wavelet analysis of the lightcurves using the IDL function.
+". Evenly sampled lightcurves are required as input, thus missing data points were filled with repeated values of the last valid flux measurement."," Evenly sampled lightcurves are required as input, thus missing data points were filled with repeated values of the last valid flux measurement."
+" We employed different wavelet shapes, concentrating on Morlets of order 4-6."," We employed different wavelet shapes, concentrating on Morlets of order 4-6."
+" For TW Hya no signal with a significance above is found in the wavelet analysis, while AA Tau shows a peak in the power spectrum after about 500 s, which is very likely caused by the dip in the lightcurve (Fig. 3))."," For TW Hya no signal with a significance above is found in the wavelet analysis, while AA Tau shows a peak in the power spectrum after about 500 s, which is very likely caused by the dip in the lightcurve (Fig. \ref{AATau_lc}) )."
+ We showed above that this feature is related to changing seeing conditions., We showed above that this feature is related to changing seeing conditions.
+" In TW Hya we do not detect periodic signals on the confidence level, that is power level Py=17.4."," In TW Hya we do not detect periodic signals on the confidence level, that is power level $P_N= 17.4$."
+" Knowing the variance of the noise c2, we now invert eqn.", Knowing the variance of the noise $\sigma_N^2$ we now invert eqn.
+" 4 to calculate the minimum amplitude A,, which should have been detected at that level, if present: For the parameters of both lightcurves from TW Hya Amn£F:100n."," \ref{eqn_pn} to calculate the minimum amplitude $A_m$ which should have been detected at that level, if present: For the parameters of both lightcurves from TW Hya $A_m\approx10 \sigma_N$."
+" The total noise level in the normalised lightcurves is about 107? and 5x107? in the 340 nm and 380 nm observations of TW Hya, therefore the confidence limit on Aj is 107? and 5x1074, respectively for the presence of a single, sinusoidal mode in the frequency range 0.02-3 Hz in the 340 nm band and 3 Hz in the 380 nm band."," The total noise level in the normalised lightcurves is about $10^{-2}$ and $5\times 10^{-3}$ in the 340 nm and 380 nm observations of TW Hya, therefore the confidence limit on $A_m$ is $10^{-3}$ and $5\times 10^{-4}$, respectively for the presence of a single, sinusoidal mode in the frequency range 0.02-3 Hz in the 340 nm band and 0.1-3 Hz in the 380 nm band."
+ We performed Monte-Carlo simulations to estimate the detection efficiency for a combination of many periods., We performed Monte-Carlo simulations to estimate the detection efficiency for a combination of many periods.
+" We used the time bins from the observations, keeping the irregularities introduced by the outlier rejection and simulated 1000 lightcurves, calculated Lomb-Scargle periodograms for each of them, and determined the largest peak in the power spectrum."," We used the time bins from the observations, keeping the irregularities introduced by the outlier rejection and simulated 1000 lightcurves, calculated Lomb-Scargle periodograms for each of them, and determined the largest peak in the power spectrum."
+" For random measurement uncertainties of1%,, which corresponds to the observations of TW Hya in the 340 nm filter, we distributed multiple sinusoidal waves with random frequencies between 0.02 and 3 Hz and added them to a constant component."," For random measurement uncertainties of, which corresponds to the observations of TW Hya in the 340 nm filter, we distributed multiple sinusoidal waves with random frequencies between 0.02 and 3 Hz and added them to a constant component."
+ In the case of 40 different, In the case of 40 different
+proto-cluster environments and are believed to be the ancestors of early-type central-cluster ealaxies (e.g.Veneniuisetal.2007).,proto-cluster environments and are believed to be the ancestors of early-type central-cluster galaxies \citep[e.g.][]{ven07}.
+. TizRCs are in an active stage of their evolution. with powerful radio jets cluanating from. a massive ceutral black-hole.," HzRGs are in an active stage of their evolution, with powerful radio jets emanating from a massive central black-hole."
+ These radio jets vigorously interact with their =rrounding environment (c.g.Ihuuphnrevetal.2006) aud πεrye as beacols for tracing the faint ost galaxy aud ΕΠΣ proto-cluster., These radio jets vigorously interact with their surrounding environment \citep[e.g.][]{hum06} and serve as beacons for tracing the faint host galaxy and surrounding proto-cluster.
+ This makes Πως amoung the best studied high-: objecs and ideal aboratories for investigating both f1C ornatiolji and evolution of ealaxies/clusters as well as fje. relationship beween carly star formation aud AGN activity., This makes HzRGs among the best studied $z$ objects and ideal laboratories for investigating both the formation and evolution of galaxies/clusters as well as the relationship between early star formation and AGN activity.
+ A crucial comoeat i ithe evohtionary studies of TzRCs aud their associated star formation processes IS a detailec lkuowlecexορ about the content and proorties of the 120lecular gas., A crucial component in the evolutionary studies of HzRGs and their associated star formation processes is a detailed knowledge about the content and properties of the molecular gas.
+ The most abundant molecule. molecular hwdrosgeu orIL... has stronglv forbiden rotational transitions ancl is virtually iivisible. except when it is shocked or heated to high temperatures.," The most abundant molecule, molecular hydrogen or, has strongly forbidden rotational transitions and is virtually invisible, except when it is shocked or heated to high temperatures."
+ An exceleut racer for eeas Is carbon nionoxile Or CO. which cuits strong rotational transition lines that occur prinarily hrough collisious withΠο. even at relatively low densities (e.g.Solomon&VandenBout2005).," An excellent tracer for gas is carbon monoxide or CO, which emits strong rotational transition lines that occur primarily through collisions with, even at relatively low densities \citep[e.g.][]{sol05}."
+. Brown&VandeBout(1901) «xvere the first to observe CO (referred to as CO in this paper) vevoud 2=2., \citet{bro91} were the first to observe $^{12}$ CO (referred to as CO in this paper) beyond $z=2$.
+ During that saue decade. dtensive searches ‘ailed to detec CO in IIZRCs (Evansetal.1996:vanOjikctal. 1997a).," During that same decade, intensive searches failed to detect CO in HzRGs \citep{eva96,oji97}."
+. Since tlie1. CO emission jas been detected 3i individual stiilies of powerful radio galaxies between ~25 (scereviewsby2008.audSect.L2 for iudividialreferences).," Since then, CO emission has been detected in individual studies of powerful radio galaxies between $z \sim 2-5$ \citep[see reviews by][and Sect. \ref{sec:discussion2} for individual."
+ Iu sone cases the CO has been resolved ou scales of several teus of spe (e.g.Papadopoulosetal.2000)., In some cases the CO has been resolved on scales of several tens of kpc \citep[e.g.][]{pap00}.
+. Despite these results. two iuajor limitations ο. svsteimatic searches for CO in HzRGs have )oen the very iuited velocity coverage of existing unrspectromieters (often uot much wider than he velocity rauge of the CÓ eas and/or the accuracy of the redshift) aud the fact tiat ost observatories can oulv target the higos rotational trations of CO at hieh-:.," Despite these results, two major limitations to systematic searches for CO in HzRGs have been the very limited velocity coverage of existing mm-spectrometers (often not much wider than the velocity range of the CO gas and/or the accuracy of the redshift) and the fact that most observatories can only target the higher rotational transitions of CO at $z$."
+ While hese higher CO transitions trace the deIsc and thermally excited eas (such as that iu he melear starburst/ACN region). Sugeest that more wicdely distributed. xSOPVOlUS of less dense aud sub-thermally excited eas could be largelv missed by these observaions and could thus more casily be detected in the lower transitions than often assuned (seealsoe.g.Dannerbaucretal.2009:erencesthereiu).," While these higher CO transitions trace the dense and thermally excited gas (such as that in the nuclear starburst/AGN region), suggest that more widely distributed reservoirs of less dense and sub-thermally excited gas could be largely missed by these observations and could thus more easily be detected in the lower transitions than often assumed \citep[see also e.g.][and references therein]{dan09,car10,ivi11}."
+ The bhuuimositv of the eround transition. CO(1-0). is least affected bv the excitatio1 coucditious of he gas.," The luminosity of the ground transition, CO(1-0), is least affected by the excitation conditions of the gas."
+ Therefore. observations of CO(t-0) are crucial for deriving the most accurate estimates of he overall molecular gas coutent 1 high-: galaxies.," Therefore, observations of CO(1-0) are crucial for deriving the most accurate estimates of the overall molecular gas content in $z$ galaxies."
+ Oueoi1ο developineusin broadband correlators are OVCLCOMMLο the above moentioued. linitations., Ongoing developments in broadband correlators are overcoming the above mentioned limitations.
+ As we discussed in Enoutsetal.(2011).. the uperade of the Australia Telescope Compact Array (ATCA) with tjio 2ν2GGIIz Coupact Array Broadband Dackeud (CABB:Wilsouctal.POLL) LOW aows searching for CO(1-0) ACTOSS the southern remuisphere.," As we discussed in \citet{emo11}, the upgrade of the Australia Telescope Compact Array (ATCA) with the $2 \times 2$ GHz Compact Array Broadband Backend \citep[CABB;][]{wil11} now allows searching for CO(1-0) across the southern hemisphere."
+ We are currenlv iu the process of observing an unbiased sauje of IIZRCs from i6. Molouglo Reference. Catalogue (MeCirthy.etal.1990) in CO(1-0) with the ATCA/CABB ται system between 1.5<2<3., We are currently in the process of observing an unbiased sample of HzRGs from the Molonglo Reference Catalogue \citep{mcc90} in CO(1-0) with the ATCA/CABB 7mm system between $1.7<z<3$.
+ Iun this LeOr. WO preseit the deteclon of CO(1-0) emission iu one of our sunuple sources. 00152-2099 (:—1.92).," In this Letter, we present the detection of CO(1-0) emission in one of our sample sources, 0152-209 $z$ $1.92$ )."
+ 00152-209 In a radio galaxy wili a disturbed optical 1uorphiologv (Peutericcietal.2001)., 0152-209 is a radio galaxy with a disturbed optical morphology \citep{pen01}.
+. It is bright iu the ucar- aud iuh (Sevinouretal.2107:DeDreuck 2010).. detectc (ii citepiucc)l and has a compact radio source (1.6arcsec/13speiudiameter:Peutericciet 20003.," It is bright in the near- and mid-IR \citep{sey07,bre10}, detected in \\citep{mcc91} and has a compact radio source \citep[1.6~arcsec/13~kpc in diameter;][]{pen00_radio}. ."
+". Throughout lis paper. we assume Ty=71 +. Ong0.3 and O4=0.7. correspouding to an angular distance scale of 8.3 kpc i aud η,vy distance Dj=11571 Mpe for MRC 0152-209(followingWrieht2006)."," Throughout this paper, we assume $_{0} = 71$ $^{-1}$, $\Omega_{\rm M} = 0.3$ and $\Omega_{\Lambda} = 0.7$, corresponding to an angular distance scale of 8.3 kpc $^{-1}$ and luminosity distance $D_{\rm L} = 14571$ Mpc for MRC 0152-209\citep[following][]{wri06}."
+.! AIRC 0152-209 was observed ou 25 & 26 August and 28 & 29 September 2010 with ATCA/CABB in the ost compact hybrid Πο and II75, MRC 0152-209 was observed on 25 $\&$ 26 August and 28 $\&$ 29 September 2010 with ATCA/CABB in the most compact hybrid H168 and H75
+The final step to the data reduction is to apply aperture corrections.,The final step to the data reduction is to apply aperture corrections.
+ These corrections are necessarv since there may be variations from field to field., These corrections are necessary since there may be variations from field to field.
+ The aperture corrections were determined. using the overlapping regions on adjacent fields. starting by the central pointing. which contains the cluster and the standard. stars.," The aperture corrections were determined using the overlapping regions on adjacent fields, starting by the central pointing, which contains the cluster and the standard stars."
+ ‘To eliminate spurious detections. ancl possibly galaxies. we performed. a cut in the magnitude error of the final photometric table.," To eliminate spurious detections, and possibly galaxies, we performed a cut in the magnitude error of the final photometric table."
+ This cut. eliminates sources with errors larger than those expected. for point sources at their magnitude value., This cut eliminates sources with errors larger than those expected for point sources at their magnitude value.
+ This process eliminates most spurious detections including many galaxies that could introduce uncertainties to the tidal tail detection., This process eliminates most spurious detections including many galaxies that could introduce uncertainties to the tidal tail detection.
+ The final photometric sample has approximately 152000 stars., The final photometric sample has approximately $152000$ stars.
+ The mean photometric error in the range 16xV22 is less than 0.05. which is enough for our purposes.," The mean photometric error in the range $16 \leq V \leq 22$ is less than $0.05$, which is enough for our purposes."
+ All magnitudes were corrected. for extinction. using Schlegel.Finkbeiner.&Mare(1998). dust maps., All magnitudes were corrected for extinction using \citet{schlegel} dust maps.
+ The mean reddening for the entire observed region is £E(CD132020., The mean reddening for the entire observed region is $E(B-V) = 0.20$.
+ In Figure Gaa we show the CMD for the stars within ]2'=~37pe. which corresponds to 1.5 tical radius as quoted by DeMarchi&Pulone(2007)," In Figure \ref{cmd2298}a a we show the CMD for the stars within $12 \arcmin = 
+\sim 37~pc$, which corresponds to 1.5 tidal radius as quoted by \citet{marchi2298}."
+ The structure of the CMD ts typical of an old metal-poor GC., The structure of the CMD is typical of an old metal-poor GC.
+ Note that at bright magnitudes we lose stars due to saturation. although the blue end of the extended LED is still visible.," Note that at bright magnitudes we lose stars due to saturation, although the blue end of the extended HB is still visible."
+ Larger errors on brighter magnitudes are due to saturation in the E band., Larger errors on brighter magnitudes are due to saturation in the I band.
+ In Figure 6bb we show the CMD for stars outside. 12’. this plot we choose to display only a fraction of 104 of the total number of points for clarity.," In Figure \ref{cmd2298}b b we show the CMD for stars outside $12\arcmin$, in this plot we choose to display only a fraction of $10\%$ of the total number of points for clarity."
+ Using the best determination of NGC 2298 age and metallicity (DeMarchi&Pulone2007).. we overlay the corresponding Padova isochrone to the data in Figure 6aa. The best fit occurs for a distance modulus of(ii. which closely agrees with the estimate from. HarrisT(19096).," Using the best determination of NGC 2298 age and metallicity \citep{marchi2298}, we overlay the corresponding Padova isochrone to the data in Figure \ref{cmd2298}a a. The best fit occurs for a distance modulus of $(m-M) = 15.15$ , which closely agrees with the estimate from \citet{harris}."
+. -Llowever. a CV-1) offset oft1G towards blue colours was to properly fit the isochrone to the MS and AIS'TO.," However, a (V-I) offset of 0.06 towards blue colours was applied to properly fit the isochrone to the MS and MSTO."
+" ""This olfset. rellects. the discrepancy between the reddening value from. the Schlegel.Finkbeiner.&/Marc dust maps. £(2V)=0.20. and that. found by DeMarchi&Pulone(2007)."," This offset reflects the discrepancy between the reddening value from the \citet{schlegel} dust maps, $E(B-V) = 0.20$, and that found by \citet{marchi2298}."
+ In what follows. for the sake of coherence. we use the Schlegel.Pinkbeiner.&Marc(1998) values over the entire field covered by our MOSAIC?2 data.," In what follows, for the sake of coherence, we use the \citet{schlegel}
+ values over the entire field covered by our MOSAIC2 data."
+ l'or completeness reasons we chose not to match stars in overlapping regions., For completeness reasons we chose not to match stars in overlapping regions.
+ All the analysis was done on a field bv field. basis and. when necessary. we adopt the proper area correction(e.g. border of the fields and chips) using a carefully constructed mask.," All the analysis was done on a field by field basis and, when necessary, we adopt the proper area correction (e.g. border of the fields and chips) using a carefully constructed mask."
+ This mask also takes bad. pixel regions into account when calculating anv area., This mask also takes bad pixel regions into account when calculating any area.
+ All areas calculated hereafter were obtained by a Monte-Carlo integral and accounting for the mask we built., All areas calculated hereafter were obtained by a Monte-Carlo integral and accounting for the mask we built.
+ See Balbinot(2009) for further details., See \citet{balbinot09} for further details.
+ To make and independent measurement of NGC 229s structural parameters we built its radial density profile (RDP)., To make and independent measurement of NGC 2298 structural parameters we built its radial density profile (RDP).
+ Vhe RDP was built by counting stars in racial bins out to the point where the background. is clearly. reached., The RDP was built by counting stars in radial bins out to the point where the background is clearly reached.
+ For the cluster center. we simply used. values given in the literature.," For the cluster center, we simply used values given in the literature."
+ In figure 7 we show the resulting. RDP for NGC 2298., In figure \ref{rdp} we show the resulting RDP for NGC 2298.
+ We choose not to use the cluster most central region due to incompleteness caused both by crowding ancl gaps on the CCD mosaic., We choose not to use the cluster most central region due to incompleteness caused both by crowding and gaps on the CCD mosaic.
+ Our RDP analysis covers oul to an angular distance of 25 which corresponds to TSpe., Our RDP analysis covers out to an angular distance of $25\arcmin$ which corresponds to $78~pc$.
+ Ες is approximately three times the best estimated of tidal radius found in the literature., This is approximately three times the best estimated of tidal radius found in the literature.
+" We fitted a Wing profile with fixed core radius (r,.=0:290.9pc) since our central. densities are not accurate due to crowding.", We fitted a King profile with fixed core radius $r_c = 0\arcmin.29 = 0.9pc$ ) since our central densities are not accurate due to crowding.
+" We find a larger tidal radius of à,=15.91+107. which corresponds to 50pe. which is twice larger than previously found."," We find a larger tidal radius of $r_t = 15\arcmin.91 \pm 1\arcmin.07$, which corresponds to $50~pc$ , which is twice larger than previously found."
+" Phe bac‘Keround density found is of m5,=9.50+£0.10starsfarcmin-stich", The background density found is of $\sigma_{bg} = 9.50\pm0.10 ~ stars/arcmin^{2}$.
+ The last timeNGC 2298 was observed with a laree FOV and. photometric depth was by ‘Trageretal.(1995)., The last timeNGC 2298 was observed with such a large FOV and photometric depth was by \citet{trager}.
+.. Phe authors find a tidal radius of 648 although not reaching deep magnitudes., The authors find a tidal radius of $6\arcmin.48$ although not reaching deep magnitudes.
+ Later. DeMarchi&Pulone(2007) found a tidal raclius of S'.() although using only data that cover a 3.4aremin FOV.," Later, \citet{marchi2298} found a tidal radius of $8\arcmin.0$, although using only data that cover a $3.4\times3.4 ~ arcmin$ FOV."
+ Hence their estimated tidal radius relies on data that do not reach the full extension of the eluster., Hence their estimated tidal radius relies on data that do not reach the full extension of the cluster.
+ Our determination combines the advantages of lavecr depth compared to photographic plates and larger FOV compared to LISTΑΟ., Our determination combines the advantages of larger depth compared to photographic plates and larger FOV compared to HST/ACS.
+ NGC 2298 was one ofthe first globular clusters to be bound with a high degree of depletion of low mass stars (DeAlarchi&Pulone 2007)., NGC 2298 was one of the first globular clusters to be found with a high degree of depletion of low mass stars \citep{marchi2298}.
+. The invertec mass function and the ow concentration parameter are expected. for old. globular clusters that are subject to a high degree of tidal interactions Baumeardt&Malkino(2003)., The inverted mass function and the low concentration parameter are expected for old globular clusters that are subject to a high degree of tidal interactions \citet{baumgardt}.
+. Based on the fitted isochrone. the mass range sampled w our observation of NGC 2298 is very limited 0.6<AL.de . ," Based on the fitted isochrone, the mass range sampled by our observation of NGC 2298 is very limited $0.6 \leq M_{\odot} \leq 0.79$ ."
+Theinner parts of NGC 2298 are not accessible to us to crowding., The inner parts of NGC 2298 are not accessible to us due to crowding.
+ On the outer parts the number of stars is too low at our photometric depth. giving low statistical significance to the mass function.," On the outer parts the number of stars is too low at our photometric depth, giving low statistical significance to the mass function."
+ We thus refrain. [rom making a PDMLE reconstruction with our ΓΙΟ data: deeper observations covering a wider range of masses. such as those that will be provided by DIES. are necessary to accurately determine the slope of the mass function in the outskirts of NGC 2298," We thus refrain from making a PDMF reconstruction with our CTIO data; deeper observations covering a wider range of masses, such as those that will be provided by DES, are necessary to accurately determine the slope of the mass function in the outskirts of NGC 2298."
+ We follow the same ME recipe adopted for Palomar 5 in 82 in order to investigate the presence of tidal tails associated NGC 2298.," We follow the same MF recipe adopted for Palomar 5 in \ref{findgama}
+ in order to investigate the presence of tidal tails associated NGC 2298."
+ was built using stars that are less than 37pe from NGC 2298 centre., was built using stars that are less than $37~pc$ from NGC 2298 centre.
+ As seen on figure 6.. a relatively large amount of field stars are present on the region chosen to build.," As seen on figure \ref{cmd2298}, a relatively large amount of field stars are present on the region chosen to build."
+. Phese field stars were eliminated by choosing a CMD locus that is consistent with the cluster evolutionary sequence in the same fashion as in Dalbinotetal., These field stars were eliminated by choosing a CMD locus that is consistent with the cluster evolutionary sequence in the same fashion as in \citet{balbinot09}.
+ was built using 6 fields near the edges of the observed region., was built using 6 fields near the edges of the observed region.
+ and were built. using the completeness limit of the most crowded: region., and were built using the completeness limit of the most crowded region.
+ ΕΠ conservative. approach avoids most issues that might arise from non-homogeneous »notometry., This conservative approach avoids most issues that might arise from non-homogeneous photometry.
+ In figure S. we show the resulting star count map for GC 2298 after applying the ME., In figure \ref{2298tail} we show the resulting star count map for NGC 2298 after applying the MF.
+ Several features are found o»ve the le confidence level., Several features are found above the $\sigma$ confidence level.
+ One interpretation of our Lrndings is that the extended: Northwest tail is the trailing all since its orientation is opposed to the proper motion Dinescuetal.19993..., One interpretation of our findings is that the extended Northwest tail is the trailing tail since its orientation is opposed to the proper motion \citep{dinescu}. .
+ Phe two smaller opposing structures »ind in the central East-West direction may be the result ‘tical interaction with the clisk. since they point towards je disk and NGC 2298 is closeto theGalactic plane.," The two smaller opposing structures found in the central East-West direction may be the result of tidal interaction with the disk, since they point towards the disk and NGC 2298 is closeto the Galactic plane."
+ In, In
+depend only on the recent deformation history of that parcel of ης.,depend only on the recent deformation history of that parcel of fluid.
+ 1n recent vears. numerical simulations of accretion Lows subject to the magnetorotational instability have been taken bevond the confines of the shearing box to study the cllects ol evlindrical geometry. the How across the marginally stable orbit around a black hole. or the boundary laver between a star and a disc (e.g. Armitage 1998: Lawley 2000: Llawley. Dalbus Stone 2001: Armitage 2002).," In recent years, numerical simulations of accretion flows subject to the magnetorotational instability have been taken beyond the confines of the shearing box to study the effects of cylindrical geometry, the flow across the marginally stable orbit around a black hole, or the boundary layer between a star and a disc (e.g. Armitage 1998; Hawley 2000; Hawley, Balbus Stone 2001; Armitage 2002)."
+ These advances are to be welcomed. ancl represent. significant computational achievements.," These advances are to be welcomed, and represent significant computational achievements."
+ However. | believe there is still. ancl will continue to exist. à need for a description such as the one developed: in this paper.," However, I believe there is still, and will continue to exist, a need for a description such as the one developed in this paper."
+ This is not just. because global studies of realistic thin disces are well bevond. the current capabilities of direct numerical simulations. but also because such a moclel will allow a variety of methods. analytical and less clirect numerical ones. to continue to be used in studying accretion disc vnamics while taking seriously the nature of magnetorotational turbulence.," This is not just because global studies of realistic thin discs are well beyond the current capabilities of direct numerical simulations, but also because such a model will allow a variety of methods, analytical and less direct numerical ones, to continue to be used in studying accretion disc dynamics while taking seriously the nature of magnetorotational turbulence."
+ Lt should be particularly useful in. understanding the dynamics of warped. eccentric and tically distorted. disces. non-Ixeplerian. accretion [lows close to black holes. and a variety of. time-dependent accretion phenomena.," It should be particularly useful in understanding the dynamics of warped, eccentric and tidally distorted discs, non-Keplerian accretion flows close to black holes, and a variety of time-dependent accretion phenomena."
+ ] thank Charles Ganmic and Jim Pringle for helpful discussions., I thank Charles Gammie and Jim Pringle for helpful discussions.
+ E acknowledge the support of the Roval Society through a University Research Fellowship., I acknowledge the support of the Royal Society through a University Research Fellowship.
+" The realizability constraint requires that the three eigenvalues of 2), and the three cigenvalues of Mj; remain non-negative under the evolutionary model.", The realizability constraint requires that the three eigenvalues of $R_{ij}$ and the three eigenvalues of $M_{ij}$ remain non-negative under the evolutionary model.
+ A full investigation of this issue would require a consideration of a Large number of degenerate cases., A full investigation of this issue would require a consideration of a large number of degenerate cases.
+ For simplicity. 1 demonstrate here only the weaker result that no one of the eigenvalues can become negative while the others remain positive.," For simplicity, I demonstrate here only the weaker result that no one of the eigenvalues can become negative while the others remain positive."
+ lt ds therefore assumed that. in the initial condition. AR; and Alj; are positive definite tensors.," It is therefore assumed that, in the initial condition, $R_{ij}$ and $M_{ij}$ are positive definite tensors."
+ Lf either. £3; or AM; were to develop a negative eigenvalue in the subsequent evolution. when the first such eigenvalue passed through zero the traces Z? and A (which are the sums ofthe eigenvalues) would. still be positive.," If either $R_{ij}$ or $M_{ij}$ were to develop a negative eigenvalue in the subsequent evolution, when the first such eigenvalue passed through zero the traces $R$ and $M$ (which are the sums of the eigenvalues) would still be positive."
+ Therefore one may take AM.o>0 below. without loss of generalitv.," Therefore one may take $R,M>0$ below, without loss of generality."
+ Deline the quadratic forms 2=δι and Q= Yi¥). where NX; and Y; are differentiable vector. fields aclvected according to the time-reversible equations Then ∖∖⋎↓↥∢⋅↓⋅⋖⋅∐∐∣∶∩∕↖∣∣∣∩∣⋠↓≱∖⇂↓↥∢⊾∟⋜↧⋏∙≟↓⋅⋜⋯⋏∙≟⊲⋯⊔⇂↕⊔↓⋖⋅− ln t," Define the quadratic forms $P=R_{ij}X_iX_j$ and $Q=M_{ij}Y_iY_j$ , where $X_i$ and $Y_i$ are differentiable vector fields advected according to the time-reversible equations Then where${\rm D}/{\rm D}t=\partial_t+u_i\partial_i$ is the Lagrangian time-derivative."
+he initial condition. Z?0 everywhere for all VY; and Q>0 everywhere for all 3;.," In the initial condition, $P>0$ everywhere for all $X_i$ and $Q>0$ everywhere for all $Y_i$ ."
+ According to equation. (A3)).," According to equation \ref{dpdt}) ),"
+of polars is their high levels of polarisation.,of polars is their high levels of polarisation.
+ However. until now. no yolarimetry data was avaiable on this object.," However, until now, no polarimetry data was available on this object."
+ Our »olarimery shows circular polarisaion of up to LO per cent., Our polarimetry shows circular polarisation of up to 10 per cent.
+ Such high levels of polarisation together with the detection of only one photometric period. conlirm that CP Tuc is a »olar arr not an IP (IPs show comlex amplitude spectra)., Such high levels of polarisation together with the detection of only one photometric period confirm that CP Tuc is a polar and not an IP (IPs show complex amplitude spectra).
+ The one unusual feature that CP Tuc exhibits for a volar is 110 prominent X-rav dip wjiich lasts approximately wll the orbital evele., The one unusual feature that CP Tuc exhibits for a polar is the prominent X-ray dip which lasts approximately half the orbital cycle.
+ Misaki et al (1996) suggested that this dip is caused. by obseuration of the accretion region by the accretion stream., Misaki et al (1996) suggested that this dip is caused by obscuration of the accretion region by the accretion stream.
+ This conclusion was reached. mainly due o the dip having a greater depth in soft. N-rays (1dkeV) compared to hard. N-rays (413keV) which is indicative of ;whoto-electrie absorption (Fig. 1))., This conclusion was reached mainly due to the dip having a greater depth in soft X-rays (1–4keV) compared to hard X-rays (4–13keV) which is indicative of photo-electric absorption (Fig. \ref{folded}) ).
+ In contrast. the much ooader dips seen in the EUV are erev.," In contrast, the much broader dips seen in the EUV are grey."
+ Lt is very unusual or à polar to have an absorption clip lasting such a Large raction ofthe binary orbit., It is very unusual for a polar to have an absorption dip lasting such a large fraction of the binary orbit.
+ In hard X-ravs. absorption dips asting ~O.1 of the binary orbit have been seen in a number of systems (eg EE Iri: Done. Osborne Beardmore 1995). rut dips lasting 70.5r of the binary orbit have not.," In hard X-rays, absorption dips lasting $\sim$ 0.1 of the binary orbit have been seen in a number of systems (eg EF Eri: Done, Osborne Beardmore 1995), but dips lasting $\sim$ 0.5 of the binary orbit have not."
+ Our modelling of our optical polarimetry data suggests hat the bulk of the optical [ux originates from an extencded accretion region far [from the visible magnetic pole., Our modelling of our optical polarimetry data suggests that the bulk of the optical flux originates from an extended accretion region far from the visible magnetic pole.
+ It is reasonable to assume that the polarised optical Dux and the A-ray Lux originate close to the shock region above the white dwarf., It is reasonable to assume that the polarised optical flux and the X-ray flux originate close to the shock region above the white dwarf.
+ Thus. the dip seen in A-ravs between ὁ ~0.90.2 is simply due to the bulk of the primary accretion region being self-eclipsed by the white dwarf according to our polarisation mocelling.," Thus, the dip seen in X-rays between $\phi\sim$ 0.9–0.2 is simply due to the bulk of the primary accretion region being self-eclipsed by the white dwarf according to our polarisation modelling."
+ However. even at the deepest point of the X-ray dip a significant X-ray signal is present with hard. X-ravs ing more dominant than soft. N-ravs.," However, even at the deepest point of the X-ray dip a significant X-ray signal is present with hard X-rays being more dominant than soft X-rays."
+ ὃν examining Fig., By examining Fig.
+ 4. it can be seen that the trailing edee of the primary. accretion region is still just. visible a 10 deepest point of the X-ray dip (although this maybe not be so clear in the reproduction of Fig. 4))., \ref{globes} it can be seen that the trailing edge of the primary accretion region is still just visible at the deepest point of the X-ray dip (although this maybe not be so clear in the reproduction of Fig. \ref{globes}) ).
+ In models where 3 ds >50° this is more obvious., In models where $\beta$ is $>50^{\circ}$ this is more obvious.
+ On the other han 10 larger (ancl leading) edge of the region is sell-eclipsec or the phases centered on the X-ray clip., On the other hand the larger (and leading) edge of the region is self-eclipsed for the phases centered on the X-ray dip.
+ Lf the leading edge X the region is predominately emitting in soft. X-rays au 10 trailing edge predominates in hard N-ravs then this can account for the energy dependance of the X-ray dip., If the leading edge of the region is predominately emitting in soft X-rays and the trailing edge predominates in hard X-rays then this can account for the energy dependance of the X-ray dip.
+ The structure of the aceretion region in polars depends on how and where the accreting Blow attaches to the magnetic field lines., The structure of the accretion region in polars depends on how and where the accreting flow attaches to the magnetic field lines.
+ A magnetic field can be thought of as acting as a density filter in which less dense material from the aceretion [low is threaded. first while censer material is threaded. futher along its trajectory., A magnetic field can be thought of as acting as a density filter in which less dense material from the accretion flow is threaded first while denser material is threaded futher along its trajectory.
+ Vhis results in a variation in the local specilie accretion rate over the region and hence a range of X-ray οποιον., This results in a variation in the local specific accretion rate over the region and hence a range of X-ray energy.
+ Our model suggests that (at. least in the case of CP Tuc) the trailing edge of the primary accretion region has a higher mean X-ray temperature than the leading edge., Our model suggests that (at least in the case of CP Tuc) the trailing edge of the primary accretion region has a higher mean X-ray temperature than the leading edge.
+ Until this stage we do not. know how our spin phase convention relates to the binary orbital phase., Until this stage we do not know how our spin phase convention relates to the binary orbital phase.
+ We can use the optical spectroscopy of Thomas Reinsch (1996) to roughly locate the position of the secondary. star., We can use the optical spectroscopy of Thomas Reinsch (1996) to roughly locate the position of the secondary star.
+ Thomas Reinsch obtained low resolution (32 A)) spectroscopy of CP Tuc and were able to resolve the Ho line into broad and narrow componens., Thomas Reinsch obtained low resolution (32 ) spectroscopy of CP Tuc and were able to resolve the $\alpha$ line into broad and narrow components.
+ In low resolution svectra such as these the broad comporncnt is likely to originatο from the accretion stream close. to rw white dwarf., In low resolution spectra such as these the broad component is likely to originate from the accretion stream close to the white dwarf.
+ On he other hane the narrow componen is likely to originate in the stream much further from the white dwarf. perhaps c‘loser to (or at) the secondary star.," On the other hand the narrow component is likely to originate in the stream much further from the white dwarf, perhaps closer to (or at) the secondary star."
+ Tje maximum red shit of the broad. and narrow componen sare expected to occur just before inferior conjunction (the point where the secondary is closest to the observer) and in that order., The maximum red shift of the broad and narrow components are expected to occur just before inferior conjunction (the point where the secondary is closest to the observer) and in that order.
+ The spectroscopic. phase 0.0. of. Thomas lteinsch (1996) corresponds to our polarimetric phase 0.08., The spectroscopic phase 0.0 of Thomas Reinsch (1996) corresponds to our polarimetric phase 0.08.
+ From Fig., From Fig.
+ 3 of Thomas Reinsch (1996) the broad component has a maximum blue shift at our ó 70.5 while it reaches maximum. red shift at our ὁ —1.0., 3 of Thomas Reinsch (1996) the broad component has a maximum blue shift at our $\phi\sim$ 0.5 while it reaches maximum red shift at our $\phi\sim$ 1.0.
+ Examining Fig. 4..," Examining Fig. \ref{globes},"
+ we can see that αἱ ὦ 0.40.5 the magnetic field. configuration is such the magnetic field lines feeding the accretion region appearing over the limb of the white dwarf will result in the accreting material being maximally blue shifted., we can see that at $\phi\sim$ 0.4–0.5 the magnetic field configuration is such the magnetic field lines feeding the accretion region appearing over the limb of the white dwarf will result in the accreting material being maximally blue shifted.
+ Similarly at ó ~O.S0.9 our model predicts the accreting. material will have maximum red shift., Similarly at $\phi\sim$ 0.8--0.9 our model predicts the accreting material will have maximum red shift.
+ Bearing in mind the uncertainties in the relative phasing. our model results are reasonably consistent with the spectroscopy of Thomas Iteinsch (1996).," Bearing in mind the uncertainties in the relative phasing, our model results are reasonably consistent with the spectroscopy of Thomas Reinsch (1996)."
+ Let us compare this result to the low resolution spectroscopy. of the celipsing polar MN Ίνα (Ramsay Wheatley 1998) where o=0.0 defined. the eclipse center., Let us compare this result to the low resolution spectroscopy of the eclipsing polar MN Hya (Ramsay Wheatley 1998) where $\phi$ =0.0 defined the eclipse center.
+ Ramsay Wheatley (1998) found a maximum red shift for the broad. component at ó ~0.9 and maximum bluc shift for the broad. component at ó —0.4., Ramsay Wheatley (1998) found a maximum red shift for the broad component at $\phi\sim$ 0.9 and maximum blue shift for the broad component at $\phi\sim$ 0.4.
+ Therefore to obtain, Therefore to obtain
+frequencies of the system (the index j running over the number of planets) and the coefficients M and N are the secular amplitudes.,frequencies of the system (the index $j$ running over the number of planets) and the coefficients $M$ and $N$ are the secular amplitudes.
+ In Paper 1 we concluded that close encounters between the giant planets is the only known mechanism that can explain the current amplitude of Ms5 i.c. the excitation of Jupiter's eccentricity associated with the gs secular frequency.," In Paper I we concluded that close encounters between the giant planets is the only known mechanism that can explain the current amplitude of $M_{5,5}$ i.e. the excitation of Jupiter's eccentricity associated with the $g_5$ secular frequency."
+" Other possible mechanisms, such as the crossing of multiple mean-motion resonances between Jupiter and Saturn, only excite Ms."," Other possible mechanisms, such as the crossing of multiple mean-motion resonances between Jupiter and Saturn, only excite $M_{5,6}$."
+ Resonance crossings between Saturn and Uranus (which also kick Jupiter's orbit) are not strong enough to pump Mss to its current value.," Resonance crossings between Saturn and Uranus (which also kick Jupiter's orbit) are not strong enough to pump $M_{5,5}$ to its current value."
+" However, a Neptune-mass planet (presumably Neptune itself or Uranus) encountering Saturn in general excites Ms to values comparable to the current one."," However, a Neptune-mass planet (presumably Neptune itself or Uranus) encountering Saturn in general excites $M_{5,5}$ to values comparable to the current one."
+" In principle, encounters of a planet with Jupiter do not need to have occurred."," In principle, encounters of a planet with Jupiter do not need to have occurred."
+ Planet-planet encounters are consistent with the giant planet evolution model of Thommes et al. (, Planet-planet encounters are consistent with the giant planet evolution model of Thommes et al. (
+"1999) and with the so-called ""Nice model” (Tsiganis et al, 2005; Gomes et al.,","1999) and with the so-called “Nice model” (Tsiganis et al, 2005; Gomes et al.,"
+ 2005)., 2005).
+ In Paper Η we focused our attention on the terrestrial planets., In Paper II we focused our attention on the terrestrial planets.
+" We showed that if the terrestrial planets were initially on quasi-circular, nearly-coplanar orbits, the divergent migration of Jupiter and Saturn must have been fast, otherwise the orbits of the terrestrial planets would have become too eccentric and/or inclined."," We showed that if the terrestrial planets were initially on quasi-circular, nearly-coplanar orbits, the divergent migration of Jupiter and Saturn must have been fast, otherwise the orbits of the terrestrial planets would have become too eccentric and/or inclined."
+" Their excitation during giant planets migration is primarily caused by the crossing of the secular resonances gs=g» and gs=σι. which pump M», and Mj."," Their excitation during giant planets migration is primarily caused by the crossing of the secular resonances $g_5=g_2$ and $g_5=g_1$, which pump $M_{2,k}$ and $M_{1,k}$."
+ These resonances occur because gs decreases while the orbital separation between Jupiter and Saturn increases., These resonances occur because $g_5$ decreases while the orbital separation between Jupiter and Saturn increases.
+" More precisely, assuming that the orbital separation between Jupiter and Saturn"," More precisely, assuming that the orbital separation between Jupiter and Saturn"
+spectrum extraction with XSELECT and dead-time correction with HXDDTCOR (version 1.50).,spectrum extraction with XSELECT and dead-time correction with HXDDTCOR (version 1.50).
+ The cosmic background was created with a model provided by the team using flat response (ae_hhxd_ppinflate2_220080129.rsp) and then combineda with the internal background model provided by the team (ae401094010_hhxd_ppinbgd.evt)., The cosmic background was created with a model provided by the team using a flat response 20080129.rsp) and then combined with the internal background model provided by the team pinbgd.evt).
+ The resulting combination is used for the background subtraction., The resulting combination is used for the background subtraction.
+" The response matrix used for the analysis is the one proposed by the team for the time of our observation, ae_hhxd_ppinxinome2_220080129.rsp."," The response matrix used for the analysis is the one proposed by the team for the time of our observation, 20080129.rsp."
+ The count rates of IGR J16318—4848 are 0.1437+0.001ctss! for the combined XISs and 0.6108+0.004ctss! for the HXD PIN diodes.," The count rates of IGR $-$ 4848 are $\mathrm{0.1437\pm0.001}\,\mathrm{cts}\,\mathrm{s}^{-1}$ for the combined XISs and $\mathrm{0.6108\pm0.004}\,\mathrm{cts}\,\mathrm{s}^{-1}$ for the HXD PIN diodes."
+" For the analysis with XSPEC we rebinned the spectrum to a minimum of 250 and ccounts per bin for the XIS and the PIN, respectively."," For the analysis with XSPEC we rebinned the spectrum to a minimum of 250 and counts per bin for the XIS and the PIN, respectively."
+ The uncertainties for all fits are quoted at the level for a single parameter of interest., The uncertainties for all fits are quoted at the level for a single parameter of interest.
+" In order to account for flux cross calibration issues among the instruments, in all spectral fits a multiplicative constant was introduced."," In order to account for flux cross calibration issues among the instruments, in all spectral fits a multiplicative constant was introduced."
+" Although we detected a soft excess in the spectrum below 5 keV (Fig. 1)),"," Although we detected a soft excess in the spectrum below 5 keV (Fig. \ref{fig:excess}) ),"
+ we did not include it in the modeling because itis most probably due to a serendipitous source at a distance -30” from IGR J16318—48482003)., we did not include it in the modeling because it is most probably due to a serendipitous source at a distance $\simeq 30''$ from IGR $-$ 4848.
+". The presence of this source could not be confirmed here because of the lower angular resolution of the XISs compared toXMM-Newton,, even when using an optimal attitude solution for by measuring the attitude directly through following the location of IGR J16318—4848 on the XIS chips."," The presence of this source could not be confirmed here because of the lower angular resolution of the XISs compared to, even when using an optimal attitude solution for by measuring the attitude directly through following the location of IGR $-$ 4848 on the XIS chips."
+" In order to describe the kkeV broad-band spectrum of the source we fit the spectral continuum with an absorbed cutoff powerlaw, taking also into account non-relativistic Compton scattering."," In order to describe the keV broad-band spectrum of the source we fit the spectral continuum with an absorbed cutoff powerlaw, taking also into account non-relativistic Compton scattering."
+" Photoabsorption was modeled with a revised version of the TBabs model2006),, using the interstellar medium abundances summarized by(2000)."," Photoabsorption was modeled with a revised version of the TBabs model, using the interstellar medium abundances summarized by."
+. This model describes the continuum extremely well (Fig. 3))., This model describes the continuum extremely well (Fig. \ref{fig:spec}) ).
+" In addition to the continuum, strong fluorescent"," In addition to the continuum, strong fluorescent"
+the decay. products from resonances. Polarizatious ofhvperons having acommon cliquark,"with the projectile, such as $\Sigma^+, \Sigma^0$ and $\Lambda$, in unpolarized $pA$ collisions hardly depend"
+associated with the sliocks.,associated with the shocks.
+ IIcatiug of the ICM may also result frou gas. “sloshing” iu the cluster's central potential well (CÁscasibar Abukeviteh 2006: Zullone et 22010)., Heating of the ICM may also result from gas “sloshing” in the cluster's central potential well (Ascasibar Markevitch 2006; ZuHone et 2010).
+ Cold. fonts. regions where the surface brightuess chauges sharply but the pressure does not. can result from these sloshine motions driven initially by cluster or sub-cluster merecrs.," Cold fronts, regions where the surface brightness changes sharply but the pressure does not, can result from these sloshing motions driven initially by cluster or sub-cluster mergers."
+ The sloshing can produce a spiral distribution of cool cluster gas reaching iuto the cluster ceuter CClarke et 22001. Lagana et 22010). aud is also seen iu ealaxv groups (Randall et 220092).," The sloshing can produce a spiral distribution of cool cluster gas reaching into the cluster center Clarke et 2004, Laganá et 2010), and is also seen in galaxy groups (Randall et 2009a)."
+ Tere. we present a very deep observation of the cool core cluster Abell 2052.," Here, we present a very deep observation of the cool core cluster Abell 2052."
+ The only other cool core clusters observed to similar or ereater depth with are Perseus (Fabian ot 22006) and M7που (Million et 22010. Werner ct 22010).," The only other cool core clusters observed to similar or greater depth with are Perseus (Fabian et 2006) and M87/Virgo (Million et 2010, Werner et 2010)."
+ Abell 2052 is a moderately rich cluster at a redshift of 2=0.03519 (Ocecrle Till 2001)., Abell 2052 is a moderately rich cluster at a redshift of $z=0.03549$ (Oegerle Hill 2001).
+ Its central radio source. 3C 317. is hosted by the central cD galaxy. UGC 090799.," Its central radio source, 3C 317, is hosted by the central cD galaxy, UGC 09799."
+ Abell 2052 was previously observed in the Nav with Einstein White. Jones. Forman 1997).ROSAT (Peres ct 11998. Rizza ot 22000). (White 2000). (Blauntou et 22001. 2003. 2009. 2010). (Tamra et 22008). and (de Plaa et 22010).," Abell 2052 was previously observed in the X-ray with $Einstein$ (White, Jones, Forman 1997), (Peres et 1998, Rizza et 2000), (White 2000), (Blanton et 2001, 2003, 2009, 2010), (Tamura et 2008), and (de Plaa et 2010)."
+ In addition. we present archival racio observations from the Very Large Array (VLÀ).," In addition, we present archival radio observations from the Very Large Array (VLA)."
+" We asstuune JJ,=70 lau st 1, Q3;=0.3. aud Q4—0.7 (1%=0.7059 kpe at 2= 0.03519) throuehout."," We assume $H_{\circ}=70$ km $^{-1}$ $^{-1}$ , $\Omega_{M}=0.3$, and $\Omega_{\Lambda}=0.7$ $1\arcsec = 0.7059$ kpc at $z = 0.03549$ ) throughout."
+ Errors are given at the lo level unless otherwise stated., Errors are given at the $1\sigma$ level unless otherwise stated.
+ Abell 2052 was observed with for a total of 662 ksec in Cycles 1. 6. and 10 from 9000 2009.," Abell 2052 was observed with for a total of 662 ksec in Cycles 1, 6, and 10 from 2000 – 2009."
+ A summary of the observations is given in Table 1., A summary of the observations is given in Table 1.
+ All observations were performed with the ACIS-S as the primary iustrmment., All observations were performed with the ACIS-S as the primary instrument.
+ The ACTS-S was chosen for its ereater response at low energies than the ACTS-1. vielding a higher count rate for this relatively cool. AEx=3 keV cluster.," The ACIS-S was chosen for its greater response at low energies than the ACIS-I, yielding a higher count rate for this relatively cool, $kT\approx3$ keV cluster."
+" The nominal roll angles used for the observations varied from 99° to 265"" providing coverage at large radii at a range of azimuthal angeles around the cluster.", The nominal roll angles used for the observations varied from $99^{\circ}$ to $265^{\circ}$ providing coverage at large radii at a range of azimuthal angles around the cluster.
+ The eveuts were telemietered in Very. Faint mode for all but the Cycle 1 observation. where the events were telemoetered in Faint moce.," The events were telemetered in Very Faint mode for all but the Cycle 1 observation, where the events were telemetered in Faint mode."
+ The data were processed in the standard manner. using CIAO L2 and CALDB 1.2.2.," The data were processed in the standard manner, using CIAO 4.2 and CALDB 4.2.2."
+ After cleaning. and filtering for background flares (using the 2.5 7.0 keV ranee for the DI chips aud he 0.3. 12.0 keV range for the FI chips). the total exposure roiiaining for the eleven data sets was 657 ssec.," After cleaning, and filtering for background flares (using the 2.5 – 7.0 keV range for the BI chips and the 0.3 – 12.0 keV range for the FI chips), the total exposure remaining for the eleven data sets was 657 ksec."
+" Background corrections were made using the blauk-sky backerouud fields. iucliding the new ""period E ckeround files for the more recent data."," Background corrections were made using the blank-sky background fields, including the new “period E” background files for the more recent data."
+ For cach target events file. a corresponding backeround events file was created. normalizing by the ratio of counts in the 10. 12 seV. enerev range for the source aud. background files to set the scaling.," For each target events file, a corresponding background events file was created, normalizing by the ratio of counts in the 10 – 12 keV energy range for the source and background files to set the scaling."
+ We have used the NRÀO data achive to extract observations of 3C 317 at L8 aud 1.1 GIIz., We have used the NRAO data achive to extract observations of 3C 317 at 4.8 and 1.4 GHz.
+ A παν of the data sets is presented iu Table 2.., A summary of the data sets is presented in Table \ref{tbl:radio}.
+ The archival radio data were calibrated and veduced with the NRAO Astronomical huage Processing System (AIPS)., The archival radio data were calibrated and reduced with the NRAO Astronomical Image Processing System (AIPS).
+ huages were produced through the standard Fourier transform) decouvolution method for each frequency and configuration., Images were produced through the standard Fourier transform deconvolution method for each frequency and configuration.
+ Several loops of imagine and sclf&calibration were undertaken for cach data set to reduce the effects of phase and amplitude errors in the data., Several loops of imaging and self-calibration were undertaken for each data set to reduce the effects of phase and amplitude errors in the data.
+ The final radio nuage at [8 11 was obtained through combining the three VLA coufiguratious aud two observing frequencies., The final radio image at 4.8 GHz was obtained through combining the three VLA configurations and two observing frequencies.
+ We have also produced a combined configuration iniage at 1.I Cz but do uot show it as the structure is remarkably similar to the Ls CITIz inae., We have also produced a combined configuration image at 1.4 GHz but do not show it as the structure is remarkably similar to the 4.8 GHz image.
+ We have mace a spectral iudex map between 1.1 CIIz and LS GIIz to compare the spectral features more directly with the N-rav structure., We have made a spectral index map between 1.4 GHz and 4.8 GHz to compare the spectral features more directly with the X-ray structure.
+ This map was created from the combined coufleuration data at cach frequency., This map was created from the combined configuration data at each frequency.
+ The we-coverage of both data sets was matched aud both frequencies were tmaged with a ccireular beam., The $uv$ -coverage of both data sets was matched and both frequencies were imaged with a circular beam.
+ We have blauked all pixels in the spectral index map that were lower than the 5o level ou either of the input maps., We have blanked all pixels in the spectral index map that were lower than the $\sigma$ level on either of the input maps.
+ Alereed N-rav naages of the source were created., Merged X-ray images of the source were created.
+ The nuages were corrected using merged background miases and exposure πας., The images were corrected using merged background images and exposure maps.
+ An image showing the combined data from all CCD chips used in the analysis (ACIS SI. S2. S35. I2. and I3) in the 0.3.10.0 keV. baud is show in refüe:allehips..," An image showing the combined data from all CCD chips used in the analysis (ACIS S1, S2, S3, I2, and I3) in the $0.3 - 10.0$ keV band is shown in \\ref{fig:allchips}."
+ This figure illustrates the different roll angeles that were used when the observations were performed., This figure illustrates the different roll angles that were used when the observations were performed.
+ The exteuded cluster enüssion ids visible. ax well as numerous point sources which appear more exteuded while increasingly off-axis due to the larger PSF in these regions.," The extended cluster emission is visible, as well as numerous point sources which appear more extended while increasingly off-axis due to the larger PSF in these regions."
+ The nuage has beeu smoothed with a Cassia., The image has been smoothed with a Gaussian.
+ An uusinoothed image in the 0:3.2.0 keV baud of the ceutral region of the cluster is shown in reffieumsimoo.., An unsmoothed image in the $0.3 - 2.0$ keV band of the central region of the cluster is shown in \\ref{fig:unsmoo}.
+ The image reveals exquisite detail related to the interaction of the AGN with the ICAL, The image reveals exquisite detail related to the interaction of the AGN with the ICM.
+ Point sources. including the central ACN. are visible.," Point sources, including the central AGN, are visible."
+ Cavitics or bubbles to the N and S of the ACN are secu. as well as outer cavities to the NW and SE.," Cavities or bubbles to the N and S of the AGN are seen, as well as outer cavities to the NW and SE."
+ The inner cavities are bounded by bright. dense. rins. aud the NW cavity is surrounded by a very narrow. filamentary loop.," The inner cavities are bounded by bright, dense, rims, and the NW cavity is surrounded by a very narrow, filamentary loop."
+ A filament extends into the N bubble., A filament extends into the N bubble.
+ A shock is seen exterior to the bubbles aud rims. aud a probable second shock is visible to the NE.," A shock is seen exterior to the bubbles and rims, and a probable second shock is visible to the NE."
+ A three-color image of the ceutral 6/56«6156 region of A2052 is clisplaved in reffie:3color.., A three-color image of the central $6\farcm56 \times 6\farcm56$ region of A2052 is displayed in \\ref{fig:3color}. .
+ The nuage has been slightly suoothect. using a 1755 Caussian. aud the scaling for each of the three colors is logarithinic.," The image has been slightly smoothed, using a 5 Gaussian, and the scaling for each of the three colors is logarithmic."
+ Red represents the soft (0.31.0 keV) baud. exeen is medium (1.03.0 keV). and blue represeuts the hard band (3.010.0 keV).," Red represents the soft $0.3 - 1.0$ keV) band, green is medium $1.0 - 3.0$ keV), and blue represents the hard band $3.0 - 10.0$ keV)."
+ The bright tims surrounding the inner bubbles appear cool., The bright rims surrounding the inner bubbles appear cool.
+ The first shock exterior to the bubble rius i slishtlv elliptical with the major axis iu the north-south direction (Blauton et 22009)., The first shock exterior to the bubble rims is slightly elliptical with the major axis in the north-south direction (Blanton et 2009).
+ It is visible as a discoutinuity i the surface brightuess. aud the bluish color is consistent with a tempcrative rise in this region.," It is visible as a discontinuity in the surface brightness, and the bluish color is consistent with a temperature rise in this region."
+ Outside of the inner shock. a second surface brightuess discoutimuty is seen as erecuish emission in this figure.," Outside of the inner shock, a second surface brightness discontinuity is seen as greenish emission in this figure."
+ The discontiuuity is sharper to the NEand more extended to the SW., The discontinuity is sharper to the NEand more extended to the SW.
+ This feature may represent ashock or a cold frout., This feature may represent ashock or a cold front.
+ These features will be explored in further detail iu, These features will be explored in further detail in
+matter cistribution (Navarro οἱ al 1996. LHIuss ct al 1998. Moore et al 1998).,"matter distribution (Navarro et al 1996, Huss et al 1998, Moore et al 1998)."
+ llowever on galaxies scale he evolution is g£overnec not only by the DM. but also by the eas. whose dynamics regulate on large scale the formaion of erand design spira arms and extended: thin. clisks. and on smaller scale. the ormation of stars and the interaction between stars arc he multi-phase interstellar medium. (ISM). (Phornton ο al 1998).," However on galaxies scale the evolution is governed not only by the DM, but also by the gas, whose dynamics regulate on large scale the formation of grand design spiral arms and extended thin disks, and on smaller scale the formation of stars and the interaction between stars and the multi-phase interstellar medium (ISM) (Thornton et al 1998)."
+ To understand how galaxies formed. and evolve it is therefore necessary to couple gravitational forces anc ivdro-dvnamics., To understand how galaxies formed and evolved it is therefore necessary to couple gravitational forces and hydro-dynamics.
+" This can be done semi-analitveally or numerically,", This can be done semi-analitycally or numerically.
+ Semi-analvtical models. of galaxy. formation (ναπιαπα et al 1993). although successful in reproducing many galaxy properties (Baugh et al 1998). are based. on several ad-hoc assumptions about the interaction between eas and dark matter. and the way stars form from gas and interact with the 1981.," Semi-analytical models of galaxy formation (Kauffmann et al 1993), although successful in reproducing many galaxy properties (Baugh et al 1998), are based on several ad-hoc assumptions about the interaction between gas and dark matter, and the way stars form from gas and interact with the ISM."
+ Many. physical processes. [ike thermal shocks. pressure forces and clissipation are required. together with gravity. to realistically model the formation ancl evolution of galaxies.," Many physical processes like thermal shocks, pressure forces and dissipation are required, together with gravity, to realistically model the formation and evolution of galaxies."
+ This can be done more properly by means of numerical simulations. in which the equations of motion of a large numbers of particles and/or grid cells are integrated under gravitational ancl hvero-dyvnamical forces.," This can be done more properly by means of numerical simulations, in which the equations of motion of a large numbers of particles and/or grid cells are integrated under gravitational and hydro-dynamical forces."
+ In numerical astrophysics gravity can. be. computed using different methods. which can generally be divided in particle ancl ericl technique.," In numerical astrophysics gravity can be computed using different methods, which can generally be divided in particle and grid technique."
+ The simplest. particle method (PP) directly sums up the pairwise contribution between all the particles. but has the disadvantage to increase the computational time as /Nδν making impossible to hancle simulations with more than LO! particles. (Aarseth 1985).," The simplest particle method $PP$ ) directly sums up the pairwise contribution between all the particles, but has the disadvantage to increase the computational time as $N\times N$, making impossible to handle simulations with more than $10^{4}$ particles (Aarseth 1985)."
+ Although this method. can be used. for other. purposes (flor example. open star clusters. Aarseth 1998). it cannot » appropriate for testing cosmic structure and/or galaxy ormation theories.," Although this method can be used for other purposes (for example, open star clusters, Aarseth 1998), it cannot be appropriate for testing cosmic structure and/or galaxy formation theories."
+ X recent development. dis. represented w GRAPE (Ανν Pipl) boards. where the PP force computation is performed with a special-purpose hardware (Llut Alaking 1999).," A recent development is represented by GRAPE (GRAvity PipE) boards, where the $PP$ force computation is performed with a special-purpose hardware (Hut Makino 1999)."
+ Tree. codes (Barnes Ilut. 1986. Appel 1985) reduce the scaling to ΑΝfogN putting the system in a hierarchical structure. and computing the forces rom distant particle groups in an approximate fashion.," Tree codes (Barnes Hut 1986, Appel 1985) reduce the scaling to $N\times logN$ putting the system in a hierarchical structure, and computing the forces from distant particle groups in an approximate fashion."
+ The 'article-Mesh. (PM) method is based on a grid evaluation of the newtonian gravitation potential using Fast Fourier “Transforms (FET)., The Particle-Mesh $PM$ ) method is based on a grid evaluation of the newtonian gravitation potential using Fast Fourier Transforms (FFT).
+ lt scales as NVfogN as well. but ias the drawback that the spatial resolution is limited. by he cell size (lockney Eastwood. 1981).," It scales as $N\times logN$ as well, but has the drawback that the spatial resolution is limited by the cell size (Hockney Eastwood 1981)."
+ The resolution imitation can be alleviated using a hybrid. method. called I7. which uses the eric estimate of the potential only or distant regions. whereas the neighboring contributions are computed with the 22 method.," The resolution limitation can be alleviated using a hybrid method, called $P^{3}M$, which uses the grid estimate of the potential only for distant regions, whereas the neighboring contributions are computed with the $PP$ method."
+ However the best καν ο circumvent resolution problems is to use adaptive grids. which deform and adapt to regions of dilferent. density (livedra. Pearce and Couchman JOT).," However the best way to circumvent resolution problems is to use adaptive grids, which deform and adapt to regions of different density (Hydra, Pearce and Couchman 1997)."
+ Lvclro-dynamical forces. are caleulatecl adopting a Lagrangian>o or Eulerian formalism., Hydro-dynamical forces are calculated adopting a Lagrangian or Eulerian formalism.
+ The use of ogrids is natural for Eulerian codes., The use of grids is natural for Eulerian codes.
+ In these codes the values of hyelro-dvnamical quantities are estimated: inside the eric Cells. whereas [luxes are evaluated across the cell. borders. as in the Piecewise Parabolic Method (PPM) (Woodward Colella. 1984: Brian et al 1995. Cheller οἱ al 19982).," In these codes the values of hydro-dynamical quantities are estimated inside the grid cells, whereas fluxes are evaluated across the cell borders, as in the Piecewise Parabolic Method (PPM) (Woodward Colella 1984; Brian et al 1995, Gheller et al 1998a)."
+ The spatial resolution can be improved. using Adaptive Mesh volinement. (AAR. Bevan Norman 1995).," The spatial resolution can be improved using Adaptive Mesh Refinement (AMR, Bryan Norman 1995)."
+ On the other mud. Lagrangian codes start from. the hyvdro-dynamical conservation laws in Lagrangian formalism (Landau Lifchitz 1971). and mostly utilize particles to map the luid properties.," On the other hand, Lagrangian codes start from the hydro-dynamical conservation laws in Lagrangian formalism (Landau Lifchitz 1971), and mostly utilize particles to map the fluid properties."
+ “Phis is the case of the Smoothed Particle Ivdro-dynamies (SPLD) technique developed by Lucy (1977) and Cingold Monaghan (1977)., This is the case of the Smoothed Particle Hydro-dynamics (SPH) technique developed by Lucy (1977) and Gingold Monaghan (1977).
+ The advantage of this echnique is the great [exibility and adaptivity (Llernequist katz 1989: Steinmetz Mülller 1993: Carraro et al 1998a), The advantage of this technique is the great flexibility and adaptivity (Hernquist Katz 1989; Steinmetz Mülller 1993; Carraro et al 1998a).
+ In the last decade a variety of cliflerent combinations of gravity solvers ancl hyelro-cvnamical methods appeared., In the last decade a variety of different combinations of gravity solvers and hydro-dynamical methods appeared.
+ Brick. there are pure grid. codes like ALPPAL (Brvan Norman 1995). erid codes combined with particle codes ike PPALSPL (Exrard 1988). or pure particle codes. like GRAPE-SPLHI (Steinmetz 1996) or Tree-SPLLE (Llerneuist katz 1989).," Briefly, there are pure grid codes like $PM-PPM$ (Bryan Norman 1995), grid codes combined with particle codes like $P^{3}M-SPH$ (Evrard 1988), or pure particle codes, like GRAPE-SPH (Steinmetz 1996) or Tree-SPH (Hernquist Katz 1989)."
+ Basically Lagrangian particle codes provide a xtter resolution in the over-dense regions. but exhibit a xo»orer shocks resolution and. under-dense regions resolution compared with Eulerian codes (Ixang et al 1994).," Basically Lagrangian particle codes provide a better resolution in the over-dense regions, but exhibit a poorer shocks resolution and under-dense regions resolution compared with Eulerian codes (Kang et al 1994)."
+ However most astro-physical phenomena occur in high density regions. in particular galaxies ancl clusters are over-dense regions.," However most astro-physical phenomena occur in high density regions, in particular galaxies and clusters are over-dense regions."
+ This fact renders lagrangian code more favourite to study astro-physical problems. provided. the involved dvnanmies is not strongly dominated by shocks.," This fact renders lagrangian code more favourite to study astro-physical problems, provided the involved dynamics is not strongly dominated by shocks."
+ Carraro ct al (1998a) developed a pure particle cocle. combining Barnes Ilut (1986) octo-tree with SPILL. and applving this code to the formation of a spiral galaxy like the Milkv. Was.," Carraro et al (1998a) developed a pure particle code, combining Barnes Hut (1986) octo-tree with SPH, and applying this code to the formation of a spiral galaxy like the Milky Way."
+ The code is similar to Llernquist lxatz (1989) TIreeSPLHI., The code is similar to Hernquist Katz (1989) TreeSPH.
+ Te uses SPII to solve the hyero-dynamical equations (see also Carraro et al 1998b: Lia Carraro 1999)., It uses SPH to solve the hydro-dynamical equations (see also Carraro et al 1998b; Lia Carraro 1999).
+ In SPELL a fluid is sampled using particles. there is no resolution limitation due to the absence of evicls. and great. Dexibility thanks to the use of a time and space dependent smoothing length.," In SPH a fluid is sampled using particles, there is no resolution limitation due to the absence of grids, and great flexibility thanks to the use of a time and space dependent smoothing length."
+ Shocks are captured by adopting an artificial viscosity. tensor. and the neighbors search is. performed. using the octo-tree.," Shocks are captured by adopting an artificial viscosity tensor, and the neighbors search is performed using the octo-tree."
+ The octo-tree. combined with SPL. allows a global time scaling of logN.," The octo-tree, combined with SPH, allows a global time scaling of $N\times logN$ ."
+ X good advantage of such codes is that it is easy to introduce new physics. like cooling and racliative processes. maenctic fields and so forth.," A good advantage of such codes is that it is easy to introduce new physics, like cooling and radiative processes, magnetic fields and so forth."
+ Finally the kernel. which is utilized to perform hyelro-cdynamical quantities estimates. can be made adaptive by using anisotropic smoothing lengths (Shapiro et al 1996).," Finally the kernel, which is utilized to perform hydro-dynamical quantities estimates, can be made adaptive by using anisotropic smoothing lengths (Shapiro et al 1996)."
+ Lt ds widely recognized that. TreeSPLHI codes. although deficient in some aspects. can give reasonable answers in many astrophysical situations. like in simulations of [fragmentation and star formation in giant molecular clouds. (GMC) (Monaghan. Lattanzio 1991). explosions (Bravo CGarcia-Senz 1995). merging of galaxies (Mihos Llernquist 1994). ealaxies and clusters formation (ναι Gunn 1991. Ixatz White 1993) and Lyman alpha forest (Llernquist ct al 1996).," It is widely recognized that TreeSPH codes, although deficient in some aspects, can give reasonable answers in many astrophysical situations, like in simulations of fragmentation and star formation in giant molecular clouds (GMC) (Monaghan Lattanzio 1991), explosions (Bravo Garcia-Senz 1995), merging of galaxies (Mihos Hernquist 1994), galaxies and clusters formation (Katz Gunn 1991, Katz White 1993) and Lyman alpha forest (Hernquist et al 1996)."
+ Galaxy formation in particular requires a huge cvnamical range (Dave et al 1997)., Galaxy formation in particular requires a huge dynamical range (Davé et al 1997).
+ In fact an ideal galaxy, In fact an ideal galaxy
+in the rate at which coronal gas cools. and the density and A-ray luminosity of the corona increase rapidly.,"in the rate at which coronal gas cools, and the density and X-ray luminosity of the corona increase rapidly."
+ Ehe higher they eet. the lower the chance that at the next major merger a cool-gas filament. will reach the galactic centre and renew thecusp.," The higher they get, the lower the chance that at the next major merger a cool-gas filament will reach the galactic centre and renew the."
+. Phus there is a direct connection between the process discussed. here and. the stark contrast. in. the detectability with X-ravs of the coronae of star-forming and earlv-tvpe galaxies (Rasmussenctal., Thus there is a direct connection between the process discussed here and the stark contrast in the detectability with X-rays of the coronae of star-forming and early-type galaxies \citep{Rasmussen09}.
+"2009) The progressive reduction. with cosmic time in the tvpical halo mass of galaxies making the transition from the blue cloud to the red sequence (Codownsizing""). follows from. the decrease in the abundance of cold infalling gas with both cosmic time and halo mass-scale that is expected on analytic erouncds and observed in ab-initio cosmological simulations."," The progressive reduction with cosmic time in the typical halo mass of galaxies making the transition from the blue cloud to the red sequence (“downsizing”), follows from the decrease in the abundance of cold infalling gas with both cosmic time and halo mass-scale that is expected on analytic grounds and observed in ab-initio cosmological simulations."
+ lt would be interesting to adel this idea to semi-analytic models of the evolution of the galaxy population., It would be interesting to add this idea to semi-analytic models of the evolution of the galaxy population.
+ We thank an anonymous referee for his/her valuable comments., We thank an anonymous referee for his/her valuable comments.
+ Most of the numerical simulations were performed using the BCX svstem at CINECA. Bologna. with CPU time assigned uncer the INAP-CINECA agreement 2005-2014V.," Most of the numerical simulations were performed using the BCX system at CINECA, Bologna, with CPU time assigned under the INAF-CINECA agreement 2008-2010."
+ EM gratefully acknowledges support from the Marco Polo program. University of Bologna.," FM gratefully acknowledges support from the Marco Polo program, University of Bologna."
+V363 Aur is normally observed αἱ V—14.2 but has been observed to be as faint as V = 15.,V363 Aur is normally observed at V=14.2 but has been observed to be as faint as V = 15.
+ It has never been seen at a low brightness state thus suggesting its classification as a member of the UX UMa subclass., It has never been seen at a low brightness state thus suggesting its classification as a member of the UX UMa subclass.
+ The FUY spectrum reveals a steeply rising continuum toward shorter wavelengths and a line spectrum dominated by strong emission features due to C IV and IIe II., The FUV spectrum reveals a steeply rising continuum toward shorter wavelengths and a line spectrum dominated by strong emission features due to C IV and He II.
+ The He H strength is strong for a non-magnetic nova-like. raising the possibility that the white dwar! could be magnetic.," The He II strength is strong for a non-magnetic nova-like, raising the possibility that the white dwarf could be magnetic."
+ Absorption features due to $i HE (1300). C II (1335) and O V (1371) are also seen along with the C IV and He II emission.," Absorption features due to Si III (1300), C II (1335) and O V (1371) are also seen along with the C IV and He II emission."
+ On the basis of V363 Aur having an optically thick. steady state accretion disk. Hoare ancl Drew (1991) used (he Zanstra method to determine a boundary laver temperature of TOOOOK (o 100000Ix using both the He II 1640 flix and the He II 4686 flux.," On the basis of V363 Aur having an optically thick, steady state accretion disk, Hoare and Drew (1991) used the Zanstra method to determine a boundary layer temperature of 70000K to 100000K using both the He II 1640 flux and the He II 4686 flux."
+ For the model fitting to V363 Aur. we adopted a white cwarf mass Mwd=0.86 M... inclination i = 73 degrees. and a reddening value of E(B-V) = 0.3 (see Table 1).," For the model fitting to V363 Aur, we adopted a white dwarf mass $\mathrm{wd} = 0.86$ $M_{\sun}$, inclination i = 73 degrees, and a reddening value of E(B-V) = 0.3 (see Table 1)."
+ We carried out model accretion disk fitting using optically Chick. steady. state disk models.," We carried out model accretion disk fitting using optically thick, steady state disk models."
+" Our accretion disk model to (he de-reddened spectium SWP25335 consisted of a white dwarf mass M4=0.8.M. (ου)= 8.26). an inclination i = 75 degrees. a 47=2.24 and vielded an accretion rate M=1x10""M. /vr."," Our best-fitting accretion disk model to the de-reddened spectrum SWP25335 consisted of a white dwarf mass $M_{wd} = 0.8 M_{\odot}$ $log(g)= 8.26$ ), an inclination i = 75 degrees, a $\chi^2 = 2.24$ and yielded an accretion rate $\dot{M} = 1 \times 10^{-9}M_{\odot}$ /yr."
+ This best-fit gave a distance of 134 pc Lor V363 Aur., This best-fit gave a distance of 134 pc for V363 Aur.
+ This best-fitting aceretion disk model is displaved in Fie.1., This best-fitting accretion disk model is displayed in Fig.1.
+ RZ Gru is twpically observed al V—12.3 but has been observed as faint as 13.4., RZ Gru is typically observed at V=12.3 but has been observed as faint as 13.4.
+ Thus. it is alwavs in a high brightness state with excursions in brightness of a magnitude or less. also characteristic of the UN. UMa subclass of nova-like variables.," Thus, it is always in a high brightness state with excursions in brightness of a magnitude or less, also characteristic of the UX UMa subclass of nova-like variables."
+ The IUE spectrum reveals a strong continuum sleeply rising toward shorter wavelengths with deep absorption features due to € IV (blended 1543. 1550 components). οἱ IV. (blended 1393. 1402 components). N V (blended 1238. 1242 components)," The IUE spectrum reveals a strong continuum steeply rising toward shorter wavelengths with deep absorption features due to C IV (blended 1548, 1550 components), Si IV (blended 1393, 1402 components), N V (blended 1238, 1242 components)."
+ The N V. Si IV and C IV. resonance lines all have blue-shifted. absorption and weak. sharp P Cvgni absorption flanking the short wavelength sile of the profiles.," The N V, Si IV and C IV resonance lines all have blue-shifted absorption and weak, sharp P Cygni absorption flanking the short wavelength side of the profiles."
+ These features mav manifest a hot. outflowing wind arising fom the inner accretion disk and boundary laver between the star aud disk.," These features may manifest a hot, outflowing wind arising from the inner accretion disk and boundary layer between the star and disk."
+ The wind absorption al C IV and N V are very deep as expected for à UN UMa system viewed at low inclination where a larger disk surface area is emitting continuum photons which are absorbed by the outflowing wind., The wind absorption at C IV and N V are very deep as expected for a UX UMa system viewed at low inclination where a larger disk surface area is emitting continuum photons which are absorbed by the outflowing wind.
+ Relatively little is known about RZ Gru., Relatively little is known about RZ Gru.
+ We adopted a reddening of E(B-V) = 0.03 (Bruch and Eneel 1994) and adopted a nominal white dwarf mass of 0.6.V. (log(g)= 8) in (he absence of anv published may value., We adopted a reddening of E(B-V) = 0.03 (Bruch and Engel 1994) and adopted a nominal white dwarf mass of $0.6 M_{\odot}$ $log(g)= 8$ ) in the absence of any published may value.
+ From Table 1. the published orbital inclination is 61 degrees.," From Table 1, the published orbital inclination is 61 degrees."
+ With these parameters. we carried oul accretion disk model lits keeping the accrelion rate as a [ree parameter.," With these parameters, we carried out accretion disk model fits keeping the accretion rate as a free parameter."
+ Our best fitting model accretion disk to WD18138 is displaved in Fig.2., Our best fitting model accretion disk to SWP18138 is displayed in Fig.2.
+ The best-fit parameters are given in Table 3., The best-fit parameters are given in Table 3.
+ The model accretion disk had i = GO degrees. M;=0.55.M. (ου)= 7.71). à A?=4.314 and vielded a distance of," The model accretion disk had i = 60 degrees, $M_{wd} = 0.55 M_{\odot}$ $log(g)= 7.71$ ), a $\chi^2 = 4.314$ and yielded a distance of"
+excluding their detections of ULIRGs in Έναblobs!.. see also Francis et al.,"excluding their detections of ULIRGs in $\alpha$, see also Francis et al."
+ 2001) were bright in the observed infrared. and that they were ULIRGs.," 2001) were bright in the observed infrared, and that they were ULIRGs."
+ The percentage of ULIRGs in the >=23 sample presentedhere is I478%.. if only the objects in the deep survey are considered (7/50).," The percentage of ULIRGs in the $z = 2.3$ sample presentedhere is $ 14^{+8}_{-5}$, if only the objects in the deep survey are considered (7/50)."
+ This number is a lower estimate. as the sensitivity of the deep COSMOS MIPS data only reaches logLjg=12.4L. at z=2.3.," This number is a lower estimate, as the sensitivity of the deep COSMOS MIPS data only reaches $\log L_{IR} = 12.4 \, \mathrm{L}_{\odot}$ at $z = 2.3$."
+ In the z=0.3 sample. the results are complete in infrared luminosity. and the percentage of ULIRGs in the sample is 20176€ (6/30). where we again have included 6 unclassified objects but not the AGN.," In the $z = 0.3$ sample, the results are complete in infrared luminosity, and the percentage of ULIRGs in the sample is $20^{+12}_{-8}$ (6/30), where we again have included 6 unclassified objects but not the AGN."
+ We return to the robustness also of this result against AGN inclusion in Sec., We return to the robustness also of this result against AGN inclusion in Sec.
+ ?? below., \ref{sec:ulirgcolour} below.
+ In Fig., In Fig.
+ 4. the ULIRG fractions are plotted às a function of redshift., \ref{fig:frac} the ULIRG fractions are plotted as a function of redshift.
+ The data points are few. but there is a clear trend for the ULIRG fraction to grow from high redshifts up to the present day.," The data points are few, but there is a clear trend for the ULIRG fraction to grow from high redshifts up to the present day."
+ There even seems to be a hint that the transition from almost zero ULIRG fraction to the current value ts rather sudden., There even seems to be a hint that the transition from almost zero ULIRG fraction to the current value is rather sudden.
+ For the sole purpose of illustrating the time and steepness of the transition. we overplot in Fig.," For the sole purpose of illustrating the time and steepness of the transition, we overplot in Fig."
+" 4. a hyperbolic tangents function of the form: In this equation. ÜF is the ULIRG fraction. 8 represents the steepness of the transition. and z;, is the transition redshift."," \ref{fig:frac} a hyperbolic tangents function of the form: In this equation, $UF$ is the ULIRG fraction, $\theta$ represents the steepness of the transition, and $z_{tr}$ is the transition redshift."
+ The function plotted in Fig., The function plotted in Fig.
+" 4 has UFy=0.20 (present-day fraction). @=2.28. and Z5,=2.52."," \ref{fig:frac} has $UF_0 = 0.20$ (present-day fraction), $\theta = 2.28$, and $z_{tr} = 2.52$."
+ This means that the transition redshift for the ULIRG fraction is z:~2.5; the ULIRG fraction is negligible at higher redshifts. whereas at lower redshifts it approaches ~25%.," This means that the transition redshift for the ULIRG fraction is $z \sim 2.5$; the ULIRG fraction is negligible at higher redshifts, whereas at lower redshifts it approaches $\sim 25$."
+. For the definition of the Lya-selected ULIRG subsamples we used the common definition based on Ly»., For the definition of the $\alpha$ -selected ULIRG subsamples we used the common definition based on $L_{IR}$.
+ IR-selected ULIRG samples are known to have a significant (15-50%)) fraction of galaxies with AGN components. but it is also known that their total luminosity is often dominated by star formation rather than by the AGNs (Veilleux et al.," IR-selected ULIRG samples are known to have a significant ) fraction of galaxies with AGN components, but it is also known that their total luminosity is often dominated by star formation rather than by the AGNs (Veilleux et al."
+ 2009)., 2009).
+ Nevertheless. to make sure that our Ένα selection did not bias the results. we chose the most conservative approach and therefore removed all X-ray detected objects from the samples.," Nevertheless, to make sure that our $\alpha$ selection did not bias the results, we chose the most conservative approach and therefore removed all X-ray detected objects from the samples."
+" We now ask the question of whether a less conservative approach is possible. and if that would change the conclusions,"," We now ask the question of whether a less conservative approach is possible, and if that would change the conclusions."
+ We therefore searched for an additional test to certify that a candidate is a ULIRG., We therefore searched for an additional test to certify that a candidate is a ULIRG.
+ Using the ULIRG template SEDs of Vega et al. (, Using the ULIRG template SEDs of Vega et al. (
+2008). we find that a useful index at =2.3 Is optical R band minus Ch4 of IRAC.,"2008), we find that a useful index at $z=2.3$ is optical $R$ band minus Ch4 of IRAC."
+ All Vega templates have R—Ch>2.7. see Fig. 5..," All Vega templates have $R-Ch4>2.7$, see Fig. \ref{fig:colours}."
+ The points in the plot show all the LAE candidates at from Nilsson et al. (, The points in the plot show all the LAE candidates at $z=2.3$ from Nilsson et al. (
+2009) with >3c detections in R and Cha.,2009) with $> 3\sigma$ detections in $R$ and $Ch4$.
+ The black points are those detected with MIPS. whereas the green points are not MIPS-detected.," The black points are those detected with MIPS, whereas the green points are not MIPS-detected."
+ X-ray detected objects are marked with a pink ring., X-ray detected objects are marked with a pink ring.
+ Four ULIRGs are not detected in Ch and are shown with upper limits., Four ULIRGs are not detected in $Ch4$ and are shown with upper limits.
+ It can be seen from this plot that the flux selection with MIPS very cleanly samples the region of colour space where ULIRGs are expected to be. except the very brightest AGN that are too blue for this selection.," It can be seen from this plot that the flux selection with MIPS very cleanly samples the region of colour space where ULIRGs are expected to be, except the very brightest AGN that are too blue for this selection."
+ Selecting ULIRGs based on the infrared flux and the R—Ch4 colour. and including all objects that have the correct colours within Io. changes the z=2.3 sample by increasing the number of ULIRGs from seven of a total 50 LAE candidates 1 the first selection to eight of a total 60 candidates in the second selection.," Selecting ULIRGs based on the infrared flux and the $R-Ch4$ colour, and including all objects that have the correct colours within $1\sigma$, changes the $z=2.3$ sample by increasing the number of ULIRGs from seven of a total 50 LAE candidates in the first selection to eight of a total 60 candidates in the second selection."
+ The new fraction ts κ... in perfect agreement with the initial result.," The new fraction is $13^{+7}_{-5}$ , in perfect agreement with the initial result."
+ The ULIRG fraction at z=0.3 was obtained using the Deharveng et al. (, The ULIRG fraction at $z=0.3$ was obtained using the Deharveng et al. (
+2008) sample but excluding the AGN reported by Cowie et al. (,2008) sample but excluding the AGN reported by Cowie et al. (
+2009).,2009).
+ If the AGN are included.," If the AGN are included,"
+(Evrarcl.Metzler.&Navarro1996:Horner.Mushotzky.Seharl1999) we found that the A548b cluster has a mass of zz2.7xLONAL. for h=0.7 or 3.7xLOMA. for h=0.5.,"\citep*{evr96,hor99}
+ we found that the A548b cluster has a mass of $\approx 2.7 \times 10^{14}
+M_{\odot}$ for $h=0.7$ or $3.7 \times 10^{14} M_{\odot}$ for $h=0.5$."
+ Since very lew clusters with lower temperatures have been found to possess radio halos. the threshold mass mav be on the order of ~101...," Since very few clusters with lower temperatures have been found to possess radio halos, the threshold mass may be on the order of $\sim 10^{14} M_{\odot}$."
+" We choose 10M, to be the threshold mass Mj for cluster mergers (to form radio halos.", We choose $10^{14} M_{\odot}$ to be the threshold mass $M_{\mathrm{th}}$ for cluster mergers to form radio halos.
+" For comparison. we also consider different AM, (5xLOMAL. and 5xLOMA.) in our caleulation."," For comparison, we also consider different $M_{\mathrm{th}}$ $5 \times 10^{13} M_{\odot}$ and $5 \times 10^{14}
+M_{\odot}$ ) in our calculation."
+ The condition for the disruption of cluster cores have been determined [rom dynamical simulation., The condition for the disruption of cluster cores have been determined from dynamical simulation.
+" Salvador-Solé.Solanes.&Manrique(1998) found that a cluster with mass | experiencing a mereine process woulddisrupt ils core structure when (he relative mass increase AV/M exceeds a certain threshold Ay,=(ZAM/M)usguga0.6.", \citet*{sal98} found that a cluster with mass $M$ experiencing a merging process woulddisrupt its core structure when the relative mass increase $\triangle\!M / M$ exceeds a certain threshold $\triangle_{\mathrm{m}} \equiv (\triangle\!M / M)_{\mathrm{threshold}} = 0.6$.
+" We use these two conditions. My=101. and A,,=0.6. as the criteria for cluster mergers to [form radio halos."," We use these two conditions, $M_{\mathrm{th}}=10^{14} M_{\odot}$ and $\triangle_{\mathrm{m}}=0.6$, as the criteria for cluster mergers to form radio halos."
+ The Press-Schechter mass function (P9) is used to model the cluster number density and its evolution., The Press-Schechter mass function (PS) is used to model the cluster number density and its evolution.
+ The comoving number censitv of clusters in the mass range AJ~A+dAl al lime / is given by where py is the present mean density of the universe. 0.(/) is the eritical density threshold for a spherical perturbation to collapse by the time /. ancl 66V) is the present rms density [Iuctuation smoothed over a region of mass AM.," The comoving number density of clusters in the mass range $M \sim M+dM$ at time $t$ is given by where $\rho_{0}$ is the present mean density of the universe, $\delta_{c}(t)$ is the critical density threshold for a spherical perturbation to collapse by the time $t$ , and $\sigma(M)$ is the present rms density fluctuation smoothed over a region of mass $M$."
+ We adopt the expressions of 0.(/) summarized in Randall&Sarazin(2001) for different cosmological models., We adopt the expressions of $\delta_{c}(t)$ summarized in \citet{ran01} for different cosmological models.
+ For 2(M ). we use an approximate [formula proposed by Ixitavama&Suto(1996b). [or the cold dark model fluctuationspectrum," For $\sigma(M)$ , we use an approximate formula proposed by \citet{kit96b} for the cold dark model fluctuationspectrum"
+The harsh radiation environment above the Earth's atmosphere gradually cegracles all clectronic equipment. including the sensitive Charee-Couplecl Device. (CCD) imaging detectors used in the (LIST) Advanced Camera for Surveys. (ACS)/Wide Field Channel (WEO).,"The harsh radiation environment above the Earth's atmosphere gradually degrades all electronic equipment, including the sensitive Charge-Coupled Device (CCD) imaging detectors used in the (HST) Advanced Camera for Surveys (ACS)/Wide Field Channel (WFC)."
+ The detectors work by collecting photoclectrons in a potential well at cach pixel., The detectors work by collecting photoelectrons in a potential well at each pixel.
+ At the end of an exposure. these electrons are then transferred. row by row. to an amplifier at the edge of the device. where they are counted.," At the end of an exposure, these electrons are then transferred, row by row, to an amplifier at the edge of the device, where they are counted."
+ Lowever. radiation damage to the silicon lattice creates charge traps that temporarily capture electrons. and. release. them. only after a characteristic delay.," However, radiation damage to the silicon lattice creates charge traps that temporarily capture electrons and release them only after a characteristic delay."
+" Any electrons captured. during the transfer to the reacout register can reemerge. several pixels later. as a spurious ""Charge Transfer. Inelficiency (οΕΙ) trail behind every bright source."," Any electrons captured during the transfer to the readout register can reemerge, several pixels later, as a spurious “Charge Transfer Inefficiency” (CTI) trail behind every bright source."
+ CPL trailing is particularly troublesome because the amount of [lux trailed is a nonlinear function of the flux. size and shape of a source.," CTI trailing is particularly troublesome because the amount of flux trailed is a nonlinear function of the flux, size and shape of a source."
+ The effect. is thereforeeof a convolution., The effect is therefore a convolution.
+" A multitude of ac-hoe schemes have been invented to estimate (ancl subtract) the ellect of οΕΙ from catalogues of object photometry. astrometry and shape (e.g.Chiabergeefa£.2009:Cawleyαἱ2002:Rhodesclal, 2007)."," A multitude of ad-hoc schemes have been invented to estimate (and subtract) the effect of CTI from catalogues of object photometry, astrometry and shape \citep[\eg][]{chiaberge,cawley02,rhodes07}."
+. Llowever. since CTI moves electrons. around fairly preclictably at the image level. the ideal approach for correction is to directly shullle those electrons back to where they belong.," However, since CTI moves electrons around fairly predictably at the image level, the ideal approach for correction is to directly shuffle those electrons back to where they belong."
+ Since detector reacout is the last. process to happen during data acquisition. this can be conveniently carried. out as the first. process in data processing.," Since detector readout is the last process to happen during data acquisition, this can be conveniently carried out as the first process in data processing."
+ While no algorithm can undo the nonlinear movement of electrons in à single step. Bristow(2003a) and Piatekοἱa£(2005) proposed an iterative algorithm. to remove trailing by repeatedly (re-)adding new trailing.," While no algorithm can undo the nonlinear movement of electrons in a single step, \cite{bristow03im} and \cite{piatek05} proposed an iterative algorithm to remove trailing by repeatedly (re-)adding new trailing."
+ This requires a mocel of the (forward) readout: process., This requires a model of the (forward) readout process.
+" A’ physically-motivated model was developed. for ACS/WEC by Masseyefαἱ, (2010).. using measurements from trails behind warm pixels in science imagine."," A physically-motivated model was developed for ACS/WFC by \cite{m10}, using measurements from trails behind warm pixels in science imaging."
+ This letter updates the Masseyοἱa£(2010) CCD readout model ancl pixel-based CPL correction., This letter updates the \cite{m10} CCD readout model and pixel-based CTI correction.
+ In refsec:imethod.. we account for an additional species of charge trap with a long characteristic release time. and measure the density of traps in the detector up to mid-2010.," In \\ref{sec:method}, we account for an additional species of charge trap with a long characteristic release time, and measure the density of traps in the detector up to mid-2010."
+ Ins refsecmplementation.. we implement the improved. CTI correction algorithm and evaluate its performance.," In \\ref{sec:implementation}, we implement the improved CTI correction algorithm and evaluate its performance."
+ In refsec:conclusions.. we discuss the overall. performance of the detectors. in light of changes to their. operational temperature and the long period during which they were ollline.," In \\ref{sec:conclusions}, we discuss the overall performance of the detectors, in light of changes to their operational temperature and the long period during which they were offline."
+ We use a cvoluume-criven” CCD readout niocdel. whose first inerecdient is the rate at which electrons fill up the potential well in a pixel.," We use a “volume-driven” CCD readout model, whose first ingredient is the rate at which electrons fill up the potential well in a pixel."
+" A cloud. of n, electrons. grows in size as electrons are added. and a cloud with a larger cross-sectional area will be exposed to more charge traps when it is swept through the silicon lattice during readout."," A cloud of $n_e$ electrons grows in size as electrons are added, and a cloud with a larger cross-sectional area will be exposed to more charge traps when it is swept through the silicon lattice during readout."
+ As first suggested, As first suggested
+We again used the data and memberships from ?..,We again used the data and memberships from \cite{1953ApJ...117..313J}.
+ The € BipV diagram. especially the region where the gradient is reversed. shows that there is variable extinction to this cluster.," The $U-B$ $B-V$ diagram, especially the region where the gradient is reversed, shows that there is variable extinction to this cluster."
+ We therefore dereddened the data on ἃ star-bv-star basis. which limits us to D.V«0.0.," We therefore dereddened the data on a star-by-star basis, which limits us to $B-V<0.0$."
+ Before fitting we also excluded LHertzsprung 371 (which appears to be reddened), Before fitting we also excluded Hertzsprung 371 (which appears to be reddened).
+ For the PAIS ages we require à set of consistent ages. and we therefore adopt the ages of ?.. with the following exceptions.," For the PMS ages we require a set of consistent ages, and we therefore adopt the ages of \cite{2008MNRAS.386..261M}, with the following exceptions."
+ We take the PAIS age for. NOGC2547 [rom ? as 38.5Mvr. which is derived from isochrone fitting. though also agrees with the Lithium depletion age.," We take the PMS age for NGC2547 from \cite{2006MNRAS.373.1251N} as 38.5Myr, which is derived from isochrone fitting, though also agrees with the Lithium depletion age."
+ The age for 162602 (25Msr) is taken from ?.. which again is based on isochrone fitting to the PAIS stars.," The age for IC2602 (25Myr) is taken from \cite{1997ApJ...479..776S}, which again is based on isochrone fitting to the PMS stars."
+ We use a PAIS age of 4.5Myr for Cop ΟΕ. which is from Littlefair at al (in prep) out is based on the system of ?..," We use a PMS age of 4.5Myr for Cep OB3b, which is from Littlefair at al (in prep), but is based on the system of \cite{2008MNRAS.386..261M}."
+ The position on the sky of our AIS sample is shown in ligure 7.. along with the positions of the sample of ?.. which represents stars in the ONC itself. and the llanking fields of 7..," The position on the sky of our MS sample is shown in Figure \ref{onc_fig}, along with the positions of the sample of \cite{1997AJ....113.1733H}, which represents stars in the ONC itself, and the flanking fields of \cite{2004AJ....128..787R}."
+ Given the distribution of stars. it is clear that the PAIS age we should use is that of the Hanking fields.," Given the distribution of stars, it is clear that the PMS age we should use is that of the flanking fields."
+ Although ? calculate this. they do so on the assumption that the ONC is 470pe away. whilst à more modern estimate is 400pe (0.andreferences therein)..," Although \cite{2004AJ....128..787R} calculate this, they do so on the assumption that the ONC is 470pc away, whilst a more modern estimate is 400pc \citep[][and references therein]{2008MNRAS.386..261M}."
+ In the Vo/(V—Do diagram ? place the [anking fields 0.3 mags above NGC2264., In the $V_0$ $(V-I)_0$ diagram \cite{2004AJ....128..787R} place the flanking fields 0.3 mags above NGC2264.
+ Correcting the distance to 400pc will bring the Ilanking fields PAIS to the same magnitude as that of NGC2264. and therefore to an age of 3Myvr on the scale of ?..," Correcting the distance to 400pc will bring the flanking fields PMS to the same magnitude as that of NGC2264, and therefore to an age of 3Myr on the scale of \cite{2008MNRAS.386..261M}."
+ We collect together our. measurements of the ages of the eroups and clusters in Table 1.. along with the other xwameters from our fits.," We collect together our measurements of the ages of the groups and clusters in Table \ref{ages_table}, along with the other parameters from our fits."
+ For completeness we include the distances. though as these are derived from. two-parameter its we emphasise that those of 2 are to be preferred.," For completeness we include the distances, though as these are derived from two-parameter fits we emphasise that those of \cite{2008MNRAS.386..261M}
+ are to be preferred."
+ We xot. PAIS against AIS age in Figure S.., We plot PMS against MS age in Figure \ref{ages}.
+ Whilst the PAIS ancl AIS ages for individual clusters may agree to within he uncertainties. the average of the AIS ages is significantly older than the average of the PAIS ages.," Whilst the PMS and MS ages for individual clusters may agree to within the uncertainties, the average of the MS ages is significantly older than the average of the PMS ages."
+ LE we take only hose clusters and associations less than IOMvr old. the MS ages are. on average. a factor two larger.," If we take only those clusters and associations less than 10Myr old, the MS ages are, on average, a factor two larger."
+ Phe issue is clearly which of these age scales is correct., The issue is clearly which of these age scales is correct.
+ Explanations as to why the MS ages may be. incorrect fall into two groups. those associated. with the statistical techniques ancl those associated with the models.," Explanations as to why the MS ages may be incorrect fall into two groups, those associated with the statistical techniques and those associated with the models."
+ We can rule out problems with the fitting procedure by comparing our ages with those obtained by ον, We can rule out problems with the fitting procedure by comparing our ages with those obtained by \cite{1993A&AS...98..477M}.
+ They use similar isochrones to the ones presented. here ancl measure the age of the Pleiades as 100. Myr and the environs of the ONC as 4Alve., They use similar isochrones to the ones presented here and measure the age of the Pleiades as 100 Myr and the environs of the ONC as 4Myr.
+ Both jese ages are compatible with those we measure. suggesting our technique eives similar ages to “hy eve fitting.," Both these ages are compatible with those we measure, suggesting our technique gives similar ages to “by eye” fitting."
+ Equally importantlv. we match the lithium depletion age for the NCGO2547 (3436Myr27) and are very close to the depletion age for the Pleiades (125130Mvr?)..," Equally importantly, we match the lithium depletion age for the NGC2547 \citep[34--36Myr][]{2005MNRAS.358...13J}
+ and are very close to the depletion age for the Pleiades \citep[125--130Myr][]{1998ApJ...499L.199S}."
+ To check that our uncertainties are at least reasonable we tested how the result changes if one star is removed from. each fit., To check that our uncertainties are at least reasonable we tested how the result changes if one star is removed from each fit.
+ As one might expect. the brightest star in the fit provides the tightest limits on the age.," As one might expect, the brightest star in the fit provides the tightest limits on the age."
+ We therefore removed the brightest star from each dataset and replotted Figure S.., We therefore removed the brightest star from each dataset and replotted Figure \ref{ages}.
+ As Figure 9. shows. the result remains clear. though as one might expect the error bars are larger. and one more dataset (the ONC) returns an upper limit for the age.," As Figure \ref{ages_less} shows, the result remains clear, though as one might expect the error bars are larger, and one more dataset (the ONC) returns an upper limit for the age."
+ This strongly, This strongly
+"Pendleton, Hakkila Meegan 1998; Paciesas et al.","Pendleton, Hakkila Meegan 1998; Paciesas et al."
+ 1999; Meegan et al., 1999; Meegan et al.
+" 2000; Johnson, Meegan Hakkila 2000; Brainerd et al."," 2000; Johnson, Meegan Hakkila 2000; Brainerd et al."
+ 2000; Hakkila et al., 2000; Hakkila et al.
+ 2003a; Hakkila et al., 2003a; Hakkila et al.
+ 2003b)., 2003b).
+" Model-dependent studies of the BATSE triggering threshold for different Ep,op; — based on the observational data as well as comparison with other detectors — has also been presented by Band (2006), Band (2004), Band (2003) Band (2002)."," Model-dependent studies of the BATSE triggering threshold for different $\epo$ – based on the observational data as well as comparison with other detectors – has also been presented by Band (2006), Band (2004), Band (2003) Band (2002)."
+" A precise determination of the trigger threshold, however, involves modeling the triggering algorithm of BATSE and its observational efficiency, taking into account the sensitivity of BATSE LADs at different angles of the incident photons."," A precise determination of the trigger threshold, however, involves modeling the triggering algorithm of BATSE and its observational efficiency, taking into account the sensitivity of BATSE LADs at different angles of the incident photons."
+ BATSE had a relatively simple triggering algorithm compared to other instruments such as BAT onboard SWIFT and HETE-II (e.g. Band 2003)., BATSE had a relatively simple triggering algorithm compared to other instruments such as BAT onboard SWIFT and HETE-II (e.g. Band 2003).
+" It generally triggered on a burst if the count rate in the second most brightly illuminated detector exceeded a threshold specified in units of standard deviations — nominally 5.50 for normal incidence on a single detector — above background rate that was determined for each detector over a given time interval, usually set at 17.4 sec."," It generally triggered on a burst if the count rate in the second most brightly illuminated detector exceeded a threshold specified in units of standard deviations – nominally $5.5\sigma$ for normal incidence on a single detector – above background rate that was determined for each detector over a given time interval, usually set at 17.4 sec."
+" However, as indicated by Band (2003), the requirement that a trigger occurs when the flux in the second most brightly illuminated detector reaches 5.50, raises the significance threshold from 5.50 — for a normal incidence — to 5.960—7.780 above the background, depending on the cosine of the angle between the normal axis of the detector and the source."," However, as indicated by Band (2003), the requirement that a trigger occurs when the flux in the second most brightly illuminated detector reaches $5.5 \sigma$, raises the significance threshold from $5.5\sigma$ – for a normal incidence – to $5.96\sigma-7.78\sigma$ above the background, depending on the cosine of the angle between the normal axis of the detector and the source."
+" BATSE was programmed to trigger on any of three time scales: 64 ms, 256 ms and 1024 ms and the trigger energy range was generally set to 50 KeV - 300 KeV. To simulate BATSE, we use BATSE's Detector Response Matrices (DRMs) and public GRB data available through the online HEASARC archives[."," BATSE was programmed to trigger on any of three time scales: 64 ms, 256 ms and 1024 ms and the trigger energy range was generally set to 50 KeV - 300 KeV. To simulate BATSE, we use BATSE's Detector Response Matrices (DRMs) and public GRB data available through the online HEASARC archives."
+" The fluences for four BATSE channels covering the energy range of 20—2000 KeV -- which we treat as bolometric — were taken from the “Current BATSE"" catalog."," The fluences for four BATSE channels covering the energy range of $20-2000$ KeV – which we treat as bolometric – were taken from the “Current BATSE"" catalog."
+ Jimenez et al. (, Jimenez et al. (
+2001) have shown that the fluences resulting from the processing pipeline used to create the BATSE catalog are slightly — up to a factor of 2 — larger than the fluences obtained by fitting high-resolution spectra.,2001) have shown that the fluences resulting from the processing pipeline used to create the BATSE catalog are slightly – up to a factor of 2 – larger than the fluences obtained by fitting high-resolution spectra.
+" By simulation, using a wide range of E,,;5 (1—2000 KeV) with typical photon indices a=—1.1 and 8=—2.3, we find the ratio of the bolometric fluence (given the rest-frame energy range of 1—10,000 KeV) as defined in G07 and the fluence in BATSE energy range to be very close to unity (on average « 1.2)."," By simulation, using a wide range of $\epo$ $1-2000$ KeV) with typical photon indices $\alpha = -1.1$ and $\beta = -2.3$, we find the ratio of the bolometric fluence (given the rest-frame energy range of $1-10,000$ KeV) as defined in G07 and the fluence in BATSE energy range to be very close to unity (on average $< 1.2$ )."
+ Comparing this ratio with the results of Jimenez et al. (, Comparing this ratio with the results of Jimenez et al. (
+"2001), extension of the energy band to compute bolometric fluence via spectral fits appears to be unnecessary for the current analysis (Band Preece 2005; Friedman Bloom 2005).","2001), extension of the energy band to compute bolometric fluence via spectral fits appears to be unnecessary for the current analysis (Band Preece 2005; Friedman Bloom 2005)."
+" In some previous analyses (e.g. N08 G08), the trigger threshold limits were generally obtained using the Band model with typical photon indices fixed to a=—1.1 B=—2.3 as a representation of the spectra of the whole sample of BATSE LGRBs."," In some previous analyses (e.g. N08 G08), the trigger threshold limits were generally obtained using the Band model with typical photon indices fixed to $\alpha=-1.1$ $\beta=-2.3$ as a representation of the spectra of the whole sample of BATSE LGRBs."
+" However, we will show in that these types of simulations are strongly affected by model biases and circularity problems."," However, we will show in \ref{sec:sbep} \ref{sec:pbep} that these types of simulations are strongly affected by model biases and circularity problems."
+" Instead, we rely on the light curves of the GRBs in their original forms as detected by BATSE detectors and use them to find the lower limits for the trigger threshold on planes by decreasing monotonically the amount of received photon counts in all energy channels for each individual burst."," Instead, we rely on the light curves of the GRBs in their original forms as detected by BATSE detectors and use them to find the lower limits for the trigger threshold on planes by decreasing monotonically the amount of received photon counts in all energy channels for each individual burst."
+" The analysis of background count rates in the light curves of 2145 BATSE GRBs gives average minimum required peak fluxes of 0.66-0.86, 0.34-0.43, 0.18-0.24 phcm?s! for the three timescales 64, 256 and 1024 ms respectively, given the aforementioned BATSE significance threshold."," The analysis of background count rates in the light curves of 2145 BATSE GRBs gives average minimum required peak fluxes of 0.66-0.86, 0.34-0.43, 0.18-0.24 $ph\,cm^{-2}s^{-1}$ for the three timescales 64, 256 and 1024 ms respectively, given the aforementioned BATSE significance threshold."
+ These values correspond well to the lowest observed peak fluxes of BATSE GRBs on these timescales., These values correspond well to the lowest observed peak fluxes of BATSE GRBs on these timescales.
+ We also include all three timescales in the trigger criteria of our simulation., We also include all three timescales in the trigger criteria of our simulation.
+" This is particularly important for SGRBs, as they are mainly triggered on 64 ms and 256 ms timescales."," This is particularly important for SGRBs, as they are mainly triggered on 64 ms and 256 ms timescales."
+ The entire simulation algorithms in this work are written in FORTRAN., The entire simulation algorithms in this work are written in FORTRAN.
+" Figure shows the plane of 1900 BATSE bursts with continuous light curves and known durations and fluences, chosen from the complete BATSE catalog."," Figure \ref{fig6} shows the plane of 1900 BATSE bursts with continuous light curves and known durations and fluences, chosen from the complete BATSE catalog."
+" Interestingly, the distributions of both SGRBs (blue dots) LGRBs (red dots) look similar to each other and exhibit similar slopes on both their left and right sides (the dotted lines in Figure [I))."," Interestingly, the distributions of both SGRBs (blue dots) LGRBs (red dots) look similar to each other and exhibit similar slopes on both their left and right sides (the dotted lines in Figure \ref{fig6}) )."
+" The only difference is that SGRBs are shifted towards the upper left of the plane, possibly indicating that the duration of the burst has an important effect on its detection and determines the position of the GRB on the plots of E,,5, versus the bolometric fluence (Spo)."," The only difference is that SGRBs are shifted towards the upper left of the plane, possibly indicating that the duration of the burst has an important effect on its detection and determines the position of the GRB on the plots of $\epo$ versus the bolometric fluence $\sbol$ )."
+" Moreover, the scatter of SGRBs on this plot appears to be smaller than the scatter of LGRBs."," Moreover, the scatter of SGRBs on this plot appears to be smaller than the scatter of LGRBs."
+ SGRBs also show a slightly tighter correlation (Kendall, SGRBs also show a slightly tighter correlation (Kendall
+metallicity ==00.03 (Cirardietal.2000).. is shown in Fig.e 3.,"metallicity 0.03 \citep{girardi00}, is shown in Fig. \ref{HR}."
+0 Also shown are isochrones computed.1 by Marigoinjetal.(2008) lor the same Z and for [ive ages. Le. log/ ==SS.87. 8.89. 8.91. 8.03 and 8.95 (/ in vears).," Also shown are isochrones computed by \citet{marigo08} for the same Z and for five ages, i.e. $\log t$ 8.87, 8.89, 8.91, 8.93 and 8.95 $t$ in years)."
+ The non- metallicity follows from our abundance analysis. and Z=0.03 is the nearest metallicity in the models computed by Cirardietal.(2000) andAlarigoctal.(2008).," The non-solar metallicity follows from our abundance analysis, and Z=0.03 is the nearest metallicity in the models computed by \citet{girardi00} and\citet{marigo08}."
+.. The location of the star indicates a mass z2 MM; and an age ol£z +440 Aves., The location of the star indicates a mass $\approx$ $_\odot$ and an age of $t \approx$ $\pm$ 40 Myrs.
+ In Table 8— we summarized all the astrophisical quantities for 227411., In Table \ref{phot} we summarized all the astrophisical quantities for 27411.
+ lt is well known that the vast majority. of Am stars are binaries., It is well known that the vast majority of Am stars are binaries.
+ There are lew racial velocity measurements of I1D227411 in the Literature., There are few radial velocity measurements of 27411 in the literature.
+ We could only find three measurements in. Buscombe(1963): |17. 134 and 53 km s!.," We could only find three measurements in \citet{buscombe63}: +17, +34 and $-$ 53 km $^{-1}$."
+" The scatter in these values does. indeed. suggest that 227411 is à binary,"," The scatter in these values does, indeed, suggest that 27411 is a binary."
+ Lf so. then the companion is probably much fainter than the primary. otherwise one may expect to see some evidence in the spectrum.," If so, then the companion is probably much fainter than the primary, otherwise one may expect to see some evidence in the spectrum."
+ In Fig., In Fig.
+ 4. we compare the LPE abundances derived from model atmospheres computed: using ΑΔΕΙΑ and AFLASI2., \ref{A12_A9} we compare the LTE abundances derived from model atmospheres computed using ATLAS9 and ATLAS12.
+ Lt is evident that there is good agreement in the abundances derived with the two different codes., It is evident that there is good agreement in the abundances derived with the two different codes.
+ Phere are some very small dilferences for Na. V. Zr and Nel. but these are only 0.1 dex or less.," There are some very small differences for Na, V, Zr and Nd, but these are only 0.1 dex or less."
+ Thus we may use the faster ΑΓΙΑΟ code with confidence., Thus we may use the faster ATLAS9 code with confidence.
+ In Fig., In Fig.
+ 5. we show the abuncances relative to solar standard abuncdances (Cirevesseetal. 2010)., \ref{pattern} we show the abundances relative to solar standard abundances \citep{grevesse10}.
+.. Phe chemical pattern cisplaved here is typical of that observed in Am stars. ic. an underabundance of C. N. O. Ca and Se and à general increasing overabundance for heavy elements.," The chemical pattern displayed here is typical of that observed in Am stars, i.e. an underabundance of C, N, O, Ca and Sc and a general increasing overabundance for heavy elements."
+ The atmospheric abundance of Li is interesting in rw context of dillusion., The atmospheric abundance of Li is interesting in the context of diffusion.
+. Burkhart&Coupry(1991) and Burkhartetal.(2005). find that. in general. the Li abundance in Am stars is close to the cosmic value ogNSN&9.04 dex) although some Am/bin stars appear to have an underabundance of Li.," \citet{burkhart91} and \citet{burkhart05} find that, in general, the Li abundance in Am stars is close to the cosmic value $\log N_{\rm Li}/N_{\rm Tot} \approx -9.04$ dex), although some Am/Fm stars appear to have an underabundance of Li."
+ Normal A-type stars in the range 7000.«μμ<8500 IX appear to have a higher Li abundance. Le. Ίουτιων6S.64. dex. Burkhart&Coupry 1995)..," Normal A-type stars in the range $7000 < T_{\rm eff} < 8500$ K appear to have a higher Li abundance, i.e. $\log N_{\rm Li}/N_{\rm Tot} \approx -8.64$ dex, \citep{burkhart95}. ."
+ To determine the Li abundance. we used the A6707A line. taking into account the hyperfine structure GXndersen.Gustafson&Lambert1984).," To determine the Li abundance, we used the $\lambda$ 6707 line, taking into account the hyperfine structure \citep{andersen84}."
+. The abundance that gives the best fit is logNuiIN ==- S.42-E0.10. which is closer to the abundance in normal A stars and agrees with the average Li abundance of three cluster Am/Em stars observed by Fossatietal. (2007). ücher.," The abundance that gives the best fit is $\log N_{\rm Li}/N_{\rm Tot}$ $-$ $\pm$ 0.10, which is closer to the abundance in normal A stars and agrees with the average Li abundance of three cluster Am/Fm stars observed by \citet{fossati07}. ."
+Michaud&Turcotte(2000). precliet abundances isa function of stellar age ancl ellective temperature using iir moclels of dilfusion., \citet{richer00} predict abundances as a function of stellar age and effective temperature using their models of diffusion.
+ Fig., Fig.
+ 14 of their paper allows us o estimate the predicted. abundances for a star with a given “Par up to a maximum age of zz 6670 Alves., 14 of their paper allows us to estimate the predicted abundances for a star with a given $_{\rm eff}$ up to a maximum age of $\approx$ 670 Myrs.
+ We find mt age of 227411 to be about SLO+40 Myr. which is considerably. older than the maximunm age of models in ucher.Michaud.&Turcotte (2000).. but we will assume mt models of 670 Myr still giveafair approximation of the abundances.," We find that age of 27411 to be about $810 \pm 40$ Myr, which is considerably older than the maximum age of models in \citet{richer00}, , but we will assume that models of 670 Myr still giveafair approximation of the abundances."
+ For Tr 77400 Ix. which is our best estimate or 227411. the models by Richer.Michaud&Turcotte predicts underabundances ranging from -0.3 dex to," For $_{\rm eff}$ 7400 K, which is our best estimate for 27411, the models by \citet{richer00} predicts underabundances ranging from -0.3 dex to"
+ou the other haud. were taken from spectra of the Type T-P SN 1999011 by Ehuliuci (2003) at 9. 10. 113. and 168 davs since explosion. again normalized to their V-band light curves at 1. 10. 100. and 200 davs. with an assiunied distance of 7.8 Mpc.,"on the other hand, were taken from spectra of the Type II-P SN 1999em by Elmhamdi (2003) at 9, 10, 113, and 168 days since explosion, again normalized to their V-band light curves at 1, 10, 100, and 200 days, with an assumed distance of 7.8 Mpc."
+ The inost striking feature from this comparison is that despite chormous kinetic cnereies of ~50<103 eres. the peak optical luminosities of are similar to those of other SNe. even falling below the Ia and II curves in many cases.," The most striking feature from this comparison is that despite enormous kinetic energies of $\sim 50 \times 10^{51}$ ergs, the peak optical luminosities of are similar to those of other SNe, even falling below the Ia and II curves in many cases."
+ This is because the higher ejecta mass produces a large optical depth and most of the internal energy of the eas is converted into kinctic enerey by adiabatic expausion., This is because the higher ejecta mass produces a large optical depth and most of the internal energy of the gas is converted into kinetic energy by adiabatic expansion.
+ Finthernore. the overall spectral shapes of the PPSN curves are not unlike those of more usual cases.," Furthermore, the overall spectral shapes of the PPSN curves are not unlike those of more usual cases."
+ Iu fact. pair-production supernovae spend most of their lives iu the sale temperature range as other SNe aud therefore exhibit a similar range in colors throughout their evolution.," In fact, pair-production supernovae spend most of their lives in the same temperature range as other SNe and therefore exhibit a similar range in colors throughout their evolution."
+ Lastly. iav of line features iu the observed spectra would also be xeseut in more detailed models ofPPSNe.," Lastly, many of line features in the observed spectra would also be present in more detailed models of."
+. Iu particular. )ecause of their hydrogen euvelopes. PPSN spectra should contain hydrogen lines similar to those in Twpe-II SNe.," In particular, because of their hydrogen envelopes, PPSN spectra should contain hydrogen lines similar to those in Type-II SNe."
+" Clearly. then. will not be obviously distinguishable roni their more usual counterparts ""at first elance.”"," Clearly, then, will not be obviously distinguishable from their more usual counterparts “at first glance.”"
+ A closer comparison between spectra. however. uncovers wo kev features that are uniquely characteristic toPPSNe.," A closer comparison between spectra, however, uncovers two key features that are uniquely characteristic to."
+. The first of these is a dramatically exteuded intrinsic decay time. which is especially noticeable in the uodels with the strongest enricluneut of CNO in the cuvelope.," The first of these is a dramatically extended intrinsic decay time, which is especially noticeable in the models with the strongest enrichment of CNO in the envelope."
+ This is due to the long adiabatic cooling times of supergiant progenitors. whose radii are ~20 AU. but whose expansion velocities are simular or even less than those of other SNe.," This is due to the long adiabatic cooling times of supergiant progenitors, whose radii are $\sim 20$ AU, but whose expansion velocities are similar or even less than those of other SNe."
+ Second. are the ouly objects that show an extremely late rise at times =LOO davs.," Second, are the only objects that show an extremely late rise at times $\geq 100$ days."
+ This is cue to energy released by the decay of °° Co. which uulike in the Type Ia case. takes mouths to dominate over the internal euerev imparted by the initial shock.," This is due to energy released by the decay of $^{56}$ Co, which unlike in the Type Ia case, takes months to dominate over the internal energy imparted by the initial shock."
+ In, In
+of stars on the TP-AGB is known as a function of time.,of stars on the TP-AGB is known as a function of time.
+" There now exist sufficient data in the literature for Galactic TP-AGB stars to constrain such colors, and we adopt the mid-IR bands [14.6], and [21.3] from the work of Guandalini[8.8],etal.[12.5],(2006) and (2008)."," There now exist sufficient data in the literature for Galactic TP-AGB stars to constrain such colors, and we adopt the mid-IR bands [8.8], [12.5], [14.6], and [21.3] from the work of \cite{guandalini2006} and \cite{guandalini2008}."
+". Using 200 Galactic TP-AGB C stars from Guandalinietal. and Bussoetal. and conservatively excluding (2006)stars redder than (2007),,K—[8.8]>10 mag to reduce any bias towards IR-Iuminous post-AGB objects, we compute mean colors for the population of dust-enshrouded Galactic TP-AGB stars of (K—X)co=5.66+1.16,6.41c1.15,6.311.06,Cand6.870.86 mag for the four bandpasses, respectively (Vegamagnitudes;zeropointsfromGuandalinietal. 2006)."," Using $200$ Galactic TP-AGB C stars from \cite{guandalini2006} and \cite{busso2007}, , and conservatively excluding stars redder than $K-[8.8]>10$ mag to reduce any bias towards IR-luminous post-AGB objects, we compute mean colors for the population of dust-enshrouded Galactic TP-AGB C stars of $\langle K-X\rangle_C = 5.66\pm 1.16, 6.41\pm 1.15,
+6.31\pm 1.06, {\rm\ and\ } 6.87\pm 0.86$ mag for the four bandpasses, respectively \citep[Vega magnitudes; zeropoints from][]{guandalini2006}."
+". We do not claim that all stars with sufficient mass have these colors upon joining the TP-AGB, but that they are dusty and IR luminous, such that when averaged over timescales longer than the TP-AGB lifetimes (106yrs;e.g.Marigo&Girardi2007) this color is à good approximation for the ensemble."," We do not claim that all stars with sufficient mass have these colors upon joining the TP-AGB, but that they are dusty and IR luminous, such that when averaged over timescales longer than the TP-AGB lifetimes \citep[$10^6$ yrs; e.g.][]{marigo2007} this color is a good approximation for the ensemble."
+" Using Guandalini&Busso(2008),, we compute ensemble colors of mag for M stars."," Using \cite{guandalini2008}, we compute ensemble colors of mag for TP-AGB M stars."
+" Formal errors in these colors are smaller than the systematic uncertainties, which have been listed above."," Formal errors in these colors are smaller than the systematic uncertainties, which have been listed above."
+ These were computed by varying the color bounds of the TP-AGB and by using several different estimators., These were computed by varying the color bounds of the TP-AGB and by using several different estimators.
+ These uncertainties dominate the errors in subsequent modeling., These uncertainties dominate the errors in subsequent modeling.
+" TP-AGB ensemble colors are critical, and use of the Galactic samples assumes that these stars adequately sample the phases of TP- evolution."," TP-AGB ensemble colors are critical, and use of the Galactic samples assumes that these stars adequately sample the phases of TP-AGB evolution."
+" We now combine Figure 1((a) with the AGBknowledge that the C star ensemble color should be used for SSP ages within 0.2«t1.5 Gyr, and the M ensemble color for all other ages."," We now combine Figure \ref{fig:agbf}( (a) with the knowledge that the C star ensemble color should be used for SSP ages within $0.2 < t < 1.5$ Gyr, and the M ensemble color for all other ages."
+" Whether stars have sufficient mass to become C stars depends on mass-loss and the efficiency of dredge-up, both of which are sensitive to metallicity (Marigo&Girardi2007)."," Whether stars have sufficient mass to become C stars depends on mass-loss and the efficiency of dredge-up, both of which are sensitive to metallicity \citep{marigo2007}."
+". For the purposes of making initial estimates of the contribution of the TP-AGB in the mid-IR, we ignore these important effects."," For the purposes of making initial estimates of the contribution of the TP-AGB in the mid-IR, we ignore these important effects."
+"'The colors given in the previous section imply that for ages less than 1.5 Gyr, the TP-AGB emits more than","The colors given in the previous section imply that for ages less than 1.5 Gyr, the TP-AGB emits more than"
+On this plot. RXN0532+62 lies uear the short-periocd eucl. where systems appear o be evolving through the period αμα.,"On this plot, RX0532+62 lies near the short-period end, where systems appear to be evolving through the period minimum."
+ These occupy a relatively wide range of e values., These occupy a relatively wide range of $\epsilon$ values.
+ RN0532+62 ts toward the top of this range. as if it lias evolved iuto the turnarouud region relatively recently.," RX0532+62 is toward the top of this range, as if it has evolved into the turnaround region relatively recently."
+ Patterson(2001) fits the empirical relation between e aud the mass ratio g as e(q)=0.2164., \citet{patt01} fits the empirical relation between $\epsilon$ and the mass ratio $q$ as $\epsilon(q)=0.216q$.
+ Usine our e value. we determine --0.071(19). which is typical for dwarf novae with similarperiods!.," Using our $\epsilon$ value, we determine $q=0.074(19)$, which is typical for dwarf novae with similar."
+. We conclude that RNO532+62 is a typical SU UMa star lying near the minimum period for secoudary systems.Acknowledginents:, We conclude that RX0532+62 is a typical SU UMa star lying near the minimum period for hydrogen-rich secondary systems.:
+ We thauk Holly Sheets for taking the 2006 January spectra., We thank Holly Sheets for taking the 2006 January spectra.
+ The National Science Foundation funded this research through award AST-0307113 and an REU supplement to that award., The National Science Foundation funded this research through award AST-0307413 and an REU supplement to that award.
+ Travel for Ann Ixapusta was made possible by a generous gift [rom C‘lauclia aud Jay Weecl., Travel for Ann Kapusta was made possible by a generous gift from Claudia and Jay Weed.
+Opticaland / photometry of IC 1613 were obtained [rom the OGLE-II microlensing survey (Udalskietal.2001).,Optical and photometry of IC 1613 were obtained from the OGLE-II microlensing survey \citep{uda01}.
+. These data are used to aid in identilvineg the stellar (wpe of the stars detected in our IRAC data. which is challenging with IRAC data alone (see ??)).," These data are used to aid in identifying the stellar type of the stars detected in our IRAC data, which is challenging with IRAC data alone (see \ref{photometry}) )."
+" The optical data were acquired in the context of detecting objects that are variable (due to either microlensing events or intrinsic stellar variability) by taking a large number of frames, each with relatively short integration times."," The optical data were acquired in the context of detecting objects that are variable (due to either microlensing events or intrinsic stellar variability) by taking a large number of frames, each with relatively short integration times."
+ The photometry presented here is ihe mean of the individual photometric observations. including those objects found to be variable.," The photometry presented here is the mean of the individual photometric observations, including those objects found to be variable."
+ The field of view of the Udalskietal.(2001) optical data covers the entire IRAC 3.6. 4.5. 5.8. and 8.0 ccoverage area. except lor a small triangular region 70.9 square arcmin in size in the edge of the IRAC coverage area (see Figure 1)).," The field of view of the \citet{uda01} optical data covers the entire IRAC 3.6, 4.5, 5.8, and 8.0 coverage area, except for a small triangular region $\sim$ 0.9 square arcmin in size in the south-eastern edge of the IRAC coverage area (see Figure \ref{coverage}) )."
+ To estimate the number of foreground stars in both the optical and IR data we used the Milky Wav stellar population svuthesis model of Robinetal.(2003).., To estimate the number of foreground stars in both the optical and IR data we used the Milky Way stellar population synthesis model of \citet{rob03}.
+ Sources were modeled for the Galactic coordinates of IC 1613 in a one square degree field to reduce statistical error., Sources were modeled for the Galactic coordinates of IC 1613 in a one square degree field to reduce statistical error.
+ The Robinetal.(2003) moclel provides magnitudes in the Johnson-Cousins L-band. whose central wavelength is close to that of the TRAC 3.6 bband.," The \citet{rob03} model provides magnitudes in the Johnson-Cousins L-band, whose central wavelength is close to that of the IRAC 3.6 band."
+ In the region of IC 1613 where we have coverage in the optical and all four IRAC channels. we expect 15 foreground stars in our HX color-magnitude diagram (CMD).," In the region of IC 1613 where we have coverage in the optical and all four IRAC channels, we expect 15 foreground stars in our IR color-magnitude diagram (CMD)."
+" These 15 foreground stars are expected to have —7.5«Mà5,< —16 and [3.6]— H.5| colors near Zero.", These 15 foreground stars are expected to have $-$ $<$ $_{3.6 \mu m}<-$ 16 and $-$ [4.5] colors near zero.
+ IRAC 3.6. 4.5. 5.8. and 8.0 iimages of IC 1613 are shown in Figure 2..," IRAC 3.6, 4.5, 5.8, and 8.0 images of IC 1613 are shown in Figure \ref{irac}."
+ In general. (he stars in (hese images are uniformly distributed throughout the images. consistent. with the regular elliptical distribution of red eiants discussed by Daade&Gaposchkin(1963).. and dubbed Daade's Population 1I ‘sheet’ by Sandage(1971).," In general, the stars in these images are uniformly distributed throughout the images, consistent with the regular elliptical distribution of red giants discussed by \citet{baa63}, and dubbed Baade's Population II `sheet' by \citet{san71}."
+. In addition to the smooth distribution of stars. a concentration of bright stars is seen in (he large star forming complex in the northeast.," In addition to the smooth distribution of stars, a concentration of bright stars is seen in the large star forming complex in the northeast."
+ The only. diffuse emission detected in these images is 8.0 eenussion from hot (5350 IX) dust aud/or polvevclic aromatic hydrocarbons., The only diffuse emission detected in these images is 8.0 emission from hot $\sim$ 350 K) dust and/or polycyclic aromatic hydrocarbons.
+ This emission, This emission
+Among the 240 extrasolar planets that have been detected so [ar with à semi.major axis larger than 0.1 astronomical unit (au). about 100 have an eccentricityv e-0.3.,"Among the 240 extrasolar planets that have been detected so far with a semi–major axis larger than 0.1 astronomical unit (au), about 100 have an eccentricity $e > 0.3$."
+ Five of them even have e>O.S., Five of them even have $e>0.8$.
+ Such large eccentricities. which cannot be the result. of dise.planet interaction (Papaloizou et al.," Such large eccentricities, which cannot be the result of disc–planet interaction (Papaloizou et al."
+ 2001). are probably produced by planet interactions. either through scattering or secular perturbation (see Ford Rasio 2008 and references therein). that occur after the disc clissipates (Juric “Tremaine 2008. Chatterjee et. al.," 2001), are probably produced by planet–planet interactions, either through scattering or secular perturbation (see Ford Rasio 2008 and references therein), that occur after the disc dissipates (Juric Tremaine 2008, Chatterjee et al."
+ 2008. Ford Rasio 2008).," 2008, Ford Rasio 2008)."
+ Llere. we show that high eccentricities can be pumped if the orbit of the planet is inclined with respect to the disc.," Here, we show that high eccentricities can be pumped if the orbit of the planet is inclined with respect to the disc."
+ Phe process involved is an extension of the Ixozai mechanism. in which a planet is perturbed by ai distant companion on an inclined. orbit (Ixozai 1962).," The process involved is an extension of the Kozai mechanism, in which a planet is perturbed by a distant companion on an inclined orbit (Kozai 1962)."
+ While the Ixozai effect has always been studied for the case in which the companion is lar away [from the planet. the process investigated here is shown to be ellicient even if the orbit of the planet crosses the disc.," While the Kozai effect has always been studied for the case in which the companion is far away from the planet, the process investigated here is shown to be efficient even if the orbit of the planet crosses the disc."
+ Phe classical Ixozai effect has of course been very well studied., The classical Kozai effect has of course been very well studied.
+ Here. we show that some significant dillerences occur when the classical scenario is extended to apply to à disc.," Here, we show that some significant differences occur when the classical scenario is extended to apply to a disc."
+ In section 2. we review the Ixozai ellect. and. show that the same behaviour is expected. whether the planet is perturbed by a distant companion or by a ring of material orbiting [far away.," In section \ref{sec:Kozai} we review the Kozai effect, and show that the same behaviour is expected whether the planet is perturbed by a distant companion or by a ring of material orbiting far away."
+ In section 3... we present the results of numerical simulations of the interaction between a planet on an inclined orbit ancl a disc.," In section \ref{sec:simulations}, we present the results of numerical simulations of the interaction between a planet on an inclined orbit and a disc."
+ We show that. provided most of the mass in the disc is beyond the orbit. and. the initial inclination is Larger than some critical value. the eravitational potential from the disc causes the eecentricity and the inclination of the planets orbit to oscillate with time.," We show that, provided most of the mass in the disc is beyond the orbit, and the initial inclination is larger than some critical value, the gravitational potential from the disc causes the eccentricity and the inclination of the planet's orbit to oscillate with time."
+ This may occur even if the orbit. crosses the. disc., This may occur even if the orbit crosses the disc.
+ ln section 3 we summarise our [indings. and cliscuss under which conditions this mechanism could operate.," In section \ref{sec:discussion} we summarise our findings, and discuss under which conditions this mechanism could operate."
+ The important result is that a planet on an inclined orbit with respect to the disc and located in or within the planet formation region may have its eccentricity puniped. up to high values by the interaction with the disc., The important result is that a planet on an inclined orbit with respect to the disc and located in or within the planet formation region may have its eccentricity pumped up to high values by the interaction with the disc.
+ This is of astronomical interest. since inclinations are beginning to be measured for extrasolar planets.," This is of astronomical interest, since inclinations are beginning to be measured for extrasolar planets."
+" We consider a planet of mass AZ, orbiting around a star of mass M, which is itself suurounced: by a ring of material of mass ιο.", We consider a planet of mass $M_p$ orbiting around a star of mass $M_\star$ which is itself surrounded by a ring of material of mass $M_{\rm disc}$.
+ Phe ring is in the equatorial plane of the star whereas the orbit of the planet is inclined. with respect to this plane., The ring is in the equatorial plane of the star whereas the orbit of the planet is inclined with respect to this plane.
+ The motion of the planet is dominated. by, The motion of the planet is dominated by
+normalization of the UV spectrum is known and equal to the HM05 spectrum is simply theaverage spectrum); the true normalization(which and shape are still quite uncertain (c.f.??)..,"normalization of the UV spectrum is known and equal to the HM05 spectrum (which is simply the spectrum); the true normalization and shape are still quite uncertain \citep[c.f.][]{scott-etal-02-UV-background, shull-etal-99-UV-background}."
+" We have no ability to determine the true shape or normalization of the radiation field oour data do not show any density diagnostics), so we (i.e.can only proceed using the canonical UV background field; nevertheless, this should be kept in mind."," We have no ability to determine the true shape or normalization of the radiation field our data do not show any density diagnostics), so we can only proceed using the canonical UV background field; nevertheless, this should be kept in mind."
+ A plot of the predicted metal line column densities as a function of aand iis shown in Figure 6.., A plot of the predicted metal line column densities as a function of and is shown in Figure \ref{fig: photo-plot}.
+" To construct this plot, we took the metallicity to be [M/H]——0.7, in the middle of the allowed range."," To construct this plot, we took the metallicity to be $\MH = -0.7$, in the middle of the allowed range."
+ The vertical dashed lines mark the range of density/ionization parameter permitted by the tto ccolumn density ratio., The vertical dashed lines mark the range of density/ionization parameter permitted by the to column density ratio.
+" At these densities, the fraction of hydrogen in neutral form is extremely low (10-45--91)."," At these densities, the fraction of hydrogen in neutral form is extremely low $10^{-4.5 -
+ -5.1}$ )."
+ To obtain the observed neutral hydrogen column density of logNg1=14.6 at this ionization fraction and physical density would require a cloud with a length scale L~0.1—1.2Mpc.," To obtain the observed neutral hydrogen column density of $\logNHI = 14.6$ at this ionization fraction and physical density would require a cloud with a length scale $L \sim
+0.1-1.2\Mpc$."
+" A cloud with this physical size is unlikely to be a distinct cloud in a galaxy halo, but the length scale is consistent with absorption arising in a large-scale filamentary structure (i.e. the or a galaxy group (see,oforexample?). fforest"," A cloud with this physical size is unlikely to be a distinct cloud in a galaxy halo, but the length scale is consistent with absorption arising in a large-scale filamentary structure (i.e. the forest) or a galaxy group \citep[see, for
+example][]{schaye-01-Lya-forest}."
+")Intragroup gas has separately been suggested as a possible origin for highly ionized species such as citepmulchaey-etal-96-hot-gas and a few examples have been observed, both in the Local Group and the low-z IGM (??).."," Intragroup gas has separately been suggested as a possible origin for highly ionized species such as \\citep{mulchaey-etal-96-hot-gas} and a few examples have been observed, both in the Local Group and the $z$ IGM \citep{tripp-savage-00-PG0953+415_z0.1423, nicastro-etal-03-nature}."
+" In order to assess a possible group origin, we examined galaxies in the field near uusing available SDSS data."," In order to assess a possible group origin, we examined galaxies in the field near using available SDSS data."
+ We took the photometric redshift for galaxies in the surrounding region as given by ?.., We took the photometric redshift for galaxies in the surrounding region as given by \citet{cunha-etal-09-SDSS-photoz}.
+" 'The closest potential group candidate is a set of three galaxies ~1Mpc to the NNW: all are brighter than 0.4L* if they are at the redshift of the absorber, and have photometric redshift probability distributions p(z) broadly consistent with z=0.35."," The closest potential group candidate is a set of three galaxies $\sim 1\Mpc$ to the NNW: all are brighter than $0.4\,\Ls$ if they are at the redshift of the absorber, and have photometric redshift probability distributions $p(z)$ broadly consistent with $z=0.35$."
+" The 0.4L* level is typically used in structure-finding algorithms, exploiting the red sequence of galaxies in large virialized structures such as groups and clusters (see,e.g.?);; the fact that iis also iis coincidental."," The $0.4\,\Ls$ level is typically used in structure-finding algorithms, exploiting the red sequence of galaxies in large virialized structures such as groups and clusters \citep[see, e.g.][]{koester-etal-07-maxBCG}; the fact that is also is coincidental."
+" With such a small number of galaxies, the radius of such a system would be only rao000.4Mpc (?) and would not encompass227."," With such a small number of galaxies, the radius of such a system would be only $r_{200} \sim
+0.4\Mpc$ \citep{hansen-etal-09-clusters-groups} and would not encompass."
+"19.. There are a number of galaxies in the vicinity of wwhich have photometric redshifts broadly consistent with z—0.35, but all are fainter than 0.4L*."," There are a number of galaxies in the vicinity of which have photometric redshifts broadly consistent with $z=0.35$, but all are fainter than $0.4\,\Ls$."
+" These galaxies lie primarily to the S. We thus conclude that pprobably does not exist in a galaxy group, and that the ddoes not therefore arise from the group environment."," These galaxies lie primarily to the S. We thus conclude that probably does not exist in a galaxy group, and that the does not therefore arise from the group environment."
+ 'This conclusion is also consistent with the above finding that CIE models do no adequately describe the data., This conclusion is also consistent with the above finding that CIE models do no adequately describe the data.
+ We are thus left with the possibility that the aarises in a filament., We are thus left with the possibility that the arises in a filament.
+ Spectroscopic confirmation of these faint galaxies is required to further investigate the filamentary structure hypothesis., Spectroscopic confirmation of these faint galaxies is required to further investigate the filamentary structure hypothesis.
+" To check the physical consistency of the filament scenario, we compared the measured properties of the ccloud to the IGM equation of state predictions of ?.."," To check the physical consistency of the filament scenario, we compared the measured properties of the cloud to the IGM equation of state predictions of \citet{ricotti-etal-00-IGM-eqn-state}."
+" At a density ng——4.4—4.9, the cosmic overdensity is 6~12—87."," At a density $\hden = -4.4 - -4.9$, the cosmic overdensity is $\delta \simeq 12-37$."
+" Using this overdensity and the effective IGM equation of state (EOS), T=To(1-- from ?,, we can derive the implied thermal Doppler0)?! parameter forI."," Using this overdensity and the effective IGM equation of state (EOS), $T = T_0
+(1+\delta)^{\gamma-1}$ from \citet{ricotti-etal-00-IGM-eqn-state}, we can derive the implied thermal Doppler parameter for."
+. We interpolate between redshifts in their Table 6 to obtain EOS parameters Το=7.1x10? K and y—1=0.75 for z=0.4., We interpolate between redshifts in their Table 6 to obtain EOS parameters $T_0 = 7.1 \times 10^3$ K and $\gamma - 1 = 0.75$ for $z = 0.4$.
+ This resulting EOS then predicts temperatures in the range logT'=4.6—5.0K and b;=(2kT/m)°>28—43kms!., This resulting EOS then predicts temperatures in the range $\logT = 4.6-5.0\K$ and $b_{t} = (2kT/m)^{0.5} = 28-43\kms$.
+" 'To compare to this predicted Doppler parameter with the measurements, we performed a curve of growth analysis on the 3 strongest Lyman series lines in the ccloud."," To compare to this predicted Doppler parameter with the measurements, we performed a curve of growth analysis on the 3 strongest Lyman series lines in the cloud."
+ The results are shown in Figure 7.., The results are shown in Figure \ref{fig: cog}.
+ The vertical dashed lines show the Doppler parameter predicted by the IGM EOS., The vertical dashed lines show the Doppler parameter predicted by the IGM EOS.
+" The curve of growth shows a broad region of allowed Doppler parameters, which are consistent with the theory predictions."," The curve of growth shows a broad region of allowed Doppler parameters, which are consistent with the theory predictions."
+" We conclude that the absorber has a density, ionization state, cloud size, and line broadening consistent with a filament of the large scale structure."," We conclude that the absorber has a density, ionization state, cloud size, and line broadening consistent with a filament of the large scale structure."
+" While we observe that some galaxies in the region around hhave photometric redshift estimates consistent with 0.35, a more detailed spectroscopic survey of the scale galaxy distribution around iis necessary to confirm this suggestion."," While we observe that some galaxies in the region around have photometric redshift estimates consistent with $z=0.35$ , a more detailed spectroscopic survey of the large-scale galaxy distribution around is necessary to confirm this suggestion."
+" In addition, the principle source of error in constraining the physical conditions of the ccloud is the large range of column densities allowed by the fits to the lline."," In addition, the principle source of error in constraining the physical conditions of the cloud is the large range of column densities allowed by the fits to the line."
+ A higher resolution and/or higher signal-to-noise, A higher resolution and/or higher signal-to-noise
+mechanisms in creating a GRB. alter the onset of the supernova on the basis of X-ray line- in GRDOII211.,"mechanisms in creating a GRB, after the onset of the supernova on the basis of X-ray line-emissions in GRB011211."
+ However. ταν line-enissions produced in radiativelv. powered supernovae allow the same time-of-onset of the GRD and the supernova. obviating the need for anv delay mechanism (vanPuttenοἱal.2004).," However, X-ray line-emissions produced in radiatively powered supernovae allow the same time-of-onset of the GRB and the supernova, obviating the need for any delay mechanism \citep{mvp04}."
+. Rotating black holes are described by Kerr (Ixerr.1963)., Rotating black holes are described by Kerr \citep{ker63}.
+.. In. core-collapse of massive stars. rotating black holes nucleate by accumulation of mass and angular momentum from infalling matter.," In core-collapse of massive stars, rotating black holes nucleate by accumulation of mass and angular momentum from infalling matter."
+ The Ixerr solution describes the constraint for a black hole of mass M. and angular momentum 4j. where G is Newton's constant and eis the velocity of light.," The Kerr solution describes the constraint for a black hole of mass $M$ and angular momentum $J_H$ , where $G$ is Newton's constant and $c$ is the velocity of light."
+ Table I summarizes the kev quantities of Nery black holes., Table I summarizes the key quantities of Kerr black holes.
+ Quite eenerallv. initial collapse of à rotating core produces a torus (tees.Ruffini&Wheeler1914:Duez.Shapiro&Yo 2004).. which initially satisfies Jy>GAZ/c. Thus. the nucleation of black holes takes place through afirsf-order phase-Uransition: a torus forms of increasing mass by accumulation of matter. diluting its angular momentum until it satisfies (2)) ancl collapses into an extremal black hole.," Quite generally, initial collapse of a rotating core produces a torus \citep{ree74,due04}, which initially satisfies $J_T>GM_T^2/c.$ Thus, the nucleation of black holes takes place through a phase-transition: a torus forms of increasing mass by accumulation of matter, diluting its angular momentum until it satisfies \ref{EQN_JM}) ) and collapses into an extremal black hole."
+ The alternative of a second-order phase transition which initially forms a sub-solar mass black hole. requires rapid shedding of excess angular momentum by eravitational radiation.," The alternative of a second-order phase transition which initially forms a sub-solar mass black hole, requires rapid shedding of excess angular momentum by gravitational radiation."
+ However. limited mass-densities in core-collapse probably render (his mechanism inellective in competition with mixing on the free-fall timescale of the core.," However, limited mass-densities in core-collapse probably render this mechanism ineffective in competition with mixing on the free-fall timescale of the core."
+ Nevertheless. gravitational radiation emitted [rom a non-axisvmmetre torus prior to the nucleation of the black hole is potentially. interesting (Rees.Rullini&WheelerShapiro&Yo 2004).," Nevertheless, gravitational radiation emitted from a non-axisymmetric torus prior to the nucleation of the black hole is potentially interesting \citep{ree74,due04}."
+. Gravilational radiation in the formation of black holes through a fist-order phase (ransilion is important in non-spherical collapse. even when ils energy emissions are small relative to the initial mass of the black hole.," Gravitational radiation in the formation of black holes through a first-order phase transition is important in non-spherical collapse, even when its energy emissions are small relative to the initial mass of the black hole."
+ The Bekenstein gravitational radiation-recoil mechanism operates already in the presence of initial asphericities of about 10.7. producing a recoil of 300km/s or less.," The Bekenstein gravitational radiation-recoil mechanism operates already in the presence of initial asphericities of about $10^{-3}$, producing a recoil of 300km/s or less."
+ The radius of the accretion disk or torus around a newly formed stellar mass black hole is Ry~ I0*cm., The radius of the accretion disk or torus around a newly formed stellar mass black hole is $R_T\sim 10^7$ cm.
+ A torus of a few tenths of a solar mass forms by accumulation of matter spiralling in. compressed by a [actor e(r/riseo)! as it stalls against ihe angular momentum barrier outside the inner most stable circular orbit (ISCO).The time-ol-collapse of stellar matter from a radius r is about the Iree-Iall timescale.," A torus of a few tenths of a solar mass forms by accumulation of matter spiralling in, compressed by a factor $\sim(r/r_{ISCO})^4$ as it stalls against the angular momentum barrier outside the inner most stable circular orbit (ISCO).The time-of-collapse of stellar matter from a radius $r$ is about the free-fall timescale,"
+converting most of their mass to stars.,converting most of their mass to stars.
+" Thus, the relevant for high-z giant clumps, as derived in the old-star limit, is €= 1-€ = 0.086"," Thus, the relevant for $z$ giant clumps, as derived in the old-star limit, is e = 1 - = 0.086."
+":""eg....(9) Since €=0.5, we expect the resulting1/4M; stellar systems to remain gravitationally bound."," Since $\calE \ga 0.5$, we expect the resulting stellar systems to remain gravitationally bound."
+" Clumps with 3.1—1 have €«0.5 and suffer significant disruption only if their masses are below ~10° Mo, and they reach €«0.3 and undergo complete disruption only at masses €2x10? Mo."," Clumps with $\Sigma_{-1}=1$ have $\calE < 0.5$ and suffer significant disruption only if their masses are below $\sim 10^6$ $\msun$, and they reach $\calE < 0.3$ and undergo complete disruption only at masses $\la 2\times 10^5$ $\msun$."
+" Thus the observed ~105—10? Mo giant clumps in high—z galaxies should survive disruption unless eg,_2>>1.", Thus the observed $\sim 10^8-10^9$ $\msun$ giant clumps in $-z$ galaxies should survive disruption unless $\epsfftwo \gg 1$.
+" Even if we are maximally conservative and assume that the smaller clumps have X..,=0.5, tthat their surface densities do not exceed the mean surface densities of their host galaxies, our estimated maximum mass for disruption only increases by a factor of 2."," Even if we are maximally conservative and assume that the smaller clumps have $\Sigma_{-1} = 0.5$, that their surface densities do not exceed the mean surface densities of their host galaxies, our estimated maximum mass for disruption only increases by a factor of 2."
+" Equation (9)) shows that the most important parameter in determining the survival of giant clumps is how quickly, normalized to their free-fall times, they turn themselves into stars."," Equation \ref{eq:expulsion}) ) shows that the most important parameter in determining the survival of giant clumps is how quickly, normalized to their free-fall times, they turn themselves into stars."
+" Only if they do so with a very high rate efficiency, eg,»Z;10, do we expect significant gas expulsion."," Only if they do so with a very high rate efficiency, $\epsfftwo \ga 10$, do we expect significant gas expulsion."
+" In the local universe, one measures eg by determining the star formation rate of an object or a population of objects, and comparing this to the objects’ gas mass divided by the free-fall time computed for their density."," In the local universe, one measures $\epsff$ by determining the star formation rate of an object or a population of objects, and comparing this to the objects' gas mass divided by the free-fall time computed for their density."
+" This procedure was first applied by ? to the population of giant molecular clouds in the Milky Way, and has subsequently been extended to other objects by ? and ?.."," This procedure was first applied by \citet{zuckerman74} to the population of giant molecular clouds in the Milky Way, and has subsequently been extended to other objects by \citet{krumholz07e} and \citet{evans09a}."
+" These measurements give e#,_2~1 over a very wide range of star-forming environments, from small star clusters in the Milky Way to entire starburst galaxies."," These measurements give $\epsfftwo \sim 1$ over a very wide range of star-forming environments, from small star clusters in the Milky Way to entire starburst galaxies."
+" The results are subject to considerable uncertainty, but values eg,9210 are strongly excluded by the data."," The results are subject to considerable uncertainty, but values $\epsfftwo \ga 10$ are strongly excluded by the data."
+" However, we lack comparable data for star formation at high redshifts."," However, we lack comparable data for star formation at high redshifts."
+" In this section, we therefore turn to the question of the likely value of eg in high-z giant clumps."," In this section, we therefore turn to the question of the likely value of $\epsff$ in $z$ giant clumps."
+" Unfortunately, we cannot easily apply the direct measurement procedure used to determine eg in the local universe to the high-z clumps, because, while we can evaluate star-formation rates using Ha luminosities, we are limited in our ability to measure the corresponding gas properties."," Unfortunately, we cannot easily apply the direct measurement procedure used to determine $\epsff$ in the local universe to the $z$ clumps, because, while we can evaluate star-formation rates using $\alpha$ luminosities, we are limited in our ability to measure the corresponding gas properties."
+"For example, ? uuse stellar population synthesisto estimate masses and ages for the stellar populations seen in giant clumps in high-z galaxies, and they then compare the ages T to the clump dynamical times, defined as tayn=(αρ)z 0.5tg, where p is taken to be thestellar mass density.","For example, \citet{elmegreen09c} use stellar population synthesisto estimate masses and ages for the stellar populations seen in giant clumps in $z$ galaxies, and they then compare the ages $\tau$ to the clump dynamical times, defined as $t_{\rm dyn} \equiv (G\rho)^{-1/2}\approx 0.5 t_{\rm ff}$ , where $\rho$ is taken to be the mass density."
+"1/3 They find typical values T/tayn~1—10, corresponding to 7T/tg~2—20, with a factor of ~10 scatter."," They find typical values $\tau/t_{\rm dyn}\sim 1-10$, corresponding to $\tau/t_{\rm ff}\sim 2-20$, with a factor of $\sim 10$ scatter."
+" It is tempting to identify 7 with taep and simply estimate eg~tg/T0.05—0.5, but this is likely to be significant overestimate."," It is tempting to identify $\tau$ with $t_{\rm dep}$ and simply estimate $\epsff \sim t_{\rm ff}/\tau \sim 0.05-0.5$, but this is likely to be a significant overestimate."
+" Based on dynamical mass estimates,a ? estimate that gas comprises of the total mass within the disc radius."," Based on dynamical mass estimates, \citet{genzel08a} estimate that gas comprises of the total mass within the disc radius."
+ Estimates based on direct CO measurements indicate gas fractions of 45—60% at z~2 (?7)..," Estimates based on direct CO measurements indicate gas fractions of $45-60\%$ at $z \sim 2$ \citep{daddi09a,tacconi10a}."
+" It is likely to be an even larger fraction of the mass within the dense, rapidly star-forming giant clumps, which are self-gravitating and lack a dark matter component."," It is likely to be an even larger fraction of the mass within the dense, rapidly star-forming giant clumps, which are self-gravitating and lack a dark matter component."
+" If the stars measured by ? ccomprise only a small fraction f. of the total mass in clump, then eg will be reduced by a factor of roughly 15 relative to the previous estimate — one power of f. to account for the gas that has not yet formed stars, and another factor of [0:5 because the free-fall time will be shorter than the value ? eestimate based on the stars alone."," If the stars measured by \citeauthor{elmegreen09c} comprise only a small fraction $f_*$ of the total mass in clump, then $\epsff$ will be reduced by a factor of roughly $f_*^{1.5}$ relative to the previous estimate – one power of $f_*$ to account for the gas that has not yet formed stars, and another factor of $f_*^{0.5}$ because the free-fall time will be shorter than the value \citeauthor{elmegreen09c} estimate based on the stars alone."
+" To give a sense of the possible magnitude of the error, note that ? estimate f«—0.2 for a giant clump in BX 482 (and we argue in Section 4.2 that f. is probably even smaller), and this value of f. would be sufficient to lower the estimate of eg by a factor of 10, to 5x107?—5x 10?. A secondary worry is that, as ?ppoint out, since the clumps are selected using rest- blue light there is a strong bias against selecting older, redder clumps, causing an underestimate of 7."," To give a sense of the possible magnitude of the error, note that \citet{murray09a} estimate $f_* = 0.2$ for a giant clump in BX 482 (and we argue in Section \ref{murraymass} that $f_*$ is probably even smaller), and this value of $f_*$ would be sufficient to lower the estimate of $\epsff$ by a factor of 10, to $5\times 10^{-3} - 5\times 10^{-2}$ A secondary worry is that, as \citeauthor{elmegreen09c}point out, since the clumps are selected using rest-frame blue light there is a strong bias against selecting older, redder clumps, causing an underestimate of $\tau$ ."
+ In, In
+their spectroscopic redshift measurement possible and the use of a photometric technique unnecessary.,their spectroscopic redshift measurement possible and the use of a photometric technique unnecessary.
+ The total fraction of infrared selected galaxies whose radio-to-FIlt. SED shows a clear sign of energetic AGN is only a few percent in the local universe ).. and (his fraction appears to remain roughly (he same at higher redshifts.," The total fraction of infrared selected galaxies whose radio-to-FIR SED shows a clear sign of energetic AGN is only a few percent in the local universe , and this fraction appears to remain roughly the same at higher redshifts."
+ Another important source of uncertainty is poor sampling of the SED., Another important source of uncertainty is poor sampling of the SED.
+ Sampling only the radio part of the SED offers virtually no redshift information. ancl the same is true if only (he subi part of the SED is measured.," Sampling only the radio part of the SED offers virtually no redshift information, and the same is true if only the submm part of the SED is measured."
+ On the other hand. because we are fitting the SED features changing over logarithmic scales. only a few well placed SED data points are needed to derived the redshifts 1999).," On the other hand, because we are fitting the SED features changing over logarithmic scales, only a few well placed SED data points are needed to derived the redshifts ."
+. More than one measurement on both side of the SED trough is highly desirable for a successful photometric redshift determination., More than one measurement on both side of the SED trough is highly desirable for a successful photometric redshift determination.
+" A related problem is the calibration and the relative weighting of individual SED measurements,", A related problem is the calibration and the relative weighting of individual SED measurements.
+ Measurement uncertainties lound in the literature often do not fully account for the overall calibration accuracy. as demonstrated by the magnitude of the scatter and the sizes of the error bars plotted for the SED measurements of Arp 220 in Figure 1.., Measurement uncertainties found in the literature often do not fully account for the overall calibration accuracy as demonstrated by the magnitude of the scatter and the sizes of the error bars plotted for the SED measurements of Arp 220 in Figure \ref{fig:a220sed}.
+ For example. (he IRAS 60 jan and LOO jii measurements found in the Point Source Catalog (PSC) and the Faint Source Catalog (FSC) are known to have svstemie calibration differences of order while the flux densities reported in the FSC often include uncertainties of only a [ew percent.," For example, the IRAS 60 $\mu$ m and 100 $\mu$ m measurements found in the Point Source Catalog (PSC) and the Faint Source Catalog (FSC) are known to have systemic calibration differences of order while the flux densities reported in the FSC often include uncertainties of only a few percent."
+ For deriving the SED template [rom (he 23 inlrared selected starbursts. we adopted flux densities in the FSC! but increased measurement uncertainties to (unless the reported uncertainty. was Lhuwger).," For deriving the SED template from the 23 infrared selected starbursts, we adopted flux densities in the FSC but increased measurement uncertainties to (unless the reported uncertainty was larger)."
+ This re-calibration of the measurement uncertainty is more (han cosmetic since the relative weight of the data points are directly reflected in the 4? minimization., This re-calibration of the measurement uncertainty is more than cosmetic since the relative weight of the data points are directly reflected in the $\chi^2$ minimization.
+ We found this re-weighting to be critically important in utilizing the subnmim measurements., We found this re-weighting to be critically important in utilizing the submm measurements.
+ Similarly. all radio continuum measurements are assigned at least uncertainty in order to account for the overall uneertainty in the flix calibration.," Similarly, all radio continuum measurements are assigned at least uncertainty in order to account for the overall uncertainty in the flux calibration."
+ No effort was made to re-calibrate any of (he subniun data points since most measurements generally curry relatively large fractional uncertainties. but some are clearly under-estimates based on the SED plots such as shown in Figure 1..," No effort was made to re-calibrate any of the submm data points since most measurements generally carry relatively large fractional uncertainties, but some are clearly under-estimates based on the SED plots such as shown in Figure \ref{fig:a220sed}."
+ The same problems also plague the reported SED measurements of many subi galaxies. and they. will inevitably impact the accuracy of the photometric redshifts derived from these data.," The same problems also plague the reported SED measurements of many submm galaxies, and they will inevitably impact the accuracy of the photometric redshifts derived from these data."
+ One way to improve the situation in the future would be employing a frequency. selective bolometer Chat can make simultaneous measurements of several submun bands with accurate relative calibration between the measurement bancs., One way to improve the situation in the future would be employing a frequency selective bolometer that can make simultaneous measurements of several submm bands with accurate relative calibration between the measurement bands.
+presence within maguetised. collapsing cores may lave consequences for star formation dynamics. through their influence on the ambipolay diffusion rate.,"presence within magnetised, collapsing cores may have consequences for star formation dynamics, through their influence on the ambipolar diffusion rate."
+ Molecular anions. including CTlCN . have been considered as plausible carriers for at least some of the unideuti&cd diffuse interstellar bands).," Molecular anions, including $_2$ $^-$, have been considered as plausible carriers for at least some of the unidentified diffuse interstellar bands."
+. Amon abuudances are theorised to be sensitive to clectron attachinent aud photodetachiment rates (see 2007)). aud may therefore provide a useful tool for the determunation of accurate electron densities and cosmüc rav/ploto- ionisatiou rates in astrophlivsical euvironuents.," Anion abundances are theorised to be sensitive to electron attachment and photodetachment rates (see ), and may therefore provide a useful tool for the determination of accurate electron densities and cosmic ray/photo- ionisation rates in astrophysical environments."
+ Preseutlv. many ofthe key reactions aud rates relevaut to molecular anion chemistry. such as those of radiative electron. attaclunent. have not been well-studied in the laboratory. and so are poorly constrained erossly approximated inm chemical models.," Presently, many of the key reactions and rates relevant to molecular anion chemistry, such as those of radiative electron attachment, have not been well-studied in the laboratory, and so are poorly constrained or grossly approximated in chemical models."
+ Caven the current lack of laboratory experimueuts and detailed (quantum) theoretical calculations. the oulv wav to constrain these paramctors and further our understanding of the role of anions in astrophysics is by radio observations and complementary chemical modeling of molecular auious In various astroplvsical cuvirouuents.," Given the current lack of laboratory experiments and detailed (quantum) theoretical calculations, the only way to constrain these parameters and further our understanding of the role of anions in astrophysics is by radio observations and complementary chemical modelling of molecular anions in various astrophysical environments."
+ This Letter reports new detections of the carbon chain anion Col in the vicinity of a voung. low-1uass protostar in the deuse molecular cloud L1251A (in the Cepheus complex). aud in the prestellar core L1512 (iu the Taurus-Auriga complex).," This Letter reports new detections of the carbon chain anion $_6$ $^-$ in the vicinity of a young, low-mass protostar in the dense molecular cloud L1251A (in the Cepheus complex), and in the prestellar core L1512 (in the Taurus-Auriga complex)."
+ Sources were selected. frou: a set of sixteen dense clouds for which Πουν /=109 maps had previously been obtained using the Onsala 20-1 telescope between 2005 and 2007 (some of which were published by 200633., Sources were selected from a set of sixteen dense clouds for which $_3$ N $J=10-9$ maps had previously been obtained using the Onsala 20-m telescope between 2005 and 2007 (some of which were published by ).
+ The close chemical relationship between polyvucs and cvanopolvues (see 1)). suggests that CU and 1 should be abundant in dense molecular clouds. close to where the Που oeaks.," The close chemical relationship between polyynes and cyanopolynes (see ), suggests that $_4$ H and $_6$ H should be abundant in dense molecular clouds, close to where the $_3$ N peaks."
+ Therefore. to iuprove the chances of successtully detecting CIL Coll and ο compared with previous stinveys16]||eupQ9.. we arected the strongest ουν ciission peaks in hese anaps. which. due to chemical differentiation are not eoncrally coicident with the moaiuDodyCitationEud66locatious of peak LI]buctüUG.bor0T.eissiou Youn commonly observed molecules such as NID. NoIT! or CO (for which maps already. exist in the howe," Therefore, to improve the chances of successfully detecting $_4$ H, $_6$ H and $_6$ $^-$ compared with previous surveys, we targeted the strongest $_3$ N emission peaks in these maps, which, due to chemical differentiation, are not generally coincident with the locations of peak emission from commonly observed molecules such as $_3$, $_2$ $^+$ or CO (for which maps already exist in the literature)."
+ Ousala HC4N 7=10.9 inaps of L1251À. and L15ha are shown in Figure l1., Onsala $_3$ N $J=10-9$ maps of L1251A and L1512 are shown in Figure \ref{fig:map}.
+ The adopted coordinates our anion and carbon chain searches L125dA: RA(2000) = 22:30: 10.1.DEC(2000) = |DEx16: omnL151The RA(2000) = 5moOL0r.l. DEC(2000) = ©:of»," The adopted coordinates for our anion and carbon chain searches were: L1251A: RA(2000) = 22:30:40.4, DEC(2000) = +75:13:46; L1512: RA(2000) = 5:04:07.1, DEC(2000) = 32:43:09."
+" L1I251A position lies LO” from the centre the Class Ü protostar L1251À IRS3. aud the L1512 position lies 26 from the ceutre of DU.a (2 120""7-wide) pre-stellar core2005)."," The L1251A position lies $40''$ from the centre of the Class 0 protostar L1251A IRS3, and the L1512 position lies $26''$ from the centre of a $\approx120''$ -wide) pre-stellar core."
+. The positious observed by (scho did not detect CoM in either cloud). are uot coiekleut with our observed positions.," The positions observed by (who did not detect $_6$ $^-$ in either cloud), are not coincident with our observed positions."
+ Observations were carried out iu the mouths of April and June 2010 using the NRAO LOO-neter Creen Banka, Observations were carried out in the months of April and June 2010 using the NRAO 100-meter Green Bank.
+nd The Ίνα receiver was used with 50 AMIIz bandwidth 8192 channels (corresponding to achanuuel spaciug of z:6.065 1)). in each of our spectral windows.," The Ka receiver was used with 50 MHz bandwidth and 8192 channels (corresponding to achannel spacing of $\approx0.065$ ), in each of four spectral windows."
+" In the middle of the observed requencey range (28 GIIz). the telescope beam FWIIM was 267"" and the beam effcieucv factor was 0.88."," In the middle of the observed frequency range (28 GHz), the telescope beam FWHM was $26''$ and the beam efficiency factor was 0.88."
+ We performed deep iuteerations to search for cussion ines from ΟΠ. οσα. ο and Ποστ. Additional observations were obtained of IICSN. C35 aud CIT4CCIT.," We performed deep integrations to search for emission lines from $_4$ H, $_6$ H, $_6$ $^-$ and $_7$ N. Additional observations were obtained of $_5$ N, $_3$ S and $_3$ CCH."
+ For the compact. spatially isolated source L1512. beam switching (with 78” throw) was used. aud for the nore extended source L1215À. frequency switching was used.," For the compact, spatially isolated source L1512, beam switching (with $78''$ throw) was used, and for the more extended source L1215A, frequency switching was used."
+ Pointing was checked every one to two hours aud was typically accurate to within 5”., Pointing was checked every one to two hours and was typically accurate to within $5''$.
+ Total svste1 eniperatures were iu the rauge 10 to 60 Is. New maps of IIC4N J=10.9 eission in LI251À aud ΟΙ =98 cinission iu L1512 were obtaiued using he Onsala 20 11 telescope in May 2010 (shown in Figure 1))., Total system temperatures were in the range 40 to 60 K. New maps of $_3$ N $J=10-9$ emission in L1251A and $_4$ H $N=9-8$ emission in L1512 were obtained using the Onsala 20 m telescope in May 2010 (shown in Figure \ref{fig:map}) ).
+ The basic observation aud reduction tecliuiques were xeseuted by(2006)., The basic observation and reduction techniques were presented by.
+. The CoM and CSI spectra observed in L1251À aud L1512are shown in Figure 2.., The $_6$ $^-$ and $_6$ H spectra observed in L1251A and L1512are shown in Figure \ref{fig:spectra}.
+ To improve the signalto-noise ratio for LI251À. the ο 7=10.9 and J—1l 10spectra have been averaged in velocity space.," To improve the signal-to-noise ratio for L1251A, the $_6$ $^-$ $J=10-9$ and $J=11-10$ spectra have been averaged in velocity space."
+ The observed spectral lines were least-squares fitted using sinele Gaussians. the parameters for which are even in Table 1..," The observed spectral lines were least-squares fitted using single Gaussians, the parameters for which are given in Table \ref{tab:lines}."
+ The five hyperfine peaks of the IIC'SN J=32 transition are wellvesolved in our spectra. from which we derive gas rotational fenüperatures 1). of6.20 I& for L1251Àexcitation audSavage 8.7|etd0.7 Is for L1512Thes (ising Equations 1l aud 2 of 2," The five hyperfine peaks of the $_3$ N $J=3-2$ transition are well-resolved in our spectra, from which we derived gas rotational excitation temperatures $T_{ex}$ ) of $6.2\pm0.3$ K for L1251A and $8.7\pm0.7$ K for L1512 (using Equations 1 and 2 of )."
+002)). temperatures were used iu the caleulation of the column deusities of the observed species assumius LTE and optically thin ciissiou (see for example. Equation 2 of 2002)).," These temperatures were used in the calculation of the column densities of the observed species assuming LTE and optically thin emission (see for example, Equation 2 of )."
+ Where uniltiple transitious were observed for a given species. the average of the calculated cohuun densities was taken. except for IIEC3N. for which only the optically thin AF’=0 transitions were used.," Where multiple transitions were observed for a given species, the average of the calculated column densities was taken, except for $_3$ N, for which only the optically thin $\Delta F=0$ transitions were used."
+ The HIC5N and other recorded spectra will be presented in a future article., The $_3$ N and other recorded spectra will be presented in a future article.
+ Coluun densities are given in Table 2, Column densities are given in Table \ref{tab:colds}.
+ Large abundances of carbou-chain-bearing species were observed in both LI251À aud L1512., Large abundances of carbon-chain-bearing species were observed in both L1251A and L1512.
+ The measured 1 and Coll colin deusities (see Table 2)) are similar to those in iuterstcllar clouds with the highest kuowu carbon chain abuudances in the Galaxy 2, The measured $_4$ H and $_6$ H column densities (see Table \ref{tab:colds}) ) are similar to those in interstellar clouds with the highest known carbon chain abundances in the Galaxy .
+009).. The:NUM ofcarbon chaius is also lighliehted iu L1512 by the large CIL;CCTII column deusity. which is comparable to that found in L1527 aud TMC-1 2008)).," The abundance of carbon chains is also highlighted in L1512 by the large $_3$ CCH column density, which is comparable to that found in L1527 and TMC-1 ."
+. The €5S cohunn deusitics in both. LI251À and L1512 are somewhat less than previously observed m carbou-chain-producing regions’ by (2006)., The $_3$ S column densities in both L1251A and L1512 are somewhat less than previously observed in `carbon-chain-producing regions' by .
+. Possible explanatious for this include depletion outo dust or reduced clemental sulphur abundances., Possible explanations for this include depletion onto dust or reduced elemental sulphur abundances.
+ The derivedIT columnu densities of 2740.7&1027 2 im C'sL1251À and 2.04+0.5&1029 cu? , The derived$_6$ $^-$ column densities of $\pm0.7\times10^{10}$ $^{-2}$ in L1251A and $\pm0.5\times10^{10}$ $^{-2}$ 
+discovered in the solar surface rotation rate (Howard LaBonte 1980: Ulrich et al.,discovered in the solar surface rotation rate (Howard LaBonte 1980; Ulrich et al.
+ 1983: snocderass 1992)., 1988; Snodgrass 1992).
+ On the other hand. at high latitudes (245°) these bands move poleward with lime (Antia Basu 2001: Ulrich. 2001).," On the other hand, at high latitudes $> 45^\circ$ ) these bands move poleward with time (Antia Basu 2001; Ulrich 2001)."
+ With acciumnulation of more seismic data il became clear that these zonal flows penetrate through most of the convection zone (Vorontsov el al., With accumulation of more seismic data it became clear that these zonal flows penetrate through most of the convection zone (Vorontsov et al.
+ 2002; Basu Antia 2003. 2006: Howe et al.," 2002; Basu Antia 2003, 2006; Howe et al."
+ 2006; Antia et al., 2006; Antia et al.
+ 2003)., 2008).
+ The zonal-flow pattern is correlated to the magnetic butterfly diagram (Antia et al., The zonal-flow pattern is correlated to the magnetic butterfly diagram (Antia et al.
+ 2008; Sivaraman et al., 2008; Sivaraman et al.
+ 2005) and is thought to be intimately connected with the solar dvnaimo., 2008) and is thought to be intimately connected with the solar dynamo.
+ It is believed lo arise from a nonlinear interaction between differential rotation and magnetic field (e.g. Covas οἱ al.," It is believed to arise from a nonlinear interaction between differential rotation and magnetic field (e.g., Covas et al."
+ 2000: Rempel 2006)., 2000; Rempel 2006).
+ Hence the properties of solar zonal flows provide a strong constraint on the solar dvnamo models., Hence the properties of solar zonal flows provide a strong constraint on the solar dynamo models.
+ This makes it crucial to detect. and understand evele-lo-cvele changes of solar zonal flows., This makes it crucial to detect and understand cycle-to-cycle changes of solar zonal flows.
+ This would provide the necessary input in came models to understand ancl predict the variation in activity levels [rom cycle to evele., This would provide the necessary input in dynamo models to understand and predict the variation in activity levels from cycle to cycle.
+ In this work we study. the temporal variations in solar rotation rate with particular emphasis on the differences between (he minima of eveles 23 and 24., In this work we study the temporal variations in solar rotation rate with particular emphasis on the differences between the minima of cycles 23 and 24.
+ We also study. the changes in the outer shear laver of the Sun., We also study the changes in the outer shear layer of the Sun.
+ We also use the available data (ο estimate the length of solar evele 23 using a time-shilted auto-correlation., We also use the available data to estimate the length of solar cycle 23 using a time-shifted auto-correlation.
+ This enables us to identilv the epoch of latest minimum from (he point of view of solar dvnanmies and (his may be useful in prediction of activity during the next cvele., This enables us to identify the epoch of latest minimum from the point of view of solar dynamics and this may be useful in prediction of activity during the next cycle.
+ We use data obtained by the GONG (Lill et al., We use data obtained by the GONG (Hill et al.
+ 1996) and MDI (Schou 1999) projects for this work., 1996) and MDI (Schou 1999) projects for this work.
+ These data sets consist of the mean frequency ancl (he splitting coefficients ofdifferent. (η.£) multiplets.," These data sets consist of the mean frequency and the splitting coefficients ofdifferent $(n,\ell)$ multiplets."
+ Onlv the odcd-order splitting coellicients are needed to determine the rotation rate in the solar interior (e.g.. Ritzwoller Lavelv 1991).," Only the odd-order splitting coefficients are needed to determine the rotation rate in the solar interior (e.g., Ritzwoller Lavely 1991)."
+ We use 145 data sets from GONG. each set covering a period of LOS davs.," We use 145 data sets from GONG, each set covering a period of 108 days."
+ The first set starts on 1995. May 7 and the last set ends on 2009. October 31. with a spacing of 36 davs between consecutive data sels.," The first set starts on 1995, May 7 and the last set ends on 2009, October 31, with a spacing of 36 days between consecutive data sets."
+ Thus these sets cover about a vear of data leading to the minimum of evele 23 as well as several months of data following the minimum of evele 24., Thus these sets cover about a year of data leading to the minimum of cycle 23 as well as several months of data following the minimum of cycle 24.
+ The MDI data consist of 69 non-overlapping sets each obtained from observations taken over a period of 72 davs., The MDI data consist of 69 non-overlapping sets each obtained from observations taken over a period of 72 days.
+ The first set begins on 1996. May 1: and the last set ends on 2010. April 29.," The first set begins on 1996, May 1 and the last set ends on 2010, April 29."
+ The MDI data start close to the minimum of cevele 23 and do not cover preceding period. while (hese data cover a period of about a vear following the minimum of evcle 24.," The MDI data start close to the minimum of cycle 23 and do not cover preceding period, while these data cover a period of about a year following the minimum of cycle 24."
+ We use the two dimensional Reeularizecd Least Squares (2D RLS) inversion technique as described by Antia οἱ al. (, We use the two dimensional Regularized Least Squares (2D RLS) inversion technique as described by Antia et al. (
+1998) to infer the rotation rate in the solar interior [rom each of the available data sets.,1998) to infer the rotation rate in the solar interior from each of the available data sets.
+ The first eight odd-order splitting coellicients. (4.603..... (45. Were," The first eight odd-order splitting coefficients, $a_1,a_3,\ldots,a_{15}$ , were"
+Large separation lenses. on the other hand. were modeled by ? using the triaxial NFW model based on simulations by ?..,"Large separation lenses, on the other hand, were modeled by \cite{oguri-keeton04} using the triaxial NFW model based on simulations by \cite{jing02}."
+ With an inner slope a=1.0 (pure NFW). the image multiplicities are shown to be remarkably ditferent from those of galaxy lensing: cusp systems have the highest probability. followed by quads. with doubles being least common.," With an inner slope $\alpha = 1.0$ (pure NFW), the image multiplicities are shown to be remarkably different from those of galaxy lensing: cusp systems have the highest probability, followed by quads, with doubles being least common."
+" However. because the overall lensing probabilities are orders of magnitude smaller than that of galaxy lensing. such dark-matter dominated lenses are relatively rare—the only observed lens in this regime is a quad at 14.62"". found in the Sloan Digital Sky Survey (2)."," However, because the overall lensing probabilities are orders of magnitude smaller than that of galaxy lensing, such dark-matter dominated lenses are relatively rare—the only observed lens in this regime is a quad at $14.62''$, found in the Sloan Digital Sky Survey \cite{inada03}) )."
+" In contrast to the paucity of confirmed. large-separation lenses. a significant fraction of observed. gravitational lenses lie in the intermediate region between these two extremes. with image separations roughly between |"" and 10""."," In contrast to the paucity of confirmed large-separation lenses, a significant fraction of observed gravitational lenses lie in the intermediate region between these two extremes, with image separations roughly between $1''$ and $10''$."
+ The total lensing probability in this region has been studied for the case of spherical mass distributions (2... 2... 25). but the distribution of image multiplicities for nonspherical lenses has not yet been derived.," The total lensing probability in this region has been studied for the case of spherical mass distributions \cite{oguri06}, \cite{ma03}, \cite{kochanek01}) ), but the distribution of image multiplicities for nonspherical lenses has not yet been derived."
+ We note that none of the single-component elliptical models mentioned thus far can be applied in this regime. since both the dark matter and the central galaxy contribute significantly at these separations.," We note that none of the single-component elliptical models mentioned thus far can be applied in this regime, since both the dark matter and the central galaxy contribute significantly at these separations."
+ Furthermore. any lensing model including baryons and dark matter must take into account the moditication of dark matter halos due to baryon cooling. which increases the concentration of the inner parts ofhalos (2.. 25).," Furthermore, any lensing model including baryons and dark matter must take into account the modification of dark matter halos due to baryon cooling, which increases the concentration of the inner parts ofhalos \cite{kochanek01}, \cite{keeton01}) )."
+ Although a multi-component model was employed by ? using galaxies with ellipticity. the dark matter was modeled by spherical NFW halos. ignoring adiabatic contraction and the riaxiality of dark matter halos.," Although a multi-component model was employed by \cite{moeller07} using galaxies with ellipticity, the dark matter was modeled by spherical NFW halos, ignoring adiabatic contraction and the triaxiality of dark matter halos."
+ In this paper we derive a distribution of image multiplicities hat will be approximately valid for all observed image separations. including the intermediate region.," In this paper we derive a distribution of image multiplicities that will be approximately valid for all observed image separations, including the intermediate region."
+ However. to accurately model his region. we must combine the triaxial NFW model. adiabatic contraction. and a central lens galaxy.," However, to accurately model this region, we must combine the triaxial NFW model, adiabatic contraction, and a central lens galaxy."
+ The lensing properties hen become sensitive to the ellipticities. concentrations. and orientations of both galaxy and halo. in addition to the fraction of the total mass contained in the galaxy.," The lensing properties then become sensitive to the ellipticities, concentrations, and orientations of both galaxy and halo, in addition to the fraction of the total mass contained in the galaxy."
+ This requires a more complicated calculation of lensing probabilities than those in the literature., This requires a more complicated calculation of lensing probabilities than those in the literature.
+ Despite the complexity of the models considered in this work. several important factors (considered. prevoiously in isolation) have to be neglected.," Despite the complexity of the models considered in this work, several important factors (considered prevoiously in isolation) have to be neglected."
+ In particular. we highlight the absence of dark matter substructure in our model and the effect of the lens galaxy environment.," In particular, we highlight the absence of dark matter substructure in our model and the effect of the lens galaxy environment."
+ We will later discuss the expected impact of these factors., We will later discuss the expected impact of these factors.
+ We tind several interesting features that suggest new methods of constraining Jens mass distributions., We find several interesting features that suggest new methods of constraining lens mass distributions.
+ For example. the ratio of image multiplicities is found to be sensitive to the alignment of the mayor axes of galaxies with that of their host halos.," For example, the ratio of image multiplicities is found to be sensitive to the alignment of the major axes of galaxies with that of their host halos."
+ Further. the transition from small- to large-separation lenses is marked by a “crossing” of lens probabilities when the expected lensing rate of quads overtakes that of doubles. and likewise for cusps.," Further, the transition from small- to large-separation lenses is marked by a “crossing” of lens probabilities when the expected lensing rate of quads overtakes that of doubles, and likewise for cusps."
+ Understanding the precise location of these turnover points may provide further means of constraining galaxy-halo mass distributions once the sample of observed lens systems becomes sutticiently large., Understanding the precise location of these turnover points may provide further means of constraining galaxy-halo mass distributions once the sample of observed lens systems becomes sufficiently large.
+ The paper is organized as follows., The paper is organized as follows.
+ In $2? we will present a general form for the image separation distribution. and discuss the relevant model parameters.," In \ref{sec:imgdist} we will present a general form for the image separation distribution, and discuss the relevant model parameters."
+ Section ?? will describe the lens model we used and outline our approach toward applying the adiabatic contraction model to triaxial halos.," Section \ref{sec:lensmodel}
+ will describe the lens model we used and outline our approach toward applying the adiabatic contraction model to triaxial halos."
+ Section 2? covers the details of the calculation. and results are given in 322..," Section \ref{sec:calc} covers the details of the calculation, and results are given in \ref{sec:results}."
+ Finally we summarize and discuss our results in $??.., Finally we summarize and discuss our results in \ref{sec:discussion}. .
+" The differential image separation distribution is usually defined as the probability that a source object at redshift z, will be multiply imaged with an image separation Ο.", The differential image separation distribution is usually defined as the probability that a source object at redshift $z_s$ will be multiply imaged with an image separation $\Theta$.
+ This is written as πα 4. Here Bo is the biased lensing cross section (explained in refsec:cale)) in units of solid angle. and z is the comoving number density of halos.," This is written as \cite{turner84}, \cite{schneider06}) ), Here $B\sigma$ is the biased lensing cross section (explained in \\ref{sec:calc}) ) in units of solid angle, and $n$ is the comoving number density of halos."
+" The lensing optical depth is given by the integral of this distribution over all image separations: in this paper we will simply refer to it as the total lens probability. though always with respect to a given source redshift z,."," The lensing optical depth is given by the integral of this distribution over all image separations; in this paper we will simply refer to it as the total lens probability, though always with respect to a given source redshift $z_s$."
+ The i; are the parameters relevant to lensing. whose distributions p(y;) are integrated over. and which depend on the particular lens model chosen.," The $\chi_i$ are the parameters relevant to lensing, whose distributions $p(\chi_i)$ are integrated over, and which depend on the particular lens model chosen."
+ In this study we take our model to be an adiabatically-contracted NFW dark matter halo with a central early-type galaxy represented by a Hernquist protile., In this study we take our model to be an adiabatically-contracted NFW dark matter halo with a central early-type galaxy represented by a Hernquist profile.
+" The lensing parameters are then where c, is the triaxial concentration parameter (defined below) of the original dark matter halo before adiabatic contraction. afc.bíc are the axis ratios of the halo. (9.6) are the orientation angles of the halo. /; is the baryon mass fraction. xi=Εν gives the scale radius of the Hernquist profile. gj, is the projected axis ratio of the Hernquist profile. and 4, is the angle between the projected major axis of the Hernquist profile and the projected major axis of the dark matter halo."," The lensing parameters are then where $c_{e}$ is the triaxial concentration parameter (defined below) of the original dark matter halo before adiabatic contraction, $a/c, b/c$ are the axis ratios of the halo, $\theta,\phi$ ) are the orientation angles of the halo, $f_b$ is the baryon mass fraction, $x_b = r_b / r_s$ gives the scale radius of the Hernquist profile, $q_H$ is the projected axis ratio of the Hernquist profile, and $\theta_q$ is the angle between the projected major axis of the Hernquist profile and the projected major axis of the dark matter halo."
+ The total mass of the lens was not included among the parameters listed above because for à given image separation O. lens redshift z and set of lensing parameters 4;. the mass is uniquely fixed.," The total mass of the lens was not included among the parameters listed above because for a given image separation $\Theta$, lens redshift $z$ and set of lensing parameters $\chi_i$, the mass is uniquely fixed."
+ In this section we deseribe the model for the lenses telliptical galaxies and clusters)., In this section we describe the model for the lenses (elliptical galaxies and clusters).
+ We split the discussion into four parts dealing with the dark matter halo. the baryonic component. etfect of baryon cooling on the dark matter halo and finally correlations between the luminous and dark matter mass distributions.," We split the discussion into four parts dealing with the dark matter halo, the baryonic component, effect of baryon cooling on the dark matter halo and finally correlations between the luminous and dark matter mass distributions."
+ Throughout this paper we assume a concordance cosmology with iA34=0.3. O4=0.7. h=0.71. and oy=0.9 and use the (?)) mass function of dark matter halos.," Throughout this paper we assume a concordance cosmology with $\Omega_M = 0.3$, $\Omega_\Lambda = 0.7$, $h = 0.71$, and $\sigma_8 = 0.9$ and use the Seth-Tormen \cite{sheth02})) mass function of dark matter halos."
+ We assume the density profile of a triaxial dark matter halo to be (2). where The triaxial concentration parameter is detined in ?. as R./Rg. where R. is defined so that the mean density within an," We assume the density profile of a triaxial dark matter halo to be \cite{jing02}) ), where The triaxial concentration parameter is defined in \cite{jing02} as $c_e = 
+R_e/R_0$ , where $R_e$ is defined so that the mean density within an"
+numerical integrations indicate that the first sign of a Revnolds number larger than unity appears [or an object with a radius of 10 em.,numerical integrations indicate that the first sign of a Reynolds number larger than unity appears for an object with a radius of 10 cm.
+" Since for an object at (1.25.0.125) AU, p.* varies between 0.512 and 1. the graphs of 7?—42. as shown in Figure 12. indicate that [or a centimeter-sized object. 47 is negative."," Since for an object at (1.25,0.125) AU, ${\hat r}^{-3}$ varies between 0.512 and 1, the graphs of ${{\hat r}^2}-{{\hat \omega}^2}$, as shown in Figure 12, indicate that for a centimeter-sized object, ${\hat \omega}^2$ is negative."
+ A negative value for w? is an indication of an overdamped motion., A negative value for ${\hat \omega}^2$ is an indication of an overdamped motion.
+ Increasing the density of the solid while keeping the gas temperature constant. «7 becomes larger and approaches zero (critically damped) indicating that the overdamped nature of the vertical motion of the particle weakens aud it may even change into an underdamped motion where « is positive.," Increasing the density of the solid while keeping the gas temperature constant, ${\hat \omega}^2$ becomes larger and approaches zero (critically damped) indicating that the overdamped nature of the vertical motion of the particle weakens and it may even change into an underdamped motion where ${\hat \omega}^2$ is positive."
+ From equation (23). the overdamping nature of the vertical motion becomes even more pronounced when the radius of the object is smaller than 1 cm.," From equation (23), the overdamping nature of the vertical motion becomes even more pronounced when the radius of the object is smaller than 1 cm."
+ That means. the rate of verlical migration of a solid decreases by decreasing its size (Figure 9).," That means, the rate of vertical migration of a solid decreases by decreasing its size (Figure 9)."
+ The left column of Figure 12 also shows that [or a constant value of the solid's density. increasing (he gas temperature results in lower values for 47.," The left column of Figure 12 also shows that for a constant value of the solid's density, increasing the gas temperature results in lower values for ${\hat \omega}^2$."
+ That means. at higher temperatures. the overdamping of the vertical motion of the object becomes more pronounced (top eraph of Figure 13).," That means, at higher temperatures, the overdamping of the vertical motion of the object becomes more pronounced (top graph of Figure 13)."
+ This can be explained noting (hat from equation (5) and as shown in Figure 15. at anv radial distance r. the vertical distribution of the gas maximizes al 2=0 (ie. on the midplane).," This can be explained noting that from equation (5) and as shown in Figure 15, at any radial distance $r$, the vertical distribution of the gas maximizes at $z=0$ (i.e., on the midplane)."
+ As shown in Figure 14. the width of the gas distribution funetion along {he z-axis increases by increasing the gas temperature.," As shown in Figure 14, the width of the gas distribution function along the $z$ -axis increases by increasing the gas temperature."
+ As a result of this broadening. the resistive effect of gas drag on the vertical motion of solids appear in a larger range of variable 2 which results in a longer Gime lor descending toward (he midplane.," As a result of this broadening, the resistive effect of gas drag on the vertical motion of solids appear in a larger range of variable $z$ which results in a longer time for descending toward the midplane."
+ In (his ease. equation (20) can be written as Ip 5pl sinilar to (he previous case. in (he vicinitv of the midplane. equation (24) can be," In this case, equation (20) can be written as 1pt 5pt Similar to the previous case, in the vicinity of the midplane, equation (24) can be"
+"In thisLetter, we analyze an HVC along the sight line to the blazar Mrk 421 and set a tight lower limit to the metallicity of Complex M. Savageetal.(2005) detected the HVC in1,i, and aabsorption lines in theFUSE spectrum, but they focused on the origin of the extended red wing of aabsorption.","In this, we analyze an HVC along the sight line to the blazar Mrk 421 and set a tight lower limit to the metallicity of Complex M. \citet{sav05} detected the HVC in, and absorption lines in the spectrum, but they focused on the origin of the extended red wing of absorption."
+" Throughout, we quote errors at lo and upper limits at 3c confidence levels."," Throughout, we quote errors at $1\sigma$ and upper limits at $3\sigma$ confidence levels."
+" We adopt atomic data from Morton(2003) and solar abundances from Asplundetal. (2009)., in particular, log(X/H)c-1228.43, 7.83, 8.69, 7.51, 7.12, and 7.50 for C, N, O, Si, S, and Fe, respectively."," We adopt atomic data from \citet{mor03} and solar abundances from \citet{asp09}, , in particular, $\log (X/H)_\odot + 12 =8.43$, $7.83$, $8.69$, $7.51$, $7.12$, and $7.50$ for C, N, O, Si, S, and Fe, respectively."
+" Our data come from far ultraviolet spectroscopic observations of the background target Mrk 421 (L,b=179.83°,65.03°; z= 0.03) obtained with the Cosmic Origins Spectrograph (COS; Ostermanetal. 2011)) aboard theTelescope (HST) and with theExplorer (FUSE)."," Our data come from far ultraviolet spectroscopic observations of the background target Mrk 421 $l,b=179.83^\circ, 65.03^\circ$; $z=0.03$ ) obtained with the Cosmic Origins Spectrograph (COS; \citealt{ost11}) ) aboard the ) and with the )."
+" Mrk 421 is one of the Guaranteed Time Observation (GTO; PI: Green) targets of COS, which has two medium-resolution gratings G130M and G160M, covering wavelengths 1150—1450 aand 1405-1775 wwith a spectral resolution ~18,000."," Mrk 421 is one of the Guaranteed Time Observation (GTO; PI: Green) targets of COS, which has two medium-resolution gratings G130M and G160M, covering wavelengths $1150-1450$ and $1405-1775$ with a spectral resolution $\sim 18,000$."
+" TheFUSE data cover 905-1187 wwith a resolution ~20,000."," The data cover $905-1187$ with a resolution $\sim20,000$."
+ COS observed Mrk 421 on 2009 November 24 (PID 11520) with total exposures of 1.7 ks (G130M) and 1.8 ks (G160M)., COS observed Mrk 421 on 2009 November 24 (PID 11520) with total exposures of 1.7 ks (G130M) and 1.8 ks (G160M).
+FUSE observed the source on 2000 December 1 (PID P101) and 2003 January 19 (PID Z010) with a total exposure of 85.0 ks., observed the source on 2000 December 1 (PID P101) and 2003 January 19 (PID Z010) with a total exposure of 85.0 ks.
+ The observations and calibrated data were retrieved from the Multimission Archive (MAST2)) at the Space Telescope Science Institute., The observations and calibrated data were retrieved from the Multimission Archive ) at the Space Telescope Science Institute.
+ The COS data were further processed following the procedure described in Danforthetal. (2, The COS data were further processed following the procedure described in \citet{dan10}.
+"010).. theFUSE LiFla and SiC2a data were used, reduced Onlyfurther as described in Danforthetal.(2006)."," Only the LiF1a and SiC2a data were used, reduced further as described in \citet{dan06}."
+". Individual exposures were weighted by their exposures and coadded to form final spectra, whose signal-to-noise ratios per resolution element range from 25-35 (COS) and 5-15 (FUSE)."," Individual exposures were weighted by their exposures and coadded to form final spectra, whose signal-to-noise ratios per resolution element range from 25–35 (COS) and 5--15 )."
+ We find «596 flux fluctuation in the COS spectra arising from unremoved small fixed-pattern features., We find $<5\%$ flux fluctuation in the COS spectra arising from unremoved small fixed-pattern features.
+" These fluctuations, compared to systematic uncertainties in wavelength calibration, are not important in measurements described below."," These fluctuations, compared to systematic uncertainties in wavelength calibration, are not important in measurements described below."
+" The wavelengths of these spectra were calibrated with respect to the 221 cm emission data from Leiden/Argentine/Bonn (LAB), which have an angular resolution of ~0.6?, spectral resolution of 1.3kms!, and rms brightness-temperature noise of 70—90 mK (Kalberlaetal.2005).."," The wavelengths of these spectra were calibrated with respect to the 21 cm emission data from Leiden/Argentine/Bonn (LAB), which have an angular resolution of $\sim0.6^\circ$, spectral resolution of $1.3~{\rm km~s^{-1}}$, and rms brightness-temperature noise of 70–90 mK \citep{kal05}."
+" Four LAB pencil beams adjacent to the Mrk 421 sight line were extracted, averaged, and weighted by the inverse square of their angular separations from the sight line."," Four LAB pencil beams adjacent to the Mrk 421 sight line were extracted, averaged, and weighted by the inverse square of their angular separations from the sight line."
+ An IVC component is clearly resolved from the zero-velocity component in the final emission spectrum., An IVC component is clearly resolved from the zero-velocity component in the final emission spectrum.
+" We use a Gaussian fit to the IVC and obtain its LSR velocity, column density, and Doppler width: vin?2-60kms!, MOY=6.82x10?cm, and bi=V20217.5kms""! (Fig."," We use a Gaussian fit to the IVC and obtain its LSR velocity, column density, and Doppler width: $v_{\rm HI}^{(\rm IVC)} = -60~{\rm km~s^{-1}}$, $N_{\rm HI}^{(\rm IVC)} = 6.82\times10^{19}~{\rm cm^{-2}}$, and $b_{\rm HI}^{(\rm IVC)} = \sqrt2\sigma=17.5~{\rm km~s^{-1}}$ (Fig."
+ 2aa)., \ref{fig:MRKa}a a).
+" The IVC can also be resolved in lines of À1259.529 (COS), A1020.699(FUSE LiFla), and A929.517 (FUSE SiC2a) — see 2"," The IVC can also be resolved in lines of $\lambda1259.529$ (COS), $\lambda 1020.699$ LiF1a), and $\lambda 929.517$ SiC2a) – see Figures."
+ and 3.., \ref{fig:MRKa} and \ref{fig:MRKb}. .
+ We shift the ultraviolet spectra to align theirFigures. IVC velocities with the 21 cm data., We shift the ultraviolet spectra to align their IVC velocities with the 21 cm data.
+" In the shifted spectra, an HVC at 131kms! is visible inSin,OLNL andCumu aat >3c significance, and inIv,,vi, and aat 2-3 c significance."," In the shifted spectra, an HVC at ${\rm -131~km~s^{-1}}$ is visible in, and at $>3\sigma$ significance, and in, and at 2–3 $\sigma$ significance."
+ All line centroids are well aligned within ~15kms! (Figs., All line centroids are well aligned within $\sim15~{\rm km~s^{-1}}$ (Figs.
+ 2 and 3))., \ref{fig:MRKa} and \ref{fig:MRKb}) ).
+" The HVC, however, is not visible in the 21 cm data (Fig."," The HVC, however, is not visible in the 21 cm data (Fig."
+ 2aa)., \ref{fig:MRKa}a a).
+ We use the model developed in Yaoetal.(2010) to measure absorption of the HVC., We use the model developed in \citet{yao10} to measure absorption of the HVC.
+ This model uses ionic column density (N) and Doppler width (bp) as fitting parameters and adopts a Voigt function to approximate the absorption-line profile., This model uses ionic column density $N$ ) and Doppler width $b_D$ ) as fitting parameters and adopts a Voigt function to approximate the absorption-line profile.
+ This model can beused to fit individual absorption lines or to jointly analyze multiple absorption lines., This model can beused to fit individual absorption lines or to jointly analyze multiple absorption lines.
+ We include instrumental line spread functions (LSFs) in our spectral fit by using the wavelength-dependent LSF of COS and a Gaussian with FWHM 20kms'! across the entire wavelength range ofFUSE observations., We include instrumental line spread functions (LSFs) in our spectral fit by using the wavelength-dependent LSF of COS and a Gaussian with FWHM $20~{\rm km~s^{-1}}$ across the entire wavelength range of observations.
+" Because hhas multiple transitions (Fig. 2),"," Because has multiple transitions (Fig. \ref{fig:MRKa}) ),"
+ we jointly analyze these lines to reduce statistical uncertainty., we jointly analyze these lines to reduce statistical uncertainty.
+ This method is equivalent to a curve-of-growth analysis but propagates errors automatically., This method is equivalent to a curve-of-growth analysis but propagates errors automatically.
+ The consistent Doppler widths measured in aand (11-14 s) indicate that non-thermal broadening dominates the line width., The consistent Doppler widths measured in and (11–14 ) indicate that non-thermal broadening dominates the line width.
+" To measure column densities or upper limits for ions whose HVC components are unresolved or whose line widths are unconstrained, we fix their line centroids and their widths to those ofu."," To measure column densities or upper limits for ions whose HVC components are unresolved or whose line widths are unconstrained, we fix their line centroids and their widths to those of."
+. Our results are presented in Table 1.., Our results are presented in Table \ref{tab:results}.
+" Because aand iin the HVC are severely blended with the IVC components, we have notmeasured their column densities."," Because and in the HVC are severely blended with the IVC components, we have notmeasured their column densities."
+ We are able to measure Ny in the high Lyman lines of the HVC by subtracting out aabsorption from the IVC and low-velocity gas., We are able to measure $N_{\rm HI}$ in the high Lyman lines of the HVC by subtracting out absorption from the IVC and low-velocity gas.
+"Both components are saturated, but only the IVC (-60 s7!)) contributes to the absorption in higher Lyman lines atthe(—131 s.) velocity of the HVC (Fig. 3)).","Both components are saturated, but only the IVC $-60$ ) contributes to the absorption in higher Lyman lines atthe$-131$ ) velocity of the HVC (Fig. \ref{fig:MRKb}) )."
+" Since we have well-determined values of bo, and NÀY© from 21 cm data 2)), we fix its absorptionvro, in the total"," Since we have well-determined values of $v_{\rm HI}^{(\rm IVC)}$ , $b_{\rm HI}^{(\rm IVC)}$ , and $N_{\rm HI}^{(\rm IVC)}$ from 21 cm data \ref{sec:obs}) ), we fix its absorption in the total"
+H-boaad spectra probably requires (wo separate evelotron features in this baucdpass.,-band spectra probably requires two separate cyclotron features in this bandpass.
+ If. we assume (he phasing and geometry from Bailev et al. (, If we assume the phasing and geometry from Bailey et al. (
+1952). the phase 0.0 exclotron features that we see in the {Ρας are presumably associated with (he pole that accretes during the hieh state.,"1982), the phase 0.0 cyclotron features that we see in the -band are presumably associated with the pole that accretes during the high state."
+ This suggests that the single. dominant cevleotron feature at phase 0.5 might be from the other pole.," This suggests that the single, dominant cylcotron feature at phase 0.5 might be from the other pole."
+ While the minimum light A-band spectrum is fairly consistent with the shape of a cool star spectrum. with little evidence for evelotron emission. (he maximum light spectrum is Clearly contaminated by at least one. possibly two evelotron features.," While the minimum light -band spectrum is fairly consistent with the shape of a cool star spectrum, with little evidence for cyclotron emission, the maximum light spectrum is clearly contaminated by at least one, possibly two cyclotron features."
+ Obviously. (here is a evelotron harmonic that peaks below 2.0 jan. The rapid change in slope of the red half of the A-band spectrum. and the dip near 2.07 jan. suggest the possibility of an additional. weaker cyclotron feature al phase 0.5.," Obviously, there is a cyclotron harmonic that peaks below 2.0 $\mu$ m. The rapid change in slope of the red half of the $K$ -band spectrum, and the dip near 2.07 $\mu$ m, suggest the possibility of an additional, weaker cyclotron feature at phase 0.5."
+ To better examine (he phase 0.5 evelotvon features. we have subtracted (he phase 0.0 spectrum from the phase 0.50 spectrum. to produce the residual spectrum shown in Fig.," To better examine the phase 0.5 cyclotron features, we have subtracted the phase 0.0 spectrum from the phase 0.50 spectrum, to produce the residual spectrum shown in Fig."
+ 8., 8.
+ The result shows (wo strong evclotron features. one that peaks in the /-band near 1.58 yam. and one that (sve estimate) peaks near 1.93 jan. The possibili: of a second evelotvon feature in the. A-band seems more likely after this subtraction.," The result shows two strong cyclotron features, one that peaks in the -band near 1.58 $\mu$ m, and one that (we estimate) peaks near 1.93 $\mu$ m. The possibility of a second cyclotron feature in the -band seems more likely after this subtraction."
+ If the secondary star has a similar spectrum at phases 0.0 ancl 0.5. Chen any water vapor feature should disappear upon subtraction.," If the secondary star has a similar spectrum at phases 0.0 and 0.5, then any water vapor feature should disappear upon subtraction."
+ lenoring the false IL I Dr-5 emission feature. there continues to be a modest dip near 2.07 san and a small rise that peaks near 2.15 yan and that falls off slowly bevond 2.2 jum. Regardless of the exact interpretation of the spectra presented here. it is clear (hat the model proposed by Harrison et al. (," Ignoring the false H I $\gamma$ emission feature, there continues to be a modest dip near 2.07 $\mu$ m and a small rise that peaks near 2.15 $\mu$ m and that falls off slowly beyond 2.2 $\mu$ m. Regardless of the exact interpretation of the spectra presented here, it is clear that the model proposed by Harrison et al. ("
+2003) is incorrect.,2003) is incorrect.
+ The photometric modulations in the A-band are certainly being driven by the changing strength of the evelotron feature at the blue end of this bandpass (ancl possibly aided by a broader. ancl weaker feature that peaks," The photometric modulations in the -band are certainly being driven by the changing strength of the cyclotron feature at the blue end of this bandpass (and possibly aided by a broader, and weaker feature that peaks"
+a resolution of at least 25 bins with a minimum noise of 100 per bin.,a resolution of at least 25 bins with a minimum signal-to-noise of 100 per bin.
+ Our goal is to make high precision measurements of the projected metallicity distribution while maintaining a spatial resolution that optimizes the measurement of the spatial distribution of the high metallicity gas for comparison to the radio sources., Our goal is to make high precision measurements of the projected metallicity distribution while maintaining a spatial resolution that optimizes the measurement of the spatial distribution of the high metallicity gas for comparison to the radio sources.
+" The clusters that meet these criteria are A133, A262, Perseus, MS0735.6+7421, Hydra A, M87, A1835, A2029, A2199, and A2597."," The clusters that meet these criteria are A133, A262, Perseus, MS0735.6+7421, Hydra A, M87, A1835, A2029, A2199, and A2597."
+ Each cluster observation was processed using version 4.1.2 of CIAO and the calibration database., Each cluster observation was processed using version 4.1.2 of CIAO and the calibration database.
+" Background flares were excluded using standard filtering techniques and point sources were identified in the combined image of each cluster usingwavedetect, and removed."," Background flares were excluded using standard filtering techniques and point sources were identified in the combined image of each cluster using, and removed."
+ Blank-sky background files were normalized to the source count rate in the 9.5-12 keV band., Blank-sky background files were normalized to the source count rate in the 9.5-12 keV band.
+" We present example metallicity maps for three clusters 262, Hydra A, Abell 1835) in Figure 1."," We present example metallicity maps for three clusters (Abell 262, Hydra A, Abell 1835) in Figure 1."
+ The combined(Abell exposures for each cluster were binned using a weighted Voronoi tessellation algorithm (Cappellari&Copin2003;Diehl&Statler 2006).," The combined exposures for each cluster were binned using a weighted Voronoi tessellation algorithm \citep{cap03,die06}."
+". With approximately 10000 net counts or greater per bin, spectra were extracted and weighted response matrices were created using standard CIAO routines."," With approximately 10000 net counts or greater per bin, spectra were extracted and weighted response matrices were created using standard CIAO routines."
+" A single temperature plasma model with absorption was fit over the energy range 0.5-7 keV. (WABSxMEKAL)Temperature, abundance, and normalization were allowed to vary with the column density frozen at the values quoted by Dickey&Lockman"," A single temperature plasma model with absorption $\times$ MEKAL) was fit over the energy range 0.5-7 keV. Temperature, abundance, and normalization were allowed to vary with the column density frozen at the values quoted by \citet{dic90}."
+ The internal ratio between metals were set to the(1990)..solar photospheric values of Grevesse&Sauval (1998)., The internal ratio between metals were set to thesolar photospheric values of \citet{gre98}.
+" Allowing the a-elements to vary independently does not show any significant effect on the outcome of the iron abundance measurement, therefore we have chosen to use the simple model."," Allowing the $\alpha$ -elements to vary independently does not show any significant effect on the outcome of the iron abundance measurement, therefore we have chosen to use the simple model."
+" The spectra are reasonably well fit by the single temperature model, with the exception of the central region of A262."," The spectra are reasonably well fit by the single temperature model, with the exception of the central region of A262."
+ This region is better fit with a multi-temperature model which finds a higher central metallicity value (seeKirkpatricketal. 2009a).., This region is better fit with a multi-temperature model which finds a higher central metallicity value \citep[see][]{cck09a}.
+" However, this does not affect the interpretation of our results as the central regions of the clusters are not important for this analysis."," However, this does not affect the interpretation of our results as the central regions of the clusters are not important for this analysis."
+ Metallicity profiles have been measured for each cluster., Metallicity profiles have been measured for each cluster.
+ Spectra were extracted from regions occupied by radio emission and/or cavities in semi-annular shaped bins., Spectra were extracted from regions occupied by radio emission and/or cavities in semi-annular shaped bins.
+" Each bin size was determined by requiring a minimum signal-to-nose of 140, which achieves uncertainties of approximately 0.05 Zo."," Each bin size was determined by requiring a minimum signal-to-nose of 140, which achieves uncertainties of approximately 0.05 $_\sun$."
+ A high signal to noise ratio is required to find deviations between the profiles along and orthogonal to the radio cavity system., A high signal to noise ratio is required to find deviations between the profiles along and orthogonal to the radio cavity system.
+" 'These orthogonal profiles, which we consider to represent the undisturbed region of the cluster, are extracted using the same bin width as the radio cavity system profile."," These orthogonal profiles, which we consider to represent the undisturbed region of the cluster, are extracted using the same bin width as the radio cavity system profile."
+ T'he S/N for these regions are usually greater than 140 due to larger angular sizes., The S/N for these regions are usually greater than 140 due to larger angular sizes.
+ All profiles are fit using the same model as the 2D analysis., All profiles are fit using the same model as the 2D analysis.
+" The radial profiles for A262, Hydra A, and A1835 are presented in Figure 2."," The radial profiles for A262, Hydra A, and A1835 are presented in Figure 2."
+ The red squares represent the metallicity of the radio cavity system and the blue triangles represent the undisturbed metallicity profile., The red squares represent the metallicity of the radio cavity system and the blue triangles represent the undisturbed metallicity profile.
+ 'The dotted vertical line indicates the maximum radius at which a significant enhancement in metallicity has been detected., The dotted vertical line indicates the maximum radius at which a significant enhancement in metallicity has been detected.
+ We refer to this as the iron radius (Rg)., We refer to this as the iron radius $R_{\rm Fe}$ ).
+ This radius is concisely defined as the location of the radial bin furthest from the cluster center where the one sigma error bars of the profiles along and orthogonal to the radio axis do not overlap., This radius is concisely defined as the location of the radial bin furthest from the cluster center where the one sigma error bars of the profiles along and orthogonal to the radio axis do not overlap.
+ Regions interior to the iron radius show a systematic enhancement in iron abundance over the undisturbed regions., Regions interior to the iron radius show a systematic enhancement in iron abundance over the undisturbed regions.
+" For regions beyond the iron radius, the abundance profiles are indistinguishable."," For regions beyond the iron radius, the abundance profiles are indistinguishable."
+ From top to bottom the three panels are arranged in order of increasing jet power and increasing iron radius., From top to bottom the three panels are arranged in order of increasing jet power and increasing iron radius.
+ This shows that higher power AGN are able to launch metal enriched gas to greater distances., This shows that higher power AGN are able to launch metal enriched gas to greater distances.
+" In Figure 3, the iron radius is plotted against jet power derived using X-ray cavity data taken from Raffertyetal. for all ten clusters in the sample."," In Figure 3, the iron radius is plotted against jet power derived using X-ray cavity data taken from \citet{raf06} for all ten clusters in the sample."
+" The powers quoted (2006)there were derived from the total energy of the cavities, 4pV, and their buoyant rise time."," The powers quoted there were derived from the total energy of the cavities, $4pV$ , and their buoyant rise time."
+" In systems with multiple generations of cavities, the average power"," In systems with multiple generations of cavities, the average power"
+etections in the 1.16mm map with their deboosted. [Lux ensities in the mm and 12mm maps.,"detections in the $\,$ mm map with their deboosted flux densities in the $\,$ mm and $\,$ mm maps."
+ Our approach to obtain empirically cehboosted Hux densities is to fit a —unction that relates the Dux. densities of secure detections in the mam and 12mm maps to the lux densities and noise values in those maps.," Our approach to obtain empirically deboosted flux densities is to fit a function that relates the flux densities of secure detections in the $\,$ mm and $\,$ mm maps to the flux densities and noise values in those maps."
+ We then use the derived. formula to estimate enipirically deboosted flux densities for the secure nm cletections from. the measured mm lux densities and noises.," We then use the derived formula to estimate empirically deboosted flux densities for the secure $\,$ mm detections from the measured $\,$ mm flux densities and noises."
+ We find where Syuouaer and σ are in mw., We find where $S_{\mathrm{measured}}$ and $\sigma$ are in mJy.
+ For the secure Ακος L1mm and MAMDO 12mm detections. the residuals between the deboosted. Dux. densities from this relation and the deboostec (lux densities in Pereractal.(2008) and Grevectal.(2008). have a standard. deviation olf my. an error well below the Lux density noise values in all mnm maps.," For the secure AzTEC $\,$ mm and MAMBO $\,$ mm detections, the residuals between the deboosted flux densities from this relation and the deboosted flux densities in \citet{perera08} and \citet{greve08} have a standard deviation of $\,$ mJy, an error well below the flux density noise values in all mm maps."
+ This formula is only valid in the range of SNR. covered by the AzPEC and NLAMDO cletections. so we do not deboost the [ux density of source 1 (a 156 source).," This formula is only valid in the range of SNR covered by the AzTEC and MAMBO detections, so we do not deboost the flux density of source 1 (a $\sigma$ source)."
+ Table 1. lists the cdeboosted [ux densities for the secure 1.16 mu detections using this relation.," Table \ref{table:sources} lists the deboosted flux densities for the secure $\,$ mm detections using this relation."
+ For the main purposes of this paper. (lux deboosting is not necessary since we stack the 1.16 mm Ilux densities of sources we know to exist from other observations.," For the main purposes of this paper, flux deboosting is not necessary since we stack the $\,$ mm flux densities of sources we know to exist from other observations."
+ Fig., Fig.
+ 3 shows the comparison between deboosted. [Lux densities for secure. mim. and. mim. detections.," \ref{fig:deboosted1.1} shows the comparison between deboosted flux densities for secure $\,$ mm and $\,$ mm detections."
+ The combined mim. map recovers the majority. of secure detections identified in the AZPEC πάν map the arrows pointing down show that there are 4 secure detections in the AZTEC map that are not secure detections in the combined map.," The combined $\,$ mm map recovers the majority of secure detections identified in the AzTEC $\,$ mm map – the arrows pointing down show that there are 4 secure detections in the AzTEC map that are not secure detections in the combined map."
+ Fig., Fig.
+ 4 shows the comparison between deboosted. [Lux densities for secure mm ancl mm detections.," \ref{fig:deboosted1.2} shows the comparison between deboosted flux densities for secure $\,$ mm and $\,$ mm detections."
+ There are Lt secure detections in the NLAMDO nim map that are not coincident with secure detections in the combined 1.16 mm map (the down arrows in the right panel).," There are 14 secure detections in the MAMBO $\,$ mm map that are not coincident with secure detections in the combined $\,$ mm map (the down arrows in the right panel)."
+ Llowever. the upper limits to the [lux densities in the combined map are within the scatter about the solid line.," However, the upper limits to the flux densities in the combined map are within the scatter about the solid line."
+ We conclude. based on the comparisons in Figs.," We conclude, based on the comparisons in Figs."
+ 3. and 4.. that our method of combining the AzPEC mam and ALAMDO mm maps is ellective.," \ref{fig:deboosted1.1} and \ref{fig:deboosted1.2}, that our method of combining the AzTEC $\,$ mm and MAMBO $\,$ mm maps is effective."
+ The combined. 1.16 mum map has 13 new secure detections (Γαία 12): we co not expect the new detections to be in the individual maps.," The combined $\,$ mm map has 13 new secure detections (Table \ref{table:sources}) ); we do not expect the new detections to be in the individual maps."
+ Galaxies with detected: 0n emission. compose the most iomiogeneous set οἱ clusty galaxies whose mam Πας density can be stacked with significant. results.," Galaxies with detected $\,$$\micron$ emission compose the most homogeneous set of dusty galaxies whose mm flux density can be stacked with significant results."
+ We use the jum catalogue [rom the GOODS-N /MEPS survey. which ias a uniform depth of la~ 55s in the regions of interest (Chary et al.," We use the $\,$$\micron$ catalogue from the GOODS-N /MIPS survey, which has a uniform depth of $1\sigma \sim 5$ $\mu$ Jy in the regions of interest (Chary et al."
+ in preparation): the jm Ηχος are measured at the positions of LRAC sources. so this catalogue pushes o faint 24imi Hluxes.," in preparation); the $\,$$\micron$ fluxes are measured at the positions of IRAC sources, so this catalogue pushes to faint $\,$$\micron$ fluxes."
+". We only stack ο3a 24in sources with So, 254..."," We only stack $\ge 3\sigma$ $\,$$\micron$ sources with $S_{24} > 25$ $\mu$ Jy."
+ At Hux densities above 50 idv. the catalogue is 99 per cent complete: for 25«Sey SOyedy. he catalogue is 83 per cent complete (Magnellietal.2009).," At flux densities above $50$ $\mu$ Jy, the catalogue is 99 per cent complete; for $25 < S_{24} < 50$ $\mu$ Jy, the catalogue is 83 per cent complete \citep{magnelli09}."
+. Completeness corrections to our results are negligible. so we do not apply them.," Completeness corrections to our results are negligible, so we do not apply them."
+ We exclude sources that lie in the region of the Tl.16mm map with le lmJv: in this region. the noise is non-uniform.," We exclude sources that lie in the region of the $\,$ mm map with $1\sigma > 1$$\,$ mJy; in this region, the noise is non-uniform."
+ Vhe final 24jin catalog for stacking has 2484 sources in 0.068 deg.," The final $\,$$\micron$ catalog for stacking has 2484 sources in 0.068 $^{2}$."
+ To decompose the contribution to the mm background [rom 24uu sources as a function. of redshift. we require either a photometric or spectroscopic redshift [or each 24; source.," To decompose the contribution to the mm background from $\,$$\micron$ sources as a function of redshift, we require either a photometric or spectroscopic redshift for each $\,$$\micron$ source."
+ We start by matching a source with a spectroscopic redshift from the catalogues of Darger.Cowie.&Wang(2008) and Stern et al. (, We start by matching a source with a spectroscopic redshift from the catalogues of \citet*{barger08} and Stern et al. (
+in preparation) to each δι source.,"in preparation) to each $\,$$\micron$ source."
+ The match radius. 0.7 aresec. is chosen by maximizing the number of unique matches while minimizing the number of multiple matches.," The match radius, 0.7 arcsec, is chosen by maximizing the number of unique matches while minimizing the number of multiple matches."
+ We find spectroscopic redshifts for 1026 (41 per cent of the) jm sources.," We find spectroscopic redshifts for 1026 (41 per cent of the) $\,$$\micron$ sources."
+ Η no (or multiple) coincident sources with spectroscopic redshifts are found. we resort to the photometric redshift source catalogue of DBrodwin et al. (," If no (or multiple) coincident sources with spectroscopic redshifts are found, we resort to the photometric redshift source catalogue of Brodwin et al. ("
+in preparation) to find a source match.,in preparation) to find a source match.
+ Photometric redshifts are constrained with deep CDVIULIN imaging. and provide redshift estimates or S72. or 35 per cent. of the jun sources.," Photometric redshifts are constrained with deep $UBVRIzJK$ imaging, and provide redshift estimates for 872, or 35 per cent, of the $\,$$\micron$ sources."
+ Photometric redshift’ uncertainties are small compared to our redshift yins. since we are interested. in the contribution to the xwkeround from galaxies in large redshift bins.," Photometric redshift uncertainties are small compared to our redshift bins, since we are interested in the contribution to the background from galaxies in large redshift bins."
+" 1 no. (or multiple) coincident sources with photometric redshifts are ound. we assign the 24m source to a ""redshift unknown? jn in the stacking analysis."," If no (or multiple) coincident sources with photometric redshifts are found, we assign the $\,$$\micron$ source to a `redshift unknown' bin in the stacking analysis."
+ Of the 2484 jin sources. SSS (24 per cent) have no spectroscopic or photometric recdshift estimate available.," Of the 2484 $\,$$\micron$ sources, 588 (24 per cent) have no spectroscopic or photometric redshift estimate available."
+ Our stacking5 procedure depends on two fundamental properties of the (sub)mm maps: Stacking is the process of averaging the lux density. at some wavelength mm) of sources detected at another wavelength.," Our stacking procedure depends on two fundamental properties of the (sub)mm maps: Stacking is the process of averaging the flux density, at some wavelength $\,$ mm), of sources detected at another wavelength."
+ To resolve the (sub)jmun background. we want o stack a catalogue of sources whose emission. correlates strongly with mu emission. and we want this catalogue o have a large number of sources.," To resolve the (sub)mm background, we want to stack a catalogue of sources whose emission correlates strongly with $\,$ mm emission, and we want this catalogue to have a large number of sources."
+ A catalogue that meets hese requirements has galaxies selected: on dust. emission at both low and high redshift., A catalogue that meets these requirements has galaxies selected on dust emission at both low and high redshift.
+ We do not expect a sample of stellar mass selected sources (for example. at 3.6 jm) to » ellicient at isolating the galaxies responsible for the mm xickeground. because jin sources are a mix of dusty and," We do not expect a sample of stellar mass selected sources (for example, at $\,$$\micron$ ) to be efficient at isolating the galaxies responsible for the mm background, because $\,$$\micron$ sources are a mix of dusty and"
+because inclined binary systems do not necessarily have the same observed orientation as their circumprimary disks.,because inclined binary systems do not necessarily have the same observed orientation as their circumprimary disks.
+" This may explain why our significance, with an improved data-set, is of the same order as ? found with a smaller data-set and a more simplistic model."," This may explain why our significance, with an improved data-set, is of the same order as \citet{DB2006} found with a smaller data-set and a more simplistic model."
+ Improving the current statistics requires further enlarging the sample., Improving the current statistics requires further enlarging the sample.
+" We find that if the data followed the model distribution exactly, differentiating between the two scenarios at a 3σ level would require a sample of approximately 50 objects."," We find that if the data followed the model distribution exactly, differentiating between the two scenarios at a $\sigma$ level would require a sample of approximately 50 objects."
+" Nonetheless, the current data are consistent with the co-planar hypothesis and furthermore, this scenario is favoured over random alignments."," Nonetheless, the current data are consistent with the co-planar hypothesis and furthermore, this scenario is favoured over random alignments."
+ Our use of spectropolarimetry assumes that the spectropolarimetric signatures of HAe/Be stars can be used to trace the orientations of their circumstellar disks., Our use of spectropolarimetry assumes that the spectropolarimetric signatures of HAe/Be stars can be used to trace the orientations of their circumstellar disks.
+"?). However, the proposal of ? and ? that the spectropolarimetric signatures of HAe/Be stars are due to optical pumping and not scattering in a disk implies that these signatures do not have to trace disks."," However, the proposal of \citet{Harrington2007} and \citet{Kuhn2007} that the spectropolarimetric signatures of HAe/Be stars are due to optical pumping and not scattering in a disk implies that these signatures do not have to trace disks."
+ We address this issue here and do this by comparing spectropolarimetric observations of young stellar objects with independent measurements of the orientation of their disks., We address this issue here and do this by comparing spectropolarimetric observations of young stellar objects with independent measurements of the orientation of their disks.
+ To, To
+where ka=ky+ke and ky||ko.,where $\bf k_3=k_1+k_2$ and $\bf k_1 \parallel k_2$.
+" O=—2y—5x10."" is the mean temperature SZ distortion."," $\bar{\Theta}=-2
+\bar{y} \sim 5 \times 10^{-6}$ is the mean temperature SZ distortion."
+ OQ? and O° are [roin the corresponding Gaussian term in the correlation function. as explained in Eq. (8)).," $\bar{\Theta}^2$ and $\bar{\Theta}^3$ are from the corresponding Gaussian term in the correlation function, as explained in Eq. \ref{pressure}) )."
+ As a reminder. Cr) is the geometric function defined in Eq. (3)).," As a reminder, $C(x)$ is the geometric function defined in Eq. \ref{eqn:cx}) )."
+ Though the unsmoothed Q5 is uot directly observable due to the abseuce of the window fuuction. from the theory viewpoint. it represents the raw non-gaussianity properties.," Though the unsmoothed $\Theta_3$ is not directly observable due to the absence of the window function, from the theory viewpoint, it represents the raw non-gaussianity properties."
+ Our calculations find that O4~105 (Table 1))., Our calculations find that $\Theta_3 \sim 10^{6}$ (Table \ref{tble:tbar}) ).
+ We defined the skewness parameter (16)) as the ratio of differential measurements. which might 'esult. from interferometers in the Raleigh-Jeaus regime.," We defined the skewness parameter \ref{eqn:t3}) ) as the ratio of differential measurements, which might result from interferometers in the Raleigh-Jeans regime."
+ For comparison. Coorayetal. aud Cooray(2000b) cousicered the absolute temperature distortion aud clefinecl the skewuess ↽≻⋜↕↓⋅⋜⋃∐≺↵↕≺↵⋅⋜↕⊳∖⋌∶∖↭⇢∣∖↭−⋰∖−⋅∖↾∖↕∐↥⋯∐↥−↓⋅≺↵≺↽↓⋯↵∐∢∶∖⊽↓∐↥∩⊔∐⋜," For comparison, \citet{Cooray00a} and \citet{Cooray00b} considered the absolute temperature distortion and defined the skewness parameter as $\langle\Theta^3\rangle/\langle
+\Theta^2\rangle^2$."
+↕↓∩∐↕∐≺↵⋜↕∣≻⊳∖∩↥∐≺↵⋅∣∕∢∙⋯∐≺∣≻≺↵∐↕≺↵⋜↕⊳∖⋯⋅≺↵≺⇂⋜↕↕PENASEEIg2A - ep ⋅⋅⋅ ≺↵⋜↕∢∙∐↥↽≻∩↥∐↕∐⊳∖↥↽≻⋜↕∢∙↩⋅⋜⋯≺⊓↕⋜↕↕≺⇂≺↵∐∐∏∩↥∩↥∎⊳∖↕⊊↩∖∖↽∐≺↵⊳∖⊳∖∢∙⋯∐≺↥∐↥↽∐⋅↕∐∢∙↥↥↽≻↥≺↵≻≺↲∩," With multi-frequency information the absolute $y$ could be measured at each point in space, and that definition of skewness could in principle be observed."
+∣≻⊳∖≺↵↕⋅∖⊽≺↵≺⋅⊺∐≺↵↕⋅≺↵⊳∖⋃∐⊳∖⋜⋃⋅≺↵ ↕∩↕↕≺⋉⊓⇂∎∐⇀≺↵↕⋅≺↵∐↕∶↕∐⋯⊔⋅⋯⋯⇂↩↥⋅⋌∶↕↭⋮⋡⋝↕⋤∶↙∕∣↗⋌∶↕≺↦⋟⊋⋟∶⊋∿↽∐⋗↱↗⋅∖∖↽↥∏∢∙∐↥⊳∖⊳∖∐∐∐⋜⋃⋅↕∩↕↕≺↵↓⋅≺↵⊳∖⋃∐∩∎⋖∫⊂⊲∩∩↕⋅⋜≹⊽∖⇁⊇∪∪∩∣∪ ∖∖↽∐∐↕∐≺↵∐↕⋜∟∖↥∐∐⊔∐∖⊽∐⋅↥⋜≹↥∎∠↩≺⊔∐⋜↕⊳∖⊳∖∿⋥∪↓⊽↗⋀⊔⋅⋅⊡∐⋅↕∐≺↵↥⋅∖∖↽≺∐⋅↕⊔∙∩∐⊳∖," The results are not too different: In our model, $\langle\Theta^3\rangle/\langle 
+\Theta^2\rangle^2 \sim 10^5$, which is similar to the result of \citep{Cooray00b} with the maximum virialized mass $\sim
+10^{15}M_{\odot}$."
+↥≺⇂≺↵↕⋅↥∑≟∖∖↽↥∐≺⇂∩∖∖↽↥∎⋯∐∙⋃∩∐∐∐≺↵∏∐∑∸ ⋜↕∐≺⇂↕∐↩∑∸↩∐≺↵↕⋅⋜↕↥∢∙⋜↕↥∢∙⋃↥��∏↕∩∐∩⊓∐≺↵∺∠≻↕⊳∖↥↽≻≺↵∢∙⋃⋅⋯∐∢∙⋜↕∐∣⋈↵≺⇂∩∐≺↲↥∎∩∐∩∖∖↽↥∐∑∸↕∐≺↵∐↕≺↲↕∐⋯⇂∩↥∎⊂⊲∩∩↓⋅⋜↕⊽∖⇁≺↵↕ (2000a)., Further work considering window function filtering and the general calculation of the SZ bispectrum can be done following the method of \citet{Cooray00a}.
+. Iu the SZ effect. the redshift distribution of intergalactic gas is lost.," In the SZ effect, the redshift distribution of intergalactic gas is lost."
+ Taking advantage of the Cross correlations with other surveys baving redshift tuformation. we may be able to statistically extract IGM 3-D correlations and evolution.," Taking advantage of the cross correlations with other surveys having redshift information, we may be able to statistically extract IGM 3-D correlations and evolution."
+ Because tlie SZ ellect is mostly contributed by uouliuear structures at z~0.5—2. it should have a stroug cross correlation with galaxies at that reclshilt range.," Because the SZ effect is mostly contributed by nonlinear structures at $z\sim 0.5-2$, it should have a strong cross correlation with galaxies at that redshift range."
+ The general idea is to use a galaxy survey with coarse redshift information. where the distance information ouly needs enough accuracy to resolve the time evolution of the correlation [uuction. for example Az~)0.2. which can be reached by photometric redshift surveys.," The general idea is to use a galaxy survey with coarse redshift information, where the distance information only needs enough accuracy to resolve the time evolution of the correlation function, for example $\Delta z \sim
+0.2$, which can be reached by photometric redshift surveys."
+ Oue then correlates each redshift bin with the SZ map. and obtains the relative contribution of that SZ slice to the total projected map.," One then correlates each redshift bin with the SZ map, and obtains the relative contribution of that SZ slice to the total projected map."
+ Several questious that pose themselves in this procedure are: How big a survey area aud depth does one need?, Several questions that pose themselves in this procedure are: How big a survey area and depth does one need?
+ What fraction of the total SZ luctuatious cau be accounted using cross-correlatious?, What fraction of the total SZ fluctuations can be accounted using cross-correlations?
+ Are there optimal procedures of cross-correlatine?, Are there optimal procedures of cross-correlating?
+ Since SZ fluctuations are a non-linear function ol the ealaxy field. what fraction of the sigual will ever be accounted for by this approach?," Since SZ fluctuations are a non-linear function of the galaxy field, what fraction of the signal will ever be accounted for by this approach?"
+ Do we expect the SZ-galaxy cross-correlation coefficient to depend on redshilt?, Do we expect the SZ-galaxy cross-correlation coefficient to depend on redshift?
+ How big an uncertainty uuieht arise from a time aud scale depeudent galaxy bias?, How big an uncertainty might arise from a time and scale dependent galaxy bias?
+ In this section we will address each of these questions using our model. aud obtain quantitative estimates.," In this section we will address each of these questions using our model, and obtain quantitative estimates."
+procedure is shown in the Figure 3..,procedure is shown in the Figure \ref{fig:qso_ex}.
+" There are 43 observations in the light curve and time counts from the first epoch in Stripe 82 (run 94, fpC frame 12)."," There are 43 observations in the light curve and time counts from the first epoch in Stripe 82 (run 94, fpC frame 12)."
+ The number of epochs per quasar in our sample varies between 20 and 56 with the mean of 36., The number of epochs per quasar in our sample varies between 20 and 56 with the mean of 36.
+ The time sampling of our quasar light curves is relatively sparse (see Figure 3)) and insufficient for the use of spectral methods such as the frequency-power spectrum analysis., The time sampling of our quasar light curves is relatively sparse (see Figure \ref{fig:qso_ex}) ) and insufficient for the use of spectral methods such as the frequency-power spectrum analysis.
+" Following the common practice in AGN variability studies, we focus on the information contained in the structure function."," Following the common practice in AGN variability studies, we focus on the information contained in the structure function."
+ A number of structure function definitions are in use differing in relatively minor details that are not important for this study., A number of structure function definitions are in use differing in relatively minor details that are not important for this study.
+" We adopt a structure function (hereafter SF) in the form: where r;=2(t;) some random quantity at moment t;, the sum is taken over all measurements for which t;—t;—T and N(r) is the number of such data points."," We adopt a structure function (hereafter SF) in the form: where $x_i=x(t_i)$ some random quantity at moment $t_i$, the sum is taken over all measurements for which $t_j-t_i=\tau$ and $N(\tau)$ is the number of such data points."
+" In short, the structure function is a time dependent standard deviation of some random process."," In short, the structure function is a time dependent standard deviation of some random process."
+ SF of white noise is flat and equal to v26?., SF of white noise is flat and equal to $\sqrt{2\sigma^2}$.
+" Therefore, the SF of a random process is also flat on time-scales, where the amplitude of the intrinsic variability is comparable with the measurement errors, as well as on the longest timescales where the r.m.s."," Therefore, the SF of a random process is also flat on time-scales, where the amplitude of the intrinsic variability is comparable with the measurement errors, as well as on the longest timescales where the r.m.s."
+ stops changing as the length of the time interval keeps increasing., stops changing as the length of the time interval keeps increasing.
+ In the context of quasar variability the random quantity is the magnitude (essentially the relative luminosity) m; measured at time t; and T is taken in the quasar rest frame., In the context of quasar variability the random quantity is the magnitude (essentially the relative luminosity) $m_i$ measured at time $t_i$ and $\tau$ is taken in the quasar rest frame.
+ For a given quasar we can write: where we converted to flux units., For a given quasar we can write: where we converted to flux units.
+" The flux consists of a stationary component f and a time dependent variable part óf;, and στ is the r.m.s."," The flux consists of a stationary component $\bar{f}$ and a time dependent variable part $\delta f_i$ , and $\sigma_\tau$ is the r.m.s."
+ calculated for points separated by time interval τ., calculated for points separated by time interval $\tau$.
+" From here it is obvious that for two quasars at the same redshift and with equal c2 the brightest one will have smaller variability amplitude, S(T), at all time scales."," From here it is obvious that for two quasars at the same redshift and with equal $\sigma_{\tau}^2$ the brightest one will have smaller variability amplitude, $S(\tau)$, at all time scales."
+ 'The number of time intervals that can be constructed from an average of 36 epochs for a single quasar is 36x(36—1)/2630 and results in à very noisy structurefunction?., The number of time intervals that can be constructed from an average of 36 epochs for a single quasar is $36\times(36-1)/2 = 630$ and results in a very noisy structure.
+". Therefore, we focus our analysis on the ensemble quasar SF, i.e. we assume that the ensemble variability reflects on average the variability of a single quasar."," Therefore, we focus our analysis on the ensemble quasar SF, i.e. we assume that the ensemble variability reflects on average the variability of a single quasar."
+ The final sample contains 7562 light curves., The final sample contains 7562 light curves.
+" All possible pairs of measurements within each light curve are collected for a total of 4,796,468 magnitude differences."," All possible pairs of measurements within each light curve are collected for a total of 4,796,468 magnitude differences."
+" After reducing the time lag to the quasar rest frame using spectroscopic redshifts from the SDSS Quasar Catalog (Schneideretal.2010), the data are binned in time."," After reducing the time lag to the quasar rest frame using spectroscopic redshifts from the SDSS Quasar Catalog \citep{2010AJ....139.2360S}, the data are binned in time."
+" The bin size was adjusted to reach 20,000 measurements per bin."," The bin size was adjusted to reach 20,000 measurements per bin."
+ The resulting ensemble SF consists of 239 bins and is shown in Figure 4 (blue points)., The resulting ensemble SF consists of 239 bins and is shown in Figure \ref{fig:sf} (blue points).
+" Error bars were estimated by dividing each timescale bin into 20 smaller bins, evaluating the mean SF in each, and calculating the r.m.s."," Error bars were estimated by dividing each timescale bin into 20 smaller bins, evaluating the mean SF in each, and calculating the r.m.s."
+ scatter of the resulting values (see also deVriesetal. (2005)))., scatter of the resulting values (see also \cite{2005AJ....129..615D}) ).
+ In order to verify that we detect quasarvariability and guard against the possibility of contamination, In order to verify that we detect quasarvariability and guard against the possibility of contamination
+"st which solves the cooling flow problem (ChanandChu2007) and Ty,=1289. we get which agrees with the recent observation: a=4.022:0.32 and 3=8.1340.06 al. 2002)..","$^{-1}$ which solves the cooling flow problem \citep{Chan}
+ and $\Gamma_{3 \nu}=128 \Gamma$, we get which agrees with the recent observation: $\alpha=4.02 \pm 0.32$ and $\beta=8.13 \pm 0.06$ \citep{Tremaine}. ."
+ As an order of magnitude estimate. sin@~mp/mi. where mj is the neutrino Dirac mass.," As an order of magnitude estimate, $\sin \theta \sim m_D/m_s$, where $m_D$ is the neutrino Dirac mass."
+" In the standard see-saw mechanism. mj; is the geometric mean of the light neutrino niass-scale nm, and m. (Mohapatraa"," In the standard see-saw mechanism, $m_D$ is the geometric mean of the light neutrino mass-scale $m_{\nu}$ and $m_s$ \citep{Mohapatra}."
+"ndSenjanovic1980).. Therelore. sin9~ym,/m.."," Therefore, $\sin \theta 
+\sim 
+\sqrt{m_{\nu}/m_s}$."
+ From Eq. (, From Eq. (
+"1). we get For Pa,~10UU I. n,~30 keV for my,~1 eV. which is consistent with our assumption (n»>10 keV).","1), we get For $\Gamma_{3 \nu} \sim 10^{-17}$ $^{-1}$, $m_s \sim 30$ keV for $m_{\nu} \sim 1$ eV, which is consistent with our assumption $m_s \ge 10$ keV)."
+" We asstume (he existence of a degenerate neutrino halo (i.~ keV) at the proto-galactic center with a parameter. D4,. which is universal and can be inferred[roii observation of cluster hot eas (ChanandChu 2007)."," We assume the existence of a degenerate neutrino halo $m_s \sim$ keV) at the proto-galactic center with a parameter, $\Gamma_{3 \nu}$, which is universal and can be inferredfrom observation of cluster hot gas \citep{Chan}. ."
+". Without anv assumptions of the protogalaxy ancl existence of any intermediate mass blackohles. a24 ancl &8 require the total decay rate to be Ta,—(52x1)x100.Fs fo which is consistent with the observational data from cooling flow clusters (ChanandChu2007)."," Without any assumptions of the protogalaxy and existence of any intermediate mass blackohles, $\alpha \approx 4$ and $\beta \approx 8$ require the total decay rate to be $\Gamma_{3 \nu}=(5 \pm 1) \times
+10^{-17}$ $^{-1}$, which is consistent with the observational data from cooling flow clusters \citep{Chan}."
+.. Also. we have reviewed several models to explain the Apu—0 relation.," Also, we have reviewed several models to explain the $M_{BH,f}-\sigma$ relation."
+ These models require several assumptions or [ree parameters which may not be true Lor all galaxies., These models require several assumptions or free parameters which may not be true for all galaxies.
+ For example. Ixings model assumes (he isothermal density prolile (pc r2) for all galaxies during all the time of the blackhole formation.," For example, King's model assumes the isothermal density profile $\rho \sim 
+r^{-2}$ ) for all galaxies during all the time of the blackhole formation."
+ If the density profile changes into the form per+. then σαwr which is not a constant.," If the density profile changes into the form $\rho \sim r^{-1}$, then $\sigma \propto \sqrt{r}$ which is not a constant."
+ This problem also exists in (he Ballistic model which is based on the isothermal distribution of matter., This problem also exists in the Ballistic model which is based on the isothermal distribution of matter.
+ On the other hand. in the self-similar model. (here are several [ree parameters which are moclel dependent.," On the other hand, in the self-similar model, there are several free parameters which are model dependent."
+ Also one cannot obtain the proportionality constant2 of the Mt;σ relation.," Also one cannot obtain the proportionality constant$\beta$ of the $M_{BH,f}- \sigma$ relation."
+" We have considered. a wide range of 7=0.1— 10000. f,=0.6—3 andf»=0.60--3 encompassing almost all possibilities in galaxies."," We have considered a wide range of $\tau=0.1-10000$ , $f_1=0.6-3$ and$f_2=0.6-3$ encompassing almost all possibilities in galaxies."
+ In our model. we assume that (he supermassive," In our model, we assume that the supermassive"
+F9 ΑΠ) and F11 ). it forms a substructure that we interpret as rotational splitting.,"F9 ) and F11 ), it forms a substructure that we interpret as rotational splitting."
+ The derived equatorial rotational velocity amounts to 15 kms. nαυ], The derived equatorial rotational velocity amounts to 15 $^{\rm -1}$.
+ the case of models obtained with standard chemical composition. mode instability is incorrectly. reproduced: instead. a reduction in metallicity (from 00.0134 to 0.010) and ar increase in helium abundance (from 00.2466 to 0.29) has to be adopted to fit the observed instability range.," In the case of models obtained with standard chemical composition, mode instability is incorrectly reproduced; instead, a reduction in metallicity (from 0.0134 to 0.010) and an increase in helium abundance (from 0.2466 to 0.29) has to be adopted to fit the observed instability range."
+ In fact. the reduction of metallicity implies that the observed frequencies are fitted by models at higher effective temperature. where higher order modes are driven.," In fact, the reduction of metallicity implies that the observed frequencies are fitted by models at higher effective temperature, where higher order modes are driven."
+ If the assumption about slow rotation holds true. then we are able to test the physical conditions of the envelope within the frame of the existing atomic physics data in HD 144277.," If the assumption about slow rotation holds true, then we are able to test the physical conditions of the envelope within the frame of the existing atomic physics data in HD 144277."
+(Figure 5)).,(Figure \ref{fig3}) ).
+" With this diagram, the +1 ambiguity over the value of e, and therefore radial order, can be resolved."," With this diagram, the $\pm1$ ambiguity over the value of $\epsilon$, and therefore radial order, can be resolved."
+ The extension of this relationship to higher Tog shows promise for resolving the mode identification problem in F starsreferred to in Section ?? (White et al., The extension of this relationship to higher $T_\mathrm{eff}$ shows promise for resolving the mode identification problem in F starsreferred to in Section \ref{sec3} (White et al.
+ in prep)., in prep).
+ 'The authors gratefully acknowledge theKepler Science Team and all those who have contributed to theMission for their tireless efforts which have made these results possible., The authors gratefully acknowledge the Science Team and all those who have contributed to the for their tireless efforts which have made these results possible.
+ Funding for theMission is provided by NASA's Science Mission Directorate., Funding for the is provided by NASA's Science Mission Directorate.
+" TRW is supported by an Australian Postgraduate Award, a University of Sydney Merit Award, an Australian Astronomical Observatory PhD Scholarship and a Denison Merit Award."," TRW is supported by an Australian Postgraduate Award, a University of Sydney Merit Award, an Australian Astronomical Observatory PhD Scholarship and a Denison Merit Award."
+ TRB and DS acknowledge the support of the Australian Research Council., TRB and DS acknowledge the support of the Australian Research Council.
+ SH acknowledges financial support from the Netherlands Organisation for Scientific Research., SH acknowledges financial support from the Netherlands Organisation for Scientific Research.
+ This research was supported by grant AYA2010-17803 from the Spanish National Research Plan., This research was supported by grant AYA2010-17803 from the Spanish National Research Plan.
+ NCAR is supported by the National Science Foundation., NCAR is supported by the National Science Foundation.
+the possible redshifts of the NLAMDO sources. although for 9 of the LO sources. they are cousistent with Ξ11 witlin the uncertainties.,"the possible redshifts of the MAMBO sources, although for 9 of the 10 sources, they are consistent with $z$ =4.1 within the uncertainties."
+" e One of the faint MAMDO sources is possibly related to an ERO with RoN= 6.1. which has a likely spectroscopic redshift of :—1.18. iuplviung a star formation rate up to 1600 yr +,"," $\bullet$ One of the faint MAMBO sources is possibly related to an ERO with $R-K$ =6.1, which has a likely spectroscopic redshift of $z$ =1.18, implying a star formation rate up to 1600 $_{\odot}$ $^{-1}$."
+ e None of the 114 spectroscopically coufinued celitters show detectable millimetre cussion. acd supposing they are at z—L1. none of the 1.2 mun sources show an excess cecluission. indicating no apparent overlap between both populations.," $\bullet$ None of the 14 spectroscopically confirmed emitters show detectable millimetre emission, and supposing they are at $z$ =4.1, none of the 1.2 mm sources show an excess emission, indicating no apparent overlap between both populations."
+ e The iean star formation rate of the 11 spectroscopically coufirmed cclitters. as derived from the deep ldmaeing is two orders of magnitudes lower than the upper lianits derived from the stacked VLA aud MAÀAMDO maps. illustrating that the racio and nuu maps probe imch more actively star-forming ealaxics.," $\bullet$ The mean star formation rate of the 14 spectroscopically confirmed emitters, as derived from the deep imaging is two orders of magnitudes lower than the upper limits derived from the stacked VLA and MAMBO maps, illustrating that the radio and mm maps probe much more actively star-forming galaxies."
+ Our iuultiwaveleugth observations of this +=1.1 proto-cluster sugeest that the cexcess technique docs not detect the massive starburst conrpanons within 2 Alpe., Our multi-wavelength observations of this $z$ =4.1 proto-cluster suggest that the excess technique does not detect the massive starburst companions within 2 Mpc.
+ Their high SER could be tregecred by iuteractious between different. proto-cluster members. as iudicated by the diffuse nature of the optical identifications of MO5 aud MO2.," Their high SFR could be triggered by interactions between different proto-cluster members, as indicated by the diffuse nature of the optical identifications of M05 and M02."
+l galaxies.,1 galaxies.
+ This rather steep function iupics that the illumination - and therefore the line cussion - ix strongly concentrated iu the inner regions of the ACN. with ~DO of the line coming from within 2¢IR. and SO from within LOOR..," This rather steep function implies that the illumination - and therefore the line emission - is strongly concentrated in the inner regions of the AGN, with $\sim 50$ of the line coming from within $20 R_{\rm g}$ and 80 from within $100 R_{\rm g}$."
+ Such an Cluissivity is roughly consistent with hat of a ceutrallv-illuninated disk., Such an emissivity is roughly consistent with that of a centrally-illuminated disk.
+" No universal enassivitv kaw was found. howeer. sugecsting that there is no single ecometry,"," No universal emissivity law was found, however, suggesting that there is no single geometry."
+ The emissiou line usually peaks at 6. keV. and for low-inclination disks this inplies a low state of jonization. typically < FoXN.," The emission line usually peaks at 6.4 keV, and for low-inclination disks this implies a low state of ionization, typically $<$ Fe."
+ If the peak is due to a contribution from a narrow hne from another source. however. no constraint can be placed ou the hue rest energv (N97). and theretore lonization.," If the peak is due to a contribution from a narrow line from another source, however, no constraint can be placed on the line rest energy (N97), and therefore ionization."
+ Standard accretiou disks are expected to be very dense. aud cam therefore remain cool despite the intense illumination.," Standard accretion disks are expected to be very dense, and can therefore remain cool despite the intense illumination."
+ Nonetheless. siguificant ionization nüeht be expected especially in the central regions. where the X-rays are most strouelv concentrated (e.g. Matt. Fabian Ross 1993).," Nonetheless, significant ionization might be expected especially in the central regions, where the X-rays are most strongly concentrated (e.g. Matt, Fabian Ross 1993)."
+ A range of ionization sates throughout the disk is certainly cousisteut wit ithe curent data., A range of ionization sates throughout the disk is certainly consistent with the current data.
+ The iuner disk could be highly ionized. with the eravitational alc Doppler shifts dominating aud makine it difficult to tie down the rest enerev.," The inner disk could be highly ionized, with the gravitational and Doppler shifts dominating and making it difficult to tie down the rest energy."
+"data (Scherreretal.1995)., and also perform inversions for one selected active region to examine how this effect may affect inverted interior structures.","data \citep{sch95}, and also perform inversions for one selected active region to examine how this effect may affect inverted interior structures."
+ These results are presented in Section 2.., These results are presented in Section \ref{sec2}.
+" We then present in Section 3 the showerglass effect measurements, and in Section 4. the second-skip travel times measured trom annulus-annulus cross-correlations, which demonstrate that oscillation signals inside sunspots do not introduce significant extra phase shifts to time-distance measurements."," We then present in Section \ref{sec3}
+ the showerglass effect measurements, and in Section \ref{sec4} the second-skip travel times measured from annulus-annulus cross-correlations, which demonstrate that oscillation signals inside sunspots do not introduce significant extra phase shifts to time-distance measurements."
+ Discussions and conclusions follow in Section 4 and 5.., Discussions and conclusions follow in Section \ref{sec4} and \ref{sec5}.
+" However, it should be borne in mind that the motivation of this paper is not to compare directly measurements from time-distance and helioseismic holography techniques, but to assess how some measurement effects that were reported from holography studies affect our typical time-distance measurements."," However, it should be borne in mind that the motivation of this paper is not to compare directly measurements from time-distance and helioseismic holography techniques, but to assess how some measurement effects that were reported from holography studies affect our typical time-distance measurements."
+" To study the inclined magnetic field effect, we select the same sunspot in active region AR9026 as studied by Schunkeretal.(2005) during the similar observational periods."," To study the inclined magnetic field effect, we select the same sunspot in active region AR9026 as studied by \citet{sch05} during the similar observational periods."
+" This sunspot appeared from the solar east limb on June 2, 2000 at the latitude of 19:5 in the Northern hemisphere."," This sunspot appeared from the solar east limb on June 2, 2000 at the latitude of $19\fdg5$ in the Northern hemisphere."
+" Our time-distance analysis is different from the holography analysis (Schunkeretal.2005). in several aspects, including observational duration, annulus radius and width (or pupil size). frequency bands selection, and the use of phase-speed filtering."," Our time-distance analysis is different from the holography analysis \citep{sch05} in several aspects, including observational duration, annulus radius and width (or pupil size), frequency bands selection, and the use of phase-speed filtering."
+" Time-distanee helioseismology employs speed filtering procedure (c.g.,Duvalletal.1996). which is designed to keep all the acoustic waves that have the same first bounce distances and similar wave propagation speed (therefore, traveling similar distances with similar times), and filter out all other acoustic waves."," Time-distance helioseismology employs phase-speed filtering procedure \citep[e.g.,][]{duv96}, which is designed to keep all the acoustic waves that have the same first bounce distances and similar wave propagation speed (therefore, traveling similar distances with similar times), and filter out all other acoustic waves."
+" Applying phase-speed filtering makes the time-distance analysis with short annulus radius and short annulus width possible, and significantly improves the signal to noise ratio of time-distance measurement."," Applying phase-speed filtering makes the time-distance analysis with short annulus radius and short annulus width possible, and significantly improves the signal to noise ratio of time-distance measurement."
+" We selected three dates for analysis, and for each date, only observations of the first 512 minutes of the day are used."," We selected three dates for analysis, and for each date, only observations of the first 512 minutes of the day are used."
+ The p-modes frequency range covers approximately 2.5—5.5 mHz after filtering. though the acoustic power peaks at approximately 3.5 mHz.," The $p$ -modes frequency range covers approximately $2.5 - 5.5$ mHz after filtering, though the acoustic power peaks at approximately $3.5$ mHz."
+" The acoustic travel times are measured between a central point and the surrounding annulus, with the annulus size range of 3.7—8.7 Mm."," The acoustic travel times are measured between a central point and the surrounding annulus, with the annulus size range of $3.7 - 8.7$ Mm."
+" Following the data analysis procedure described by Zhao,Kosovichev.&Duvall (2001). we compute the acoustic travel times inside the sunspot penumbra for the selected observation periods."," Following the data analysis procedure described by \citet{zha01}, , we compute the acoustic travel times inside the sunspot penumbra for the selected observation periods."
+" Figure 1. presents the mean travel times, which are averages of the outgoing and ingoing travel times, as functions of the penumbral azimuthal angle for different dates."," Figure \ref{fg1} presents the mean travel times, which are averages of the outgoing and ingoing travel times, as functions of the penumbral azimuthal angle for different dates."
+" The azimuthal angle starts from Q in the West of the sunspot umbra center, and increases counter-clockwise inside the sunspot penumbra."," The azimuthal angle starts from $0\degr$ in the West of the sunspot umbra center, and increases counter-clockwise inside the sunspot penumbra."
+" Basically, the mean acoustic travel times are longer inside the sunspot penumbra than these in the quiet Sun for this specific annulus range that we use."," Basically, the mean acoustic travel times are longer inside the sunspot penumbra than these in the quiet Sun for this specific annulus range that we use."
+It can be found that for each,It can be found that for each
+where yo=yl)—o).,where $\varphi_- \equiv \varphi(t=-\infty)$.
+ The energv and angular momentum of the photon can be determined from & aud ((2.1)) with /2—2x., The energy and angular momentum of the photon can be determined from $b$ and \ref{eqNewton}) ) with $t=-\infty$.
+ E=1J=b (, E =; J = b .
+3) The stationary condition *=0 has one solution because £z0. and that is the distance (ry) of the closest approach of the photon to the mass AL.," The stationary condition $\dot r =0$ has one solution because $E >0$ , and that is the distance $r_0$ ) of the closest approach of the photon to the mass $M$."
+ y= —GM + (GAAP +?(4) since J>0. the azimuthal angle increases with time (45>0).," r_0 = -GM + (G^2 M^2 + Since $J > 0$, the azimuthal angle increases with time $\dot \varphi > 0$ )."
+ The equation of the orbit can be found by a simple integration of ((2.1)) alter eliminating d/., The equation of the orbit can be found by a simple integration of \ref{eqNewton}) ) after eliminating $dt$.
+ For the incoming photon. I DEία42.pUTM and the equation of the orbit is given by £3 Hem ) )..(6) where B=GALJ? and A=(J74BYP," For the incoming photon, = ^r - } (1 + -, and the equation of the orbit is given by - _- = ) ), where $B \equiv GMJ^{-2}$ and $A \equiv (J^{-2} + B^2)^{1/2}$."
+ IE py=un). the photon trajectory is reflection symmetric wilh respect to the periastron. and its deflection angle is given by = 2 )—m(0) ," If $\varphi_0 \equiv \varphi(r_0)$, the photon trajectory is reflection symmetric with respect to the periastron, and its deflection angle is given by = _0 _-) - = 2 ) -."
+The equation of the orbit can berewritten interms of jy and results in the standard equation of a conic., The equation of the orbit can berewritten interms of $\varphi_0$ and results in the standard equation of a conic.
+Magrini et al. (,Magrini et al. (
+2009) also noted a flattening in the oxygen gradient in their sample of M33 planetary nebulae.,2009) also noted a flattening in the oxygen gradient in their sample of M33 planetary nebulae.
+ Figure 7 shows the present day gas surface density distribution for the galaxies studied here., Figure \ref{gas} shows the present day gas surface density distribution for the galaxies studied here.
+ For the Milky Way we note that the model MW-B describes better the present surface gas density of the disk with a smoother distribution while MW A] and MW - À2 present a remarkable break in the profile (probably associated with the constant efficiency in the SER)., For the Milky Way we note that the model MW-B describes better the present surface gas density of the disk with a smoother distribution while MW - A1 and MW - A2 present a remarkable break in the profile (probably associated with the constant efficiency in the SFR).
+ For M31 the models with constant star formation efficiency can reproduce an exponential profile but fail to explain the peak located at about 12kpe of distance from the galaxy center., For M31 the models with constant star formation efficiency can reproduce an exponential profile but fail to explain the peak located at about 12kpc of distance from the galaxy center.
+ The models with variable efficiency (M31-B and M31-Bk1.25) can instead reproduce this peak as shown in figure 7.., The models with variable efficiency (M31-B and M31-Bk1.25) can instead reproduce this peak as shown in figure \ref{gas}.
+ Both models for M33 are capable of explaining the present day gas distribution of Boissier et al. (, Both models for M33 are capable of explaining the present day gas distribution of Boissier et al. (
+2007) for R> 4 kpe but overestimate the gas surface density m the inner regions.,2007) for $>$ 4 kpc but overestimate the gas surface density in the inner regions.
+ In Figure 9.. we show the predicted stellar. density distributions along the three disks.," In Figure \ref{stars}, we show the predicted stellar density distributions along the three disks."
+ The Milky Way models are in good agreement with the observed values (shaded area)., The Milky Way models are in good agreement with the observed values (shaded area).
+ All this area was obtained by assuming an exponential distribution (see also CMR2001) and using a scale length radius Ryvers equal to 2.5 kpe (Freudenreich 1998) together with a stellar density for the solar annulus equal to 35M..pce7* (Gilmore et al., All this area was obtained by assuming an exponential distribution (see also CMR2001) and using a scale length radius $R_{Dstars}$ equal to 2.5 kpc (Freudenreich 1998) together with a stellar density for the solar annulus equal to $35 M_{\odot}pc^{-2}$ (Gilmore et al.
+ 1998). in order to scale the distribution to the observed values.," 1998), in order to scale the distribution to the observed values."
+ The predictions for M31 stellar density profile show a shallower distribution. whereas for M33 our models present à steeper distribution (even without taking into account the inner part of the M33 disk which is also overestimated as a consequence of the gas profile - see figure 7)).," The predictions for M31 stellar density profile show a shallower distribution, whereas for M33 our models present a steeper distribution (even without taking into account the inner part of the M33 disk which is also overestimated as a consequence of the gas profile - see figure \ref{gas}) )."
+ In Figure 8. we show the predicted distributions of the SER in the three disks compared to observational data., In Figure \ref{sfr} we show the predicted distributions of the SFR in the three disks compared to observational data.
+ All the models can reproduce observational trends and the ditferences between, All the models can reproduce observational trends and the differences between
+also useful.,also useful.
+ Uaiman&IIolder(2003) show that we cau distinguish how ο. times the universe is reionized bv using EE spectrum with Planck seusitivity in 36 level., \cite{hai03} show that we can distinguish how much times the universe is reionized by using EE spectrum with Planck sensitivity in $\sigma$ level.
+ Since the seusitivitv ofWALAP is not enough. we iust await the launch of to reveal the reionization history by the CAIB mncasurement.," Since the sensitivity of is not enough, we must await the launch of to reveal the reionization history by the CMB measurement."
+ Moreover. ictal absorption lines like ο A1302 can be useful to prove the reionizatiou history (Oh2002:Furlanctto&Loeb 2003).," Moreover, metal absorption lines like O $\lambda1302$ can be useful to prove the reionization history \citep{oh02,fur03}."
+. Tow:er. the expected equivalent width is very small as XA.," However, the expected equivalent width is very small as $\la 5$."
+ We require the spectroscopy with a resolving power ~5000., We require the spectroscopy with a resolving power $\sim 5000$.
+ Unfortunately. observations with a ground-based telescope iiw be diffieult because a huge imnber of Earth's atinospherie ΟΠ lines conceal the metal lines.," Unfortunately, observations with a ground-based telescope may be difficult because a huge number of Earth's atmospheric OH lines conceal the metal lines."
+ Ouly JWST (James Webb Space will rave such a high spectral resolution :unong the future space telescopes having the NIR spectrograph., Only JWST (James Webb Space will have such a high spectral resolution among the future space telescopes having the NIR spectrograph.
+ Thus. w«| need to await its launch to discuss the metal absorption nes.," Thus, we need to await its launch to discuss the metal absorption lines."
+ Technique using the hydrogen 21 €i absorption liue is also proposed(c.g. Aladau et al.," Technique using the hydrogen 21 cm absorption line is also proposed (e.g., Madau et al."
+ 1997)., 1997).
+ Ilowever. the brightness of GRD afterglows is too unt to use them as a backeround light source for the esent and future radio facilities because the absorption line is verv weak. (Furlanetto&Loeb2002).," However, the brightness of GRB afterglows is too faint to use them as a background light source for the present and future radio facilities because the absorption line is very weak \citep{fur02}."
+. Thus. we need to look for otler sources.," Thus, we need to look for other sources."
+ Although Carilli et al. (, Although Carilli et al. (
+2002) proposed Iuninous ieli- radio-loud quasars as a candidate of the backerouud source.z It is very uncertain that they exist at 2—20.,"2002) proposed luminous $z$ radio-loud quasars as a candidate of the background source, it is very uncertain that they exist at $z\sim 20$."
+ The dispersion measure in GRB radio afterglows may © promising iu near future (loka2003:Tnoue2003).," The dispersion measure in GRB radio afterglows may be promising in near future \citep{iok03,ino03}."
+. If we observe the radio afterglow at about 100 NIIz within about 1000 s after the burst occurrence. the delav of he arrival time of the low frequeney photous may be detectable by the Square Nilometer Array.," If we observe the radio afterglow at about 100 MHz within about 1000 s after the burst occurrence, the delay of the arrival time of the low frequency photons may be detectable by the Square Kilometer Array."
+ As well as the ucasureineut of the red damping wing of the Lya line. his technique is seusitive to eji20.1 (nec ThyZ 107).," As well as the measurement of the red damping wing of the $\alpha$ line, this technique is sensitive to $x_{\rm HI}\ga 0.1$ (i.e., $\tau_{\rm Ly\alpha}\ga 10^5$ )."
+ Thus. this techuique aud our NIR color method are also complementary cach other.," Thus, this technique and our NIR color method are also complementary each other."
+ The disadvantage of the NIR color method is the coarse redshift resolution., The disadvantage of the NIR color method is the coarse redshift resolution.
+ As discussed in section L. the neutral fraction obtained by the color method is au averaged oue in the redshift range of ipu“2Xty.," As discussed in section 4, the neutral fraction obtained by the color method is an averaged one in the redshift range of $z_{\rm Ly\alpha,in} \leq z \leq z_{\rm S}$."
+ Thus. we cannot determine the so-called reiouizatiou redslitt.," Thus, we cannot determine the so-called reionization redshift."
+ Oulv the redshift range iu which the reionization occurred is obtained., Only the redshift range in which the reionization occurred is obtained.
+ However. even such a rough redshift resolution is enough to discuss whether the universe was reionized once. twice. or niore.," However, even such a rough redshift resolution is enough to discuss whether the universe was reionized once, twice, or more."
+ To know the detailed historv of the reionization. the spectroscopy is needed.," To know the detailed history of the reionization, the spectroscopy is needed."
+ For this purpose. the Japanese astronomical satellite cau be useful.," For this purpose, the Japanese astronomical satellite can be useful."
+ The satellite will have the spectroscopic seusitivitv of ~30 gr Jv around ~2gan. which can detect the CRB afterelows of no20 fey 10Ὁ (see Fig.2a).," The satellite will have the spectroscopic sensitivity of $\sim 30$ $\mu$ Jy around $\sim 2$, which can detect the GRB afterglows of $z_{\rm S}\sim 20$ if $x_{\rm HI}\sim 10^{-6}$ (see Fig.2a)."
+ SIRTF (Spece IufraBRed Telescope may not be useful because it has the sensitivity oulv in the wavelength longer than 5jan., SIRTF (Spece InfraRed Telescope may not be useful because it has the sensitivity only in the wavelength longer than 5.
+.. Iu future. JWST is πιο promising.," In future, JWST is much promising."
+" Finally, we note that the narrow-band photometry nay be useful because its narrow transmission width provides us with a moderate resolution of redshift keeping a higher sensitivity than that of spectroscopy."," Finally, we note that the narrow-band photometry may be useful because its narrow transmission width provides us with a moderate resolution of redshift keeping a higher sensitivity than that of spectroscopy."
+ We thank παήτο Ioka for valuable couuncuts., We thank Kunihito Ioka for valuable comments.
+" This work is supported in part by the Coant-n-Aid for Scieutific Research of the Japauese Ministry of Education. Culture. Sports. Science aud Technology. No.11017212. (TN) and No.11201021. CEN). and also supported. bv a. Craut-iu-Aid for the 21st Century COE ""Center for Diversity aud Universality iu Physics."," This work is supported in part by the Grant-in-Aid for Scientific Research of the Japanese Ministry of Education, Culture, Sports, Science and Technology, No.14047212 (TN) and No.14204024 (TN), and also supported by a Grant-in-Aid for the 21st Century COE ”Center for Diversity and Universality in Physics”."
+ AKI aud RY thank the Besearch Fellowships of the Japan Society for the Promotion of Scicuce for Young Scieutists., AKI and RY thank the Research Fellowships of the Japan Society for the Promotion of Science for Young Scientists.
+We assune that the enmissiou region is a honiogeneous sphere with radius R.,We assume that the emission region is a homogeneous sphere with radius $R$.
+ The magnetic fields are assumed to lave isotropic orieuntatious., The magnetic fields are assumed to have isotropic orientations.
+ Several electron populations with isotropic momentum distributions and power-lav uonmientuni spectra are assumed to exist simiultaucouslv in the enüssion region., Several electron populations with isotropic momentum distributions and power-law momentum spectra are assumed to exist simultaneously in the emission region.
+ Initially. thev are assuned to be »owoer-Imwv distibutious: The index / labels the different spectra. which have rormalizations of C and spectral iudices s;;," Initially, they are assumed to be power-law distributions: The index $i$ labels the different spectra, which have normalizations of $C_i$, and spectral indices $s_i$."
+ Or) is he Teaviside step-function. sohat for cmeusiouless electron momenta p—μιςο) outside pi;Oppad he electron spectra vanish.," $\Theta(x)$ is the Heaviside step-function, sothat for dimensionless electron momenta $p = P_\e
+/(m_\e \,c)$ outside $p_{1,i}<p<p_{2,i}$ the electron spectra vanish."
+ The freedom to have several xopulations allows to model complex spectra with brokeu oower-Iwwvs or spectral discoutinuitics., The freedom to have several populations allows to model complex spectra with broken power-laws or spectral discontinuities.
+ We also allow that cach electron population is located within a differcut naguctic field streneth [δει so that simall-scale multi-phase uedia can be modeled.," We also allow that each electron population is located within a different magnetic field strength $B_i$, so that small-scale multi-phase media can be modeled."
+" After he supply of fresh clectrous has stopped. the electron spectra evolve according to svuchrotrou. IC', aud adiabatic enerev losses: where ay=icr(noc). V is the source volue. ap sud ver are the magnetic aud CAIB eneres* deusities respectively."," After the supply of fresh electrons has stopped, the electron spectra evolve according to synchrotron, IC, and adiabatic energy losses: where $a_{0} = \frac{4}{3}\, \st/(m_\e\, c)$, $V$ is the source volume, $u_B$ and $u_C$ are the magnetic and CMB energy densities respectively."
+ We do not consider bronisstrahlung aud Coulonib losses iu view of the very low particle density iun most of the cousidered sources., We do not consider bremsstrahlung and Coulomb losses in view of the very low particle density in most of the considered sources.
+ We asstune that the expansion of the source volume can be described by a power-law iu BE where2 Citft) denotes the compression factor., We assume that the expansion of the source volume can be described by a power-law in time: where $\tilde{C}(t)$ denotes the compression factor.
+ Tn this case the evolution of the ποιοιτα. of au electron residiug in a magnetic field streneth D;C? can be even analytically (ναίetal.Eufblin&CGopal-Ixishua.2001. 6.9.):: where the initial momentum is py aud q;(f) can be read as the inverse of the maximal possible electron momentum p;(t) in population /: (C=Cit) for brevity).," In this case the evolution of the momentum of an electron residing in a magnetic field strength $B_i(t) = B_i\,\tilde{C}^{2/3}$ can be given analytically \cite[e.g.]{1997MNRAS.292..723K,2001A&A...366...26E}: where the initial momentum is $p_0$ and $q_i(t)$ can be read as the inverse of the maximal possible electron momentum $p_{i}^*(t)$ in population $i$: $\tilde{C} = \tilde{C}(t)$ for brevity)."
+" The electron spectra evolve according to where pijtf)=pitf) aud poil(t)=(pr,jef.We ορ.)"," The electron spectra evolve according to where $p_{1,i}(t)= p_i(p_{1,i},t)$ and $p_{2,i}(t) =
+p_i(p_{2,i},t)$."
+ also write f(p) for f(p.t)=for brevity in the following aud also drop the explicit time depeudence iu our notation of the resulting photon spectra.," We also write $f(p)$ for $f(p,t)=\sum_i \,f_i(p,t)$ for brevity in the following and also drop the explicit time dependence in our notation of the resulting photon spectra."
+ Similar evolution fonxuulae can be derived for av arbitrary initial electron spectrum., Similar evolution formulae can be derived for any arbitrary initial electron spectrum.
+" Since it js used in our exaniples. we eive here the definition of a relativistic thermal distribution: with A, denoting the modified Bessel function. 04 the thermal electron παπαος: density. Ah=ewPAT. the normalized thermal beta-paranmeter. aud Zi the electron teniperature."," Since it is used in our examples, we give here the definition of a relativistic thermal distribution: with $K_\nu$ denoting the modified Bessel function, $n_{\e,\th}$ the thermal electron number density, $\beta_{\rm th} = \me\, c^2/kT_\e$ the normalized thermal beta-parameter, and $T_\e$ the electron temperature."
+ The pitch-angle averaged svuchrotron energy losses of an ultra-relativistic electrou located im maguctic fields of streugth B are radivted in the form of photous at the characteristic svuchrotron frequency term A =S3ef(27oe)d introduced for convenieuce.," The pitch-angle averaged synchrotron energy losses of an ultra-relativistic electron located in magnetic fields of strength $B$ are radiated in the form of photons at the characteristic synchrotron frequency The term $\Lambda = 3\, e/(2\,\pi\,\me\,c)$ is introduced for convenience."
+ The enüssvitv of a sinele electron located in field streugth £2; is therefore in the monochromatic approximation. with dr) denoting Dirac’s delta function.," The emissivity of a single electron located in field strength $B_i$ is therefore in the monochromatic approximation, with $\delta(x)$ denoting Dirac's delta function."
+ The source deuxitv of svuchrotron plotous produced in the frequency interval dv by our electron populations is Eq., The source density of synchrotron photons produced in the frequency interval $d\nu$ by our electron populations is Eq.
+ BS is the general form and Eq., \ref{eq:qsync1} is the general form and Eq.
+ B7 is valid for the power-law electron populationsgiven in Eq. Al.., \ref{eq:qsync2} is valid for the power-law electron populations given in Eq. \ref{eq:finitial}. .
+ Here mit)=ABA(t)pyft) and ο)=AB;(1)psgU) ave the lower and upper frequency cutofts.," Here $\nu_{1,i}(t) = \Lambda\,B_i(t)\,p_{1,i}^2(t)$ and $\nu_{2,i}(t) = \Lambda\,B_i(t)\,p_{2,i}^2(t)$ are the lower and upper frequency cutoffs."
+ The svuchrotrou absorption cocfitcicut within the the monochromatic approxiuation (Eq. D3)), The synchrotron absorption coefficient within the the monochromatic approximation (Eq. \ref{eq:Psync}) )
+ can be calculated from the eeneric form given in Bwbicki Lightumau 97 aN:, can be calculated from the generic form given in Rybicki Lightman \cite*{1979rpa..book.....R}: :
+such SDSS galaxies (5.3> random).,such SDSS galaxies (5.3 random).
+ If neither FIRST source is within 3 aresec. then it is possible that the two FIRST components are two lobes of the same extended radio source with no core.," If neither FIRST source is within 3 arcsec, then it is possible that the two FIRST components are two lobes of the same extended radio source with no core."
+ As discussed earlier. McMahon shortcitememO02 found that in this case the two FIRST sources often have the following properties: () they have comparable flux densities: (1) the flux-weighted mean position of the two sources Is close to that of the optical galaxy: (11) the sizes of the two sources are comparable (ie.," As discussed earlier, McMahon \\shortcite{mcm02} found that in this case the two FIRST sources often have the following properties: (i) they have comparable flux densities; (ii) the flux-weighted mean position of the two sources is close to that of the optical galaxy; (iii) the sizes of the two sources are comparable (ie."
+ both are lobes. not a core and a lobe).," both are lobes, not a core and a lobe)."
+ None of these conditions on their own is sufficient to classify the source as a double without including a lot of false detections. but the Monte—Carlo simulations show that the combination of the three can be quite powerful.," None of these conditions on their own is sufficient to classify the source as a double without including a lot of false detections, but the Monte--Carlo simulations show that the combination of the three can be quite powerful."
+ Galaxies with two FIRST matches were accepted as double radio sources if the ratio of the radio source sizes. multiplied by the ratio of the radio source flux densities. multiplied by the offset in arcsec of the flux-weighted mean position. is less than 5.," Galaxies with two FIRST matches were accepted as double radio sources if the ratio of the radio source sizes, multiplied by the ratio of the radio source flux densities, multiplied by the offset in arcsec of the flux–weighted mean position, is less than 5."
+ Figure 6 shows the result of this analysis: 116 double sources are selected among SDSS galaxies. but only 1.2 for random positions. implying that the reliability is about99%.," Figure \ref{firstdoubles} shows the result of this analysis: 116 double sources are selected among SDSS galaxies, but only 1.2 for random positions, implying that the reliability is about."
+. It should be noted here that in cases where the galaxy is associated with a single FIRST component. and the other component (or components. for the cases with 3 or more matches discussed below) is genuinely unassociated. the NVSS flux will overestimate the radio luminosity due to the contaminating source.," It should be noted here that in cases where the galaxy is associated with a single FIRST component, and the other component (or components, for the cases with 3 or more matches discussed below) is genuinely unassociated, the NVSS flux will overestimate the radio luminosity due to the contaminating source."
+ However. only in very rare cases (e.," However, only in very rare cases (ie."
+ a faint point source lying nearby a much brighter source) would such a correction be, a faint point source lying nearby a much brighter source) would such a correction be
+on August 9 than they were in July but with lower amplitudes the same model provides a good fit. and a QPO does still appear to be required.,"on August 9 than they were in July but with lower amplitudes the same model provides a good fit, and a QPO does still appear to be required."
+ We did have to fix the widths of the QPOs on August 9 to the average values of the July observations to adequately constrain the fit., We did have to fix the widths of the QPOs on August 9 to the average values of the July observations to adequately constrain the fit.
+ The most important quantity for our later correlation analysis is the characteristic period of the primary QPO., The most important quantity for our later correlation analysis is the characteristic period of the primary QPO.
+ This is ss for July 6. ss for July 7. and 4.0ss for August 9.," This is s for July 6, s for July 7, and s for August 9."
+ Since the Argos time-resolution is only ss. we would not expect the QPOs to be resolved in the July optical data and their main effect is likely to be indistinguishable from an excess of white noise.," Since the Argos time-resolution is only s, we would not expect the QPOs to be resolved in the July optical data and their main effect is likely to be indistinguishable from an excess of white noise."
+ The ss ΟΡΟ in August could have a much more significant effect on our analysis., The s QPO in August could have a much more significant effect on our analysis.
+ Cross correlation. functions. (CCFs) for the July 6 and 7 observations near the outburst peak are shown in Fig. 5.., Cross correlation functions (CCFs) for the July 6 and 7 observations near the outburst peak are shown in Fig. \ref{ArgosCCFFig}.
+ Although correlations were not obvious in the lighteurve they are clearly present at high significance. with a strong response at or close to zero lag. and an extended tail lasting for a few seconds.," Although correlations were not obvious in the lightcurve they are clearly present at high significance, with a strong response at or close to zero lag, and an extended tail lasting for a few seconds."
+ The obvious interpretation of the correlation is that it arises from thermal reprocessing either in an accretion disk or on the surface of the companion star., The obvious interpretation of the correlation is that it arises from thermal reprocessing either in an accretion disk or on the surface of the companion star.
+ The lag of the companion star response modulates with orbital phase with an average of approximately the binary separation in. lightseconds., The lag of the companion star response modulates with orbital phase with an average of approximately the binary separation in lightseconds.
+ With. the parameters discussed in Section 7.. the binary separation. would be about 5llightseconds and so the companion peak lag would modulate in the ss range.," With the parameters discussed in Section \ref{TransferFunctionSection}, the binary separation would be about lightseconds and so the companion peak lag would modulate in the s range."
+ Since our observations span hhrs each night and the orbital period is estimated to be hhrs 2008).. we cover about half an orbit on each night.," Since our observations span hrs each night and the orbital period is estimated to be hrs \citep{Zurita:2008a}, we cover about half an orbit on each night."
+ On both nights we observe a peak lag close to zero. so it is unlikely that the companion is dominating the response. which instead appears to originate from the accretion disk.," On both nights we observe a peak lag close to zero, so it is unlikely that the companion is dominating the response, which instead appears to originate from the accretion disk."
+ Additionally. black hole systems are expected to usually have small mass ratios. and the presence of likely superhumps (Zuritaetal.2007) supports this expectation in this case.," Additionally, black hole systems are expected to usually have small mass ratios, and the presence of likely superhumps \citep{Zurita:2007a} supports this expectation in this case."
+ If the mass ratio is indeed small. then the companion star will subtend only a small solid angle as seen by the X-ray source. and hence would be expected to reprocess much less light than the accretion disk.," If the mass ratio is indeed small, then the companion star will subtend only a small solid angle as seen by the X-ray source, and hence would be expected to reprocess much less light than the accretion disk."
+ We should be somewhat cautious. however. as similar CCFs were seen in the UV in XTE JI1184480 (although not in the optical). where they may be associated with jet emission.," We should be somewhat cautious, however, as similar CCFs were seen in the UV in XTE J1118+480 (although not in the optical), where they may be associated with jet emission."
+ This seems less likely in this ease. however. as the correlations are observed at the peak of an apparently normal outburst. and the optical spectrum is blue and can be fitted by that of an accretion disk (Stilletal.2005)..," This seems less likely in this case, however, as the correlations are observed at the peak of an apparently normal outburst, and the optical spectrum is blue and can be fitted by that of an accretion disk \citep{Still:2005a}."
+ In addition. the optical auto-correlation function is found to be broader than the X-ray one as expected for thermal reprocessing. and unlike that seen in XTE JH1184480.," In addition, the optical auto-correlation function is found to be broader than the X-ray one as expected for thermal reprocessing, and unlike that seen in XTE J1118+480."
+ We cannot conclusively rule this possibility out. however. especially as," We cannot conclusively rule this possibility out, however, especially as"
+2008).. have poiuted out that radio sources Gadio galaxies aud radio-oud ACN) could also make a siguificaut contzibution to the sub-11Jx population.,", have pointed out that radio sources (radio galaxies and radio-loud AGN) could also make a significant contribution to the sub-mJy population."
+ It has been further suggested (Jarvis&Rawl-ines2001:Simpsonetal.2006) that radio-quiet AGN (see below) might become relevaut at sub-anJxy. levels.," It has been further suggested \citep{jar04,sim06} that radio-quiet AGN (see below) might become relevant at sub-mJy levels."
+ The Chandra Deep Field Southp. (CDFES) area is part of adje Great Observatories Origins ep Survey (GOODS) nd as such is one of the most inteuscly studied region of the sky., The Chandra Deep Field South (CDFS) area is part of the Great Observatories Origins Deep Survey (GOODS) and as such is one of the most intensely studied region of the sky.
+ Theh seusitivitv X-ray observations are available from Chandra (Cdacconietal.2002:Lelunerct2005:Luoetal. 2008).," High sensitivity X-ray observations are available from Chandra \citep{gia02,leh05,luo08}."
+. The GOODS aud Galaxy Evolution from Morphology aud Spectral cnerey distributions (GENS) απας iuagiug programs using the TST Advanced Camera for Surveys (ACS) eive sensitive high resolution optical images (Cdavaliscoctal.2001:Rixal. 2001).," The GOODS and Galaxy Evolution from Morphology and Spectral energy distributions (GEMS) multiband imaging programs using the HST Advanced Camera for Surveys (ACS) give sensitive high resolution optical images \citep{gia04,rix04}."
+. Ground based maging aud spectroscopy are available frou the ESO 2.211 and Sia telescopes (seereferencesinMainuierietal.2008.Paper IT)., Ground based imaging and spectroscopy are available from the ESO 2.2m and 8m telescopes \citep[see references in][Paper II]{mai08}.
+. We observed the CDES with the NRAO Very. Large Array (VLA) for 50 hours at 1.1 CGIIz mostly in the Bud configurations mn October 1999 and February 2001. and for 32 hours at5 GIIz mostly in the € and CuB configurations between June aud October 2001 (detailsaregiveninWellermanunctal.2008.Paper D.," We observed the CDFS with the NRAO Very Large Array (VLA) for 50 hours at 1.4 GHz mostly in the BnA configurations in October 1999 and February 2001, and for 32 hours at 5 GHz mostly in the C and CnB configurations between June and October 2001 \citep[details are given in][Paper I]{kel08}."
+. The effective aueular resolution was 3.5” and the miuinuun root-uecan-square (rius) noise was as low as 8.5 qr]. per beau at 1.1 GIIz and7 pJy per beam at5 NGIIz., The effective angular resolution was $3.5''$ and the minimum root-mean-square (rms) noise was as low as $8.5~\mu$ Jy per beam at 1.4 GHz and $7~\mu$ Jy per beam at 5 GHz.
+" A total of 266 radio sources were catalogued at CGIIz. 198 of which are in a complete sample with signal-to-noise ratio (SNR) ercater than 5 audlocated within 15/ of the field center,"," A total of 266 radio sources were catalogued at 1.4 GHz, 198 of which are in a complete sample with signal-to-noise ratio (SNR) greater than 5 andlocated within $^{\prime}$ of the field center."
+ The corresponding flux deusitv ranges frou 12 Jv at the field center to 125 µ.]ν near the field cage., The corresponding flux density ranges from 42 $\mu$ Jy at the field center to 125 $\mu$ Jy near the field edge.
+ These deep radio observations complement the lareer area. but less sensitive. lower resolution observatious of the CDFS discussed by Afousoetal.(2006) aud by Norrisetditional.(2006).," These deep radio observations complement the larger area, but less sensitive, lower resolution observations of the CDFS discussed by \cite{afo06} and by \cite{nor06}."
+. In a to our deep radio data. our set of ancillary data is quite unique aud allows us to πια iw light on the nature of the subauJv radio source »»pulation.," In addition to our deep radio data, our set of ancillary data is quite unique and allows us to shed new light on the nature of the sub-mJy radio source population."
+ These data iuclude reliable optical/near-IR identifications for 9154 of the radio sources. optica uorphological classification for 6154 of the sample. redshift iuforuatiou for 7 of the objects.with and N-rav etectious for 33%VK of the sources. upper limits or all but three of the others.," These data include reliable optical/near-IR identifications for $94\%$ of the radio sources, optical morphological classification for $61\%$ of the sample, redshift information for $73\%$ of the objects, and X-ray detections for $33\%$ of the sources, with upper limits for all but three of the others."
+ This is the largest aud nost complete sample of piJy sources iu terms of redshift information., This is the largest and most complete sample of $\mu$ Jy sources in terms of redshift information.
+ The purpose of this paper is to study the composition of the racdio-detected VLA-CDFS sample in terms of source properties by using all available information., The purpose of this paper is to study the composition of the radio-detected VLA-CDFS sample in terms of source properties by using all available information.
+" More etails on the VLA-CDFS can be found in Paper 1. which addresses the radio data. in Paper IT. which discusses the optical aud. near IR counterparts to the observed radio sources, and in Tozzietal.(2008) (Paper IT). which studies the X-ray properties."," More details on the VLA-CDFS can be found in Paper I, which addresses the radio data, in Paper II, which discusses the optical and near IR counterparts to the observed radio sources, and in \cite{toz08} (Paper III), which studies the X-ray properties."
+ The sample radio-optical data are described iu 2. while 3 deals with the XN-rav data.," The sample radio-optical data are described in 2, while 3 deals with the X-ray data."
+ | analyzes the overall properties of our sample and derives the radio number counts for various classes together with sole inodoel predictions. while 5 discusses our results.," 4 analyzes the overall properties of our sample and derives the radio number counts for various classes together with some model predictions, while 5 discusses our results."
+ Finally. 6 summnarizes our conclusions.," Finally, 6 summarizes our conclusions."
+" Throughout this paper spectral indices are defined by 5S,x»'"" aud the values Hy=70 lau | ο, Q4=0.3. and Q4=0.7 have been used."," Throughout this paper spectral indices are defined by $S_{\nu} \propto \nu^{-\alpha}$ and the values $H_0 = 70$ km $^{-1}$ $^{-1}$, $\Omega_{\rm M} = 0.3$, and $\Omega_{\rm \Lambda} =
+0.7$ have been used."
+ This paper supersedes the proliniuaryv results presented by Padovanietal.(2007))., This paper supersedes the preliminary results presented by \cite{pad07b}.
+. Before we turn to the racio sources. we first define the vpical ranges iu radio poses and radio-to-optical flux density ratio for the typeο of objects which are most likely o be part of our survey: s (normal. radio. and star ormüueg). and active galactic uuclei (AGN).," Before we turn to the radio sources, we first define the typical ranges in radio power and radio-to-optical flux density ratio for the type of objects which are most likely to be part of our survey: galaxies (normal, radio, and star forming), and active galactic nuclei (AGN)."
+" Normal galaxies. with radio powers <εν100"" W +. are thought to have their radio cussion dominated wosvuchrotron radiation frou interstellar relativistic clectrous (Philipsetal.1986:Sadler1950)."," Normal galaxies, with radio powers $\la 4 \times10^{20}$ W $^{-1}$, are thought to have their radio emission dominated by synchrotron radiation from interstellar relativistic electrons \citep{phi86,sad89}."
+ Radio galaxies are cousidered to lave gigahertz radio over zo1072 W ! (eg.Sadleretal.1989:Led-ow&Owen1996). associated with relativistic jets. siloparsec scale morphology. which extends well bevoud he region of optical huninositw. and. at least at low redshift. imostlv elliptical hosts.," Radio galaxies are considered to have gigahertz radio power $\ga 10^{22}$ W $^{-1}$ \citep[e.g.,][]{sad89,led96}, associated with relativistic jets, kiloparsec scale morphology, which extends well beyond the region of optical luminosity, and, at least at low redshift, mostly elliptical hosts."
+" This radio power is at he faint eud of the luminosity function of radio galaxies (e.g...Urry&Padovani1995). aud. therefore. represeuts a ""natural threshold for radio-oudness"" iu galaxies. although recent papers have suggested the presence of nou-therimal. njet cussion in carly type galaxies as faint as 1079 WV Ἐν se. Balmaverde&Capetti(2006)."," This radio power is at the faint end of the luminosity function of radio galaxies \citep[e.g.,][]{up95} and, therefore, represents a `natural' threshold for ""radio-loudness"" in galaxies, although recent papers have suggested the presence of non-thermal, jet emission in early type galaxies as faint as $10^{20}$ W $^{-1}$: e.g., \cite{bal06}."
+ Fanarott&Riley(1971). recognized that radio galaxies separate iuto two distiuct luminosity classes. cach with its own characteristic radio morphology.," \cite{fan74}
+ recognized that radio galaxies separate into two distinct luminosity classes, each with its own characteristic radio morphology."
+ Iligbh-Iuninosity Fanaroff-Rilev (FR) IIs have racio lobes with pronineut hot spots aud bright outer edges. while in low-huuinosity FR Is the radio emission is more diffuse.," High-luminosity Fanaroff-Riley (FR) IIs have radio lobes with prominent hot spots and bright outer edges, while in low-luminosity FR Is the radio emission is more diffuse."
+ The Iuuinositv distinction is fairly sharp at 175 AUIz. with FR Is and FR Os lying below and above. respectively. the fiducial luminosity -ϱμπα W !.," The luminosity distinction is fairly sharp at 178 MHz, with FR Is and FR IIs lying below and above, respectively, the fiducial luminosity $P_{\rm 178 MHz} \approx 10^{25}/(H_0/70)^2$ W $^{-1}$."
+" At higher radio frequencies the division is at PiΠο],ή24.orlAH/70)?. W ft. with some depeudeuey also on optical huuinositv (Owen&White1991). aud therefore a somewhat large overlap."," At higher radio frequencies the division is at $P_{\rm 1.4 GHz} \approx 2 \times 10^{24}/(H_0/70)^2$ W $^{-1}$, with some dependency also on optical luminosity \citep{ow91} and therefore a somewhat large overlap."
+ Finally. SFC cau also be relatively strong radio ciuitters but are hosted by spiral and inregular galaxies at the relative low redshifts of our sample (see ο).," Finally, SFG can also be relatively strong radio emitters but are hosted by spiral and irregular galaxies at the relative low redshifts of our sample (see \ref{CDFS-VLA}) )."
+" These star formine radio sources dominate the local (2< 0.1) radio buuinositv function below Pywp,c1077 W | but reach only Prien,c104 W | (egtoSadleretal.2002:Bestetal. 2005).. as compared Price,m100 W t or nore for the powerful radio-galaxies."," These star forming radio sources dominate the local $z < 0.4$ ) radio luminosity function below $P_{\rm 1.4 GHz} \approx 10^{23}$ W $^{-1}$ but reach only $P_{\rm
+ 1.4 GHz} \approx 10^{24}$ W $^{-1}$ \citep[e.g.,][]{sad02,be05}, as compared to $P_{\rm 1.4 GHz} \approx 10^{26}$ W $^{-1}$ or more for the powerful radio-galaxies."
+ In terms of the rest-frame ratio of radio to optical flux density. which we define as R=loe(SyοSy) (where Sy is the V-baud flux deusitv). SEG. have relatively small values.," In terms of the rest-frame ratio of radio to optical flux density, which we define as $R = \log (S_{\rm 1.4GHz}/S_{\rm
+ V})$ (where $S_{\rm V}$ is the V-band flux density), SFG have relatively small values."
+ For example. the majority of the ultra-Inninous infrared galudes in Vegaetal.(2008) have Syjn7Sian: which corresponds to Rw 0.1. while Machalski&Con- have found that of their normal aud starburst galaxies have Ro<1.3 (converting from their definition).," For example, the majority of the ultra-luminous infrared galaxies in \cite{veg08} have $S_{\rm 1.4GHz} \approx S_{\rm 1 \mu m}$, which corresponds to $R \approx 0.4$ , while \cite{mac99} have found that of their normal and starburst galaxies have $R < 1.3$ (converting from their definition)."
+ It is important to note that these two studies reach ouly 2—0.1. and ~0.2 respectively. aud therefore oue cannot exclude the possibility that higher redshift sources iueht have larger R values (see 773).," It is important to note that these two studies reach only $z \sim 0.1$ and $\sim 0.2$ respectively, and therefore one cannot exclude the possibility that higher redshift sources might have larger $R$ values (see \ref{CDFS-VLA}) )."
+ We assune for the tine being R=1.1 as the maxima fiducial value for SFC., We assume for the time being $R = 1.4$ as the maximum fiducial value for SFG.
+ We note that, We note that
+sources have been included so that all 55 stars have measurements of both the permitted and forbidden oxygen lines.,sources have been included so that all 55 stars have measurements of both the permitted and forbidden oxygen lines.
+ The data set presented here provides a strong test of any attempts to reconcile the indicators., The data set presented here provides a strong test of any attempts to reconcile the indicators.
+ We will begin our analysis by measuring the magnitude of the difference in these two oxygen abundance indicators when we adopt atmosphere parameters with from colors and from isochrones., We will begin our analysis by measuring the magnitude of the difference in these two oxygen abundance indicators when we adopt atmosphere parameters with from colors and from isochrones.
+ We find that the familiar pattern of O I lines giving higher O abundances than the [O I] lines reasserts itself., We find that the familiar pattern of O I lines giving higher O abundances than the [O I] lines reasserts itself.
+ Next. we examine whether changing the assumptions of the analysis. in particular the temperature scale. eliminates the measured difference.," Next, we examine whether changing the assumptions of the analysis, in particular the temperature scale, eliminates the measured difference."
+ We then use the knowledge of the behavior of the lines to create an ad hoc parameter set that. within the assumptions of the analysis. results in agreement between the indicators.," We then use the knowledge of the behavior of the lines to create an ad hoc parameter set that, within the assumptions of the analysis, results in agreement between the indicators."
+ Finally. we discuss whether this ad hoc parameter scale Is realistic when compared to other observables for the target stars.," Finally, we discuss whether this ad hoc parameter scale is realistic when compared to other observables for the target stars."
+ An exhaustive study of all the possible solutions to the oxygen abundance problem is beyond the scope of a single paper., An exhaustive study of all the possible solutions to the oxygen abundance problem is beyond the scope of a single paper.
+ We therefore concentrate on following up the apparent successes of parameter-based solutions in the recent works mentioned in the Introduction., We therefore concentrate on following up the apparent successes of parameter-based solutions in the recent works mentioned in the Introduction.
+ Our analysis follows the following assumptions: 1) The atmospheres of stars can be described by one-dimensional. plane-parallel models in LTE.," Our analysis follows the following assumptions: 1) The atmospheres of stars can be described by one-dimensional, plane-parallel models in LTE."
+ For most of this work. we will use Kurucz (1995)models?.," For most of this work, we will use Kurucz (1995)."
+. We use the MOOG stellar abundance package (Sneden1973) for the analysis., We use the MOOG stellar abundance package \citep{s73} for the analysis.
+ We adopt logn(Fe)..=7.52 and logn(O)..=8.69., We adopt $\log{\rm n(Fe)_\odot} = 7.52$ and $\log{\rm n(O)_\odot} = 8.69$.
+ The later value is based on the reanalysis of the solar [ο I] by AllendePrietoetal.(2001)... which takes in account the contamination of the 6300 [O I] line by a Ni I weak line.," The later value is based on the reanalysis of the solar [O I] by \citet{ap01}, which takes in account the contamination of the 6300 [O I] line by a Ni I weak line."
+ 2) Non-LTE effects limit the usefulness of Fe I lines (Thévenin&Idiart1999) in metal-poor stars.," 2) Non-LTE effects limit the usefulness of Fe I lines \citep{ti99}
+ in metal-poor stars."
+ Non-LTE conditions also affect the permitted O I lines. but for the purposes of this experiment we will assume the abundance corrections of Takedaetal.(2000) adequately compensate for the departures from LTE.," Non-LTE conditions also affect the permitted O I lines, but for the purposes of this experiment we will assume the abundance corrections of \citet{nlte} adequately compensate for the departures from LTE."
+ We also assume the lines of Fe II are free of non-LTE effects and will be used as the primary Fe abundance indicator., We also assume the lines of Fe II are free of non-LTE effects and will be used as the primary Fe abundance indicator.
+ 3) Within the assumptions above. we will assume the solution to the problem can be found by the application of the correct atmospheric parameters for the stars.," 3) Within the assumptions above, we will assume the solution to the problem can be found by the application of the correct atmospheric parameters for the stars."
+ This assumption ts similar to the solution put forth by King(2000)., This assumption is similar to the solution put forth by \citet{k00}.
+. The methods employed here are similar to those taken by Nissenetal.(2002) and King(2000).. but their samples contain only warm 5600 K) turn-off and subgiant stars. and only have 11 stars between them with both forbidden and permitted lines.," The methods employed here are similar to those taken by \citet{n02} and \citet{k00}, but their samples contain only warm $> 5600$ K) turn-off and subgiant stars, and only have 11 stars between them with both forbidden and permitted lines."
+ The Gecko echelle spectrograph on the Canada-France-Hawait Telescope (CFHT) delivers spectra with resolution of ~130000. with a dispersion of 0.018 per 13.5 jm pixel.," The Gecko echelle spectrograph on the Canada-France-Hawaii Telescope (CFHT) delivers spectra with resolution of $\sim 130000$, with a dispersion of $0.018$ per 13.5 $\mu$ m pixel."
+ Only one spectral order is observable at a time (selected by a narrow-band filter). meaning only a —75 region is observed per exposure with the thinned 2048x4096 pixel detector.," Only one spectral order is observable at a time (selected by a narrow-band filter), meaning only a $\sim 75$ region is observed per exposure with the thinned 2048x4096 pixel detector."
+ The Gecko data were obtained over 6 nights in April and September 2001., The Gecko data were obtained over 6 nights in April and September 2001.
+ In both runs. the spectra covered the wavelength range from 6290-6370A.. covering both the 6300 and 6363 lines.," In both runs, the spectra covered the wavelength range from 6290–6370, covering both the 6300 and 6363 lines."
+ The primary candidate list was created in a similar manner as the target list of Fulbright(2000)... using literature lists of known metal-poor stars. such as Bond(1980).. Carneyetal. (1994).. ete.," The primary candidate list was created in a similar manner as the target list of \citet{f00}, using literature lists of known metal-poor stars, such as \citet{bond}, \citet{c94}, etc."
+ The list was then culled of stars where we estimated that one of the two sets of lines would be undetectable., The list was then culled of stars where we estimated that one of the two sets of lines would be undetectable.
+ The spectra were reduced using normal routines., The spectra were reduced using normal routines.
+ Although previous work has shown that the scattered light effect in Gecko is less than156.. special care was taken in its removal.," Although previous work has shown that the scattered light effect in Gecko is less than, special care was taken in its removal."
+ A wavelength solution was applied using ThAr lamps taken at the beginning and end of the night., A wavelength solution was applied using ThAr lamps taken at the beginning and end of the night.
+ The variation in the wavelength of the telluric features between spectra is less than 100 m s7!., The variation in the wavelength of the telluric features between spectra is less than 100 m $^{-1}$.
+ Details of the individual observations are given in Table 1., Details of the individual observations are given in Table 1.
+ Because the 6300 feature lies within a band of telluric lines. it was sometimes necessary to remove these telluric lines from the spectrum.," Because the 6300 feature lies within a band of telluric lines, it was sometimes necessary to remove these telluric lines from the spectrum."
+ During the observation runs. spectra of bright. rapidly-rotating (vsini >250 km s! ). spectral type B or A stars were observed.," During the observation runs, spectra of bright, rapidly-rotating (vsini $> 250$ km $^{-1}$ ), spectral type B or A stars were observed."
+ These spectra were used to divide out the telluric features (some sample spectra are shown in the Appendix)., These spectra were used to divide out the telluric features (some sample spectra are shown in the Appendix).
+ For most stars the division was cosmetic. because the [O I] line was not contaminated and the very high resolution and dispersion of the Gecko spectrograph lessens the probability of contamination.," For most stars the division was cosmetic, because the [O I] line was not contaminated and the very high resolution and dispersion of the Gecko spectrograph lessens the probability of contamination."
+ Langer(1991) suggested circumstellar [O I] emission may play a role in the diserepancy., \citet{l91} suggested circumstellar [O I] emission may play a role in the discrepancy.
+ We did not observe any sign of stellar [O I] emission in our spectra. a point which we examine this point further in the Appendix.," We did not observe any sign of stellar [O I] emission in our spectra, a point which we examine this point further in the Appendix."
+ Additional data to measure the Fe and O I lines were obtained with the Lick 3-m and KPNO 4-m with their respective echelle spectrographs., Additional data to measure the Fe and O I lines were obtained with the Lick 3-m and KPNO 4-m with their respective echelle spectrographs.
+ The new data taken with the Hamilton spectrograph at Lick were obtained in the same way as the previous data (see Fulbright2000. and Johnson2002. for more details)., The new data taken with the Hamilton spectrograph at Lick were obtained in the same way as the previous data (see \citealt{f00} and \citealt{j02} for more details).
+ The echelle data from ΚΕΝΟ were obtained in January 2002 with a resolution of ~40000 and covers the wavelength range from 4480 to 7850A., The echelle data from KPNO were obtained in January 2002 with a resolution of $\sim 40000$ and covers the wavelength range from 4480 to 7850.
+. Included in the KPNO data is a spectrum of the asteroid Vesta. which provides a solar spectrum taken as if the Sun was a point source.," Included in the KPNO data is a spectrum of the asteroid Vesta, which provides a solar spectrum taken as if the Sun was a point source."
+ The equivalent widths (EW) values of the 6300 and 6363 [O I] lines were measured using both Gaussian fits and integrations., The equivalent widths (EW) values of the 6300 and 6363 [O I] lines were measured using both Gaussian fits and integrations.
+ The EW of the 6300 line for BD +23 3130 was adopted from Cayreletal.(2001)., The EW of the 6300 line for BD +23 3130 was adopted from \citet{c01}.
+. The EW values for the permitted O | lines were measured in the non-Gecko data or taken from literature sources or à combination of both., The EW values for the permitted O I lines were measured in the non-Gecko data or taken from literature sources or a combination of both.
+ Table 2 gives the oxygen EW values for, Table 2 gives the oxygen EW values for
+ 1995)). 1996)). 1994). 0«z| zy~1 1999)), \markcite{tri95}) \markcite{ba96}) \markcite{wo94}) \markcite{sed98} $0<z<1$ $z_F\sim 1$ \markcite{fe99})
+ 1995)). 1996)). 1994). 0«z| zy~1 1999))., \markcite{tri95}) \markcite{ba96}) \markcite{wo94}) \markcite{sed98} $0<z<1$ $z_F\sim 1$ \markcite{fe99})
+cosmic variance) of the simulated mock samples (see Pollo et al.,cosmic variance) of the simulated mock samples (see Pollo et al.
+ 2006)., 2006).
+ The analvtical model is based on the halo model (see Cooray Sheth 2002. for a review). here we will brielly mention the main ingredients.," The analytical model is based on the halo model (see Cooray Sheth 2002, for a review), here we will briefly mention the main ingredients."
+ ALL mass is assumed to be bound up into dark matter halos having a range of masses which in turn lost galaxies., All mass is assumed to be bound up into dark matter halos having a range of masses which in turn host galaxies.
+ In this model the power spectrum. Pik). and/or correlation function. η). of the galaxies (which are fourier ransform pairs) can be written as the sum of two ternis.," In this model the power spectrum, P(k), and/or correlation function, $\xi(r)$, of the galaxies (which are fourier transform pairs) can be written as the sum of two terms."
+ One that dominates on non-linear scales -smaller than the size of a halo. and the other term becoming significant on arger linear scales. known as the 1-halo ancl 2-halo terms respectively.," One that dominates on non-linear scales -smaller than the size of a halo, and the other term becoming significant on larger linear scales, known as the 1-halo and 2-halo terms respectively."
+ The 1 halo termi arises from pairs of galaxies ving within the same halo. whereas pairs of galaxies Iving in dillerent halos contribute to the 2 halo term.," The 1 halo term arises from pairs of galaxies lying within the same halo, whereas pairs of galaxies lying in different halos contribute to the 2 halo term."
+" In fourier space this can be written as. where. and yer ds the number density of galaxies: ]lere. GN,|i) is the average number of galaxies occupying a halo of mass m. u(K|m) is the halo density profile in fourier space. n(m.z) is the number density. of halos of mass m. b(m.z) is the bias factor which describes the strength of halo clustering. ancl £7;,,(4.2) is the power spectrum of the mass in linear theory all at a given redshift 2."," In fourier space this can be written as, where, and $\bar n_{gal}$ is the number density of galaxies: Here, $\langle N_{g}|m\rangle$ is the average number of galaxies occupying a halo of mass $m$, $u(k|m)$ is the halo density profile in fourier space, $n(m,z)$ is the number density of halos of mass $m$, $b(m,z)$ is the bias factor which describes the strength of halo clustering, and $P_{lin}(k,z)$ is the power spectrum of the mass in linear theory all at a given redshift $z$."
+ Phe upper limit of integration. ζω approximately accounts for the halo exclusion ellect (dilferent halos cannot overlap) by suppressing the 2-halo term at small scales.," The upper limit of integration, $M_{max}$, approximately accounts for the halo exclusion effect (different halos cannot overlap) by suppressing the 2-halo term at small scales."
+ Following Zehavi ct al. (, Following Zehavi et al. (
+2004). Ανν is the mass of the ido with virial radius 7/2.,"2004), $M_{max}$ is the mass of the halo with virial radius $r/2$."
+ One can also calculate the one- term for the correlation function exactly in real space. which is the approach we have taken.," One can also calculate the one-halo term for the correlation function exactly in real space, which is the approach we have taken."
+ For more details we refer the reader to. Berlind Weinberg (2002)., For more details we refer the reader to Berlind Weinberg (2002).
+ The two-ido term has been computed. in k-space and then fourier ransformec to obtain the correlation function., The two-halo term has been computed in k-space and then fourier transformed to obtain the correlation function.
+ Then the xojected correlation function is obtained as in Equation 2.., Then the projected correlation function is obtained as in Equation \ref{Eq_wprp}.
+" Similarly to the data the upper limit is chosen to be inito (7,,,,2:20 Mpc/h) in order to avoid noise caused. by uncorrelatecl distant. pairs.", Similarly to the data the upper limit is chosen to be finite $r_{max}$ =20 $h$ ) in order to avoid noise caused by uncorrelated distant pairs.
+ We assume that the density. profiles of halos have he form described by Navarro ct al. (, We assume that the density profiles of halos have the form described by Navarro et al. (
+L997). with a halo concentration parameter e(M)=LIAL/AL.)OLS to account or the definition of halos as spheres enclosing 200. times he background density (Zehavietal.2005) and where GM.2)=ὃν (0 and ὃς. ave defined. below),"1997), with a halo concentration parameter $c(M) = 11(M/M_*)^{-0.13}$ to account for the definition of halos as spheres enclosing 200 times the background density \citep{zeh05} and where $\sigma(M_*,z) = 
+ \delta_{sc}$ $\sigma$ and $\delta_{sc}$ are defined below)."
+ The halo abundances and clustering are. described by the Sheth 'Tormen (1999) parameterization:, The halo abundances and clustering are described by the Sheth Tormen (1999) parameterization:
+"where X=(xx,JA,/Ar is the image position relative to the ceuter of the leus.","where ${\bf \tilde{x}} \equiv({\bf x} - {\bf x}_o) \lambda_o/
+\lambda_E$ is the image position relative to the center of the lens."
+ Also useful is the convergence or dimensiouless surface deusitv at the Einstein racius iu these powerlaw model: &(Ag)=1/2., Also useful is the convergence or dimensionless surface density at the Einstein radius in these power–law model: $\kappa(\lambda_E)=n/2$.
+ For a Singular Isothermal Sphere (SIS) leus η=1 aud For a point mass 7=0 and Ag=vn (7)., For a Singular Isothermal Sphere (SIS) lens $n=1$ and For a point mass $n=0$ and $\lambda_E=\sqrt{m/(\pi\Sigma_c)}$ .
+ Asan exaniple. fieure 1. shows the curvature magnification factor. C. for a SIS lens with no substructure.," As an example, figure \ref{fig:curv_mag} shows the curvature magnification factor, $C$, for a SIS lens with no substructure."
+ This quantity depends on both the position of the source and the tangent vector to one of the images (iu this case the outer image is chosen)., This quantity depends on both the position of the source and the tangent vector to one of the images (in this case the outer image is chosen).
+ The factor is ecnerally larger for jets that are radial (3i which case the outer image is more bent) or tangeutial Qvlere the opposite is tue)., The factor is generally larger for jets that are radial (in which case the outer image is more bent) or tangential (where the opposite is true).
+ Cases with C much different from, Cases with $C$ much different from
+ (e.g..Gültekinetal.2009).," \citep[e.g.,][]{silkrees1998,hopkinsetal2006}. \citep[e.g.,][]{gultekinetal2009}."
+. (Greene&Ho2007:SchulzeWisotzki2010) (e.g..Antonucci&Cohen1983: (see ," \citep{greeneho2007b,schulzewisotzki2010} \citep{wooetal2006,kollmeieretal2006,shenetal2008b,wooetal2008,
+vestergaardosmer2009,kellyetal2009}, \citep{blandfordmckee1982}. \citep[see
+the recent compilation by ][]{petersonetal2004}. \citep[e.g.,][]{antonuccicohen1983,petersonetal1983,ulrichetal1984,
+ gaskellsparke1986}, \citep[see][]{petersonetal2004}."
+LH//Ru , $1/\sqrt{R}$ 
+Collimated extragalactic [lows. exhibiting a complex pattern of internal siructures. often with high brightness temperature and/or superlununal speed. have been explored with increasing spatial and temporal resolution over the last two decades (Zensus1997).,"Collimated extragalactic flows, exhibiting a complex pattern of internal structures, often with high brightness temperature and/or superluminal speed, have been explored with increasing spatial and temporal resolution over the last two decades \citep{zen}."
+. Remarkably. it has become evident over the last [ive vears that the highly energetic [lows associated with active galactic nuclei are also found in Galactic objects with stellar mass 'engines.," Remarkably, it has become evident over the last five years that the highly energetic flows associated with active galactic nuclei are also found in Galactic objects with stellar mass `engines'."
+ In the galactic superluminals GRS 1915+105 (MirabelandRodrienez1994.1995) and GRO J165540 C]jellmineandRupen1995:Tingayοἱal.1995). the observed motions indicate jet flow speeds up lo 0.98c.," In the galactic superluminals GRS 1915+105 \citep{mr1,mr2} and GRO J1655–40 \citep{hje,tin} the observed motions indicate jet flow speeds up to $0.98c$."
+ Further. (here is compelling evidence from the observation ol optical afterglow that eamnma-ray bursts are of cosmological origin. ancl whether produced bv the mergers of compact objects. or accretion induced collapse (AIC). simple relativistic fireball models seem ruled out. the data strongly favoring highly relativistic jets (Dar1993)..," Further, there is compelling evidence from the observation of optical afterglow that gamma-ray bursts are of cosmological origin, and whether produced by the mergers of compact objects, or accretion induced collapse (AIC), simple relativistic fireball models seem ruled out, the data strongly favoring highly relativistic jets \citep{dar}."
+ Thus a detailed understanding of the dynamics of collimated relativistic [lows has wide application in astrophysics., Thus a detailed understanding of the dynamics of collimated relativistic flows has wide application in astrophysics.
+ A particularly important facet of these flows has been revealed over the last decade as the quantity and. quality of images of extragalactic relativistic jets have increased: the curvature of parsec-scale flows is (he norm. not the exception: see e.g.. Wardleetal. (1994)..," A particularly important facet of these flows has been revealed over the last decade as the quantity and quality of images of extragalactic relativistic jets have increased: the curvature of parsec-scale flows is the norm, not the exception; see e.g., \citet{war}. ."
+observations.,observations.
+ Using the minimum set (2 HWP angles) leads to large errors. especially in the case of low polarizations (~ 156)). and it is therefore strongly discouraged.," Using the minimum set (2 HWP angles) leads to large errors, especially in the case of low polarizations $\sim$ ), and it is therefore strongly discouraged."
+ The Fourier analysis shows that all harmonies with &40.4 are negligible in the wavelength range 3800-7400A.. where a rms accuracy of can be reached with a sufficient signal-to-noise.," The Fourier analysis shows that all harmonics with $k\neq$ 0,4 are negligible in the wavelength range 3800–7400, where a rms accuracy of can be reached with a sufficient signal-to-noise."
+ This can be considered as the instrumental limit attainable with CAFOS with + HWP angles. and within this spectral range.," This can be considered as the instrumental limit attainable with CAFOS with 4 HWP angles, and within this spectral range."
+ Below 3800 the polarimetric properties rapidly degrade. requiring a larger number of HWP angles.," Below 3800 the polarimetric properties rapidly degrade, requiring a larger number of HWP angles."
+ The same applies. though to a smaller extent. to the region redwards of 8200A.," The same applies, though to a smaller extent, to the region redwards of 8200."
+. For this work we have used data obtained with the B200 erism., For this work we have used data obtained with the B200 grism.
+ Because of its tilted surfaces. the grism can act as a poor linear polarizer.," Because of its tilted surfaces, the grism can act as a poor linear polarizer."
+ Since in CAFOS this 15 placed after the analyzer. the spurious polarization produced by transmission is not modulated by the HWP rotation. and hence the redundancy in the retarder-plate position effectively removes it (see Patat Romaniello 2006)).," Since in CAFOS this is placed after the analyzer, the spurious polarization produced by transmission is not modulated by the HWP rotation, and hence the redundancy in the retarder-plate position effectively removes it (see Patat Romaniello \cite{patat06}) )."
+ The exact effect produced by the grism depends on its properties., The exact effect produced by the grism depends on its properties.
+ However. the conclusions reached in this paper do not depend on the grism. provided that the data are obtained using at least 4 HWP angles.," However, the conclusions reached in this paper do not depend on the grism, provided that the data are obtained using at least 4 HWP angles."
+ In general. CAFOS appears to be perfectly suitable. for linear polarization studies aiming at accuracies of a fewO.1%.. making it a valid instrument for bright objects.," In general, CAFOS appears to be perfectly suitable for linear polarization studies aiming at accuracies of a few, making it a valid instrument for bright objects."
+ As a term of reference. an accuracy of «0.15€ per resolution element (~50 A)) was reached for SN 2010j1 (V.~ 13.5). with 4 exposures of 40 minutes each (Patat et al. 2010)).," As a term of reference, an accuracy of $\sim$ per resolution element $\sim$ 50 ) was reached for SN 2010jl $V\sim$ 13.5), with 4 exposures of 40 minutes each (Patat et al. \cite{patat10}) )."
+We note that assigning errors to the folded and averaged fluxes is not a unique procedure.,We note that assigning errors to the folded and averaged fluxes is not a unique procedure.
+ Here. we first divide the light curve into time bins of the 1/20 of the period torbital or superorbital) and obtain the average of the flux. 75;. where { is the number of he phase bin and j is the number of the time bin contributing o the ;-th phase bin. weighted by the inverse squares of their measurement error.," Here, we first divide the light curve into time bins of the 1/20 of the period (orbital or superorbital) and obtain the average of the flux, $F_{ij}$, where $i$ is the number of the phase bin and $j$ is the number of the time bin contributing to the $i$ -th phase bin, weighted by the inverse squares of their measurement error."
+ In. this way. we avoid any contribution to our folded/averaged light curves from the source variability on ime scales shorter than that corresponding to the length of our chosen phase bin.," In this way, we avoid any contribution to our folded/averaged light curves from the source variability on time scales shorter than that corresponding to the length of our chosen phase bin."
+ The resulting measurement uncertainties on fy; are rather small. and we neglect them.," The resulting measurement uncertainties on $F_{ij}$ are rather small, and we neglect them."
+ Then. we consider he average and standard deviation of all time bins contributing ο a given phase bin.," Then, we consider the average and standard deviation of all time bins contributing to a given phase bin."
+ Here. we want to take into account the actual aperiodic variabilityof the source. not just the measurement errors.," Here, we want to take into account the actual aperiodic variability of the source, not just the measurement errors."
+" Therefore. we calculate the unweighted average within he 7-th phase bin. ;=(Xp)N;. where JN; Is the number of time bins contributing to the j-th phase bin. and he corresponding rms standard deviation. σε, given by a?=YunFl ΟΕ "," Therefore, we calculate the unweighted average within the $i$ -th phase bin, $F_i = \left(\sum_{j=1}^{N_i} F_{ij}\right)/N_i$, where $N_i$ is the number of time bins contributing to the $i$ -th phase bin, and the corresponding rms standard deviation, $\sigma_i$, given by $\sigma_i^2= \left[\sum_{j=1}^{N_i} (F_i-F_{ij})^2\right]/ [N_i(N_i-1)]$ ."
+Then. we take €;=Inby and fit momwith equation. (2)).," Then, we take $G_i=\ln F_i$ and fit with equation \ref{eq:harmonics}) )."
+ Subsequently. we prewhiten the logarithmic light curves by subtracting from them the orbital variability fitted above. We now again apply the Lomb-Scargle method to the prewhitened data. (," Subsequently, we prewhiten the logarithmic light curves by subtracting from them the orbital variability fitted above, We now again apply the Lomb-Scargle method to the prewhitened data. ("
+We do not prewhiten the BATSE data. as their orbital modulation is very small. see above.),"We do not prewhiten the BATSE data, as their orbital modulation is very small, see above.)"
+obtained main (150-4) superorbital periods are presented in TableThe 2.., The obtained main $\sim$ 150-d) superorbital periods are presented in Table \ref{t:ampl}.
+ The given Lo errors have been estimated using the method described in section 4.4 of Schwarzenberg-Czerny(1991)., The given $1\sigma$ errors have been estimated using the method described in section 4.4 of \citet{sc91}.
+.. The superorbital period of Cyg | was determined by LO06 using the weighted average. 151.43 d. of periods determined for a large number of data sets.," The superorbital period of Cyg X-1 was determined by L06 using the weighted average, 151.43 d, of periods determined for a large number of data sets."
+ For our data alone we find 152.4+0.6 d. and including the and5 data from LO. we obtain a similar value. 2?=152.0+ O.4d (lo error). consistent with that of LO6.," For our data alone we find $152.4\pm0.6$ d, and including the and data from L06, we obtain a similar value, $P = 152.0\pm 0.4$ d $1\sigma$ error), consistent with that of L06."
+ Thus. for consistency with LO6. we hereafter use their ephemeris. We then fold and average within each of 20 superorbital phase bins the prewhitened data. and fit them using equation (2)) with N= L.," Thus, for consistency with L06, we hereafter use their ephemeris, We then fold and average within each of 20 superorbital phase bins the prewhitened data, and fit them using equation \ref{eq:harmonics}) ) with $N=1$."
+The errors usedin the fit are calculated in the waydescribed above., The errors used in the fit are calculated in the way described above.
+ The fits give us the amplitude of the variability. which we detine as. where fansand Jin are. respectively. the maximum and minimum fitted fluxes. which ratio is δὲστωςPus and R=(1)D/OA).," The fits give us the amplitude of the variability, which we define as, where $F_{\rm max}$and $F_{\rm min}$ are, respectively, the maximum and minimum fitted fluxes, which ratio is $R\equiv F_{\rm max}/F_{\rm min}$, and $R= (1+A)/(1-A)$."
+ For our harmonic fit with No=1. Pyuusauin=e6xp(6)dC4).," For our harmonic fit with $N=1$, $F_\mathrm{max, min}=\exp(\langle G \rangle \pm G_1)$."
+ The obtained valuesof ;V are listed in Table 2.., The obtained values of $A$ are listed in Table \ref{t:ampl}.
+ We see that the of amplitudesthe superorbital modulation in both the ASM and radio data are almost independent. of the confirming the previous result of Ld.," We see that the amplitudes of the superorbital modulation in both the ASM and radio data are almost independent of the energy, confirming the previous result of L06."
+ On the other hand. we energy.contirm the result of LO6 that the modulation amplitudes in the 20—300 keV energy range are significantly smaller than those in the 1.5-12 keV range.," On the other hand, we confirm the result of L06 that the modulation amplitudes in the 20–300 keV energy range are significantly smaller than those in the 1.5–12 keV range."
+ of Twothe resulting superorbital phase diagrams are shown in reff:sophase.., Two of the resulting superorbital phase diagrams are shown in \\ref{f:sophase}.
+ We point out that the obtained amplitudes are somewhat affected by systematic uncertainties. in particular those related to the selection criteria.," We point out that the obtained amplitudes are somewhat affected by systematic uncertainties, in particular those related to the selection criteria."
+ We have tested that indeed changing the criteria used to choose the titted parts of the light curves results in some changes of the values of 2., We have tested that indeed changing the criteria used to choose the fitted parts of the light curves results in some changes of the values of $A$.
+ To illustrate this effect. Table 2. also gives the values of 21 of LO06. who used both somewhat different selection criteria and different prewhitening method (namely. prewhitening the data with all other significant modulations. not only the orbital one).," To illustrate this effect, Table \ref{t:ampl} also gives the values of $A$ of L06, who used both somewhat different selection criteria and different prewhitening method (namely, prewhitening the data with all other significant modulations, not only the orbital one)."
+ We can see that their values differ from ours by typically 10 per cent. which can be taken as the systematic uncertainty of the values of 21.," We can see that their values differ from ours by typically 10 per cent, which can be taken as the systematic uncertainty of the values of $A$ ."
+ The largest difference is for the Ryle data. which appears to result from not applying," The largest difference is for the Ryle data, which appears to result from not applying"
+The direct. observational signature of the Newtonian virial heorem in elliptical galaxies emerged two decades ago with he discovery of the “Fundamental Plane” (Djorgovski Davis LOST. Dressler et al.,"The direct observational signature of the Newtonian virial theorem in elliptical galaxies emerged two decades ago with the discovery of the “Fundamental Plane"" (Djorgovski Davis 1987, Dressler et al."
+ LOST)., 1987).
+" Assuming that ellipticals are reasonably homologous Newtonian pressure supported systems. then they will comprise à two-parameter set: if uminosityv is proportional to mass then the virial theorem implies that £LxAge where 2,jy is the elfective radius and σ is the line-of-sight velocity dispersion."," Assuming that ellipticals are reasonably homologous Newtonian pressure supported systems, then they will comprise a two-parameter set: if luminosity is proportional to mass then the virial theorem implies that $L\propto \mathit{R_{eff}}{\sigma^2}$ where $\mathit{R_{eff}}$ is the effective radius and $\sigma$ is the line-of-sight velocity dispersion."
+ This is essentially 1ο observed. Fundamental Plane relation., This is essentially the observed Fundamental Plane relation.
+ Although the 'onnection between the Fundamental Plane (hereafter FP) uxd the virial theorem. was originally obscured by the Lact ju the observationally inferred. exponents in the above relation were not. precisely as expected. 1t is now established wt this deviation is primarily due to a svstematie variation in the mass-to-light ratio of ellipticals: M/L appears to gaowly increase with luminosity (see Faber et al.," Although the connection between the Fundamental Plane (hereafter FP) and the virial theorem was originally obscured by the fact that the observationally inferred exponents in the above relation were not precisely as expected, it is now established that this deviation is primarily due to a systematic variation in the mass-to-light ratio of ellipticals: M/L appears to slowly increase with luminosity (see Faber et al."
+ LOST [or an early realization of the relevance of the virial theorem)., 1987 for an early realization of the relevance of the virial theorem).
+ That the FP is the observational manifestation of the virial rweorem: has become indisputable with the determination. via gravitational lensing. of the mass-basec EI the “more Fundamental Plane”(Bolton ct al.," That the FP is the observational manifestation of the virial theorem has become indisputable with the determination, via gravitational lensing, of the mass-based FP, the “more Fundamental Plane""(Bolton et al."
+ 2007)., 2007).
+ Llere surface density. V. replaces surface brightness. /. in the equivalent FP relation. Royx07/1. and thereby removes the ellects of systematic and random variations in M/L. The striking aspect of the traditional FP is the small seatter about the expected relation.," Here surface density, $\Sigma$ , replaces surface brightness, $I$, in the equivalent FP relation, $\mathit{R_{eff}} \propto \sigma^2/I$, and thereby removes the effects of systematic and random variations in M/L. The striking aspect of the traditional FP is the small scatter about the expected relation."
+ Such a precise relation surely reflects the facts that elliptical galaxies coniprise a rather homologous class of objects with a fairly. isotropic velocity distribution (at least in the inner regions) and that the random. Uuctuation in M/L is relatively small., Such a precise relation surely reflects the facts that elliptical galaxies comprise a rather homologous class of objects with a fairly isotropic velocity distribution (at least in the inner regions) and that the random fluctuation in M/L is relatively small.
+ This small Ductuation in dynamical M/L is puzzling given the current view of a galaxy as a luminous barvonic component immersed in a more massive and extensive dark halo composed of non-barvonic particles., This small fluctuation in dynamical M/L is puzzling given the current view of a galaxy as a luminous baryonic component immersed in a more massive and extensive dark halo composed of non-baryonic particles.
+" Lllipticals are presumably assembled over a Hubble timescale via mergers Uwe or Cdrv. ""major or “minor” of smaller mass ageregales: it is not evident that such a fixed. proportion of dark ancl barvon matter within the visible object should survive what must be a rather chaotic process."," Ellipticals are presumably assembled over a Hubble timescale via mergers – “wet"" or “dry"", “major"" or “minor"" – of smaller mass aggregates; it is not evident that such a fixed proportion of dark and baryon matter within the visible object should survive what must be a rather chaotic process."
+" The second. scaling relation for elliptical galaxies was actually discovered long before the FP that is the relation between Luminosity and line-of-sight. velocity dispersion. of the form Lxo6"" where 3xa<5.", The second scaling relation for elliptical galaxies was actually discovered long before the FP – that is the relation between luminosity and line-of-sight velocity dispersion of the form $L\propto \sigma^\alpha$ where $3\le\alpha\le 5$.
+ Phis correlation was [irs deseribed in a qualitative way by Morgan Alavall (1957) who remarked that “kor progressively fainter Virgo cluster ellipticals the spectral lines tend to become narrower. as if there were a line width-absolute magnitude elfect for. the brighter members.”," This correlation was first described in a qualitative way by Morgan Mayall (1957) who remarked that “For progressively fainter Virgo cluster ellipticals the spectral lines tend to become narrower, as if there were a line width-absolute magnitude effect for the brighter members."""
+ This elfect was given its definite. aux quantitive form by Faber Jackson (1976) and has become known as the Faber-Jackson relation (hereafter EJ)., This effect was given its definite and quantitive form by Faber Jackson (1976) and has become known as the Faber-Jackson relation (hereafter FJ).
+ The E. relation implies that ellipticals comprise à one parameter family: given the velocity dispersion. the Luminosity (or luminous mass) follows.," The FJ relation implies that ellipticals comprise a one parameter family: given the velocity dispersion, the luminosity (or luminous mass) follows."
+ Phe observed EJ relation does have a much Larger dispersion about its mean than does the ΕΤ. ancl this has led to the often-heard assertion that the FP has superseded the FJ that the PJ represents a non-orthogonal projection of the FP onto a lower dimensional parameter space (Franx de Zecuw 1991).," The observed FJ relation does have a much larger dispersion about its mean than does the FP, and this has led to the often-heard assertion that the FP has superseded the FJ – that the FJ represents a non-orthogonal projection of the FP onto a lower dimensional parameter space (Franx de Zeeuw 1991)."
+ While it is certainly true that the FR. because of its much smaller scatter. is superior to the FJ as a distance indicator for elliptical galaxies. it is not evident that the ]J is a simple projection of the FP: nothing like Faber-Jackson is required by the virial relation.," While it is certainly true that the FP, because of its much smaller scatter, is superior to the FJ as a distance indicator for elliptical galaxies, it is not evident that the FJ is a simple projection of the FP; nothing like Faber-Jackson is required by the virial relation."
+ EJ. implies that the mass of a pressure-supported. svstem is correlated. to the velocity dispersion more-or-less independently of the size of the object., FJ implies that the mass of a pressure-supported system is correlated to the velocity dispersion more-or-less independently of the size of the object.
+" E have argued. (Sanders 1994) that this same correlation extends to the great. clusters of galaxies in the form of the gasmass - temperature relation (AL,x 77).", I have argued (Sanders 1994) that this same correlation extends to the great clusters of galaxies in the form of the gasmass - temperature relation $M_g\propto T^2$ ).
+distribution.,distribution.
+⋅⋅⋅ The↴ expression⋅ (207--H).HEN indicates that he quantity in pareutheses uM )o caleulated based ou he radial field given by Eq.A CD)., The expression $(2\Omega^2-\kappa^2)_{\sigma}$ indicates that the quantity in parentheses should be calculated based on the radial field given by Eq. \ref{campovero}) ).
+ For a relatively thin disk the quantity (207|Q2 nav be taken to )o nearly incdepeudeut of τν πο that the οταν]ational field obtained by inteerating Eq. (AL))," For a relatively thin disk the quantity $(2\Omega^2-\kappa^2)_{\sigma}+\Omega^2_{ext}$ may be taken to be nearly independent of $z$, so that the gravitational field obtained by integrating Eq. \ref{bahcallext}) )"
+ iucludes a term that is approximately hnear in :., includes a term that is approximately linear in $z$.
+ The result can be inserted iu he right haud side of Eq. (A2)).," The result can be inserted in the right hand side of Eq. \ref{hydro}) ),"
+ which λα...” where we have introduced the integrated density o.=nρε λατ! , which becomes: where we have introduced the integrated density $\sigma_z=2\int_0^z\rho(z')\mbox{d}z'$ .
+"Note that the term Q2,,: is just the vertical conrponeut of the spherical external field. as nuelt have been anticipated."," Note that the term $\Omega_{ext}^2z$ is just the vertical component of the spherical external field, as might have been anticipated."
+" If po=p((0) is the density on the equatorial plane. we candefine two scales; X—V2pre aud J|—ονοπρῃ. aud the natural dimeusiouless parameter A=[207/—42),|O2[7ExC?py. so that Eq. (AD))"," If $\rho_0=\rho(z=0)$ is the density on the equatorial plane, we can define two scales, $\Sigma=\sqrt{2\rho_0/\pi G}c$ and $H=c/\sqrt{2\pi G\rho_0}$, and the natural dimensionless parameter $A=[(2\Omega^2-\kappa^2)_{\sigma}+\Omega^2_{ext}]/4\pi G\rho_0$, so that Eq. \ref{eq:bahcall}) )"
+ becomes: to be solved uuder the couditiousg(0)=0. y/(0)=1: here we iive used the dinieusiouless variables y=a./X and ¢=II.," becomes: to be solved under the conditions$y(0)=0$, $y'(0)=1$; here we have used the dimensionless variables $y=\sigma_z/\Sigma$ and $\zeta=z/H$."
+ For cach value of A one can compute the desired deusitv profile and the associated surface deusitv 6—σι: x)., For each value of $A$ one can compute the desired density profile and the associated surface density $\sigma=\sigma_z(z=\infty)$ .
+ One can then introduce a scalehcieht P such that &=2py/. aud the value of / can be computed directly from the relation f=/Ty{.= x).," One can then introduce a scaleheight $h$ such that $\sigma=2\rho_0h$, and the value of $h$ can be computed directly from the relation $h=Hy(z=\infty)$ ."
+ The standard non sclberavitating model. where the contribution of the ons deusity p το the vertical field is negligible. corresponds to the luit equation eravitationalyf!2Acu/.," The standard non self-gravitating model, where the contribution of the gas density $\rho$ to the vertical gravitational field is negligible, corresponds to the limit equation $y''=-2A\zeta y'$."
+" Usually the analysis is carried out with the additional approximation 4lzO2teayfAxG py.dut this is not needed: by retaining the coutribution (20?H),5 one can thus keep track ft the effects associated. with the rest of he disk mass distribution. which may be significant even where the local disk density py is sunall."," Usually the analysis is carried out with the additional approximation $A\approx \Omega_{ext}^2/4\pi G\rho_0$, but this is not needed; by retaining the contribution $(2\Omega^2-\kappa^2)_{\sigma}$ one can thus keep track of the effects associated with the rest of the disk mass distribution, which may be significant even where the local disk density $\rho_0$ is small."
+" Iu any case. such anit equation viclds the well known Cassia xofile p==pyoxp(jtiDorey?/2hi). with. fzD=TP?2Aοο.82),| OF). here improved with respect to the standard expression /,zze/Oy used in the so-called I&eplerian liit."," In any case, such limit equation yields the well known Gaussian profile $\rho=\rho_0\exp(-z^2/2h_{\star}^2)$, with $h_{\star}^2=H^2/2A=c^2/[(2\Omega^2-\kappa^2)_{\sigma}+\Omega^2_{ext}]$ , here improved with respect to the standard expression $h_{\star}\approx c/\Omega_K$ used in the so-called Keplerian limit."
+" Note that 7, is slightly differcut roni the scaleheight / defined above. since fh=VESEIS"," Note that $h_{\star}$ is slightly different from the scaleheight $h$ defined above, since $h=\sqrt{\pi/2}h_{\star}$."
+ The huit of the homogeneous fully solf-eravitatiug slab correspouds to the equation y/=δή so that he vertical deusity profile is givenby (Spitzer 1912)) p=pysccl?(./h). with bhο πάσι here the scale D is the same as that defined by the relation ¢=2pyh.," The limit of the homogeneous fully self-gravitating slab corresponds to the equation $y''=-2yy'$, so that the vertical density profile is given by (Spitzer \cite{spitzer}) ) $\rho=\rho_0~\mbox{sech}^2(z/h)$, with $h=c^2/\pi G\sigma$ ; here the scale $h$ is the same as that defined by the relation $\sigma=2\rho_0h$."
+ Note hat even when no external (spherical) field is xeseut. the solution for an axisviunetric disk obtained from Eq.(À6)) is not exactly the oue derived iu the homogeneous. eravitating slab. uuless the rotation curve is flat.," Note that even when no external (spherical) field is present, the solution for an axisymmetric disk obtained from \ref{eq:spessadim}) ) is not exactly the one derived in the homogeneous, self-gravitating slab, unless the rotation curve is flat."
+ The density profile associated with Eq.(A6)) is neither Cassia: nor sech2, The density profile associated with \ref{eq:spessadim}) ) is neither Gaussian nor $\mbox{sech}^2$.
+" For practical purposes it mav be convenieut to use a simple interpolation formula for the vertical scale. which ds justified bv the followine description ""biased"" towards the Gaussian Bit."," For practical purposes it may be convenient to use a simple interpolation formula for the vertical scale, which is justified by the following description “biased” towards the Gaussian limit."
+ If we start from Eq.(A2)). naively expand the vertical field as (OD/O2)(2)~@PB/I-7).yr.aud then use Eq.(ÀA1)) to estimate the factor (074011]. y. we find an equation leading ο an unrealistic Gaussian density profile with scaleheieht 5 eiven by: The quantity py contains a dependence on / (for given 0).," If we start from \ref{hydro}) ), naively expand the vertical field as $(\partial\Phi/\partial z)(z)\sim (\partial^2\Phi/\partial z^2)_{z=0}z$,and then use \ref{poisson}) ) to estimate the factor $(\partial^2\Phi/\partial z^2)_{z=0}$ , we find an equation leading to an unrealistic Gaussian density profile with scaleheight $h$ given by: The quantity $\rho_0$ contains a dependence on $h$ (for given $\sigma$ )."
+ This sugeests the use of the following interpolation formula: Note that the exact calculation. from. 7 gives a relation h=(?/zGo)f(CA). with f(A)=i.," This suggests the use of the following interpolation formula: Note that the exact calculation, from $h=Hy(z=\infty)$ gives a relation $h=(c^2/\pi G\sigma)f(A)$, with $f(A)=[y(z=\infty)]^2$."
+ The choice of the factor (1π) iu Eq. (Α5)), The choice of the factor $(4/\pi)$ in Eq. \ref{eq:interpol}) )
+ euarautees the proper limits for A=0 and for A> x., guarantees the proper limits for $A=0$ and for $A\rightarrow\infty$ .
+ The accuracy of he interpolationformula isillustrated iu Fie. (Al))., The accuracy of the interpolationformula isillustrated in Fig. \ref{fig:interpol}) ).
+ Iu turn. since οἱ depends on / and on σ via the quantity py. for the purposes of the present paper it may be couvenicut to reexpress Eq. (AS))," In turn, since $A$ depends on $h$ and on $\sigma$ via the quantity $\rho_0$ for the purposes of the present paper it may be convenient to reexpress Eq. \ref{eq:interpol}) )"
+ as:with Go=QP(207αν]OF|fe? , as:with $a=Q^2[(2\Omega^2-\kappa^2)_{\sigma}+\Omega^2_{ext}]/\kappa^2$ .
+This leadsto Eq. (17)), This leadsto Eq. \ref{improve}) )
+ of the aad text. applicable when the," of the main text, applicable when the"
+evolution of the LGAL during the reionisation era with and without the ionising effect. of quasars.,evolution of the IGM during the reionisation era with and without the ionising effect of quasars.
+ “Phe left-hand: ane central. left-hand. panels. of Figure 5 show multi-phase ionisation maps for simulation S2 at 2=6.6. Y. 74 and SN. with and without the tonising inlluence of quasars respectively.," The left-hand and central left-hand panels of Figure \ref{128results} show multi-phase ionisation maps for simulation S2 at $z = 6.6$, 7, 7.4 and 8, with and without the ionising influence of quasars respectively."
+ Ehe right-hand panel shows the corresponding power spectra and sample variance (faint lower curves)., The right-hand panel shows the corresponding power spectra and sample variance (faint lower curves).
+ lt is evident. both in the slices shown as well as in the corresponding averaged: global ionisation fraction. that the simulated volumes with quasars are more ionised than those without at the same redshift.," It is evident, both in the slices shown as well as in the corresponding averaged global ionisation fraction, that the simulated volumes with quasars are more ionised than those without at the same redshift."
+ The resulting impact. of the quasar-generated: contribution to ionisation on the 21-cm power spectra is to decrease the overall power on all scales., The resulting impact of the quasar-generated contribution to ionisation on the 21-cm power spectra is to decrease the overall power on all scales.
+ The effect is epoch dependent since the accumulated quasar contribution to reionisation is redshift dependent (as shown in Figure 3))., The effect is epoch dependent since the accumulated quasar contribution to reionisation is redshift dependent (as shown in Figure \ref{Qevolution}) ).
+ The modification of the spectra arises due to the redistribution of Huetuation power between wave-niocles. originating from the spatial deformation of the. stellar-generated bubbles by the quasar emission.," The modification of the spectra arises due to the redistribution of fluctuation power between wave-modes, originating from the spatial deformation of the stellar-generated bubbles by the quasar emission."
+ In addition to the modification of power. the addition of quasar ionisations also modifies the slope of the power spectra.," In addition to the modification of power, the addition of quasar ionisations also modifies the slope of the power spectra."
+ This cect is most easily seen by calculating the relative amplitude between the power spectra without and with quasears. The resulting relative power for simulations S1. 52 and 83 are shown in the left-hand. panels of Figures 6. and. 7..," This effect is most easily seen by calculating the relative amplitude between the power spectra without and with quasars, The resulting relative power for simulations S1, S2 and S3 are shown in the left-hand panels of Figures \ref{comp_comp_10S_z_xHI_128} and \ref{comp_comp_k}."
+ Figure 6 shows the relative power at different redshifts as à function of wave-number., Figure \ref{comp_comp_10S_z_xHI_128} shows the relative power at different redshifts as a function of wave-number.
+ Figure 7. shows cuts of the data as a Function of redshift at three dilferent scales corresponding to wave-numbers &zz0.1. 0.4 and MÀMpe:," Figure \ref{comp_comp_k} shows cuts of the data as a function of redshift at three different scales corresponding to wave-numbers $k \approx 0.1$, 0.4 and Mpc."
+ The dependence on scale is cliscussed further in Section 6.., The dependence on scale is discussed further in Section \ref{Discussion}.
+ In the examples above. the addition of quasar lonisation lowers the neutral fraction at a fixed redshift. and hence the comparison in Equation (16)) is not made at an equivalent stage of reionisation.," In the examples above, the addition of quasar ionisation lowers the neutral fraction at a fixed redshift, and hence the comparison in Equation \ref{Deltasqrelkz}) ) is not made at an equivalent stage of reionisation."
+ Pherefore. we next isolate the elfect of quasars as à function of neutral fraction by removing the dillerence in overall power.," Therefore, we next isolate the effect of quasars as a function of neutral fraction by removing the difference in overall power."
+ This is achieved by boosting the stellar ionising cllicieney in the stellar-only simulations so as to match the mass-weighted &lobal. neutral fraction of the simulation including quasars., This is achieved by boosting the stellar ionising efficiency in the stellar-only simulations so as to match the mass-weighted global neutral fraction of the simulation including quasars.
+" The resulting ""boosted? realisations for simulation S2 (to match wy,= 0.03. 0.16. 0.28 and 0.43 in the examples) are shown in the central right-hand panels of Figure 5.. and their corresponding. power spectra and sample variance in the right-hand panels."," The resulting “boosted” realisations for simulation S2 (to match $x_{\HIeqn} = 0.03$ , 0.16, 0.28 and 0.43 in the examples) are shown in the central right-hand panels of Figure \ref{128results}, and their corresponding power spectra and sample variance in the right-hand panels."
+ The effect of quasars on the simulated 2l-cn power spectra as a function of neutral fraction. (and. Apu) is again determined by calculating the relative power. where. as mentioned. ASasoso has been calculated using the boostecl simulations so as to match the mass-weighted neutral. fraction. of the simulations including Cquasars.," The effect of quasars on the simulated 21-cm power spectra as a function of neutral fraction (and $M_{\rm h,min}$ ) is again determined by calculating the relative power, where, as mentioned, $\Delta^2_{21,\rm{no\,QSO}}$ has been calculated using the boosted simulations so as to match the mass-weighted neutral fraction of the simulations including quasars."
+" The resulting relative power for simulations S1. 52 and 53 are shown in the right-hand panels of Figures ο and τν,"," The resulting relative power for simulations S1, S2 and S3 are shown in the right-hand panels of Figures \ref{comp_comp_10S_z_xHI_128} and \ref{comp_comp_k}."
+ Figure 6 shows the relative power at different. neutral fractions as a function of wave-numboer., Figure \ref{comp_comp_10S_z_xHI_128} shows the relative power at different neutral fractions as a function of wave-number.
+ Figure 7 shows cuts of the data as a function of neutral fraction at. three different scales corresponding to wave-numbers &zz0.1. 0.4 and MMpc.," Figure \ref{comp_comp_k} shows cuts of the data as a function of neutral fraction at three different scales corresponding to wave-numbers $k \approx 0.1$, 0.4 and Mpc."
+ Ehe effect. of the quasars is scale-depencdent. indicating that the additional bias of the quasars makes fewer larger bubbles than a corresponding increase in stellar flux.," The effect of the quasars is scale-dependent, indicating that the additional bias of the quasars makes fewer larger bubbles than a corresponding increase in stellar flux."
+ Our results indicate that a LO per cent contribution to reionisation from. quasars leads to a damping of the cm power spectrum by a factor of up to 2 at. constant redshift. and by up to 30 per cent at constant neutral fraction shortly before the end of reionisation (εντ~1).," Our results indicate that a 10 per cent contribution to reionisation from quasars leads to a damping of the 21-cm power spectrum by a factor of up to 2 at constant redshift, and by up to 30 per cent at constant neutral fraction shortly before the end of reionisation $\Delta z \sim 1$ )."
+ The level of damping is both a function of scale and redshift., The level of damping is both a function of scale and redshift.
+ Firstly. in all simulations and at all scales. the effect on the relative power is more prominent at later epochs when there is a larger quasar population.," Firstly, in all simulations and at all scales, the effect on the relative power is more prominent at later epochs when there is a larger quasar population."
+ At these later epochs. the cumulative quasar contribution to rejonisation (from both newly formed. and fossil regions) is greater (sce Figure 3)).," At these later epochs, the cumulative quasar contribution to reionisation (from both newly formed and fossil regions) is greater (see Figure \ref{Qevolution}) )."
+ Moreover. the rate at which the moclification evolves is more rapid in models with more massive host halos. resulting from. rapid evolution of the halo mass function.," Moreover, the rate at which the modification evolves is more rapid in models with more massive host halos, resulting from rapid evolution of the halo mass function."
+ The scale. dependence. of the relative power can. be attributed to the ellects of halo bias., The scale dependence of the relative power can be attributed to the effects of halo bias.
+ As discussed. our results show that quasars increase the ionisation fraction at fixed redshift.," As discussed, our results show that quasars increase the ionisation fraction at fixed redshift."
+ his leads to a decrease in the amplitude of the power spectrum late in the reionisation era. since the growth of regions removes signal at a rate that counteracts the increase in power resulting from the region-induced structure in the IGM.," This leads to a decrease in the amplitude of the power spectrum late in the reionisation era, since the growth of regions removes signal at a rate that counteracts the increase in power resulting from the region-induced structure in the IGM."
+ This is the “Pall” of 21-em fluctuations described. by 2..," This is the “fall"" of 21-cm fluctuations described by \cite{lidz2008}."
+ The elfect is larger on small scales. since the addition of quasars increases the size of the regions. owing to quasars being biased. to overdense regions which reionise first.," The effect is larger on small scales, since the addition of quasars increases the size of the regions, owing to quasars being biased to overdense regions which reionise first."
+ Vhis removes more of the sources of small-scale power., This removes more of the sources of small-scale power.
+ The ellect is less significant on large scales. where removal of the source of large-scale power due to the ionisation of gas is partly counteracted by its increase from larger regions.," The effect is less significant on large scales, where removal of the source of large-scale power due to the ionisation of gas is partly counteracted by its increase from larger regions."
+ In the case of comparison at constant neutral fraction the opposite is true., In the case of comparison at constant neutral fraction the opposite is true.
+ Dy increasing the contribution of star formation to match the neutral fraction. the fraction of small-scale power removed by ionisation is the same for stars and quasars.," By increasing the contribution of star formation to match the neutral fraction, the fraction of small-scale power removed by ionisation is the same for stars and quasars."
+ As a result AZ)~1 for large &., As a result $\Delta^2_{\rm{rel}} \sim 1$ for large $k$.
+ However. quasars result in fewer larger ionised regions as they have a higher bias than the galaxies.," However, quasars result in fewer larger ionised regions as they have a higher bias than the galaxies."
+" Late in reionisation this accelerates the decline in power (during the fall) and results in AZ,l for intermediate values of &...", Late in reionisation this accelerates the decline in power (during the “fall”) and results in $\Delta^2_{\rm{rel}} > 1$ for intermediate values of $k$.
+ Thus. the level of modification of the shape of the 21-cm power spectrum can be related. to the host halo masses of the C(quasars.," Thus, the level of modification of the shape of the 21-cm power spectrum can be related to the host halo masses of the quasars."
+ Any radio interferometer is subject to. instrumental noise and has limited sensitivity. arising [rom the finite volume of the observation., Any radio interferometer is subject to instrumental noise and has limited sensitivity arising from the finite volume of the observation.
+ We consider observational parameters corresponding to the design specications of the AIWA. ancl of a hypothetical follow-up to the AIWA (which we term theMWABOOQ0).," We consider observational parameters corresponding to the design specications of the MWA, and of a hypothetical follow-up to the MWA (which we term theMWA5000)."
+ We assume a continuous.," We assume a continuous,"
+live maenituce ranges of2 mag size in the range 13<J«23.,five magnitude ranges of 2 mag size in the range $13<{\rm J}<23$.
+ The lefthand panels show histograms of raw magnitude dillerences. while the right.hand panels show the underlving distributions of random errors left over after. positionallvdependent systematic errors have been removed.," The left–hand panels show histograms of raw magnitude differences, while the right–hand panels show the underlying distributions of random errors left over after positionally--dependent systematic errors have been removed."
+ Once again. the increase in random: error. and decrease in systematic error. with decreasing brightness are apparent.," Once again, the increase in random error, and decrease in systematic error, with decreasing brightness are apparent."
+ Figure 14 shows the consistency. of our survey colours JR (with. Ro [rom both Ulx ancl ESO Schmidts) and R lin the magnitude range 16<Djli., Figure \ref{colover} shows the consistency of our survey colours J–R (with R from both UK and ESO Schmidts) and R--I in the magnitude range $16<{\rm B}_{\rm J}<17$.
+ Once again.o because we have calibrated out the systematic errors in colour as a function of position and magnitude. and. also set the colour zeropoint in every. field. (see Section 2.3.4)) the colours are consistent at a level dicateck by random. errors alone.," Once again, because we have calibrated out the systematic errors in colour as a function of position and magnitude, and also set the colour zeropoint in every field (see Section \ref{colbias}) ) the colours are consistent at a level dictated by random errors alone."
+ The RAIS values about the y=x lines in Figure 14 are ~QOL] mag. indicating single measurement random errors in the colours at the level of ~0.07 mae in this magnitude range.," The RMS values about the $y=x$ lines in Figure \ref{colover} are $\sim0.1$ mag, indicating single measurement random errors in the colours at the level of $\sim0.07$ mag in this magnitude range."
+ While this is strictly speaking a test of internal consistency. it also illustrates the true external errors in the colours (cl.," While this is strictly speaking a test of internal consistency, it also illustrates the true external errors in the colours (cf."
+ Section 3.2.3)) owing to the wav in which any systematic errors have been eliminated., Section \ref{colext}) ) owing to the way in which any systematic errors have been eliminated.
+only major differences are visible for the first cut during the nights when the WSRT was observing with the 84 and 96 m configurations.,only major differences are visible for the first cut during the nights when the WSRT was observing with the 84 and 96 m configurations.
+" The total Faraday depth of a source as a function of time (ó,,,C0) stems from the time-dependent ionospheric Faraday rotation (6;,,(0)) plus the constant contribution of the observed source (Oyj): Through Eq.", The total Faraday depth of a source as a function of time $\phi_{tot}(t)$ ) stems from the time-dependent ionospheric Faraday rotation $\phi_{ion}(t)$ ) plus the constant contribution of the observed source $\phi_{src}$ ): Through Eq.
+ 5 we calculate that has an RM=+8.190.08 rad m. which represents the most accurate determination of the RM of this pulsar to date.," \ref{phi} we calculate that has an ${\rm RM}~=~+8.19
+\pm 0.08$ rad $^{-2}$, which represents the most accurate determination of the RM of this pulsar to date."
+ From previous observations at 2 m. where no correction for the ionospheric Faraday rotation was applied. A.G. de Bruyn found a rotation measure for the pulsar of +8.5+0.5 rad m.," From previous observations at 2 m, where no correction for the ionospheric Faraday rotation was applied, A.G. de Bruyn found a rotation measure for the pulsar of $+8.5\pm0.5$ rad $^{-2}$."
+ At 350 MHz and after correcting for the rotation due to the ionosphere. ? derived a value of RM=+7.86+0.20 rad m7.," At 350 MHz and after correcting for the rotation due to the ionosphere, \citet{Brentjens} derived a value of ${\rm RM} =+7.86\pm0.20$ rad $^{-2}$."
+ After checking the validity of the model for the polarized calibrator. we computed the expected ionospheric RM variations towards the center ofA2255.. which is located in a different area of the sky. and we removed the ionospheric contribution from the data.," After checking the validity of the model for the polarized calibrator, we computed the expected ionospheric RM variations towards the center of, which is located in a different area of the sky, and we removed the ionospheric contribution from the data."
+ After performing the polarization calibration procedure described in Sects. 2.5-—2.6..," After performing the polarization calibration procedure described in Sects. \ref{bandpassandpolarisationcalibration}- \ref{ionosphericfaradayrotation},"
+ we produced the Q and U images of the target source., we produced the Q and U images of the target source.
+ For the 30 background sources observed at 2] em. we determined the polarization angle in each band and plotted it as a function of οἱ”. obtaining the values of their RM through a linear fit.," For the 30 background sources observed at 21 cm, we determined the polarization angle in each sub-band and plotted it as a function of $\lambda^2$, obtaining the values of their RM through a linear fit."
+ For the cluster observations. we made Q and U image cubes and inspected them to remove the channel maps affected by residual RFI and/or problems related to specific telescopes.," For the cluster observations, we made Q and U image cubes and inspected them to remove the channel maps affected by residual RFI and/or problems related to specific telescopes."
+ For the 18 em. 21 em. and 25 cm datasets. a total of 193. 217. and 133 channel images were used for further analysis.," For the 18 cm, 21 cm, and 25 cm datasets, a total of 193, 217, and 133 channel images were used for further analysis."
+ At 85 em we flagged 48 channels from the available 448., At 85 cm we flagged 48 channels from the available 448.
+ At 2 m we used 1523 images out of the available 1904., At 2 m we used 1523 images out of the available 1904.
+ Table 3. shows the average noise level (c) for each Q and U channel for each dataset., Table \ref{noisesQU} shows the average noise level $\sigma$ ) for each Q and U channel for each dataset.
+ This was estimated by subtracting adjacent channel maps., This was estimated by subtracting adjacent channel maps.
+ The final images were carried. through for further processing in. RM-synthesis., The final images were carried through for further processing in RM-synthesis.
+ To improve the resolution in RM space. as well as the sensitivity. we decided to combine the 18 em. 21 em. and 25 em datasets. giving 543 channel images as input to the RM-synthesis.," To improve the resolution in RM space, as well as the sensitivity, we decided to combine the 18 cm, 21 cm, and 25 cm datasets, giving 543 channel images as input to the RM-synthesis."
+ In the wide range of frequencies of each observation the primary beam width changes significantly across theband!., In the wide range of frequencies of each observation the primary beam width changes significantly across the.
+. The input images to RM-synthesis were not corrected for the primary beam attenuation., The input images to RM-synthesis were not corrected for the primary beam attenuation.
+ By doing this a constant spatial noise level across the map was retained., By doing this a constant spatial noise level across the map was retained.
+ After imaging. RM-synthesis was applied to the Q and U cubes.," After imaging, RM-synthesis was applied to the Q and U cubes."
+ This technique derotates the Q-U vectors for each pixel in each channel image. in order to compensate for a certain assumec rotation measure.," This technique derotates the Q-U vectors for each pixel in each channel image, in order to compensate for a certain assumed rotation measure."
+ After derotation. the channel images are averaged.," After derotation, the channel images are averaged."
+ RM-synthesis maximizes the sensitivity to radiatior at the assumed Faraday depth. because that emission ts added coherently.," RM-synthesis maximizes the sensitivity to radiation at the assumed Faraday depth, because that emission is added coherently."
+ All other emission will add up coherently only in part. hence the sensitivity to emission not at the assumed Faraday depth is reduced.," All other emission will add up coherently only in part, hence the sensitivity to emission not at the assumed Faraday depth is reduced."
+ RM-synthesis has already bee mentioned by ? and is known by the pulsar community (??).. but only recently it has been applied to large datasets.," RM-synthesis has already been mentioned by \citet{1966MNRAS.133...67B} and is known by the pulsar community \citep{2003A&A...398..993M,2004ApJS..150..317W}, but only recently it has been applied to large datasets."
+ A full description of how this technique works is presented by ?.. bt= in the following. we give a short overview.," A full description of how this technique works is presented by \citet{br2005}, but in the following, we give a short overview."
+ After computing the complex polarization P = Q + 1U. on[2 multiplies it by à complex phase factor to perform the rotation.," After computing the complex polarization P = Q + iU, one multiplies it by a complex phase factor to perform the rotation."
+" ? find that a general inversion of the polarization as a functio=) of wavelength squared is given by where ¢ 1s the Faraday depth. defined as where Bj is the magnetic field component along the line of sight (/) and 7, the electron density."," \citet{br2005}
+ find that a general inversion of the polarization as a function of wavelength squared is given by where $\phi$ is the Faraday depth, defined as where $_{\parallel}$ is the magnetic field component along the line of sight $l$ ) and $n_e$ the electron density."
+ In equation 6... ΕΦ) ts the emission as a function of Faraday depth. R(@) the rotation measure structure function (RMSF). vty the wavelength to which all vectors are derotated. and WGU) the weight function and represents the sensitivity as a function of wavelength squared.," In equation \ref{1}, $F(\phi)$ is the emission as a function of Faraday depth, $R(\phi)$ the rotation measure structure function (RMSF), $\lambda_0$ the wavelength to which all vectors are derotated, and $W(\lambda^2)$ the weight function and represents the sensitivity as a function of wavelength squared."
+ If the width of a channel is much smaller than the width of the band. the weight function can be approximated by à sum of ó functions. and equation 6 becomes," If the width of a channel is much smaller than the width of the band, the weight function can be approximated by a sum of $\delta$ functions, and equation \ref{1} becomes"
+"the ""constanut-anuetalicity” approximation.",the `constant-metallicity' approximation.
+ This scenario could. result from. no mnüxiug between euviroumuenuts., This scenario could result from no mixing between environments.
+ The ietallicity evolves independently aud rapidly in cach separate star-forming region with a cjwacteristic or average metallicity representing all epocjs (separate regions) of star formation., The metallicity evolves independently and rapidly in each separate star-forming region with a characteristic or average metallicity representing all epochs (separate regions) of star formation.
+ Since the TAdverse has neither complete or no musing. something between these two approximations nüght be Exaiples of differences between these are shown in Figure 9..," Since the Universe has neither complete or no mixing, something between these two approximations might be Examples of differences between these are shown in Figure \ref{fig:Z-approximations}."
+ Though the effect on the broadband colors depenucs on the SEIT and IME. the general effect is fo Increase the UV and near-IR fluxes with respect to the visible flrxes for the coustaut-Z approximation.," Though the effect on the broadband colors depends on the SFH and IMF, the general effect is to increase the UV and near-IR fluxes with respect to the visible fluxes for the $Z$ approximation."
+ This approxiuatio1 leads to lower moetallicities for voung/lot stars and higicr moetallicities for old/cool stars compared to the closed-box evolution scenario (matching Z to constant Z)., This approximation leads to lower metallicities for young/hot stars and higher metallicities for old/cool stars compared to the closed-box evolution scenario (matching $\bar{Z}$ to constant $Z$ ).
+ Some resuls for the constaut-iuetallicity approximation are shown in Figures 10. and 11 for the SDav data., Some results for the constant-metallicity approximation are shown in Figures \ref{fig:const-Z-conf-A} and \ref{fig:const-Z-conf-B} for the SBav data.
+ If we chose solar metallicity then the results are similar to the closed-box approximation except D is in he range 1.51.15 cont.), If we chose solar metallicity then the results are similar to the closed-box approximation except $\Gamma$ is in the range 1.05--1.45 conf.)
+ aud P.«1.7 at confidence., and $\Gamma<1.7$ at confidence.
+ The later confidence level holds for Z=0.025., The latter confidence level holds for $Z\la0.025$.
+ Table 2 shows a comparison between difforent publisred IMEs with multiple power-law slopes., Table \ref{tab:publ-imfs} shows a comparison between different published IMFs with multiple power-law slopes.
+ We 1se the const.uat-metallicity approximation for this comparison to avoid the additional complication introduced by metal productio jiu a closed-box model., We use the constant-metallicity approximation for this comparison to avoid the additional complication introduced by metal production in a closed-box model.
+ With this comparison. the Iro:pa. IMEs are strongly rejected.," With this comparison, the \cite*{KTG93,MS79,Scalo86} IMFs are strongly rejected."
+ This is cousisteut with their average slopes. over audAlsun.. being D1.7.," This is consistent with their average slopes, over and, being $\Gamma\ge1.7$."
+ Note that our modeling cannot strongly distinguish between different slope changes belowMziun., Note that our modeling cannot strongly distinguish between different slope changes below.
+. We have used a power law to deseribe average dust attenuation based on estimating the distributions of, We have used a power law to describe average dust attenuation based on estimating the distributions of
+emploving the self-consistent models of X-ray rellection from ionized material by Ballantvne. Ross Fabian (2001).,"employing the self-consistent models of X-ray reflection from ionized material by Ballantyne, Ross Fabian (2001)."
+ Confirming the principal result of Paper L the degree of relativistic broadening required by these data pushes one to a high value of the emissivity index and a low value of the inner disk radius.," Confirming the principal result of Paper I, the degree of relativistic broadening required by these data pushes one to a high value of the emissivity index and a low value of the inner disk radius."
+ Η one fixes the inclination of the accretion disk at ¢=28 (i.e. the value derived. from. SCA data bv Tanaka et al.," If one fixes the inclination of the accretion disk at $i=28^\circ$ (i.e., the value derived from ASCA data by Tanaka et al."
+ 1995 ancl used in Paper D. the required emissivity index and inner radius are 3=4.29(15 and iy=1.78nm ra). respectively.," 1995 and used in Paper I), the required emissivity index and inner radius are $\beta=4.29^{+0.15}_{-0.16}$ and $r_{\rm
+ in}=1.78^{+0.14}_{-0.11}\,r_{\mathrm g}$ ), respectively."
+" Allowing the inclination to be a free parameter. the best fitting values are even moreextreme (/=5644. —T.10.15 and rg,=1.30ne ra)"," Allowing the inclination to be a free parameter, the best fitting values are even moreextreme $i=56\pm 4^\circ$, $\beta=7.7\pm 0.15$ and $r_{\rm
+ in}=1.30^{+0.21}_{-0.06}\,r_{\mathrm g}$ )."
+ We proceed to discuss spectral fits to the full l0kkeV band., We proceed to discuss spectral fits to the full keV band.
+ This requires us to model the soft. X-ray spectral features. a non-trivial task since the basic physics determining the soft. X-ray spectrum. (i... absorption comission) is still the subject of much debate.," This requires us to model the soft X-ray spectral features, a non-trivial task since the basic physics determining the soft X-ray spectrum (i.e., absorption emission) is still the subject of much debate."
+ Lowever. since this work restricts itself to the medium resolution EPIC data. it ds sullicient. to model the soft. X-ray complexity with phenomenological absorption/enmission components (as oppposecl to constructing physical models of. the absorber/emitter).," However, since this work restricts itself to the medium resolution EPIC data, it is sufficient to model the soft X-ray complexity with phenomenological absorption/emission components (as oppposed to constructing physical models of the absorber/emitter)."
+ In this work. we employ two phenomenological modcls of the soft X-ray complexity.," In this work, we employ two phenomenological models of the soft X-ray complexity."
+ In tie first. model (which. we shall refer to as the “pure warm absorber). we describe the soft. N-rav structure as three simple absorption edges. with threshold. energies° ancl maximum optical1 depths! that are left. as free. parameters in. the spectral fits.," In the first model (which we shall refer to as the “pure warm absorber”), we describe the soft X-ray structure as three simple absorption edges, with threshold energies and maximum optical depths that are left as free parameters in the spectral fits."
+⋅ SinceN these edges are not of physical interest in this work. we do not report their best fitting values in ‘Tables 1- 2..," Since these edges are not of physical interest in this work, we do not report their best fitting values in Tables \ref{tab:fits1}- \ref{tab:fits3}. ."
+ Fvpical threshold, Typical threshold
+to be entirely high spikes in the galaxy noise.,to be entirely high spikes in the galaxy noise.
+ As a check of how well these cuts function to clean our catalog of objects Chat are not elobular clusters. we can apply these same euts to catalogs of objects that contain no globular clusters. ancl only conlaminating objects.," As a check of how well these cuts function to clean our catalog of objects that are not globular clusters, we can apply these same cuts to catalogs of objects that contain no globular clusters, and only contaminating objects."
+ If the variations in (he galaxy noise were purely Gaussian. (hen the number of expected noise maxima on the data image should match the number of minima.," If the variations in the galaxy noise were purely Gaussian, then the number of expected noise maxima on the data image should match the number of minima."
+ By multüiplving the galaxy sublracted image by —1. ancl searching [ου objects on (hat image. we should be able to estimate the noise contribution.," By multiplying the galaxy subtracted image by $-1$, and searching for objects on that image, we should be able to estimate the noise contribution."
+ Unfortunately. ihe asymmetry of the image histograms for these images suggest that the [lucetuations may not be Gaussian distributed. but ave rather biased towards the positive IIuctuations. as might be expected [rom the incipient detection of the brightest individual stas in MT.," Unfortunately, the asymmetry of the image histograms for these images suggest that the fluctuations may not be Gaussian distributed, but are rather biased towards the positive fluctuations, as might be expected from the incipient detection of the brightest individual stars in M87."
+ The contamination by background galaxies obviously contributes to the number of objects found., The contamination by background galaxies obviously contributes to the number of objects found.
+ A quick look at the galaxy subtracted image shows a large number of background spiral galaxies (hal are easy to identify., A quick look at the galaxy subtracted image shows a large number of background spiral galaxies that are easy to identify.
+ The IInbble Ultra Deep Field (Beckwithetal.2006) provides a dataset containing only background galaxies. ancl as such. gives a straighüforwaud wav lo estimate the contribution we expect to find [rom these galaxies.," The Hubble Ultra Deep Field \citep{UDF} provides a dataset containing only background galaxies, and as such, gives a straightforward way to estimate the contribution we expect to find from these galaxies."
+ The only complication is ensuring that the data quality matches between the two datasets., The only complication is ensuring that the data quality matches between the two datasets.
+ We first rebinned the UDF images [rom a pixel seale of 07030 to the pixel scale of our data (0705)., We first rebinned the UDF images from a pixel scale of $0\farcs030$ to the pixel scale of our data $0\farcs05$ ).
+ Next. the UDF was overlaid on our image footprint. and trimmed to match the field of view.," Next, the UDF was overlaid on our image footprint, and trimmed to match the field of view."
+ With the geometry matched. we considered the image noise.," With the geometry matched, we considered the image noise."
+ The UDF count rate was scaled to match (he exposure time of the M57 data., The UDF count rate was scaled to match the exposure time of the M87 data.
+ The AIST galaxy model was then used to generate a noise model. which assumes (hat each pixel is drawn [rom a Gaussian distribution will mean zero and variance equal to the galaxy model value.," The M87 galaxy model was then used to generate a noise model, which assumes that each pixel is drawn from a Gaussian distribution with mean zero and variance equal to the galaxy model value."
+ This noise was then scaled such that when added to the UDF image. the final image variance equaled that of the galaxy. subtracted image.," This noise was then scaled such that when added to the UDF image, the final image variance equaled that of the galaxy subtracted image."
+" This method makes it certain that objects in the UDF frame that [all “near the center of the galaxy” have more noise than objects ""at the edge.”", This method makes it certain that objects in the UDF frame that fall “near the center of the galaxy” have more noise than objects “at the edge.”
+" The reduction for objects detected in the UDF images is identical to (he M87 images. with the only exception being that the UDF filters are F606W and ΕΟΝ, Since these fillers can be converted to the same standard: V and I svstem. the effect of this difference is minimal."," The reduction for objects detected in the UDF images is identical to the M87 images, with the only exception being that the UDF filters are F606W and F775W. Since these filters can be converted to the same standard V and I system, the effect of this difference is minimal."
+ The numbers of objects detected with each of the various ents in the MST data. the inverse images. and the UDF sample are given in table 1..," The numbers of objects detected with each of the various cuts in the M87 data, the inverse images, and the UDF sample are given in table \ref{tab: cuts}."
+ It is clear (hat (he various restrictions on the data remove the majority of the contaminating objects. and vield a final nunmber of clusters consistent with what is expected based on earlier surveys.," It is clear that the various restrictions on the data remove the majority of the contaminating objects, and yield a final number of clusters consistent with what is expected based on earlier surveys."
+ The one object (hat is manually excluded from the sample is the IIST-1 knot of the jet. which due to its brightness and compact size. successfully eludes (hese cuts.," The one object that is manually excluded from the sample is the HST-1 knot of the jet, which due to its brightness and compact size, successfully eludes these cuts."
+In our previous study (VanEertenetal.2011) we presented a method to quickly caleulate light curves based on scalings on the level of the jet dynamics.,In our previous study \citep{vanEerten2011boxfit} we presented a method to quickly calculate light curves based on scalings on the level of the jet dynamics.
+ This method had the disadvantage that a full radiative transfer was still required. for each datapoint., This method had the disadvantage that a full radiative transfer was still required for each datapoint.
+ As a result. broadband: allerelow fittine-hbased on this approach still requires use of a large parallel computer in practice.," As a result, broadband afterglow fitting-based on this approach still requires use of a large parallel computer in practice."
+ The current method no longer requires radiative (ransler calculations and is therefore vastly superior in ternis of computational cost., The current method no longer requires radiative transfer calculations and is therefore vastly superior in terms of computational cost.
+ Nevertheless. in two kev aspects the BOoXx--based. framework Irom VanEertenetal.(2011) remains relevant.," Nevertheless, in two key aspects the -based' framework from \cite{vanEerten2011boxfit} remains relevant."
+ First. since the scalings there happen on the level of the dvnanmics. no parametrizalions are necessary in order to describe the transitions between different spectral regimes and smoothly connected power laws emerge naturally.," First, since the scalings there happen on the level of the dynamics, no parametrizations are necessary in order to describe the transitions between different spectral regimes and smoothly connected power laws emerge naturally."
+ Second. the scalable BOX--based blast wave dynamics dala provide a testing lab for studying the effect of various radiative processes. while the current study Cakes svnchrotron radiation as its starting point.," Second, the scalable -based blast wave dynamics data provide a testing lab for studying the effect of various radiative processes, while the current study takes synchrotron radiation as its starting point."
+ This research was supported in part by NASA through grant. NNXIOAFG2G. issued through the Astrophysics Theory Program and by the NSF through grant. AST-1009863., This research was supported in part by NASA through grant NNX10AF62G issued through the Astrophysics Theory Program and by the NSF through grant AST-1009863.
+ The software used in this work was in part developed by the DOE-supported ASCI/Alliance Center [or Astrophysical Thermonuclear Flashes at the University of Chicago., The software used in this work was in part developed by the DOE-supported ASCI/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago.
+ We thank Alexander van der Horst and Andrei Gruzinov for helpful, We thank Alexander van der Horst and Andrei Gruzinov for helpful
+"Thus. for each transition. the observed line. profile (line/continuum) is written: where 7), and 7. are the main beam and the contiuum temperatures. N. and N; are the numbers of components and hyperfinetransitions’.. ro. vo and c, are the Gaussian parameters. and a, and vy, are the LTE and optically thin line strengths relative to a reference hyperfine transition. and the velocity shift associated to each hyperfine transition.","Thus, for each transition, the observed line profile (line/continuum) is written: where $T_{\rm mb}$ and $T_{\rm c}$ are the main beam and the continuum temperatures, $N_c$ and $N_h$ are the numbers of components and hyperfine, $\tau_0$, $\upsilon_0$ and $\sigma_{\upsilon}$ are the Gaussian parameters, and $\alpha_{h}$ and $\upsilon_{0h}$ are the LTE and optically thin line strengths relative to a reference hyperfine transition and the velocity shift associated to each hyperfine transition."
+ As described in Appendix AppendixA:.. the column density associated to each velocity component is then derived assuming a single excitation temperature 7.4 for all the levels of a given molecule.," As described in Appendix \ref{AppendGauss}, the column density associated to each velocity component is then derived assuming a single excitation temperature $T_{\rm ex}$ for all the levels of a given molecule."
+ For the components which are either saturated or too weak to be singled out by the fitting procedure. limits on the column densities are derived from the integrated optical depth over the correspondingvelocity range.," For the components which are either saturated or too weak to be singled out by the fitting procedure, limits on the column densities are derived from the integrated optical depth over the correspondingvelocity range."
+ All the results are listed in Tabs. Al-, All the results are listed in Tabs. \ref{TabFitHCOp}-
+-C2 of Appendix AppendixA: and. as an example. the outcome of the multi-Gaussian decomposition. applied to the absorption lines observed toward W49N. is shown in Fig. 3..," \ref{TabFitCN} of Appendix \ref{AppendGauss} and, as an example, the outcome of the multi-Gaussian decomposition, applied to the absorption lines observed toward W49N, is shown in Fig. \ref{FigMultigauss}."
+ We find that the widths of the Gaussian components have a continuous distribution with values ranging from 0.3 tto 3.5s!.. independently of the source or the molecular species.," We find that the widths of the Gaussian components have a continuous distribution with values ranging from 0.3 to 3.5, independently of the source or the molecular species."
+ The peak optical depths range between 0.06 and 2.2 and the inferred column densities per velocity component span more than one order of magnitude on each line of sight., The peak optical depths range between 0.06 and 2.2 and the inferred column densities per velocity component span more than one order of magnitude on each line of sight.
+ The uncertainties given in Tabs. Al-, The uncertainties given in Tabs. \ref{TabFitHCOp}-
+-C2> are the formal |-σ errors derived from the diagonal elements of the covariance matrix and do not take into account the systematic errors introduced by: (1) the finite velocity resolution. (2) the uncertainty on the excitation temperatures (see Table 3)). and (3) the uncertainty on the continuum level.," \ref{TabFitCN} are the formal $\sigma$ errors derived from the diagonal elements of the covariance matrix and do not take into account the systematic errors introduced by: (1) the finite velocity resolution, (2) the uncertainty on the excitation temperatures (see Table \ref{TabTex}) ), and (3) the uncertainty on the continuum level."
+ We show in Appendix AppendixD: that the finite velocity resolution introduces an error on the column densities smaller than 12%., We show in Appendix \ref{AppendAbscissa} that the finite velocity resolution introduces an error on the column densities smaller than 12.
+. Table 3 gives the uncertainties on the excitation temperature of the lowest rotational levels of HCO™. that affects the partition function. hence the column densities.," Table \ref{TabTex} gives the uncertainties on the excitation temperature of the lowest rotational levels of $^{+}$, that affects the partition function, hence the column densities."
+ Toward G10.62-0.38. W49N. and W51. the uncertainties on T4 are small. providing uncertainties on the column densities E.," Toward G10.62-0.38, W49N, and W51, the uncertainties on $T_{\rm ex}$ are small, providing uncertainties on the column densities $<^{+50\%}_{-10\%}$."
+ Toward G34.3+0.1. the inferred column densities could be underestimated by a factor of 3.," Toward G34.3+0.1, the inferred column densities could be underestimated by a factor of 3."
+ Finally. the uncertainty ε=ὅτι on the continuum temperature (see Sect.," Finally, the uncertainty $\epsilon = \delta T_{\rm c}/T_{\rm c}$ on the continuum temperature (see Sect."
+ 2) introduces an error or on the caleulation of the optical depths: This error is larger. in most cases. than those inferred from the fitting procedure. and for e=10% the corresponding uncertainty on the column densities ranges from to when το varies between 0.5 and 2.," 2) introduces an error $\delta \tau$ on the calculation of the optical depths: This error is larger, in most cases, than those inferred from the fitting procedure, and for $\epsilon = 10$ the corresponding uncertainty on the column densities ranges from to when $\tau_0$ varies between 0.5 and 2."
+ All components with ro>2.2 are thus considered as saturated., All components with $\tau_0 > 2.2$ are thus considered as saturated.
+ The components at 16 and 4I oof the G10.62-0.38 HCO'(0-1) spectrum are the only ones that fall in this category., The components at 16 and 41 of the G10.62-0.38 $^{+}$ (0-1) spectrum are the only ones that fall in this category.
+ When all uncertainties are taken into account. the column densities are determined within a factor of 2. and within a factor of 3 for the absorption lines observed toward G34.3+0.1.," When all uncertainties are taken into account, the column densities are determined within a factor of 2, and within a factor of 3 for the absorption lines observed toward G34.3+0.1."
+ Because the line centroids of a Gaussian component observed in several transitions coincide within 0.15 km s! (except for 6 components. see Appendix AppendixA:.. Tabs.," Because the line centroids of a Gaussian component observed in several transitions coincide within 0.15 km $^{-1}$ (except for 6 components, see Appendix \ref{AppendGauss}, Tabs."
+ Al. - C2). we propose that the corresponding velocity component has a physical reality. in support of the comparison between the optical depths in the different molecular lines and the resulting column densities.," \ref{TabFitHCOp} - \ref{TabFitCN}) ), we propose that the corresponding velocity component has a physical reality, in support of the comparison between the optical depths in the different molecular lines and the resulting column densities."
+ The column densities and linewidths of all the unsaturated velocity components are displayed in Figs., The column densities and linewidths of all the unsaturated velocity components are displayed in Figs.
+ 4+ and 5 for several pars of species.," \ref{FigCorrelation1} and \ref{FigCorrelation2}
+ for several pairs of species."
+ These figures also include. for comparison. the results of previous studies of molecular absorption lines observed toward strong extragalactic mm-wave continuum sources (Lucas Liszt 1996. 2000; Liszt Lucas 2001) that mostly sample diffuse gas i1 the Solar Neighbourhood.," These figures also include, for comparison, the results of previous studies of molecular absorption lines observed toward strong extragalactic mm-wave continuum sources (Lucas Liszt 1996, 2000; Liszt Lucas 2001) that mostly sample diffuse gas in the Solar Neighbourhood."
+ Although we have not made any assumption on the ature of a Gaussian velocity component. it is remarkable that the range of the component linewidths (0.3 to 3.5 s7!)) is about the same in both data sets. and the range of column densities of most molecular species (up to two orders of magnitude) is also similar in both sets.," Although we have not made any assumption on the nature of a Gaussian velocity component, it is remarkable that the range of the component linewidths (0.3 to 3.5 ) is about the same in both data sets, and the range of column densities of most molecular species (up to two orders of magnitude) is also similar in both sets."
+ This result suggests that the gas components sampled at high galactic latitude anc toward the inner Galaxy share common kinematic properties., This result suggests that the gas components sampled at high galactic latitude and toward the inner Galaxy share common kinematic properties.
+ In addition the rotation curve in the first quadrant provides displacements ~70—100 pe/km s': therefore. if we assume that the induced spreading of the velocity centroids is responsible for the observed linewidths. a component of 3 km s! (resp.," In addition the rotation curve in the first quadrant provides displacements $\sim 70 - 100$ pc/km $^{-1}$: therefore, if we assume that the induced spreading of the velocity centroids is responsible for the observed linewidths, a component of 3 km $^{-1}$ (resp."
+ 0.3 km s~ ') would have a size of 300 pe (resp., 0.3 km $^{-1}$ ) would have a size of 300 pc (resp.
+ 30 pe)., 30 pc).
+ This would correspond to densities jj<40 cm. below our estimates (see Sect. 4.3)).," This would correspond to densities $n_H < 40$ $^{-3}$ , below our estimates (see Sect. \ref{PhysPropGas}) )."
+ Hence. the linewidths of the components are," Hence, the linewidths of the components are"
+"robust weighting (robustness = 1.0) giving a spatial resolution of19.9H17.0"" and a velocity resolution of 20.0 (alter Llanning smoothing).","robust weighting (robustness = 1.0) giving a spatial resolution of$19.9^{\prime\prime}\times
+17.0^{\prime\prime}$ and a velocity resolution of 20.0 (after Hanning smoothing)."
+ The noise in the final datacube is 0.23 mJy +., The noise in the final datacube is 0.23 mJy $^{-1}$.
+ ‘To construct the total limage. a mask was created using a datacube that was smoothed to about twice the spatial resolution and that was clipped at twice the noise of that smoothed datacube.," To construct the total image, a mask was created using a datacube that was smoothed to about twice the spatial resolution and that was clipped at twice the noise of that smoothed datacube."
+ The resulting total iis shown in Figure 1.. ancl our observations show that the iis distributed in a regular rotating ring instead of a. filled disc.," The resulting total is shown in Figure \ref{fig:optHI}, and our observations show that the is distributed in a regular rotating ring instead of a filled disc."
+" The inner radius of the ring is approximately 50"" (~ 5 kpc) and extends to 1207. which corresponds to 12 kpe. or 5 ellective radii (1 42, = 24"")."," The inner radius of the ring is approximately $50^{\prime \prime}$ $\sim$ 5 kpc) and extends to $120^{\prime
+ \prime}$, which corresponds to 12 kpc, or 5 effective radii (1 $R_e$ = $^{\prime \prime}$ )."
+ The vvelocity. Geld was derived. by fitting Gaussians to the spectra at those positions where signal is detected in the total image., The velocity field was derived by fitting Gaussians to the spectra at those positions where signal is detected in the total image.
+ Phe resulting velocity map is shown in Figure 2.., The resulting velocity map is shown in Figure \ref{fig:velmap}.
+ Typical errors on this map are 5 - 10., Typical errors on this map are 5 - 10.
+τν We find a total mass of 5.5107A7. for the &eas content of the ring. which is in agreement with Wim et al. (1988)..," We find a total mass of $5.5 \times 10^8 M_\odot$ for the gas content of the ring, which is in agreement with Kim et al. \nocite{1988ApJ...330..684K},"
+ if we correct for the dillerence in. assumed distance modulus., if we correct for the difference in assumed distance modulus.
+ The amount aid morphology of the oobserved. in NGC 2974 are not unusual for. early-type galaxies., The amount and morphology of the observed in NGC 2974 are not unusual for early-type galaxies.
+ Oosterloo et al., Oosterloo et al.
+ have found that between 5 and 10 per cent of earlv-tvpe galaxies show nimiasses well above 107AZ.. while the fraction of detections increases further for lower masses (Morganti et. al. 2006)).," \nocite{2007A&A...465..787O} have found that between 5 and 10 per cent of early-type galaxies show masses well above $10^9 M_\odot$, while the fraction of detections increases further for lower masses (Morganti et al. \nocite{2006MNRAS.371..157M}) )."
+ The majority of the lli-rich systems have the neutral hydrogen distributed in disc/ring like structures (often warped) with low surface brightness density and no or little ongoing star formation. as observed in NGC 2974.," The majority of the -rich systems have the neutral hydrogen distributed in disc/ring like structures (often warped) with low surface brightness density and no or little ongoing star formation, as observed in NGC 2974."
+ However. there is a region in the North-East of the rring where the surface density is higher. and the gas could be forming stars.," However, there is a region in the North-East of the ring where the surface density is higher, and the gas could be forming stars."
+ Jeong et al., Jeong et al.
+ published UV imaging of NGC 2974. obtained with GALEN.," \nocite{2007MNRAS.376.1021J} published UV imaging of NGC 2974, obtained with GALEX."
+ Their images reveal incleed a region of increased starformation in the North-East of the galaxy. as well as a starforming ring at the inner edges of the rring.," Their images reveal indeed a region of increased starformation in the North-East of the galaxy, as well as a starforming ring at the inner edges of the ring."
+ At least some of the most rrich structures are the results of major mergers (see e.g. Serra et al. 2006))., At least some of the most rich structures are the results of major mergers (see e.g. Serra et al. \nocite{2006A&A...453..493S}) ).
+ For the svstems with less extreme mumasses. like NGC 2974. the origin of the gas is less clear.," For the systems with less extreme masses, like NGC 2974, the origin of the gas is less clear."
+ Accretion of small companions is a possibility. but. smooth. Cold. accretion from the intergalactic medium (ICM) is an alternative scenario.," Accretion of small companions is a possibility, but smooth, cold accretion from the intergalactic medium (IGM) is an alternative scenario."
+ Maps of the stellar and ionised gas kinematics of NGC 2974. obtained with the integral-field spectrograph were presented in [msellem et al.," Maps of the stellar and ionised gas kinematics of NGC 2974, obtained with the integral-field spectrograph were presented in Emsellem et al."
+ and Sarzi οἱ al. (2006)..," \nocite{2004MNRAS.352..721E} and Sarzi et al. \nocite{2006MNRAS.366.1151S},"
+ respectively. and we refer the reader to these papers for the methocs of data reduction and extraction of the kinematics.," respectively, and we refer the reader to these papers for the methods of data reduction and extraction of the kinematics."
+ In Figure 2. we compare both the vvelocity maps of stars ane wwith the velocity map of the rring., In Figure \ref{fig:velmap} we compare both the velocity maps of stars and with the velocity map of the ring.
+ Stars and gas are well aligned. and the transition between the ionised and the neutral gas seems to be smooth. suggesting that they form one single disc.," Stars and gas are well aligned, and the transition between the ionised and the neutral gas seems to be smooth, suggesting that they form one single disc."
+ Phe twist in the velocity map of the ionised gas in the inner 4 is likely caused by the inner bar of this galaxy. (IEmisellem et al. 2003..," The twist in the velocity map of the ionised gas in the inner $4^{\prime\prime}$ is likely caused by the inner bar of this galaxy (Emsellem et al. \nocite{2003MNRAS.345.1297E},"
+ Ixrajnovié et al. 2005))., Krajnović et al. \nocite{2005MNRAS.357.1113K}) ).
+ We used kinemetry (xrajnovié ct al. 2006)), We used kinemetry (Krajnović et al. \nocite{2006MNRAS.366..787K}) )
+ to analyse the aancd VLA velocity maps., to analyse the and VLA velocity maps.
+ In our application to a eas disc. kinemetry reduces to the tilted-ring method (Begeman 1978)).," In our application to a gas disc, kinemetry reduces to the tilted-ring method (Begeman \nocite{1987PhDT.......199B}) )."
+ The velocity along each elliptical ring is expanded in Fourier components (e.g. Franx et al. 1994:, The velocity along each elliptical ring is expanded in Fourier components (e.g. Franx et al. \nocite{1994ApJ...436..642F};
+" Schoenmakers. Franx de Zecuw 1997)): where Vi, is the observed. velocity. 2 is the length of the semimajor axis of the elliptical ring. ó the azimuthal angle. measured from the projected major axis of the galaxy. Viv. the svstemic velocity of the ring and e, and s, are the coellicients of the harmonic expansion."," Schoenmakers, Franx de Zeeuw \nocite{1997MNRAS.292..349S}) ): where $V_{\mathrm{los}}$ is the observed velocity, $R$ is the length of the semimajor axis of the elliptical ring, $\phi$ the azimuthal angle, measured from the projected major axis of the galaxy, $V_{\mathrm{sys}}$ the systemic velocity of the ring and $c_n$ and $s_n$ are the coefficients of the harmonic expansion."
+" The c; term relates to the circular velocity V; in the disc. so that ej=V,sin. where 7 is the inclination of the gas disc."," The $c_1$ term relates to the circular velocity $V_c$ in the disc, so that $c_1 = V_c
+\sin i$, where $i$ is the inclination of the gas disc."
+ Assuming that motions in the ring are intrinsically circular and that the ring is infinitely thin. the inclination can be inferred. from the Hattening q of the fitted ellipse: cos£= q.," Assuming that motions in the ring are intrinsically circular and that the ring is infinitely thin, the inclination can be inferred from the flattening $q$ of the fitted ellipse: $\cos i = q$ ."
+ lf a gas disc only displavs pure circular motions. all harmonicionc terms termsother than cjin Equation (1)) are (1))zero.zer," If a gas disc only displays pure circular motions, all harmonic terms other than $c_1$ in Equation \ref{eq:expan}) ) are zero."
+ Noncircular motions. originating from e.g. inllows caused bv spiral arms or bars. or a triaxial potential. will cause," Noncircular motions, originating from e.g. inflows caused by spiral arms or bars, or a triaxial potential, will cause"
+transparency of the absorbers.,transparency of the absorbers.
+ Our parameter studies of Nyy in the lieht curve model support this interpretation (see. Fie. 5))., Our parameter studies of $N_{\rm H}$ in the light curve model support this interpretation (see Fig. \ref{nh_chi}) ).
+ An cxaimple of achromatic dipping at low energies as a consequence of hieh-Nq; absorbers iu relation to the shape of the source spectrum is Cre N-2 (?).., An example of achromatic dipping at low energies as a consequence of $N_{\rm H}$ absorbers in relation to the shape of the source spectrum is Cyg X-2 \citep{cygx2}.
+ While the οὐπα is much harder thaw in USSeo. the opacity of the absorbers is ~107 7. leading to full blockage of all light iu the 0.5-2kkeV baud.," While the continuum is much harder than in Sco, the opacity of the absorbers is $\sim 10^{23}$ $^{-2}$, leading to full blockage of all light in the keV band."
+ The harduess ratio from 0.5-1.0/1.0-2.0 band. is therefore constaut during the dips. while ouly ciission from the L5e.0kkeV baud can partially penetrate.," The hardness ratio from 0.5-1.0/1.0-2.0 band is therefore constant during the dips, while only emission from the keV band can partially penetrate."
+ Iu SSco. not οποιο] cinission at energies above kkeV is present to coustrain Nyy of the absorbers iu the wav this is possible for Cre A question of iuterest is the reformation of the accretion disk that was destroved during the mova outburst (?)..," In Sco, not enough emission at energies above keV is present to constrain $N_{\rm H}$ of the absorbers in the way this is possible for Cyg A question of interest is the reformation of the accretion disk that was destroyed during the nova outburst \citep{drake10}."
+ The transition from dipping on dav 22.9 to clean eclipse on dav 319 may. be part of the reformation process. and the concept of intersectiug opaque chuups as part of the reformation process deserves closer cousideration.," The transition from dipping on day 22.9 to clean eclipse on day 34.9 may be part of the reformation process, and the concept of intersecting opaque clumps as part of the reformation process deserves closer consideration."
+ The theoretical possibility of clumps in the ejecta has been discussed by ?.. however. being inunersed in a hieh-raciation cuviromment. the ejecta mist be highlv ionized. and as such. they do uot absorb soft N-ravs wia the photoelectric effect.," The theoretical possibility of clumps in the ejecta has been discussed by \cite{diaz10}, however, being immersed in a high-radiation environment, the ejecta must be highly ionized, and as such, they do not absorb soft X-rays via the photoelectric effect."
+ Thompson scattering in clunpy iaterial can also be ruled out as cause for the dips as that would also lead to the jue dips iu the UV/optical light curves., Thompson scattering in clumpy material can also be ruled out as cause for the dips as that would also lead to the same dips in the UV/optical light curves.
+ Chiipy ejecta can therefore not cause such deep dips iu the Noay light curve., Clumpy ejecta can therefore not cause such deep dips in the X-ray light curve.
+ A possible source of opaque clumps consisting of cooler material is the companion star. if accretion has set in again.," A possible source of opaque clumps consisting of cooler material is the companion star, if accretion has set in again."
+ Cooler plaza is opaque to soft N-rav while at the same time beime trausparcut to liebt with wavelengths longer than ((13.6ceV). the ionization energv of atomic bydroven. which would explain why no dips are seen in the simultaneous UV and optical Light curves.," Cooler plasma is opaque to soft X-ray while at the same time being transparent to light with wavelengths longer than eV), the ionization energy of atomic hydrogen, which would explain why no dips are seen in the simultaneous UV and optical light curves."
+ Absorbing atonmüc inaterial from the companion is thus strong support for the interpretation that accretion has set iu again., Absorbing atomic material from the companion is thus strong support for the interpretation that accretion has set in again.
+ This does not require the full reestablishment of the accretion disk. a ον accretion streani can als| be tmagined.," This does not require the full reestablishment of the accretion disk, a clumpy accretion stream can also be imagined."
+ The absence of dips after phase 1.25 suggests that no clhuupy material is on the far side as seen from the companion which is an argument in favor of an accretion stream. residing between the companion and the ceutral white Iu Fig. 16..," The absence of dips after phase 1.25 suggests that no clumpy material is on the far side as seen from the companion which is an argument in favor of an accretion stream, residing between the companion and the central white In Fig. \ref{trajectory},"
+ we show a theoretically. calculated initial trajectory of a reforming disk based on equations 1. aud 2 in ?.. asunuue the parameters of SSco.," we show a theoretically calculated initial trajectory of a reforming disk based on equations 1 and 2 in \cite{Flannery75_mnras}, assuming the parameters of Sco."
+" During the carly disk formation phase. aun asvnunetrical ""ring is formed."," During the early disk formation phase, an asymmetrical ""ring"" is formed."
+ The phases that correspond to different viewing aneles are inarked., The phases that correspond to different viewing angles are marked.
+ The hashed area in the ceuter encircles the optically thick parts of the cjecta., The hashed area in the center encircles the optically thick parts of the ejecta.
+ The surface can be considered the source of N-ray eniission. and no views inside are possible.," The surface can be considered the source of X-ray emission, and no views inside are possible."
+ The highly elliptical trajectory rus through the optically thick part of the ejecta. aud απ inhomoecucitics inside the plotosphere can not cause dips.," The highly elliptical trajectory runs through the optically thick part of the ejecta, and any inhomogeneities inside the photosphere can not cause dips."
+ An eccentric trajectory of infalline material cau thus be the explanation for uot seciug ay dips after phase A vemarkably similar event has been observed by ? in the accreting neutron star in 28 (0918-519.," An eccentric trajectory of infalling material can thus be the explanation for not seeing any dips after phase A remarkably similar event has been observed by \cite{zand2011}
+ in the accreting neutron star in 2S 0918-549."
+ After an X-ray burst. strong (87c1% peak-to-peak auplitude) achromatic fluctuations were seocn ax a transient phenomenon.," After an X-ray burst, strong $87\,\pm\,1$ peak-to-peak amplitude) achromatic fluctuations were seen as a transient phenomenon."
+ Since the outburst evolved mach more rapidly. the N-rav light curve taken with the Proportional Counter Arrav (PCA) ou the Rossi X-ray Timing Explorer (RNTE) may be interpreted as a fast-motion picture of the evolution iu USSco (see figure 1 in 7)).," Since the outburst evolved much more rapidly, the X-ray light curve taken with the Proportional Counter Array (PCA) on the Rossi X-ray Timing Explorer (RXTE) may be interpreted as a fast-motion picture of the evolution in Sco (see figure 1 in \citealt{zand2011}) )."
+ It was speculated by ? that the fluctuations were due to Thonpson scattering by inhomogencitics in a resctthue accretion disk that was disrupted by the effects of super-Eddingtou fluxes. much like in SSco where the blast wave has cestroved the accretion isk (?)..," It was speculated by \cite{zand2011} that the fluctuations were due to Thompson scattering by inhomogeneities in a resettling accretion disk that was disrupted by the effects of super-Eddington fluxes, much like in Sco where the blast wave has destroyed the accretion disk \citep{drake10}."
+ Furthermore. they argue that an expanding shell uav be the necessary prerequisite for the fluctuations.," Furthermore, they argue that an expanding shell may be the necessary prerequisite for the fluctuations."
+ bregulanities iu the accretion stream could be caused * the continued expansion of the nova ejecta that as stopped between days 22.9 and 31.9., Irregularities in the accretion stream could be caused by the continued expansion of the nova ejecta that has stopped between days 22.9 and 34.9.
+ The stream could be disrupted by the ejecta. or the companion star mav be distorted in the outer lavers.," The stream could be disrupted by the ejecta, or the companion star may be distorted in the outer layers."
+ A cliffereuce o 7 is the interpretation of the properties of the intervening niaterial. vielding hiehly ionized plasma.," A difference to \cite{zand2011} is the interpretation of the properties of the intervening material, yielding highly ionized plasma."
+ Dips are caused by Thompson scattering out of the ine of sight rather than photoclectric absorption., Dips are caused by Thompson scattering out of the line of sight rather than photoelectric absorption.
+ This is consistent with their observation of increases above he decay trend of the burst which can be explained wo osidewars or backward scattering of clouds located, This is consistent with their observation of increases above the decay trend of the burst which can be explained by sideways or backward scattering of clouds located
+Although Leo T contains gas. it properties as similar to those of other ultra-faint dSph galaxies and quite different from those of more massive dwarf irregular galaxies (2)..,"Although Leo T contains gas, it properties as similar to those of other ultra-faint dSph galaxies and quite different from those of more massive dwarf irregular galaxies \citep{BovillR:08}."
+ Its properties are consistent with those of pre-ionization fossils but. in the context of this model. it is difficult to understand how it could hold on to its gas.," Its properties are consistent with those of pre-ionization fossils but, in the context of this model, it is difficult to understand how it could hold on to its gas."
+ We point out that dark minihalos with large concentration are able to accrete gas from the IGM more efficiently than lower concentration halos of the same mass., We point out that dark minihalos with large concentration are able to accrete gas from the IGM more efficiently than lower concentration halos of the same mass.
+ The concentration of a halo depends on the redshift of virialization as c(2)x(1|Sen)| 2)., The concentration of a halo depends on the redshift of virialization as $c(z) \propto (1+z_{vir})/(1+z)$ .
+ Thus. as a result their growing concentration. early forming minihalos. although ofaffected by reionization feedback. may have a late phase of cold gas accretion from the IGM and form stars (perhaps for the first time) at 21.," Thus, as a result of their growing concentration, early forming minihalos, although affected by reionization feedback, may have a late phase of cold gas accretion from the IGM and form stars (perhaps for the first time) at $z<1$."
+ This model explains all the observed properties of Leo T. Having a distance of 420 kpe from the Milky Way. Leo T is just starting to fall into the Milky Way halo. thus it is likely to have evolved in isolation (fulfilling a requirement of the model) and unlikely to aave been tidally stripped.," This model explains all the observed properties of Leo T. Having a distance of $420$ kpc from the Milky Way, Leo T is just starting to fall into the Milky Way halo, thus it is likely to have evolved in isolation (fulfilling a requirement of the model) and unlikely to have been tidally stripped."
+ Recent observations suggest that Leo T ius a bimodal SEH (2). with a 12 Gyr population and a 9 Gyr population (with SF continuing until few Myrs ago). in agreement with our expectation of a late phase of accretion from he IGM.," Recent observations suggest that Leo T has a bimodal SFH \citep{DeJong:08} with a $>12$ Gyr population and a $9$ Gyr population (with SF continuing until few Myrs ago), in agreement with our expectation of a late phase of accretion from the IGM."
+" The bimodality of the SFH should be most pronounced in minihalos with the lowest c;,.. gradually disappearing as Όρη increases toward the 20 km ! threshold for. pre-reionization 5ossils. (RGOS)."," The bimodality of the SFH should be most pronounced in minihalos with the lowest $v_{cir}$, gradually disappearing as $v_{cir}$ increases toward the $ 20$ km $^{-1}$ threshold for pre-reionization fossils (RG05)."
+nc Leo T has a gas velocity. dispersion. ~-7 kms l~. hat. according to our model. implies a dark halo with maximum circular velocity of 11 km + (Eq. 41.," Leo T has a gas velocity dispersion $\sim 7$ km $^{-1}$, that, according to our model, implies a dark halo with maximum circular velocity of $11$ km $^{-1}$ (Eq. \ref{eq:vcir_c}) )."
+" Leo T has a core radius ré100 pe. Adnanfly~S. 10022) and maximum column density Vy=7«107"" 7."," Leo T has a core radius $r_{c} \sim 100$ pc, $M_{dyn}/M_g \sim 8-10$ \citep{SimonG:07, Ryan-Weber:08} and maximum column density $N_H=7\times 10^{20}$ $^{-2}$."
+ These values are all in agreement with our model as shown in 2? and right panel in Fig. 4.., These values are all in agreement with our model as shown in \ref{sec:gas} and right panel in Fig. \ref{fig:core2}. .
+ If future observations will show that the dark halo of Leo T has VU20 kms + or its SFH does not show any decline after 12 Gyr ago. our model for Leo T will be disproved.," If future observations will show that the dark halo of Leo T has $v_{cir}^{max} >20$ km $^{-1}$ or its SFH does not show any decline after $12$ Gyr ago, our model for Leo T will be disproved."
+ Minihalos smaller than the one hosting Leo T can also accrete gas from the IGM at. «1. but may not be able to form stars.," Minihalos smaller than the one hosting Leo T can also accrete gas from the IGM at $z<1$, but may not be able to form stars."
+ Sufficiently small minihalos are likely to be dark and. if they evolved in isolation. the IGM around them would be nearly metal-free.," Sufficiently small minihalos are likely to be dark and, if they evolved in isolation, the IGM around them would be nearly metal-free."
+ Cosmological simulations show that the fraction of primordial minihalos that host a luminous galaxy decreases steeply with decreasing c4: f(fPrm)omin(roi10kmsU] (?)..," Cosmological simulations show that the fraction of primordial minihalos that host a luminous galaxy decreases steeply with decreasing $v_{cir}$: $f(lum) \sim min[1, 50\%(v_{cir}/10~{\rm
+ km/s})^{4.6}]$ \citep{RicottiGS:08}."
+ A fraction of dark minihalos may beL504 able to accrete gas from the IGM but. if the gas density in their core is below the threshold required for star formation. these objects will resemble CHVCs.," A fraction of dark minihalos may be able to accrete gas from the IGM but, if the gas density in their core is below the threshold required for star formation, these objects will resemble CHVCs."
+ Additional work is needed to determine whether our model is fully consistent with the properties of known CHVCs (222).. however the typical column densities we predict in Fig. J4t(," Additional work is needed to determine whether our model is fully consistent with the properties of known CHVCs \citep{Blitz:99, Braun:99, Robishaw:02}, however the typical column densities we predict in Fig. \ref{fig:core2}( ("
+left panel) are in agreement with a subset of existing observations.,left panel) are in agreement with a subset of existing observations.
+ The typical size and size distribution of CHVCs may agree with the model assuming a mean distance of minihalos of | Mpe., The typical size and size distribution of CHVCs may agree with the model assuming a mean distance of minihalos of 1 Mpc.
+ We estimate that realistic extension of the diffuse gas in dark minihalos are <5189 kpe.," We estimate that realistic extension of the diffuse gas in dark minihalos are $\simlt
+5-10$ kpc."
+ This may exclude several known CHVC and/or set an upper limit for their distance from the Milky Way., This may exclude several known CHVC and/or set an upper limit for their distance from the Milky Way.
+ More extended extragalactic CHVCs are possible for minihalo masses AJ2107 M. that are likely to be luminous., More extended extragalactic CHVCs are possible for minihalo masses $M>10^8$ $_\odot$ that are likely to be luminous.
+ In addition this scenario for more massive minihalos has already been investigated by previous studies and found not fully satisfactory €22?2)..," In addition this scenario for more massive minihalos has already been investigated by previous studies and found not fully satisfactory \citep{Sternberg:02,
+ Maloney:03, Putman:03}."
+ More work in this area is also motivated by the ongoing survey ALFALFA (??). that may be able to discover a even fainter population of extragalactic CHVCs.," More work in this area is also motivated by the ongoing survey ALFALFA \citep{Giovanelli:05, Giovanelli:07} that may be able to discover a even fainter population of extragalactic CHVCs."
+ Hence. future work shall focus on estimating the number. flux and size distribution of gas rich minihalos in the Local Volume and the local voids. producing synthetic maps for 21 em and H-a wavelengths.," Hence, future work shall focus on estimating the number, flux and size distribution of gas rich minihalos in the Local Volume and the local voids, producing synthetic maps for 21 cm and $\alpha$ wavelengths."
+ In this appendix we derive the equations describing thedensity profile of intergalactic gas inside a dark minihalo., In this appendix we derive the equations describing thedensity profile of intergalactic gas inside a dark minihalo.
+ Although. the IGM «density. temperature and dark halo potential evolve with," Although, the IGM density, temperature and dark halo potential evolve with"
+low-luminosity transient (C13zMy2 —15) with narrow lines in emission in the spectra (v3000 km/s). and signs of a circumstellar dust excess at near-IR and mid-IR wavelengths.,"low-luminosity transient $-13 \gtrsim M_{V} \gtrsim -15$ ) with narrow lines in emission in the spectra $v \lesssim 3000$ km/s), and signs of a circumstellar dust excess at near-IR and mid-IR wavelengths."
+ Examinations of the dust properties (Prietoetal.20097; Wessonetal. 2010)) suggest (3) that the dust is carbonaceous rather than the silicate dust seen in massive stars., Examinations of the dust properties \citealt{Prieto2009}; \citealt{Wesson2010}) ) suggest (3) that the dust is carbonaceous rather than the silicate dust seen in massive stars.
+ Here we add (4) that the progenitor either does not survive or must return to its dust enshrouded state., Here we add (4) that the progenitor either does not survive or must return to its dust enshrouded state.
+ As the right panel of Fig., As the right panel of Fig.
+ 2. shows. the LBT data already rule out the presence of a massive. evolved star unless it has reconstituted an optically thick. dusty envelope.," \ref{fig:sed} shows, the LBT data already rule out the presence of a massive, evolved star unless it has reconstituted an optically thick, dusty envelope."
+ The present optical limits are somewhat stronger than those for the progenitor. and the near-IR detections already rule out RSGs more massive than 10M...," The present optical limits are somewhat stronger than those for the progenitor, and the near-IR detections already rule out RSGs more massive than $10\,M_\odot$."
+ The total luminosity is now comparable to that of the progenitor and emerges mainly in the mid-IR. but it is also continuing to rapidly fade.," The total luminosity is now comparable to that of the progenitor and emerges mainly in the mid-IR, but it is also continuing to rapidly fade."
+ Recently. Ohsawaetal.(2010) presented late-time AKARI mid-IR observations of the NGC 300 transient that show the transient is again self-enshrouded with an SED that peaks at ~3—4 jim (T.~600 K) and total bolometric luminosity ~5 times the luminosity of the mid-IR progenitor.," Recently, \cite{Ohsawa2010} presented late-time AKARI mid-IR observations of the NGC 300 transient that show the transient is again self-enshrouded with an SED that peaks at $\sim 3-4$ $\mu$ m $T \sim
+600$ K) and total bolometric luminosity $\sim 5$ times the luminosity of the mid-IR progenitor."
+" Let us first consider the possibility that the emission is again due to the progenitor,", Let us first consider the possibility that the emission is again due to the progenitor.
+ With roughly the same luminosity but double the temperature. the dust photosphere must be four times closer to the star. at Apg~40 AU*.. although we can't rule out the presence of a cooler dust component that dominates the bolometric luminosity and peaks in the mid-IR.," With roughly the same luminosity but double the temperature, the dust photosphere must be four times closer to the star, at $R_{BB} \simeq 40$ AU, although we can't rule out the presence of a cooler dust component that dominates the bolometric luminosity and peaks in the mid-IR."
+ If the optical depth is due to à constant velocity wind. this in turn requires decreasing the mass loss rate M (or opacity per unit mass) or increasing the wind velocity v by a factor of 4 relative to the progenitor.," If the optical depth is due to a constant velocity wind, this in turn requires decreasing the mass loss rate $\dot{M}$ (or opacity per unit mass) or increasing the wind velocity $v_w$ by a factor of 4 relative to the progenitor."
+ Producing the near-IR time variability is difficult in this scenario because the characteristic time scale Rgp/vy is 2 years for 100 km/s while the K band flux changed by over a factor of 2 in only 60 days., Producing the near-IR time variability is difficult in this scenario because the characteristic time scale $R_{BB}/v_w$ is $\sim 2$ years for $v_w \simeq 100$ km/s while the $K$ band flux changed by over a factor of 2 in only 60 days.
+ Thus. it seems unlikely that the system has returned to its pre-transient state.," Thus, it seems unlikely that the system has returned to its pre-transient state."
+ The rapid fading strongly suggests that the present emission is à continuation of the transient., The rapid fading strongly suggests that the present emission is a continuation of the transient.
+ Since we are now 800 days post explosion. the emission can no longer be explained as an infrared echo.," Since we are now 800 days post explosion, the emission can no longer be explained as an infrared echo."
+ At this point. echos from the transient peak are produced from a minimum distance of c£/2~70000 AU. and this is simply too distant for dust heated by a transient with a peak luminosity of order 107L.. (Botticellaetal.2009) to produce significant near-IR emission.," At this point, echos from the transient peak are produced from a minimum distance of $ct/2 \simeq 70000$ AU, and this is simply too distant for dust heated by a transient with a peak luminosity of order $10^7 \,L_\odot$ \citep{Botticella2009} to produce significant near-IR emission."
+ While the near-IR emission has roughly decayed at the rate expected from Co decay (1.023 mag/100 days. see Fig. 1))," While the near-IR emission has roughly decayed at the rate expected from $^{56}$ Co decay (1.023 mag/100 days, see Fig. \ref{fig:lc}) )"
+ the drop in the estimated bolometric luminosity is significantly slower. and it is unclear how the positron heating could be efficiently converted to near-IR emission.," the drop in the estimated bolometric luminosity is significantly slower, and it is unclear how the positron heating could be efficiently converted to near-IR emission."
+ The last possibility is shock heating of pre-existing dust., The last possibility is shock heating of pre-existing dust.
+ Botticellaetal.(2009) and Wessonetal.(2010). estimate that dust survived outside of 1000—2000 AU. which would be reached after 800 days by à shock moving at 2000—4000 km/s. Such high velocities were observed in some early line components," \cite{Botticella2009} and \cite{Wesson2010} estimate that dust survived outside of $1000-2000$ AU, which would be reached after 800 days by a shock moving at $2000-4000$ km/s. Such high velocities were observed in some early line components"
+of AL. taken as representative constant values. with radio pulsars corresponding to zero or very weak AZ. such that the disk does not quench the pulsar magnetosphere.,"of $\dot{M}$, taken as representative constant values, with radio pulsars corresponding to zero or very weak $\dot{M}$, such that the disk does not quench the pulsar magnetosphere."
+ The inner radius of the disk must lie al or bevond the light evlinder., The inner radius of the disk must lie at or beyond the light cylinder.
+ Disks around radio pulsars were first proposed by Michel Dessler (1981. 1983) ancl Michel (1988).," Disks around radio pulsars were first proposed by Michel Dessler (1981, 1983) and Michel (1988)."
+ For the ANPs Chatterjee. llernquist Naravan (2000) studied a fall-back disk with a specific time dependent rate M(/) taken to follow a self-similar thin disk solution. which entails a power-law decay of M(1) (Cannizzo. Lee Goodman 1990. Mineshige. Nomoto Shigevama 1993).," For the AXPs Chatterjee, Hernquist Narayan (2000) studied a fall-back disk with a specific time dependent mass-inflow rate $\dot{M}(t)$ taken to follow a self-similar thin disk solution, which entails a power-law decay of $\dot{M}(t)$ (Cannizzo, Lee Goodman 1990, Mineshige, Nomoto Shigeyama 1993)."
+ According to this model some neutron stars go through propeller and accretion (AXP) phases as AL(/) evolves while others. at magnetic fields <3-x1076 become radio pulsars alter a very brief initial accretion phase.," According to this model some neutron stars go through propeller and accretion (AXP) phases as $\dot{M}(t)$ evolves while others, at magnetic fields $<3 \times 10^{12} G$ become radio pulsars after a very brief initial accretion phase."
+ What are the implications of a lall-back disk for radio pulsars?, What are the implications of a fall-back disk for radio pulsars?
+ Marsclen. Lingenlfelter and Rothschild (2001a.b) argued that the presence of a propeller spin-down torque [rom a fall-back disk in conjunction wilh (he magnetic dipole radiation torque may explain the laree discrepancy (Gaensler Frail 2000) between the real (kinematic) age and the “characteristic” age P/2P corresponding to pure dipole spin-down for the pulsar D1757-24. and that -—-5AHR age discrepancies can be explained in a similar wav.," Marsden, Lingenfelter and Rothschild (2001a,b) argued that the presence of a propeller spin-down torque from a fall-back disk in conjunction with the magnetic dipole radiation torque may explain the large discrepancy (Gaensler Frail 2000) between the real (kinematic) age and the ""characteristic"" age $P/2\dot{P}$ corresponding to pure dipole spin-down for the pulsar B1757-24, and that -SNR age discrepancies can be explained in a similar way."
+ Menou. Perna and Hernquist (2001b) showed that this combination of dipole and propeller torques can explain (he braking indices η«3 of [ive voung pulsars.," Menou, Perna and Hernquist (2001b) showed that this combination of dipole and propeller torques can explain the braking indices $n<3$ of five young pulsars."
+ In this letter we apply the combined spin-down torque. using (the model of Menou. Perna and Iernquist (2001b). to the distribution of pulsars in the P— diagram.," In this letter we apply the combined spin-down torque, using the model of Menou, Perna and Hernquist (2001b), to the distribution of pulsars in the $P-\dot{P}$ diagram."
+ Presence of torques other than dipole radiation torques. in particular torques that might arise from ihe ambient medium near à voung pulsar in a supernova remnant was proposed first bv Yusilov οἱ al. (," Presence of torques other than dipole radiation torques, in particular torques that might arise from the ambient medium near a young pulsar in a supernova remnant was proposed first by Yusifov et al. ("
+1995) to explain qualitatively the measured braking indices »<3 of voung pulsars. Che discrepancy between real aud characteristic ages and the distribution of pulsars in the P— diagram.,"1995) to explain qualitatively the measured braking indices $n<3$ of young pulsars, the discrepancy between real and characteristic ages and the distribution of pulsars in the $P-\dot{P}$ diagram."
+ Gvaramadze (2001) invokes torques [rom cireumstellar chimps in the SATB to derive a πιο age lor PSR D1509-58 in agreement with the SNR age.," Gvaramadze (2001) invokes torques from circumstellar clumps in the SNR to derive a ""true"" age for PSR B1509-58 in agreement with the SNR age."
+ In Section 2. we explore (he combined torque model. using the canonical magnetic dipole torque together wilh the model of Menou. Perna and Hernq(quist (2001b) for the propeller spin-down torque. with a constant mass inflow rate M.," In Section 2, we explore the combined torque model, using the canonical magnetic dipole torque together with the model of Menou, Perna and Hernquist (2001b) for the propeller spin-down torque, with a constant mass inflow rate $\dot{M}$."
+ Evolutionary tracks in the P— diagram are presented in Section 3 along with analytical expressions for timespaus of various phases. braking indices and other properties of the tracks.," Evolutionary tracks in the $P-\dot{P}$ diagram are presented in Section 3 along with analytical expressions for timespans of various phases, braking indices and other properties of the tracks."
+ The P— diagram is divided into strips delineated by tracks of constant magnetic field and mass inflow rate., The $P-\dot{P}$ diagram is divided into strips delineated by tracks of constant magnetic field and mass inflow rate.
+ Each strip is divided into period bins and a histogram is constructed for the number of pulsars in each period bin., Each strip is divided into period bins and a histogram is constructed for the number of pulsars in each period bin.
+ A curve for the expected number of pulsars as a function of period is caleulated according to the model ancl compared with the histogram for each strip., A curve for the expected number of pulsars as a function of period is calculated according to the model and compared with the histogram for each strip.
+ The results are discussed in Section 4., The results are discussed in Section 4.
+In addition. the correlator (WIDAR) requires several major software elements:,"In addition, the correlator (WIDAR) requires several major software elements:"
+Accreting isolated neutron stars (AINS) were predicted 40 vears ago by ancl independentIv by ?.,Accreting isolated neutron stars (AINS) were predicted 40 years ago by and independently by .
+ In early 90s there was some enthusiasm clue to the launch of the ROSAT satellite. which was expected to find many sources of this kind (?).," In early 90s there was some enthusiasm due to the launch of the ROSAT satellite, which was expected to find many sources of this kind ."
+. Several populational studies have been mace (??7?777)., Several populational studies have been made .
+. Llowever. it came out that Αλ. if they exist. are very elusive (Colpi et al.," However, it came out that AINS, if they exist, are very elusive (Colpi et al."
+ 1998)., 1998).
+ The main reason is that initial (kick) velocities of NSs appeared. to be significantly larger. than it have been thought before. (2).," The main reason is that initial (kick) velocities of NSs appeared to be significantly larger, than it have been thought before ."
+.. Initial. guess that the number of Accretors is small due to low luminosity of hieh-velocity NSs was shown to be wrong., Initial guess that the number of Accretors is small due to low luminosity of high-velocity NSs was shown to be wrong.
+ In a detailed study by (hereafter Paper L) it was shown that INS with Crab-like initial parameters and constant magnetic fields spend all their lives as Ejectors (we follow the classification summarized. in 2) if their initial velocities are <100 km s, In a detailed study by (hereafter Paper I) it was shown that INS with Crab-like initial parameters and constant magnetic fields spend all their lives as Ejectors (we follow the classification summarized in ) if their initial velocities are $\ga 100$ km $^{-1}$.
+ ‘Phen. the fraction of Aceretors was mainlv determined. by the fraction of velocity NSs.," Then, the fraction of Accretors was mainly determined by the fraction of low-velocity NSs."
+ Up to the very end. of 90s. it was believed that the wast majority of NSs are born similar to the Crab pulsar.," Up to the very end of 90s, it was believed that the wast majority of NSs are born similar to the Crab pulsar."
+ Le. that they have short initial spin periods (from milliseconds to few tens of millisecond) anc magnetic fields 2~1017 €. Now it is believed. that about one hall of NSs have different initial properties," I.e., that they have short initial spin periods (from milliseconds to few tens of millisecond) and magnetic fields $B\sim 10^{12}$ G. Now it is believed, that about one half of NSs have different initial properties"
+In recent years coherent. fast pulsations have been discovered in six Low Mass X-ray Binaries (LMXBs) leading to the conclusion that they contain rapidly spinning (~2—5 ms). weakly magnetic neutron stars (see e.g. Chakrabarty 2004).,"In recent years coherent, fast pulsations have been discovered in six Low Mass X-ray Binaries (LMXBs) leading to the conclusion that they contain rapidly spinning $\sim 2-5$ ms), weakly magnetic neutron stars (see e.g. Chakrabarty 2004)."
+ X-ray light curves obtained with RXTE also revealed coherent oscillations during Type I X-ray bursts (ie. thermonuclear explosions on the neutron star surface) m fourteen LMXBs (e.g. van der Klis 2000)., X-ray light curves obtained with RXTE also revealed coherent oscillations during Type I X-ray bursts (i.e. thermonuclear explosions on the neutron star surface) in fourteen LMXBs (e.g. van der Klis 2000).
+ These findings directly confirmed evolutionary models that link the LMXB neutron stars to those of millisecond radio pulsars. the former being the progenitors of the latter (e.g. Bhattacharya van den Heuvel 1991).," These findings directly confirmed evolutionary models that link the LMXB neutron stars to those of millisecond radio pulsars, the former being the progenitors of the latter (e.g. Bhattacharya van den Heuvel 1991)."
+ A large fraction of LMXBs accrete matter at very high rates (and therefore shine as bright X-ray sources) only sporadically., A large fraction of LMXBs accrete matter at very high rates (and therefore shine as bright X-ray sources) only sporadically.
+ Among these systems are Soft X-ray Transients (SXTs). hosting old neutron stars (for a review see Campana et al.," Among these systems are Soft X-ray Transients (SXTs), hosting old neutron stars (for a review see Campana et al."
+ 1998)., 1998).
+ These systems alternate periods (weeks to months) of high X-ray luminosity with long (1-10 years) intervals of quiescence in which the X-ray luminosity drops by some 5-6 orders of magnitude., These systems alternate periods (weeks to months) of high X-ray luminosity with long (1–10 years) intervals of quiescence in which the X-ray luminosity drops by some 5–6 orders of magnitude.
+ Concerning the optical properties. SXTs in outburst are relatively bright and their emission is dominated by X-ray reprocessing in the accretion dise surrounding the neutron star.," Concerning the optical properties, SXTs in outburst are relatively bright and their emission is dominated by X-ray reprocessing in the accretion disc surrounding the neutron star."
+ In quiescence their optical lumimosity drops by as much as 6-7 mag. thereby allowing the late type compantor to be detected and studied.," In quiescence their optical luminosity drops by as much as 6–7 mag, thereby allowing the late type companion to be detected and studied."
+ Weak. residual optical emission is still present in quiescence. usually at the ~10-30% level in V. that is usually aseribed to a dim. viscously - heated accretion disc (van Paradijs McClintock 1995).," Weak, residual optical emission is still present in quiescence, usually at the $\sim 10-30\%$ level in V, that is usually ascribed to a dim, viscously - heated accretion disc (van Paradijs McClintock 1995)."
+ The large luminosity variations of SXTs correspond to vast changes t the accretion rate onto the collapsed objects., The large luminosity variations of SXTs correspond to vast changes in the accretion rate onto the collapsed objects.
+ These provide an opportunity to explore a range of accretion regimes onto the neutron star that are inaccessible to persistent sources., These provide an opportunity to explore a range of accretion regimes onto the neutron star that are inaccessible to persistent sources.
+ While in outburst accretion onto the neutron star surface takes place in à manner similar to that of persistent LMXBs. whereas when the accretion rate decreases. matter is halted at the magnetospheric boundary by a centrifugal barrier generated by the fast rotating neutron star magnetosphere (a.k.a.," While in outburst accretion onto the neutron star surface takes place in a manner similar to that of persistent LMXBs, whereas when the accretion rate decreases, matter is halted at the magnetospheric boundary by a centrifugal barrier generated by the fast rotating neutron star magnetosphere (a.k.a."
+ propeller effect: Illarionov Sunyaev 1975)., propeller effect; Illarionov Sunyaev 1975).
+ If the mass inflow rate decreases further. the neutron star can enter the radio pulsar regime and begin to expel the inflowing material from the companion star (Shaham Tavani 1991: Stella et al.," If the mass inflow rate decreases further, the neutron star can enter the radio pulsar regime and begin to expel the inflowing material from the companion star (Shaham Tavani 1991; Stella et al."
+ 1994; Campana et al., 1994; Campana et al.
+ 1995; Burdert et al., 1995; Burderi et al.
+ 2001)., 2001).
+ Some progress has recently been made in the study of the outburst decay and quiescent state of some of these systems (Asai et al., Some progress has recently been made in the study of the outburst decay and quiescent state of some of these systems (Asai et al.
+ 1996). yet the interpretation of the quiescent optical and X-ray observations is ambiguous.," 1996), yet the interpretation of the quiescent optical and X-ray observations is ambiguous."
+ The quiescent X-ray spectrum is well fit by the sum of a soft component (usually a neutron star atmosphere model with temperature &T.~0.1—0.3 keV) plus a hard tail power law with photon index in the range 1-2.," The quiescent X-ray spectrum is well fit by the sum of a soft component (usually a neutron star atmosphere model with temperature $k\,T\sim 0.1-0.3$ keV) plus a hard tail power law with photon index in the range 1–2."
+ There are cases in which only one of the two components is detected., There are cases in which only one of the two components is detected.
+ Different interpretations have been put forward., Different interpretations have been put forward.
+ Menou et al. (, Menou et al. (
+1999) suggested the presence of an Advection Dominated Accretion Flow (ADAF) dominating the residual optical emission: some of the inflowing matter would leak through the magnetosphere and aecrete onto the neutron star surface. giving rise to the soft X-ray component.,"1999) suggested the presence of an Advection Dominated Accretion Flow (ADAF) dominating the residual optical emission: some of the inflowing matter would leak through the magnetosphere and accrete onto the neutron star surface, giving rise to the soft X-ray component."
+ Campana et al. (, Campana et al. (
+1998a. b. see also Campana Stella 2000) suggestec instead that the soft component arises from the cooling of the neutron star surface heated during the repeated outburst episodes (see also Brown et al.,"1998a, b, see also Campana Stella 2000) suggested instead that the soft component arises from the cooling of the neutron star surface heated during the repeated outburst episodes (see also Brown et al."
+ 1998; Rutledge et al., 1998; Rutledge et al.
+ 1999) anc the hard tail arises from shock emission powered by an active radio pulsar., 1999) and the hard tail arises from shock emission powered by an active radio pulsar.
+ In this case (part of) the residual quiescent optical flux may come from the shock emission in the interactior between the radio pulsar wind and the matter inflowing from the companion., In this case (part of) the residual quiescent optical flux may come from the shock emission in the interaction between the radio pulsar wind and the matter inflowing from the companion.
+ We therefore decided to utilise Doppler tomography i order to study the kinematics and geometry of matter i7 such interacting binaries., We therefore decided to utilise Doppler tomography in order to study the kinematics and geometry of matter in such interacting binaries.
+ This would allow us to derive two-dimensional velocity maps of the optical lme-emitting regions. by combining spectra taken at different orbital phases (Marsh," This would allow us to derive two-dimensional velocity maps of the optical line-emitting regions, by combining spectra taken at different orbital phases (Marsh"
+instabilities in shell livdrogen burning ou the white dwiut.,instabilities in shell hydrogen burning on the white dwarf.
+ Our estimate of 10? 105MIML. of hiydrosen Όσα is 12 orders of magnitude lower than the minima. envelope mass necessary to sustain equilibrium hydrogen mining (Sala&Heornauz2005)., Our estimate of $10^{-9}$ $10^{-8}$ $_\odot$ of hydrogen burned is 1–2 orders of magnitude lower than the minimum envelope mass necessary to sustain equilibrium hydrogen burning \citep{sala05}.
+.. A late-time increase in he cuvelope hydrogen burning aud the associated mass ejection has also been invoked to explain a secondary oak iu V2362 (νο (Lvuchetal.2008)., A late-time increase in the envelope hydrogen burning and the associated mass ejection has also been invoked to explain a secondary peak in V2362 Cyg \citep{lynch08}.
+ Therefore. we conclude that the transition. plase flares are likely caused by sudden mereases m otherwise steady hydrogen miming in the white cawart cuvelope.," Therefore, we conclude that the transition phase flares are likely caused by sudden increases in otherwise steady hydrogen burning in the white dwarf envelope."
+ By analvsis of flare timines. we have shown that the intervals between consecutive fares merease as a powcr-aw function of time.," By analysis of flare timings, we have shown that the intervals between consecutive flares increase as a power-law function of time."
+ However. the length of the fares rendus essentially constaut divine the transition plase sugecsting that ouly the trigger of the flare is affected.," However, the length of the flares remains essentially constant during the transition phase suggesting that only the trigger of the flare is affected."
+ The power-law exponent is ecuerally the same for all uovae. but the normalization is larger for novae with sanaller rates of decline.," The power-law exponent is generally the same for all novae, but the normalization is larger for novae with smaller rates of decline."
+ The primary parameter that controls the rate of decline is the mass of the white dwarf: faster novae occur on more Inassive white dwarfs (Priduik&Kovetz1995)., The primary parameter that controls the rate of decline is the mass of the white dwarf: faster novae occur on more massive white dwarfs \citep{prialnik95}.
+. It is then natural to expect that with a more massive white dwarf the intervals between hvdrogeu-buruig flashes will be shorter aud also that the flare will last a shorter amount of the time., It is then natural to expect that with a more massive white dwarf the intervals between hydrogen-burning flashes will be shorter and also that the flare will last a shorter amount of the time.
+ llowever. as we have demoustrated on V1191 Aql aud GIy Por. which have essentially the same light curve but substantially different iutervals between flares. there is one or nore paranmoeters that also influence the properties of the flares.," However, as we have demonstrated on V1494 Aql and GK Per, which have essentially the same light curve but substantially different intervals between flares, there is one or more parameters that also influence the properties of the flares."
+ We suspect that chemical composition ofthe white dwarf anc its lvdrogen-burning euvelope. as well as its inass or the mass accretion rate nüeght plav a role.," We suspect that chemical composition of the white dwarf and its hydrogen-burning envelope, as well as its mass or the mass accretion rate might play a role."
+ The eradual increase of intervals between individual rebrightcnings is an unusual phenomenon in astronomy., The gradual increase of intervals between individual rebrightenings is an unusual phenomenon in astronomy.
+ We have been able to fd only one other example of suchbehavior: the X-ray aud. optical light curve Iie accretion-drziven iuillisecond pulsar SAN JisU0s.1orWe3658 during its 2000 outburst (AVijuandsetal.2001)., We have been able to find only one other example of such behavior: the X-ray and optical light curve of the accretion-driven millisecond pulsar SAX J1808.4–3658 during its 2000 outburst \citep{wijnandsetal01}.
+. plot the N-rav light curve in the right panel of Figure J, We plot the X-ray light curve in the right panel of Figure \ref{fig:lc}.
+ ILowever. duiugs the 2002 outburst the flares were more regular. with no apparent increase in finie separation (παπάς2001).," However, during the 2002 outburst the flares were more regular, with no apparent increase in time separation \citep{wijnands04}."
+. Iun other areas of physics the phenomenon of logarithmic periodicity occurs as a correction to a power-law behavior and suggests presence of discrete scale invariance (DSI.Sornette2001.," In other areas of physics the phenomenon of logarithmic periodicity occurs as a correction to a power-law behavior and suggests presence of discrete scale invariance \citep[DSI,][p.~276]{sornette04}."
+p.276). A wellknown example of DSI is the triadic Cantor set which is ecometrically identical to itself onlv under magnifications equal to powers of 3 (Soructte1998)., A well-known example of DSI is the triadic Cantor set which is geometrically identical to itself only under magnifications equal to powers of 3 \citep{sornette98}.
+. Among examples of DST aud logarithmic periodicity observed iu nature we specifically mention the rate of acoustic Cluissions as a function of time before the rupture of pressure vessels (Zhou&Sornette20022a).. and eround-water chemical anomalies precediug the 1995 Iobe earthquake (Johansenetal.2000:Zhou&Sornctte2002b).," Among examples of DSI and logarithmic periodicity observed in nature we specifically mention the rate of acoustic emissions as a function of time before the rupture of pressure vessels \citep{zhou02a}, and ground-water chemical anomalies preceding the 1995 Kobe earthquake \citep{johansen00,zhou02b}."
+. Theoretical models. which exhibit spontaneous cineregcuce of DSI. iuclude earthquake aftershock simulations based on cellular automata| (Lee&Sornette2000). aud rupture studies with a thermal fuse model (Johansen&Sornette1998).," Theoretical models, which exhibit spontaneous emergence of DSI, include earthquake aftershock simulations based on cellular automata \citep{lee00} and rupture studies with a thermal fuse model \citep{johansen98}."
+. Log-periodie echoing was also found during a critical relativistic eravitational collapse of a massless scalar field (Choptuik1993:Clop-tuiketal. 2003).," Log-periodic echoing was also found during a critical relativistic gravitational collapse of a massless scalar field \citep{choptuik93,choptuik03}."
+". Finally, we poiut out that there is uo single uuderbliug mechauisgu for the DSI iu. ecneral (Sornette2001.p.160).."," Finally, we point out that there is no single underlying mechanism for the DSI in general \citep[p.~160]{sornette04}."
+ We are grateful to Ik. Z. Stanek. T. A. Thompson. R. W. Posse. €. S. kochauck. L. Bildsten ancl D. Soruette for discussions. encouragement. advice on astroplivsical processes aud critical reading of the manuscript.," We are grateful to K. Z. Stanek, T. A. Thompson, R. W. Pogge, C. S. Kochanek, L. Bildsten and D. Sornette for discussions, encouragement, advice on astrophysical processes and critical reading of the manuscript."
+ We thank the anomvinous referce for cohunents that improved the Letter., We thank the anonymous referee for comments that improved the Letter.
+ We acknowledge the observations acquired from the AAVSO International Database. the AFOEV database. operated at CDS. France. aud from the VSNET database.," We acknowledge the observations acquired from the AAVSO International Database, the AFOEV database, operated at CDS, France, and from the VSNET database."
+canonical Freciman (1970) value of 21.65 Banag 2 Qvhich is about 20.5 nag 2).,canonical Freeman (1970) value of 21.65 $B$ -mag $^{-2}$ (which is about 20.5 $R$ -mag $^{-2}$ ).
+ Thus. M32: would just about qualify as a bulec-domunated LSB disk galaxy (Deijersbergeu. de Blok. van der Πιτ 1999: Turper Dothuu 1997).," Thus, M32 would just about qualify as a bulge-dominated LSB disk galaxy (Beijersbergen, de Blok, van der Hulst 1999; Impey Bothun 1997)."
+ It is. however. because of it’s size aud magnitude. a dwarf galaxy.," It is, however, because of it's size and magnitude, a dwarf galaxy."
+ The scale-Ieugth of the disk is 130% or 0.5 Ipc., The scale-length of the disk is $\arcsec$ or 0.5 kpc.
+" There is. however. possible evidence of a disturbance im the very outer profile: it turus downward from an expoucutial xofile at around 250""."," There is, however, possible evidence of a disturbance in the very outer profile; it turns downward from an exponential profile at around $\arcsec$."
+" That is. there appears to be a lack of light. relative to the exponeutial part of the xofile. bevoud ~250""."," That is, there appears to be a lack of light, relative to the exponential part of the profile, beyond $\sim250\arcsec$."
+ This behavior is visible in both he data of I&ent (1087) and also that of Choi et ((2002)., This behavior is visible in both the data of Kent (1987) and also that of Choi et (2002).
+" It is accompanied by a marked change iu the vchavior of the ellipticity profile of Choi et ((2002). Hattening (or even decreasing slightly) bevoud 250""."," It is accompanied by a marked change in the behavior of the ellipticity profile of Choi et (2002), flattening (or even decreasing slightly) beyond $\arcsec$."
+ This nay be a sien that material has been stripped away roni the outer disk. although it is stressed that this conchision is larecly speculative.," This may be a sign that material has been stripped away from the outer disk, although it is stressed that this conclusion is largely speculative."
+ However. if true. this irnover would be biasng the surface brightuess profile fit. ποιο the disk scale-leneth appear shorter than it was before tidal-stripping conuueuced.," However, if true, this turnover would be biasing the surface brightness profile fit, making the disk scale-length appear shorter than it was before tidal-stripping commenced."
+" Truucatius the (equally spaced) profile at 220"". to avoid the potentiallv stripped outer disk. and thereby (possibly) sampling only the original undisturbe profile gives /;233367 (1.27 kpc) and jty=22.12 for the disk."," Truncating the (equally spaced) profile at $\arcsec$, to avoid the potentially stripped outer disk, and thereby (possibly) sampling only the original undisturbed profile gives $h=336\arcsec$ (1.27 kpc) and $\mu_0=22.42$ for the disk."
+" For the bulge. r.=29"" and s=1.99 (see Figure 5: using the profile which is logaritlanically spaced Dooiu radius gives 7= 2.08)."," For the bulge, $r_{\rm e}=29\arcsec$ and $n=1.99$ (see Figure \ref{fig2}; using the profile which is logarithmically spaced in radius gives $n=2.08$ )."
+" Tere the value of log B/D=-0.09. and (c.// 0.09. which are again uot unreasonable,"," Here the value of $\log B/D$ =-0.09, and $r_{\rm e}/h$ =0.09, which are again not unreasonable."
+" Am A ""only ος). fails to provide a convincing fit to this trucated radial range.", An $r^{1/n}$ -only model fails to provide a convincing fit to this truncated radial range.
+ The surface brightness profile of M32 can be modelled remarkably well as a combination of an rv Ἰπήσο aud an expoucutial disk of stars., The surface brightness profile of M32 can be modelled remarkably well as a combination of an $r^{1/n}$ bulge and an exponential disk of stars.
+ The ellipticity profile additionally supports this iuterpretation of the data., The ellipticity profile additionally supports this interpretation of the data.
+ ILowever. could. instead. ΑΣ surface brightuess profile be πο disturbed that the outer exponential euvelope is actually due to material pulled off from what was once a one-compoucut +t n= DYellipticalgalary?," However, could, instead, M32's surface brightness profile be so disturbed that the outer exponential envelope is actually due to material pulled off from what was once a one-component $r^{1/4}$ $n=4$ ) elliptical galaxy?"
+ Thissccnarioappoarsinlibely forthe fFollepitai queasmay, This scenario appears unlikely for the following reasons.
+ Firstly. somewhat persuasive evidence comes frou the the central velocity dispersion of M32 which has beeu iieasured to be TELLIN kins | (van der Mavel ct 11998).," Firstly, somewhat persuasive evidence comes from the the central velocity dispersion of M32 which has been measured to be $\pm$ 10 km $^{-1}$ (van der Marel et 1998)."
+ This feure is in good agreement with the average value of TL lan + obtained from umnerous estimates listed in It also agrees with Gebhardt et al, This figure is in good agreement with the average value of 74 km $^{-1}$ obtained from numerous estimates listed in It also agrees with Gebhardt et al.
+« (2000) estimate of 75 lau | for the luninosity weighted velocity dispersion within one effective radius.,'s (2000) estimate of 75 km $^{-1}$ for the luminosity weighted velocity dispersion within one effective radius.
+ Receutly. it has beeu discovered that the central velocity dispersion of a bulee correlates strongly (r= 0.8) with the shape of the bulee liebt profile (as measured with the Sévrsic index: Gral et 22001b: Craham 2002).," Recently, it has been discovered that the central velocity dispersion of a bulge correlates strongly $r=0.8$ ) with the shape of the bulge light profile (as measured with the Sérrsic index: Graham et 2001b; Graham 2002)."
+ Calaxies with measured velocity dispersions less than 100 km s| are observed to have values of 1<n<2., Galaxies with measured velocity dispersions less than 100 km $^{-1}$ are observed to have values of $1<n<2$.
+ The measured value of D=1.5 for the bulee of N32 is thus exactly what one would expect frou its dynamics. not »—I.," The measured value of $n=1.5$ for the bulge of M32 is thus exactly what one would expect from its dynamics, not $n=4$."
+ Bulges (which includes elliptical galaxies) with values of 5)2L have velocity dispersious typically ercater than 100 kan st., Bulges (which includes elliptical galaxies) with values of $n\gtrsim4$ have velocity dispersions typically greater than 100 km $^{-1}$.
+ The second. related. line of reasoning comes from ADJ2/8 SAIBIT mass.," The second, related, line of reasoning comes from M32's SMBH mass."
+ Frou ground based kinematical cata. Toury (198L7) predicted a SMDBII mass of «10937. at the center of M32.," From ground based kinematical data, Tonry (1984,7) predicted a SMBH mass of $\times10^6 M_{\sun}$ at the center of M32."
+ This pioneering work has been confirmed with IIST data. οι which van der Marl ct ((1998) derived a SMBIT mass of 140.7) «109AZ... and amore recently Joseph et (2001) find a mass of ΠΛΗ...," This pioneering work has been confirmed with HST data, from which van der Marel et (1998) derived a SMBH mass of $\pm$ $\times10^6 M_{\sun}$, and more recently Joseph et (2001) find a mass of $\times10^6 M_{\sun}$."
+ The velocity dispersion (0) and SMDII ass (Adj) combination of M32 is kuown to fall ou the logMij-log a relation for (non-disturbed) ellipticals and bulges (Ferrarese Alerritt 2000: Gebhardt et 22000: [οι Ferrarese 200142)., The velocity dispersion $\sigma$ ) and SMBH mass $M_{\rm bh}$ ) combination of M32 is known to fall on the $\log M_{\rm bh}$ $\log\sigma$ relation for (non-disturbed) ellipticals and bulges (Ferrarese Merritt 2000; Gebhardt et 2000; Merritt Ferrarese 2001a).
+ Thus. uuless the process of tidal stripping modifies both the central velocity dispersion aud the SAIBIT mass in such a way that it preserves the logAfjp-log o relation. the ceutral structure and cvnamics of M32 are likely to be that of the original uuudisturbed) ealaxy.," Thus, unless the process of tidal stripping modifies both the central velocity dispersion and the SMBH mass in such a way that it preserves the $\log M_{\rm bh}$ $\log\sigma$ relation, the central structure and dynamics of M32 are likely to be that of the original undisturbed) galaxy."
+ Indeed. due to the higher densities at smaller radii. the central velocity dispersion and hence iuner nass distribution and therefore incr helt profile are expected to remain largely unaffected by the outer stripping process.," Indeed, due to the higher densities at smaller radii, the central velocity dispersion and hence inner mass distribution and therefore inner light profile are expected to remain largely unaffected by the outer stripping process."
+ Taken with the result iu the previous paragraph. this strougly suggests that the m=1.52.0 bulge profile is the original shape of the bulge.," Taken with the result in the previous paragraph, this strongly suggests that the $n=1.5-2.0$ bulge profile is the original shape of the bulge."
+" Figure 6. shows the logAf,-logy relation for bulges.", Figure \ref{fig3} shows the $\log M_{\rm bh}$ $\log n$ relation for bulges.
+ It is a variant of the relation between SAIBIT mass and central bulge concentration (C.[1/3|) shown iu. Cralam et ((2001a). (, It is a variant of the relation between SMBH mass and central bulge concentration $C_{r_{\rm e}}[1/3]$ ) shown in Graham et (2001a). (
+"The parameter C,(1/3) is a monotonicly increasing function of s. Trujillo et 22001.)","The parameter $C_{r_{\rm e}}(1/3)$ is a monotonicly increasing function of $n$, Trujillo et 2001.)"
+ If ND32 did ever have an rt?t profile. then (from Figure 6)) its SMDII mass should have ouce been aud should still be ~LOSAL.. and covtainly greater than ~LOTAL...," If M32 did ever have an $r^{1/4}$ profile, then (from Figure \ref{fig3}) ) its SMBH mass should have once been and should still be $\sim10^8 M_{\sun}$, and certainly greater than $\sim10^7 M_{\sun}$."
+ Civen the bulee value of #=1.5 agrees with the actual SMDII lnass estimate (aud with MD2's velocity dispersion). the stars composing the outer exponential euvelope are alinost certainly an excess (relative to the bulge) which have not come from a reshaped bulgc?.," Given the bulge value of $n=1.5$ agrees with the actual SMBH mass estimate (and with M32's velocity dispersion), the stars composing the outer exponential envelope are almost certainly an excess (relative to the bulge) which have not come from a reshaped ."
+. All of this is uot to sav that the outer profile of N32 has not been disrupted. Iudeed. the deficit of stars in the outer profile. causing the cdowmware kink (or break) in the disk. have been due to eravitational stripping by M3.," All of this is not to say that the outer profile of M32 has not been disrupted, Indeed, the deficit of stars in the outer profile, causing the downward kink (or break) in the disk, may have been due to gravitational stripping by M31."
+ a visille iu the B- aud £-band profiles of Choi et ((2002) aud lence less likely to be a systematic error in the profile extraction technique of EKeut (1987)., This deficit is also visible in the $B$ - and $I$ -band profiles of Choi et (2002) and hence less likely to be a systematic error in the profile extraction technique of Kent (1987).
+ Tudeed. the recently reported tidal stream of moetalaich stars around MOL is thought to have likely come frou ADJ2 and/or NGC 205 (Ibata et 22001).," Indeed, the recently reported tidal stream of metal-rich stars around M31 is thought to have likely come from M32 and/or NGC 205 (Ibata et 2001)."
+ The preseuce of a disk may also explain the iuteriiediate-age (5-15 Cx) stellar population iu M32 CCrilhuair et 11996: Davidge 2000)., The presence of a disk may also explain the intermediate-age (5-15 Gyr) stellar population in M32 Grillmair et 1996; Davidge 2000).
+ Perhaps à conuuecut ou nomenclature would be beneficial rere., Perhaps a comment on nomenclature would be beneficial here.
+ Ivormendy Cebhardt (2001) used the words ‘compact’ and fuffv when reterrine to bulges., Kormendy Gebhardt (2001) used the words `compact' and `fluffy' when referring to bulges.
+ It should )o noted that this has no refereuce at all to varving profile shape’ (or equivalently couceutration. according to the nathematical definition given im Trujillo et 22001. their equation 5 and 6).," It should be noted that this has no reference at all to varying profile `shape' (or equivalently `concentration', according to the mathematical definition given in Trujillo et 2001, their equation 5 and 6)."
+ That is. no reference to departures Toni the rt law. or structuralhomology. are implied by," That is, no reference to departures from the $r^{1/4}$ law, or structuralhomology, are implied by"
+Also the temporal indices a of Type 0. a! of Type Ia. and a! of Type II show significant similarities (see Table 3 for a summary of the relative K-S test probabilities).,"Also the temporal indices $\alpha_1^{0}$ of Type 0, $\alpha_2^{Ia}$ of Type Ia, and $\alpha_3^{II}$ of Type II show significant similarities (see Table \ref{tab_correlazioni} for a summary of the relative K-S test probabilities)."
+ We therefore group these segments as the of the GRB light curves., We therefore group these segments as the of the GRB light curves.
+ The photon indices [ have been obtained by extracting high-quality spectra within several subsequent time intervals cander the condition that 2000 photons are included in each time πεterval., The photon indices $\Gamma$ have been obtained by extracting high-quality spectra within several subsequent time intervals under the condition that $2000$ photons are included in each time interval.
+ Then. these spectra were fit by an absorbed power law uovith variable Λη.," Then, these spectra were fit by an absorbed power law with variable $N_H$ ."
+ The behaviour of the spectral index 6=T-| as a function of time for each Type of light curve Is portrayed Fig. 2.., The behaviour of the spectral index $\beta=\Gamma-1$ as a function of time for each Type of light curve is portrayed in Fig. \ref{BETAtype}.
+ We observe that for Type 0 light curves the spectral dex is almost unchanged for the full duration of the event apart from an initial evolution., We observe that for Type 0 light curves the spectral index is almost unchanged for the full duration of the event apart from an initial evolution.
+ The same situation occurs for Type la light curves., The same situation occurs for Type Ia light curves.
+ Type Ib shows variability in the spectral dices prior to the first break time., Type Ib shows variability in the spectral indices prior to the first break time.
+ This behaviour ts analogous to Type II light curves. which show variability before the first break time and no significant variations before and after the second break times.," This behaviour is analogous to Type II light curves, which show variability before the first break time and no significant variations before and after the second break times."
+ The similar behaviour of the spectral indices of the steep decays of Types Ib and II enforces their association. established in the previous section.," The similar behaviour of the spectral indices of the steep decays of Types Ib and II enforces their association, established in the previous section."
+ The other segments do not show any significant spectral evolution. and they are all consistent with B~1 (Butler2007)..," The other segments do not show any significant spectral evolution, and they are all consistent with $\tilde{\beta} \sim 1$ \citep{2007ApJ...656.1001B}."
+ Type II light curves tend to converge to a much narrower distribution of B than in other cases., Type II light curves tend to converge to a much narrower distribution of $\beta$ than in other cases.
+ In what follows. we consider the mean value of the spectral indices of all the spectra covering the same time interval as the spectral index of each segment of each light curve.," In what follows, we consider the mean value of the spectral indices of all the spectra covering the same time interval as the spectral index of each segment of each light curve."
+ This 15 an excellent approximation in all cases but the steep decays. where the average value is not always representative of the observed trend.," This is an excellent approximation in all cases but the steep decays, where the average value is not always representative of the observed trend."
+ The energy output of the different types of X-ray light curves has been calculated as the integral of the best-fitting function over the observing time. and isotropic emission Is assumed.," The energy output of the different types of X-ray light curves has been calculated as the integral of the best-fitting function over the observing time, and isotropic emission is assumed."
+ A K-correction (Bloometal.2001) has been applied in order to evaluate the energies in a common rest frame energy band (1.9—11.2 keV)., A K-correction \citep{2001AJ....121.2879B} has been applied in order to evaluate the energies in a common rest frame energy band $1.9-11.2$ keV).
+" For the prompt emission isotropic energy E;,,. we refer to Amatietal.(2008) and Krimmetal. (2009).. and to Navaetal.(2010). for those GRBs observed by Fermi//GBM whose prompt emission spectrum has been fit with a Band function."," For the prompt emission isotropic energy $E_{iso}$, we refer to \citet{2008MNRAS.391..577A} and \citet{2009ApJ...704.1405K}, and to \citet{2010arXiv1012.3968N} for those GRBs observed by /GBM whose prompt emission spectrum has been fit with a Band function."
+ Figure 3., Figure \ref{etype}.
+".a shows the distribution of the total energy output of the X-ray light curves E, as a function of Ej.", .a shows the distribution of the total energy output of the X-ray light curves $E_x$ as a function of $E_{iso}$.
+ It shows a moderate. positive correlation (with a Spearman rank p=0.56. see also Willingaleetal. 2007)).," It shows a moderate, positive correlation (with a Spearman rank $\rho=0.56$, see also \citealt{2007ApJ...662.1093W}) )."
+ We overplotted in Fig. 3., We overplotted in Fig. \ref{etype}.
+.a the total fluence in the XRT energy band (0.3—10 keV) as a function of the promptemission fluence detected by BAT (15—150 keV) for the complete sample of long GRBs observed by XRT until December 2010, .a the total fluence in the XRT energy band $0.3-10$ keV) as a function of the promptemission fluence detected by BAT $15-150$ keV) for the complete sample of long GRBs observed by XRT until December 2010
+In some cases. a dominant discrete scatterer may be common to sight-lines to a number of pulsars.,"In some cases, a dominant discrete scatterer may be common to sight-lines to a number of pulsars."
+ In. such à. case. characterization of the scattering region can be attempted using the data on the four observables which may be available for only a few pulsars and the results could. be used in determining distances to the other pulsars with nearby sight-lines.," In such a case, characterization of the scattering region can be attempted using the data on the four observables which may be available for only a few pulsars and the results could be used in determining distances to the other pulsars with nearby sight-lines."
+ This relaxes the requirement that all four observables be available for the rest of the pulsars in the set and. in fact. the knowledge of DAL ancl temporal scatter xoadening alone sullices for distance determination.," This relaxes the requirement that all four observables be available for the rest of the pulsars in the set and, in fact, the knowledge of DM and temporal scatter broadening alone suffices for distance determination."
+ We illustrate this possibility. by considering the Gun nebula region as à common scatterer for a number of pulsar sight-ines., We illustrate this possibility by considering the Gum nebula region as a common scatterer for a number of pulsar sight-lines.
+ ‘Vavlor Cordes (1993). do include. the enhanced electron density in the Gum nebula region in their model. out. state that the modelling of this component is [ar from complete since a proper modelling would require many more constraints (for both the dispersion. and. the scattering) han are presently available.," Taylor Cordes (1993) do include the enhanced electron density in the Gum nebula region in their model, but state that the modelling of this component is far from complete since a proper modelling would require many more constraints (for both the dispersion and the scattering) than are presently available."
+ Phev have. therefore. assumed he fluctuation parameter (which quantifies the amount of scattering in the medium. given the value for the electron density) associated with the Gum nebula to be equal tozero.," They have, therefore, assumed the fluctuation parameter (which quantifies the amount of scattering in the medium, given the value for the electron density) associated with the Gum nebula to be equal to."
+ In the following part of this section. we show that the data on the Vela pulsar can be exploited to this [uctuation parameter (as defined in the Tavlor Cordes model).," In the following part of this section, we show that the data on the Vela pulsar can be exploited to this fluctuation parameter (as defined in the Taylor Cordes model)."
+ In a simple exercise. we varied the assumed. number density and the fluctuation parameter associated with the Gum nebula in the framework of the Tavlor. Cordes model to obtain values of these parameters that would be consistent with the known DM (69 pe 7) ns (825 ms at 327 MllIz) as well as the derived. distance (500. pe) to the Vela pulsar.," In a simple exercise, we varied the assumed number density and the fluctuation parameter associated with the Gum nebula in the framework of the Taylor Cordes model to obtain values of these parameters that would be consistent with the known DM (69 pc $^{-3}$ ) $\tau_{sc}$ (8.25 ms at 327 MHz) as well as the derived distance (500 pc) to the Vela pulsar."
+ As shown in Table 1. (columns 1. 4 5). the electron. density in the Gun nebula region needs to be about GO higher than that suggested in the Tavlor Cordes model.," As shown in Table \ref{table:table1} (columns 1, 4 5), the electron density in the Gum nebula region needs to be about $60$ higher than that suggested in the Taylor Cordes model."
+ The value of the Ductuation. parameter is about 6.3. quite close to that for the spiral armi component in the Tavlor Cordes model.," The value of the fluctuation parameter is about $6.3$, quite close to that for the spiral arm component in the Taylor Cordes model."
+ Although the Gum nebula is quite complicated. in its structure and morphology. it may not be unreasonable to assume that the above estimates would be more or less valid for other sight-lines through the Cum region.," Although the Gum nebula is quite complicated in its structure and morphology, it may not be unreasonable to assume that the above estimates would be more or less valid for other sight-lines through the Gum region."
+ With this view. we consider another pulsar D073640 (—2354 2.6—.902 DM = 160.8 pe em 7). for which the estimated distance on the basis of the Taylor C'ordes (1993) model is 1l kpe (and therefore would be placed bevond the ‘outer’ spiral arm)," With this view, we consider another pulsar B0736–40 $l=254^{\circ}.2$, $b=-9^{\circ}.2$ DM = 160.8 pc $^{-3}$ ), for which the estimated distance on the basis of the Taylor Cordes (1993) model is $>11$ kpc (and therefore would be placed beyond the `outer' spiral arm)."
+ As part of our survey (in 1996) to measure the temporal broadening of pulse profiles at 327 MllIz (Ramachancran 1997). we observed this object and measured the temporal broadening to be 76c3 ms.," As part of our survey (in 1996) to measure the temporal broadening of pulse profiles at 327 MHz (Ramachandran 1997), we observed this object and measured the temporal broadening to be $76\pm 3$ ms."
+ Assuming the excess scattering is due to the Gum nebula. we again seek a combination of the electron density and the fluctuation parameter that would be consistent wit= the DM τος values. and the distance estimate associated with it.," Assuming the excess scattering is due to the Gum nebula, we again seek a combination of the electron density and the fluctuation parameter that would be consistent with the DM $\tau_{sc}$ values, and the distance estimate associated with it."
+ As can be seen from Table 1.. both the cases. the Vela pulsar BOT3640. are consistent with the electron density and the Ductuation. parameter values of 0.32 em5 6.3 respectively.," As can be seen from Table \ref{table:table1}, both the cases, the Vela pulsar B0736–40, are consistent with the electron density and the fluctuation parameter values of 0.32 $^{-3}$ $\sim 6.3$ respectively."
+ This. therefore. allows us to derive a distance to DOT3640 as ~4.5 kpe (with an uncertainty of about OS kpe) based on our measurement of the temporal broadening the from the Vela pulsar data.," This, therefore, allows us to derive a distance to B0736–40 as $\sim 4.5$ kpc (with an uncertainty of about 0.8 kpc) based on our measurement of the temporal broadening the from the Vela pulsar data."
+" We also find that the excess temporal broadening cannot be accounted for by even a large increase in the electron. density associated with the ""outer"" spiral arm component.", We also find that the excess temporal broadening cannot be accounted for by even a large increase in the electron density associated with the `outer' spiral arm component.
+" “The z-height of about TOO pe. based on our new estimate. implies that the ""true distance could be even shorter given the possibility that the elfective scale-height of the electron distribution may be somewhat uncler-estiniatecl in the Tavlor Cordes model."," The z-height of about 700 pc, based on our new estimate, implies that the `true' distance could be even shorter given the possibility that the effective scale-height of the electron distribution may be somewhat under-estimated in the Taylor Cordes model."
+ As one of the important implications of this distance determination. the estimated velocity of the pulsar would now be less than 1600 km/sec (using a proper motion of about 72.5 mas/vr) rather than a value 2) kmsec as was implied by the earlier distance estimate.," As one of the important implications of this distance determination, the estimated velocity of the pulsar would now be less than 1600 km/sec (using a proper motion of about 72.5 mas/yr) rather than a value $>$ 3780 km/sec as was implied by the earlier distance estimate."
+ Any more detailed modelling. of the distribution. of electron density and its Ductuation across the Cum nebula region is bevond the scope of this paper. and. needs. new measurcments of scattering parameters of a number. of pulsars that could be behind this region.," Any more detailed modelling of the distribution of electron density and its fluctuation across the Gum nebula region is beyond the scope of this paper, and needs new measurements of scattering parameters of a number of pulsars that could be behind this region."
+ We. however. have shown how distance estimates for many such pulsars could be refined. already with the help of the scatter broadening measurements.," We, however, have shown how distance estimates for many such pulsars could be refined already with the help of the scatter broadening measurements."
+ We have described. in this paper a method to. study ellectivelv the. distribution of the. electron. density Iuctuations in the interstellar medium by assuming a simple, We have described in this paper a method to study effectively the distribution of the electron density fluctuations in the interstellar medium by assuming a simple
+notably in the Small Magellanic Cloud. SAIC) from the previously established. scaling relationsDolatto- followine).,"notably in the Small Magellanic Cloud, SMC) from the previously established scaling relations ."
+. After extending parameter space bv goiug to mere metal poor galaxies. one of the next logical steps is to look for varvine cloud properties within the same galaxy.," After extending parameter space by going to more metal poor galaxies, one of the next logical steps is to look for varying cloud properties within the same galaxy."
+ Deep single dish observations revealed CO cussion far out in or even bevoud the optical disks of galaxies 2007).. including the Local Group spiral M33 2010).," Deep single dish observations revealed CO emission far out in or even beyond the optical disks of galaxies , including the Local Group spiral M33 ."
+ These observations probe molecular eas properties in a regime where conditions in the iuterstellar medi aud thus CALC properties are expected to change. because of. e.g.. chaneing metallicities or dust abundance.," These observations probe molecular gas properties in a regime where conditions in the interstellar medium and thus GMC properties are expected to change, because of, e.g., changing metallicities or dust abundance."
+ The proximity of M33 and the availability of comprehensive. homogeneous mcasurclicnts aud cousisteutly determined GALC properties across the central star forming disk makes this galaxy au ideal target for high resolution CO observations to probe GALC properties iu the outer disk.," The proximity of M33 and the availability of comprehensive, homogeneous measurements and consistently determined GMC properties across the central star forming disk makes this galaxy an ideal target for high resolution CO observations to probe GMC properties in the outer disk."
+ Iu this paper we present CO observations from he Combined Array for Besearch in Millimeter Wave Astronomy (CARMA) of 8 CAICs., In this paper we present CO observations from the Combined Array for Research in Millimeter Wave Astronomy (CARMA) of 8 GMCs.
+ The ealactoceutric distance of these clouds corresponds to approxiuatelv wo scale leneths of the molecular gas. 1... the length over which the azimuthally averaged CO cinission declines joa factor (7. or ~ O.5re5. where ros is the isophotal radius correspouding to 25 D-band magnitudes per square aresecond.," The galactocentric distance of these clouds corresponds to approximately two scale lengths of the molecular gas, i.e., the length over which the azimuthally averaged CO emission declines by a factor $e^2$, or $\sim0.5{\rm r_{25}}$ , where $_{25}$ is the isophotal radius corresponding to 25 B-band magnitudes per square arcsecond."
+ Figure 1 illustrates the location of our areeted field and shows the CO peak intensity map we obtain., Figure \ref{fig1} illustrates the location of our targeted field and shows the CO peak intensity map we obtain.
+ We introduce our observations which we use to ucasure cloud properties for8 GMCs in Section 2 and discuss the expected cuvirommental variations (molecular eas fraction. dust-to-gas ratio. metallicitv) within M33 with respect to our target field in Section 3..," We introduce our observations which we use to measure cloud properties for8 GMCs in Section \ref{data} and discuss the expected environmental variations (molecular gas fraction, dust-to-gas ratio, metallicity) within M33 with respect to our target field in Section \ref{environ}."
+ We compare our nieasuremeuts to those from the iuner disk of N33. other nearby galaxies aud the Milkv. Way in Section L.," We compare our measurements to those from the inner disk of M33, other nearby galaxies and the Milky Way in Section \ref{cloud-props}."
+ Ow CARMA observatious were targeted towards one of the outermost detected molecular gas complexes from the sinele dish observations of, Our CARMA observations were targeted towards one of the outermost detected molecular gas complexes from the single dish observations of.
+ We note that recently have TuaiuBodyCitationEud668|gardan(ü7.obtained deeper single ENish data and detected CO cinission even further out i the disk of AL33., We note that recently have obtained deeper single dish data and detected CO emission even further out in the disk of M33.
+ We observed the target position in the CO(1-0) transition using a 19-poiut mosaic in D configuration between July aud August 2008 and iu C configuration between October and November 2009 uder mostly good Ό weather conditious., We observed the target position in the CO(1-0) transition using a 19-point mosaic in D configuration between July and August 2008 and in C configuration between October and November 2009 under mostly good 3mm weather conditions.
+ The total observing fine was —27hhr aud ~30hhr in D aud C configuration. respectively.," The total observing time was $\sim$ hr and $\sim$ hr in D and C configuration, respectively."
+" The svuthesized bean size is ~LT"". corresponding to ~ ppc at the distance of M33."," The synthesized beam size is $\sim1.7\arcsec$, corresponding to $\sim7$ pc at the distance of M33."
+ We tuned the uuu receivers to the doppler-shifted CO(1-0) rest frequency in the upper sidebaud and placed two ~31 MMdBz wide bands across the line., We tuned the 3mm receivers to the doppler-shifted CO(1-0) rest frequency in the upper sideband and placed two $\sim31$ Mhz wide bands across the line.
+ This viclds an effective bandwidth of ~155kkiass+ and a velocity resolution 6Veyaame=1.27 ., This yields an effective bandwidth of $\sim$ $^{-1}$ and a velocity resolution $\delta {\rm v_{channel}}=1.27$ $^{-1}$.
+ For phase and amplitude calibration.0205|322 was observed every 20 aniuutes.," For phase and amplitude calibration, 0205+322 was observed every 20 minutes."
+ Fluxes were bootstrapped from Uranus. Neptune or MWC319 aud 3€151.3 0r 3€8E were observed for bandpass calibration.," Fluxes were bootstrapped from Uranus, Neptune or MWC349 and 3C454.3 or 3C84 were observed for bandpass calibration."
+" Radio aud optical pointing was performed every δα,", Radio and optical pointing was performed every hr.
+ We estimate the resulting calibration to be accurate within ~15%., We estimate the resulting calibration to be accurate within $\sim15\%$.
+showsData reduction was performed in MIRIAD., Data reduction was performed in MIRIAD.
+ Fieure 1 the target field with respect to the voung stars (left paucl) aud the CO(1-0) peak intensity map frou our observations (right ΟΙ}., Figure \ref{fig1} shows the target field with respect to the young stars (left panel) and the CO(1-0) peak intensity map from our observations (right panel).
+ We identify CAICS and measure their properties using he CPROPS package2006)., We identify GMCs and measure their properties using the CPROPS package.
+.. We define GAICs as contiguous regious of high signal-to-noise enission., We define GMCs as contiguous regions of high signal-to-noise emission.
+ No further decomposition appears necessary n his dataset., No further decomposition appears necessary in this dataset.
+" CPROPS uses moment methods to measure radius. line width. huninositv, and peak temperature or each cloud."," CPROPS uses moment methods to measure radius, line width, luminosity, and peak temperature for each cloud."
+ It estimates the associated uncertainties musing boostrappiug techniques and attempts to correct each measurement for biases due to finite sensitivitv and limited spectral and iugulu resolution., It estimates the associated uncertainties using boostrapping techniques and attempts to correct each measurement for biases due to finite sensitivity and limited spectral and angular resolution.
+ A key point for this study is that BOs have used exactly the same approach to measure GAC properties from a huge sample of nearby galaxies. including the iuner disk of AD.," A key point for this study is that B08 have used exactly the same approach to measure GMC properties from a large sample of nearby galaxies, including the inner disk of M33."
+This gives us a well-controlled point of comparison., This gives us a well-controlled point of comparison.
+ Table 1 lists the derived. properties for the 8 CAICS in our field., Table \ref{table1} lists the derived properties for the 8 GMCs in our field.
+ We note that more faint CO features are visible in Figure 1.. but the signal-to-noise ratio in these features is too low for reliable cloud property estimates.," We note that more faint CO features are visible in Figure \ref{fig1}, but the signal-to-noise ratio in these features is too low for reliable cloud property estimates."
+ Previous iuterferometric observations of individual GAICs in M33 focussed on GMC properties iu the iuner of N33., Previous interferometric observations of individual GMCs in M33 focussed on GMC properties in the inner of M33.
+2007). Iu the following we assess the distinguishing characteristics of the ISM at larger radi in M33 and how this may impact CMC properties., In the following we assess the distinguishing characteristics of the ISM at larger radii in M33 and how this may impact GMC properties.
+ Numerous studies have compared the radial behavior of atomic eas. molecular gas aud various star formation tracers In ΑΠΣ2007).," Numerous studies have compared the radial behavior of atomic gas, molecular gas and various star formation tracers in M33."
+. While tlie lis found to remain relatively constant across the dis- of ND33. these studies derive a range of scale lenethρα between L1 2.5kkpe for the CO emission.," While the is found to remain relatively constant across the disk of M33, these studies derive a range of scale lengths between $\sim1.4-2.5$ kpc for the CO emission."
+ We not that we adopt a value of 2 kkpe for the CO scale lenet[um throughout this paper., We note that we adopt a value of $2$ kpc for the CO scale length throughout this paper.
+ From the cdiffereut behavior of aand CO enmüsson. if is clear that the molecular gas fraction Mypo/Syy decreases as a fiction of radius.," From the different behavior of and CO emission, it is clear that the molecular gas fraction $\Sigma_{\rm H2} / \Sigma_{\rm HI}$ decreases as a function of radius."
+" Most recently. measured this fraction and find that while the eas is mostly molecular in the center. the aolecular gas fraction decreases to ouly ~10.20% at our target position at re Ekkpe (adjusted to our adopted Calactic Cconversion factor of Now=Dy107«uu?(&lausty) +),"," Most recently, measured this fraction and find that while the gas is mostly molecular in the center, the molecular gas fraction decreases to only $\sim10-20\%$ at our target position at $\approx4$ kpc (adjusted to our adopted Galactic conversion factor of $_{\rm CO}=2\times10^{20}\,{\rm cm^{-2}\, (K\,km\,s^{-1})^{-1}}$ )."
+ The stellar surface density. Moa: dno M33 d8 low conrpared to other. larger spirals," The stellar surface density, $\Sigma_{\rm star}$, in M33 is low compared to other, larger spirals."
+"2010).. As Xa, is imuportaut forthe hwvdrostatie gas pressure. which was found to correlate well with the molecular eas fraction iu nearby galaxies 2002).. this maybe an explanation for the ramsnally low"," As $\Sigma_{\rm star}$ is important forthe hydrostatic gas pressure, which was found to correlate well with the molecular gas fraction in nearby galaxies , this may be an explanation for the unusually low"
+blue edge.,blue edge.
+ A more modest shift in the blue edge for high- e-modes in slowly pulsating D stars. was. obtained by Pamwvatnykh(1999)., A more modest shift in the blue edge for high-order g-modes in slowly pulsating B stars was obtained by \citet{Pam99}.
+. Since then. extensive work on the contribution of inner-shell electrons to the OP caleulations. has brought the OP opacities into much closer agreement with those from OPAL” (Badnell&Seaton2004:Dacdnelletal. 2005).," Since then, extensive work on the contribution of inner-shell electrons to the OP calculations, has brought the OP opacities into “much closer agreement with those from OPAL” \citep{Bad04,Bad05}."
+. Directed at obtaining good atomic data [for very large numbers of atomic transitions. OPAL uses. single-configuration wavelunctions with one-electron orbitals (Rogers&Iglesias1992).," Directed at obtaining good atomic data for very large numbers of atomic transitions, OPAL uses single-configuration wavefunctions with one-electron orbitals \citep{OPAL92}."
+. OP originally obtained most of its atomic data using the R-matrix method to evaluate multi-electron. wavefunctions. resulting in. more accurate atomic data. but for fewer and simpler transitions1997).," OP originally obtained most of its atomic data using the R-matrix method to evaluate multi-electron wavefunctions, resulting in more accurate atomic data, but for fewer and simpler transitions."
+. Promptecl by substantial dilferencees between the first OP results and OPAL. primarily at high. temperature ancl densities. the introduction. of conliguration-interaction codes enabled OP to include many more bound states and radiative transition. probabilities. as well as auto-ionization probabilities ancl photo-ionization cross-sections (Seaton2005).," Prompted by substantial differences between the first OP results and OPAL, primarily at high temperature and densities, the introduction of configuration-interaction codes enabled OP to include many more bound states and radiative transition probabilities, as well as auto-ionization probabilities and photo-ionization cross-sections \citep{Sea05}."
+. For many astrophysical situations. OP anc OPAL opacities appear to be in closer agreement with one another (< 5%) (Seaton2005). than with observation.," For many astrophysical situations, OP and OPAL opacities appear to be in closer agreement with one another $<5\%$ ) \citep{Sea05} than with observation."
+ For example. helioscismology results currently require an increase of ~20% in opacity below the base of the solar convective zone (Baheallctal.2005).," For example, helioseismology results currently require an increase of $\sim20\%$ in opacity below the base of the solar convective zone \citep{Bah05}."
+. While the results presented here do not directly suggest a solution to the solar opacity problem.," While the results presented here do not directly suggest a solution to the solar opacity problem,"
+We obtain new values for the svstemic velocity ancl the velocity semi-amplitude of the secondary in agreement with previous authors.,We obtain new values for the systemic velocity and the velocity semi-amplitude of the secondary in agreement with previous authors.
+ We conclude that there is no evidence for increased heating over the inner face of the donor star during this stage of the outburst., We conclude that there is no evidence for increased heating over the inner face of the donor star during this stage of the outburst.
+ We find spin mocdulations in the ΑΗ. ratios of the lines but only tentatively in their integrated [fluxes or the continuum., We find spin modulations in the V/R ratios of the lines but only tentatively in their integrated fluxes or the continuum.
+ Spin power resolved across the line profiles extends to velocities of a large fraction of the freclall velocity of the central object.," Spin power resolved across the line profiles extends to velocities of, a large fraction of the freefall velocity of the central object."
+ The detection of kilo-seconcd QPOs across the optical emission line profiles of hhave provided an opportunity to test. the mechanism behind the unique long-timescale QPOs in this object. which are an order of magnitude longer than QPOs normally observed. in disc-accreting cataclysmic variables.," The detection of kilo-second QPOs across the optical emission line profiles of have provided an opportunity to test the mechanism behind the unique long-timescale QPOs in this object, which are an order of magnitude longer than QPOs normally observed in disc-accreting cataclysmic variables."
+ We have rejected the model of LIL which considers the direct effects of blobs orbiting at the Ixeplerian frequency of the annulus associated with a disccoverllow impact site., We have rejected the model of HL which considers the direct effects of blobs orbiting at the Keplerian frequency of the annulus associated with a disc-overflow impact site.
+ Our favoured models consider the long ΟΡΟ period to be the consequence, Our favoured models consider the long QPO period to be the consequence
+phase is gravitational instability.,phase is gravitational instability.
+ The collapse ο “the protostellar cloud alvost certaluly results iu inmost of the material iuitially lauding at large ¢isk radii. whieh is a stroiglv unstable «ιοτοι.," The collapse of the protostellar cloud almost certainly results in most of the material initially landing at large disk radii, which is a strongly unstable condition."
+ ]t seems highly plausije that the bulk of the p'OlOsellar material is thus driven inward towards the central regiors by eravitatioual torques., It seems highly plausible that the bulk of the protostellar material is thus driven inward towards the central regions by gravitational torques.
+ TIls ralses the question: why not view the T Tauri accretion process as tle not-fulls-coimpleted enc of tle eravitationally-driveu pliase?, This raises the question: why not view the T Tauri accretion process as the not-fully-completed end of the gravitationally-driven phase?
+ If eravitational itstabilities are dominant in translerriug most of tle nass trom the initia disk to the centra Lst. the ouly way for the cisx lo je highly graviatioually stable at the eud of the protostellar phase is if the MB] or some oher mechanisiu is very efficieut in driving accretion.," If gravitational instabilities are dominant in transferring most of the mass from the initial disk to the central star, the only way for the disk to be highly gravitationally stable at the end of the protostellar phase is if the MRI or some other mechanism is very efficient in driving accretion."
+ While the MRI shoLd be extremely efficient in tleiui er disk. where thermal ionization is high. ane quite plausibly also in outer disk regious. whe'e X-rays can penet‘ate and recombination rates are low. there might stilll be a dead zoje al interimeca aleradii which wieght be clriven into gravHationa instability by continually1 accumuing material from the outer disk.," While the MRI should be extremely efficient in the inner disk, where thermal ionization is high, and quite plausibly also in outer disk regions, where X-rays can penetrate and recombination rates are low, there might still be a dead zone at intermediate radii which might be driven into gravitational instability by continually accumulating material from the outer disk."
+ A majΟΥ LeasoÜi why eravitational instalilities Lave not beet strougly advocated Lor T Tauri disks is that disk masses seem to be too low., A major reason why gravitational instabilities have not been strongly advocated for T Tauri disks is that disk masses seem to be too low.
+ Disk imasses need to be of the order of 0.l of the stelar mass for graytational instajilities to operate (Pringle 1981. While typical T Tauri disk ass estituates1 have been of the order o £0.01 M.. roughly a miinimiun-mass solar nebula.," Disk masses need to be of the order of 0.1 of the stellar mass for gravitational instabilities to operate (Pringle 1981), while typical T Tauri disk mass estimates have been of the order of 0.01 $\msun$, roughly a minimum-mass solar nebula."
+ A very recent COLipreliensive survey of the Tauris star-forming region suMODeoests a median disk mass of 5x10341. (Aidrews Williauus 2005). albeit witha arge scatter.," A very recent comprehensive survey of the Taurus star-forming region suggests a median disk mass of $5 \times 10^{-3} \msun$ (Andrews Williams 2005), albeit with a large scatter."
+ This median mass estimate is an order of maguituce lower than what is required [or gravitational iustabiliv. though estimates for a small uunuber of objects ave much closer to the limiting value.," This median mass estimate is an order of magnitude lower than what is required for gravitational instability, though estimates for a small number of objects are much closer to the limiting value."
+ However. the uncertainty in disk masses estimate from dust emission is often not st[Iicientlv recognized.," However, the uncertainty in disk masses estimated from dust emission is often not sufficiently recognized."
+ Adoption of typical estimated lone-waveleleth opacities of interstellar dust esults in üunplausibly large disk masses. and the spectral iudices of observed mumn- aud sub-miu wave enilssion [rom T Tauri disks also sweest that the dust. e@raius iive evolved substantially in the ¢isk (e.g. Beckwith Sargent 1901: Andrew5 Williams 2005. aud references herein).," Adoption of typical estimated long-wavelength opacities of interstellar dust results in implausibly large disk masses, and the spectral indices of observed mm- and sub-mm wave emission from T Tauri disks also suggest that the dust grains have evolved substantially in the disk (e.g., Beckwith Sargent 1991; Andrews Williams 2005, and references therein)."
+ D'Alessio (2001) cousicle'ed power-law cisributious of dust sizes aud 5j0wed hat dus erowth to maxinun sizes greater tliui about ] nunresults iu a spectral index determied by te power law of the size distribution. 100 the maxinmuu1 size of the grains.," D'Alessio (2001) considered power-law distributions of dust sizes and showed that dust growth to maximum sizes greater than about 1 mm results in a spectral index determined by the power law of the size distribution, not the maximum size of the grains."
+ If the maximum size is much larger than 1 mua (for the dIND properties cousiclerecl by D'Alessio 2001). the imurwave opacity cau be much lower than assumed i typical observational estimates. while maiitaining a fixed spectral iudex cletermiied by the distrition of dust sizes.," If the maximum size is much larger than 1 mm (for the dust properties considered by D'Alessio 2001), the mm-wave opacity can be much lower than assumed in typical observational estimates, while maintaining a fixed spectral index determined by the distribution of dust sizes."
+ In the D'Alessio calculations. only maximum dust sizes within a factor of 3 to 10 of 1 mm provide νάνο opacities close o the standard value assumed: much sinaller or much larger maximum sizes provide much lower opacities.," In the D'Alessio calculations, only maximum dust sizes within a factor of 3 to 10 of 1 mm provide mm-wave opacities close to the standard value assumed; much smaller or much larger maximum sizes provide much lower opacities."
+ Tjus. dus erowth is more likely to yield low dust opacities than high dust opacities.," Thus, dust growth is more likely to yield low dust opacities than high dust opacities."
+ Given, Given
+We beein with Equation 2 ∖∖↽∐↕≺∢↥⊔⇂≼↲⇂∎∐∐↲⊳∖⇁⊔∐↲∐↓≀↧⊔∖⊽⊳∖⊽⋅∣∣∣↿⋅∙∩∣≼⋝∫↘⋝⇄⋝⋅↕∐⊳∖⊽↕≼⇂≼↲≀↧↴≺↕∖↽∐∐≺⇂≼↲↕⋅∪↓⋟↕⋅⋯∐∏⊳∖⇁∫↘⋝∪↕⋅↕≼↲∐∩↲≼⇂≀↧↴↥∪∐≸≟⊔∐↲∐∐≼↲∪⋟ sight: Using Equation 13. to replace X(R). we have Now we want to relate this cylindrical mass. maj4(R8). to the 3D spherical mass M(A).,"We begin with Equation \ref{eq:delta-sigma}
+ which defines the mass, $m_{cyl}(R)$, inside a cylinder of radius $R$ oriented along the line ofsight: Using Equation \ref{eq:sigma-proj} to replace $\Sigma(R)$, we have Now we want to relate this cylindrical mass, $m_{cyl}(R)$ , to the 3D spherical mass $M(R)$."
+ It is a slight abuse of notation to write M(A2). since 22 is the 2D radius in the plane of the sky. but this avoids the bulkiness of the more precise notation. A/(r=R).," It is a slight abuse of notation to write $M(R)$, since $R$ is the 2D radius in the plane of the sky, but this avoids the bulkiness of the more precise notation, $M(r=R)$."
+ Again. assuming spherical symmetry. the 3D mass can be written as the total mass inside the cvlinder minus (he mass inside the cvlinder but outside (he sphere.," Again, assuming spherical symmetry, the 3D mass can be written as the total mass inside the cylinder minus the mass inside the cylinder but outside the sphere."
+ Using evlindrical-polar coordinates wv.z.6 and spherical radius r=y?+22. with the z-axis beinge the axis of projection. we have Substituting Equation 18. for mi into 20. gives an expression for M(H2) in terms of AX and Ap.," Using cylindrical-polar coordinates $x,z,\phi$ and spherical radius $r=\sqrt{x^2+z^2}$, with the $z$ -axis being the axis of projection, we have Substituting Equation \ref{eq:m2d} for $m_{cyl}$ into \ref{eq:MR} gives an expression for $M(R)$ in terms of $\Delta\Sigma$ and $\Delta\rho$ ."
+ Since Ap is a functional of AX. this is a mapping: AX’——A/.," Since $\Delta\rho$ is a functional of $\Delta\Sigma$, this is a mapping: $\Delta\Sigma \longmapsto M$."
+ We will define (his particular form for M(G2) as M(HR). since it only uses data outside of J: since we assumed above (hat the average clensily is zero. equation 21 actually gives only the mass over-densitv. not the total mass: to properly correct for this. a term 4/3242p should be added. but this formula is (vpically used on scales where this extra contribution is negligible.," We will define this particular form for $M(R)$ as $M_{out}(R)$, since it only uses data outside of $R$: Since we assumed above that the average density is zero, equation \ref{eq:MoutR} actually gives only the mass over-density, not the total mass; to properly correct for this, a term $4/3 \pi R^3 \bar{\rho}$ should be added, but this formula is typically used on scales where this extra contribution is negligible."
+ We will now derive another formula for the mass prolile., We will now derive another formula for the mass profile.
+As we mentioned. the simplest idea is just to integrate the density.,"As we mentioned, the simplest idea is just to integrate the density."
+" We willcall (his particular form Ad;,,(22). which is given"," We willcall this particular form $M_{in}(R)$ , which is given"
+requency is the kIIz QPO which we already knew. while the second one becomes detectable as a result of the shift-zuxd-add method.,"frequency is the kHz QPO which we already knew, while the second one becomes detectable as a result of the shift-and-add method."
+ Phe significance level of the second. upper. kHz QPO is 3.10. which implies a marginal detection.," The significance level of the second, upper, kHz QPO is $\sigma$, which implies a marginal detection."
+ Note rowever that the frequency difference between the ΚΣ QPOs is 258+13 Hz. which is consistent with the value in he Z phase (see Table 1. and Homanctal. 2007b)) and herefore there are no other trials involved. in estimating his significance.," Note however that the frequency difference between the kHz QPOs is $258\pm13$ Hz, which is consistent with the value in the Z phase (see Table \ref{tab2} and \citealt{h2007}) ) and therefore there are no other trials involved in estimating this significance."
+ As a result. of this analysis. taking into account all the caveats. we find a pair of. simultaneous κ. QPOs in the atoll phase of the NIS J1701.462 which strengthens our previous suggestion that the strong ΚΣ QPOs detected in this phase are always the lower klIz As described in section 2.. we caleulate the fractional rms amplitude of the atoll QPOs in 5 dillerent. energy bands. 2-3 keV. 3-6 keV. 6-11. keV. 11-16 keV and. 16-25 keV. respectively.," As a result of this analysis, taking into account all the caveats, we find a pair of simultaneous kHz QPOs in the atoll phase of the XTE J1701–462 which strengthens our previous suggestion that the strong kHz QPOs detected in this phase are always the lower kHz As described in section \ref{data}, we calculate the fractional rms amplitude of the atoll QPOs in 5 different energy bands, 2-3 keV, 3-6 keV, 6-11 keV, 11-16 keV and 16-25 keV, respectively."
+ We find fractional rms amplitudes: of less than10.9%.. 6.9+0.3. 11.440.2. 17.5=1.2 and less than21.34... respectively (upper limits are given. at confidence).," We find fractional rms amplitudes of less than, $6.9\pm0.3$, $11.4\pm0.2$, $17.5\pm1.2$ and less than, respectively (upper limits are given at confidence)."
+ ‘Those results. as already noticed. in other sources (see Bereer et al.," Those results, as already noticed in other sources (see Berger et al."
+ 1996). show that the strength. of the variability increases as the energy. Figure 3. shows the values of the QPO quality factor as à function of frequency.," 1996), show that the strength of the variability increases as the energy Figure \ref{Q_nu} shows the values of the QPO quality factor as a function of frequency."
+ Black points represent nieasurements in the atoll phase while grey points are measurements in the Z phase., Black points represent measurements in the atoll phase while grey points are measurements in the Z phase.
+ For the Z phase. empty and filled square symbols represent. respectively. lower and upper klIz QPOs.," For the Z phase, empty and filled square symbols represent, respectively, lower and upper kHz QPOs."
+ From the plot we notice a few interesting aspects: 1) QPOs in the atoll phase are on average 10 times more coherent than in the Z phase: 2) in the atoll phase ( increases as the frequeney increases. reaching a maximum value of about 150 at ~ SLO Lz: within the errors the behaviour is consistent with what has been observed in other sources (secBarretetal. 2005a).," From the plot we notice a few interesting aspects: 1) QPOs in the atoll phase are on average 10 times more coherent than in the Z phase; 2) in the atoll phase $Q$ increases as the frequency increases, reaching a maximum value of about 150 at $\sim$ 810 Hz; within the errors the behaviour is consistent with what has been observed in other sources \citep[see][]{b2005a}."
+. No such trend. seems to be present in the Z phase: 3) comparing lower kllz QPOs in both phases it is clear that in the Z phase the lower kIlz QPO appears at lower frequencies than in the atoll phase: 4) there is an overlap in frequency around 640 Hz between Z and atoll QPOs with a significant mismatch in Q. ligure 4. shows the ΟΡΟ rms fractional amplitude as a 'unction of frequency., No such trend seems to be present in the Z phase; 3) comparing lower kHz QPOs in both phases it is clear that in the Z phase the lower kHz QPO appears at lower frequencies than in the atoll phase; 4) there is an overlap in frequency around 640 Hz between Z and atoll QPOs with a significant mismatch in $Q$ Figure \ref{r_nu} shows the QPO rms fractional amplitude as a function of frequency.
+ As in Figure 3.. black points denote measurements in the atoll phase. while grev points are QPOs in the Z phase.," As in Figure \ref{Q_nu}, black points denote measurements in the atoll phase, while grey points are QPOs in the Z phase."
+ The rms amplitude of the lower κ QPO in the atoll phase remains more or less constant around. as the QPO frequency increases from 640. 112 o TNO Lz and then drops rapidly at  SOO Llz., The rms amplitude of the lower kHz QPO in the atoll phase remains more or less constant around as the QPO frequency increases from 640 Hz to 780 Hz and then drops rapidly at $\sim$ 800 Hz.
+ There is no evidence of a similar trend for the rms of the lower or upper kllz QPOs in the Z phase., There is no evidence of a similar trend for the rms of the lower or upper kHz QPOs in the Z phase.
+ As already noticed for he quality factor. the rms amplitude values show a clear dillerence between the atoll ancl Z phases: The lower ΚΣ QPOs amplitude in the atoll phase is on average a factor of 2 higher than in the Z phase: this dilference is also apparent in the region where atoll and Z QPOs overlap in frequency between 620 Hz and 660 We test whether we may have missed. any significant kllz QPO in our analwsis.," As already noticed for the quality factor, the rms amplitude values show a clear difference between the atoll and Z phases: The lower kHz QPOs amplitude in the atoll phase is on average a factor of 2 higher than in the Z phase; this difference is also apparent in the region where atoll and Z QPOs overlap in frequency between 620 Hz and 660 We test whether we may have missed any significant kHz QPO in our analysis."
+ We caleulate upper limits to the fractional rms amplitude in frequency ranges where we did not fine ΚΙ QPOs., We calculate upper limits to the fractional rms amplitude in frequency ranges where we did not find kHz QPOs.
+ To do so we take two power density spectra. one for cach source phase. where no ΚΣ QPOs were found.," To do so we take two power density spectra, one for each source phase, where no kHz QPOs were found."
+ We fit those power density spectra using a model consisting of a constant to fit the Poissonian noise. and a Lorentzian to fit the QPO with fixed values for the centroid frequency. and the quality. factor Q.," We fit those power density spectra using a model consisting of a constant to fit the Poissonian noise, and a Lorentzian to fit the QPO with fixed values for the centroid frequency and the quality factor $Q$."
+ We fit the data and we estimate the upper limit of the amplitude of the Lorentzian using εντ=2.7. which corresponds to confidence level.," We fit the data and we estimate the upper limit of the amplitude of the Lorentzian using $\Delta \chi^{2}=2.7$, which corresponds to confidence level."
+ We repeat the same procedure. shifting eraclually the Crequeney of the Lorentzian to higher values until we cover the frequency range of For the analysis of atoll observation we use power density spectra of 256s. comparable to the time intervals over which we detected. kIIz QPOs in the atoll phase.," We repeat the same procedure shifting gradually the frequency of the Lorentzian to higher values until we cover the frequency range of For the analysis of atoll observation we use power density spectra of 256s, comparable to the time intervals over which we detected kHz QPOs in the atoll phase."
+ We use two dilferent quality factors. Q=50 and Q=20. which are comparable to the values we found for the kllz QPOs at low frequencies (sec Table 2)).," We use two different quality factors, $Q=50$ and $Q=20$, which are comparable to the values we found for the kHz QPOs at low frequencies (see Table \ref{tab1}) )."
+ We find upper limits to the fractional rms amplitude that vary within the range and within the rangeTA... respectively for Q=20 and GQ=50.," We find upper limits to the fractional rms amplitude that vary within the range and within the range, respectively for $Q=20$ and $Q=50$."
+ From these values it appears that the rms amplitude of the lower kllz QPO in the atoll phase should. decrease. at frequencies. lower nan 640 Lz., From these values it appears that the rms amplitude of the lower kHz QPO in the atoll phase should decrease at frequencies lower than 640 Hz.
+ This would be in agreement with the typical rmis-Lrequency trencl observed in most of the LAINBs (Méndezοἱal.2001: vanStraatenetal.2000.2002. 2003: DiSalvoetal.2001:: Salvoetal. 2003: Barretctal.2005a)) For the Z observation. we estimate the upper limits to the amplitude of a kllz QPO with typical atoll properties within short time intervals to check whether we could have missed such a QPO cluring the analysis.," This would be in agreement with the typical rms-frequency trend observed in most of the LMXBs \citealt*{mvf2001}; \citealt{van2000, van2002, van2003}; \citealt{ds2001}; \citealt*{dmv2003}; \citealt{b2005a}) For the Z observation, we estimate the upper limits to the amplitude of a kHz QPO with typical atoll properties within short time intervals to check whether we could have missed such a QPO during the analysis."
+ We calculate the upper limits in the frequeney range 650-940 Hz assuming a kllz QPO with Q=100 ancl setting 8 values of the integration time. 128. 256 and 512 seconds.," We calculate the upper limits in the frequency range 650-940 Hz assuming a kHz QPO with $Q=100$ and setting 3 values of the integration time, 128, 256 and 512 seconds."
+ We find that the upper limits to the fractional rms. amplitude. vary between and [for 128 seconds of integration timo. between and for 256 seconds. anc between and for 512 seconds.," We find that the upper limits to the fractional rms amplitude vary between and for 128 seconds of integration time, between and for 256 seconds, and between and for 512 seconds."
+ From this we conclude that if a kllz QPO with Q=100 and fractional rms amplitude similar to what we found for the atoll kllz QPOs was present in the Z observations. we would have detected. it significantly in intervals as short as The transient source NTE JITOLI462 is so far the only accreting neutron-star X-ray binary that changed. [rom a Z into an atoll source during an outburst.," From this we conclude that if a kHz QPO with $Q=100$ and fractional rms amplitude similar to what we found for the atoll kHz QPOs was present in the Z observations, we would have detected it significantly in intervals as short as The transient source XTE J1701–462 is so far the only accreting neutron-star X-ray binary that changed from a Z into an atoll source during an outburst."
+ Here we found that the properties of the kIIz QPOs in NTE J1701.462 in the atoll and Z phases are signilicanth: clilferent: At approximately the same frequeney (about 640. Hz). the coherence of the kllz QPOs is a factor of 2 larger. and the rms amplitude is a factor of 3 larger. in the atoll than in the Z phase (see Figure 3 ancl Figure 4. respectively).," Here we found that the properties of the kHz QPOs in XTE J1701–462 in the atoll and Z phases are significantly different: At approximately the same frequency (about 640 Hz), the coherence of the kHz QPOs is a factor of 2 larger, and the rms amplitude is a factor of 3 larger, in the atoll than in the Z phase (see Figure 3 and Figure 4, respectively)."
+ Furthermore. out of 707 observations in the Z phase. there is no single one in which the klIz QPOs have a coherence or rms amplitude similar to those seen when NTIS J1701.462," Furthermore, out of 707 observations in the Z phase, there is no single one in which the kHz QPOs have a coherence or rms amplitude similar to those seen when XTE J1701–462"
+However. the backward calculation of the motion in the Galaxy requires the knowledge of the radial velocity which cannot be measured.,"However, the backward calculation of the motion in the Galaxy requires the knowledge of the radial velocity which cannot be measured."
+ Willems et al. (, Willems et al. (
+2004. 2006). Thorsett et al (2005) and Stairs et al. (,"2004, 2006), Thorsett et al (2005) and Stairs et al. ("
+2006) therefore considered possible radial velocities drawn from either a uniform or Gaussian radial velocity distribution.,2006) therefore considered possible radial velocities drawn from either a uniform or Gaussian radial velocity distribution.
+ In this paper. we avoid the assumption of a present-day radial velocity distribution by carrying out Monte Carlo simulations of the motion of the observed DNS in the Galaxy that is forward in time instead of backward.," In this paper, we avoid the assumption of a present-day radial velocity distribution by carrying out Monte Carlo simulations of the motion of the observed DNS in the Galaxy that is forward in time instead of backward."
+ For this purpose. we distribute populations of DNS progenitors in the Galaxy according to a double exponential distribution function.," For this purpose, we distribute populations of DNS progenitors in the Galaxy according to a double exponential distribution function."
+ Each system in the populations is assigned a velocity equal to the vector sum of the local Galactic. rotational velocity and an isotropic kick velocity with a magnitude generated from the kick velocity distribution function. derived from the supernova orbital dynamies constraints for each of the observed systems.," Each system in the populations is assigned a velocity equal to the vector sum of the local Galactic rotational velocity and an isotropic kick velocity with a magnitude generated from the kick velocity distribution function, derived from the supernova orbital dynamics constraints for each of the observed systems."
+ Compared to Willems et al.(2004). Thorsett et al. (," Compared to Willems et al.(2004), Thorsett et al. ("
+2005). and Wang et al. (,"2005), and Wang et al. ("
+2006). we also relax the constraint that NS2's progenitor needs to be more massive than 2.1 M... which is a convectional limit for a helium star that explodes in a core collapse supernova.,"2006), we also relax the constraint that NS2's progenitor needs to be more massive than 2.1 $_{\sun}$, which is a convectional limit for a helium star that explodes in a core collapse supernova."
+ Instead. we restrict the progenitor to be more massive than NS2 only.," Instead, we restrict the progenitor to be more massive than NS2 only."
+ Furthermore we study alf eight known DNS in our galaxy. using the date observational parameters listed in Table 1.," Furthermore we study all eight known DNS in our galaxy, using the up-to-date observational parameters listed in Table 1."
+ The methodology of the adopted analysis is outlined in more detail in .22. while results for the individual systems and comparison with earlier studies are discussed in .33.," The methodology of the adopted analysis is outlined in more detail in 2, while results for the individual systems and comparison with earlier studies are discussed in 3."
+ In 4. we summarize our results and discuss their implications for the supernova forming the second-born neutron star in the observed DNS.," In 4, we summarize our results and discuss their implications for the supernova forming the second-born neutron star in the observed DNS."
+ Tauris & van den Heuvel (2004) reviewed the general scenario for forming DNS., Tauris $\&$ van den Heuvel (2004) reviewed the general scenario for forming DNS.
+ It starts with a binary consisting of two massive ZAMS stars € 710M ..).," It starts with a binary consisting of two massive ZAMS stars ( $\ga 10\, M_{\sun}$ )."
+ The primary star. which is initially the more massive one. will leave the main sequence first and explode as the first supernova. in which it forms the first-born neutron star (NS1) in this binary.," The primary star, which is initially the more massive one, will leave the main sequence first and explode as the first supernova, in which it forms the first-born neutron star (NS1) in this binary."
+ There may be mass transfer from the primary to the secondary star before that supernova takes place., There may be mass transfer from the primary to the secondary star before that supernova takes place.
+ As the secondary star leaves the main sequence. wind mass transfer from the secondary to ΝΡΙ can occur. leading to the formation of a X-ray binary.," As the secondary star leaves the main sequence, wind mass transfer from the secondary to NS1 can occur, leading to the formation of a X-ray binary."
+ While NS] ts accreting mass from its companion. it is spun up or “recycled”. and its spin axis is expected to align with the orbital angular momentum.," While NS1 is accreting mass from its companion, it is spun up or ""recycled"", and its spin axis is expected to align with the orbital angular momentum."
+ Later on. the X-ray binary may end up in a common envelope phase. during which NSI is engulfed by the extended envelope of its companion.," Later on, the X-ray binary may end up in a common envelope phase, during which NS1 is engulfed by the extended envelope of its companion."
+ The orbit shrinks and eireularizes rapidly due to the frictional forces acting on NSI as it moves through the envelope of the companion., The orbit shrinks and circularizes rapidly due to the frictional forces acting on NS1 as it moves through the envelope of the companion.
+ The lost orbital energy is deposited into the gas envelope. causing it to become gravitationally unbound and get ejected. leaving behind the companion's bare helium core.," The lost orbital energy is deposited into the gas envelope, causing it to become gravitationally unbound and get ejected, leaving behind the companion's bare helium core."
+ The helium star eventually explodes in a supernova and becomes NS2. orbiting the recycled. milli-second pulsar NS1 in à compact binary orbit.," The helium star eventually explodes in a supernova and becomes NS2, orbiting the recycled, milli-second pulsar NS1 in a compact binary orbit."
+ Mass transfer from the helium star to NS] can also occur prior to the second supernova explosion., Mass transfer from the helium star to NS1 can also occur prior to the second supernova explosion.
+ Here. our goal is to use the current DNS binary properties listed in Table | to constrain as best possible the magnitude of the natal kick and the mass of the progenitor of NS2 in the 8 known DNS in our galaxy at the time of the second supernova.," Here, our goal is to use the current DNS binary properties listed in Table 1 to constrain as best possible the magnitude of the natal kick and the mass of the progenitor of NS2 in the 8 known DNS in our galaxy at the time of the second supernova."
+ We also predict the unknown radial velocity of each system. as well as the total transverse velocity and NSI spin tilt angle for those systems where these properties have not been measured vet.," We also predict the unknown radial velocity of each system, as well as the total transverse velocity and NS1 spin tilt angle for those systems where these properties have not been measured yet."
+ We achieve these goals by means of Monte Carlo simulations. incorporating orbital dynamics during supernova explosions and kinematic of binaries in the Galactic potential. as detailed in what follows.," We achieve these goals by means of Monte Carlo simulations, incorporating orbital dynamics during supernova explosions and kinematic of binaries in the Galactic potential, as detailed in what follows."
+ We start our analysis by deriving a probability distribution function (PDF) for the natal kick of NS2. which depends on the semi-major axis and orbital eccentricity immediately after the second supernova (Αα and όρων respectively).," We start our analysis by deriving a probability distribution function (PDF) for the natal kick of NS2, which depends on the semi-major axis and orbital eccentricity immediately after the second supernova $A_{postSN}$ and $e_{posSN}$ respectively)."
+ As the tidal interaction between two neutron stars in DNS is extremely weak (Bildsten Cutler 1992). the major mechanism that changes the orbital semi-major axis and eccentricity of DNS ts the emission of gravitational waves.," As the tidal interaction between two neutron stars in DNS is extremely weak (Bildsten Cutler 1992), the major mechanism that changes the orbital semi-major axis and eccentricity of DNS is the emission of gravitational waves."
+ To account for gravitational radiation driven orbital evolution. we use the current observational parameters as initial conditions and integrate equations (35) and (36) in Junker Schiiffer (1992) backwards in time.," To account for gravitational radiation driven orbital evolution, we use the current observational parameters as initial conditions and integrate equations (35) and (36) in Junker Schäffer (1992) backwards in time."
+ Because of the uncertainty in the true age of the observed DNS (see Kiziltan et al., Because of the uncertainty in the true age of the observed DNS (see Kiziltan et al.
+ 2009 for details). we randomly draw the integration time 7; from a uniform distribution of ages between O and 7. where 7. 1s the characteristic age of the DNS.," 2009 for details), we randomly draw the integration time $t_i$ from a uniform distribution of ages between 0 and $\tau_c$, where $\tau_c$ is the characteristic age of the DNS."
+ If 7. is greater than the age of our galaxy. which is about 10 Gyr. we use 10 Gyr as the upper limit for ¢; instead.," If $\tau_c$ is greater than the age of our galaxy, which is about 10 Gyr, we use 10 Gyr as the upper limit for $t_i$ instead."
+" For PSR B1534412. PSR B1I913+16, and PSR J0737-3039. different upper limits are adopted for 5. as discussed in more detail in the next section."," For PSR B1534+12, PSR B1913+16, and PSR J0737-3039, different upper limits are adopted for $t_i$, as discussed in more detail in the next section."
+ Once Apovsy and paws are known from the gravitational radiation integration. we randomly draw the natal kick velocity (VY) from a uniform distribution of values between O and 2500 km/s (r.e.. we assume no prior kick distribution derived from pulsar samples analysis). and the kick direction (0.0) from an isotropic distribution.," Once $A_{postSN}$ and $e_{postSN}$ are known from the gravitational radiation integration, we randomly draw the natal kick velocity $V_k$ ) from a uniform distribution of values between 0 and 2500 km/s (i.e., we assume no prior kick distribution derived from pulsar samples analysis), and the kick direction $(\theta, \phi)$ from an isotropic distribution."
+ Here 0 is the polar angle between the natal kick velocity and the instantaneous orbital velocity of NS2’s progenitor at the time of the second supernova explosion. and o is the corresponding azimuthal angle (see Figure 1 1n Kalogera 2000 for a graphical representation).," Here $\theta$ is the polar angle between the natal kick velocity and the instantaneous orbital velocity of NS2's progenitor at the time of the second supernova explosion, and $\phi$ is the corresponding azimuthal angle (see Figure 1 in Kalogera 2000 for a graphical representation)."
+ Using the conservation laws of orbital energy and angular momentum expressed by equations (19) and (20) in Willems et al. (, Using the conservation laws of orbital energy and angular momentum expressed by equations (19) and (20) in Willems et al. (
+"2005). we obtain the orbital semi-major axis CA, s.) immediately before the second supernova and the mass of NS2's helium star (M»;).","2005), we obtain the orbital semi-major axis $A_{preSN}$ ) immediately before the second supernova and the mass of NS2's helium star $M_{2i}$ )."
+ The pre-supernova orbit is assumed to be tidally circularized at the time of the second supernova explosion., The pre-supernova orbit is assumed to be tidally circularized at the time of the second supernova explosion.
+ On the other hand. the spin tilt angle of NSI. which ts the angle between NSI axis and the post supernova orbital angular momentum axis. can be caleulated by following equation (3) in Kalogera (2000).," On the other hand, the spin tilt angle of NS1, which is the angle between NS1 axis and the post supernova orbital angular momentum axis, can be calculated by following equation (3) in Kalogera (2000)."
+ Also. using equations (12) in the same paper. we find the polar angle 0 between the natal kick and the pre-supernova orbital angular momental axis.," Also, using equations (12) in the same paper, we find the polar angle $\theta'$ between the natal kick and the pre-supernova orbital angular momental axis."
+ Immediately before the second supernova explosion. there may be mass transfer from NS2s progenitor to NSI.," Immediately before the second supernova explosion, there may be mass transfer from NS2's progenitor to NS1."
+ If the radius Rs; of NS2's progenitor is greater than its Roche lobe radius Αι. which ts approximated by equation (2) in Eggleton (1983). Roche lobe overflow (RLO) takes place.," If the radius $R_{2i}$ of NS2's progenitor is greater than its Roche lobe radius $R_L$, which is approximated by equation (2) in Eggleton (1983), Roche lobe overflow (RLO) takes place."
+ Since NS2’s progenitor is a bare helium star before the second supernova. we use equation (A10) in. Kalogera Webbink (1998) to caleulate R»;.," Since NS2's progenitor is a bare helium star before the second supernova, we use equation (A10) in Kalogera Webbink (1998) to calculate $R_{2i}$."
+" Note that the left hand side of that equation should be logRy,; instead of Ry..."," Note that the left hand side of that equation should be $\log R_{He,f}$ instead of $R_{He,f}$ ."
+" In order to have dynamically stable RLO. so that NS2's progenitor does not coalesce with NSI to form an isolated black hole. the mass ratio M»; to M, needs to be less than 3.5 (Ivanova et al."," In order to have dynamically stable RLO, so that NS2's progenitor does not coalesce with NS1 to form an isolated black hole, the mass ratio $M_{2i}$ to $M_1$ needs to be less than 3.5 (Ivanova et al."
+ 2003)., 2003).
+ Even if NS2's progenitor does not overflow its Roche lobe. there may be mass transfer by stellar wind.," Even if NS2's progenitor does not overflow its Roche lobe, there may be mass transfer by stellar wind."
+ Inaddition. M»; is," Inaddition, $M_{2i}$ is"
+"ealactic halo gas wit1 (re=O8"", 3=0.5) and cluster gas with Gy=391"". 3=0.3). rerealter called ICN2.","galactic halo gas with $r_c=98''$ , $\beta =0.5$ ) and cluster gas with $r_c=394''$, $\beta = 0.3$ ), hereafter called ICM2."
+ The surface xiehtuess edge in NCC L10£ is located at a projected clistan'e of ~I6 kkpe [rom the center of NGC 1399 where he second. (third yeta 11odel comiponeut iu he two (three) component [fit douinates., The surface brightness edge in NGC 1404 is located at a projected distance of $\sim 46$ kpc from the center of NGC 1399 where the second (third) beta model component in the two (three) component fit dominates.
+ Siuce this 'epresents the iniinu physical distauce beween the edge and the center of the dominant cluster elliptical NGC I:399 (axl the overlap between components in this region is small).la we adopt tlie «ensity beta mode rom Equation 3. with paraueters fixed by those components as representative o‘the Fornax cluster gas di our analysis of the edge.," Since this represents the minimum physical distance between the edge and the center of the dominant cluster elliptical NGC 1399 (and the overlap between components in this region is small), we adopt the density beta model from Equation \ref{eq:density} with parameters fixed by those components as representative of the Fornax cluster gas in our analysis of the edge."
+ Ouw data do not probe cluser cistaices large erough to distingui«1 between these two density modes., Our data do not probe cluster distances large enough to distinguish between these two density models.
+ Thus we consider both models in order to itvestigate how uncelalntles lu the cluster ceisity distribution at large clistauces allect our results., Thus we consider both models in order to investigate how uncertainties in the cluster density distribution at large distances affect our results.
+ In order to extract je density. distributio1 from t1le slrlace brightuess proiles. we need to uake additional assumplous about the relative three dimesional geometry of NGC 1101 with 'espect. to NGC 1399.," In order to extract the density distribution from the surface brightness profiles, we need to make additional assumptions about the relative three dimensional geometry of NGC 1404 with respect to NGC 1399."
+ Te comet-like tail tot je southeast «NX NGC 1101 seen ii Figures 1. aud 3 ds evideice [or ram p'essure stripping of gas from NGC |LIOTL., The comet-like tail to the southeast of NGC 1404 seen in Figures \ref{fig:fornax} and \ref{fig:wedgegeom} is evidence for ram pressure stripping of gas from NGC 1404.
+ Since ram pressure stripping is uore likely iu the deuser. central regions of tle cluster (i.e. near the doulnant elliptical gaAXY (GC 1390. We assume tiat NGC T1104 aud NCC 1399 both lie nea“the plane of tle sky such that he projeced distance 9.8 is representative of tjelr plvsical radial separation.," Since ram pressure stripping is more likely in the denser, central regions of the cluster (i.e. near the dominant elliptical galaxy NGC 1399), we assume that NGC 1404 and NGC 1399 both lie near the plane of the sky such that the projected distance $9'.8$ is representative of their physical radial separation."
+ The direction of the all. pointsig to the southeast aud away [roin the cli'ection of m0ion. indicates hat NGC 1101 is moving to the northwest towalc NCC 1399.," The direction of the tail, pointing to the southeast and away from the direction of motion, indicates that NGC 1404 is moving to the northwest toward NGC 1399."
+ Finav. the sharpless of the edee suggests that we are able to directly. view the stagnation region. sch that the inclination of tje motion wih 'espect to the plane of the sky shotld uot be too large (Mazzotta 2001).," Finally, the sharpness of the edge suggests that we are able to directly view the stagnation region, such that the inclination of the motion with respect to the plane of the sky should not be too large (Mazzotta 2001)."
+ These are the satje asstuuplions used in the analysis o ‘the infalliug stbeluster in Abel 3667 (Vikhlinin 2001)., These are the same assumptions used in the analysis of the infalling subcluster in Abell 3667 (Vikhlinin 2001).
+ Our analysis method closely follows that work., Our analysis method closely follows that work.
+" We calculate the surface brightless by integratiig the squareof tje density times the correspoicine X-ray emissivitv. elven in the sixth coluun of Table 1,. along the liue of sight."," We calculate the surface brightness by integrating the square of the density times the corresponding X-ray emissivity, given in the sixth column of Table \ref{tab:spectra}, along the line of sight."
+ The data are fit usi18oO a imnulti-variate chi-square minimization sdeme alowing the position ol the edge (Pedy). deusity power law index inside the edge (a). auil galaxy eectron density at the edee[n] (Aedes) to vary.," The data are fit using a multi-variate chi-square minimization scheme allowing the position of the edge $r_{\rm edge}$ ), density power law index inside the edge $\alpha$ ), and galaxy electron density at the edge $A_{\rm edge}$ ) to vary."
+ We also allow the cluster deusity normalizatior fg in Equation 3)) to be a free parameter to partially compensate [or possible local density [luctlatious in the cluster ICKL whet extracting the cluster electron density near tlie edge., We also allow the cluster density normalization $n_{0}$ in Equation \ref{eq:density}) ) to be a free parameter to partially compensate for possible local density fluctuations in the cluster ICM when extracting the cluster electron density near the edge.
+ In Figuree 5 we show the results for hree clesity models: a power law (PL density inodel for the edge region (1 kpe «ruw reqae) ἰιterior to NGC 1101 with the cluster deusity. inodelled by either [CMI or ΙΟΝΙΟ. and a single power law (SPL) model for r> Lkkpe. with no break at ro [kkpce.," In Figure \ref{fig:sbprof} we show the results for three density models: a power law (PL) density model for the edge region $4$ kpc $< r < r_{\rm edge}$ ) interior to NGC 1404 with the cluster density modelled by either ICM1 or ICM2, and a single power law (SPL) model for $r > 1$ kpc, with no break at $r \sim 4$ kpc."
+ The ICM2 model is used for the cluster density in this latter case., The ICM2 model is used for the cluster density in this latter case.
+ As tle dotted line in Figure 5 shows. the surface brielituess profile for rac TozE kkpe is well reproduce by this simple power law deusity mocel lor the galaxy (SPL+ICM2) with abest fit power law judex a1.51 (corresponding[eJ toa beta model with 6=0.17 for r22 r«). in reasonable agreemelet with the beta moclel results (9= 0.51. r.=0.57 kkpe) lor NGC 1101 by Paolillo," As the dotted line in Figure \ref{fig:sbprof} shows, the surface brightness profile for radii $\lesssim 4$ kpc is well reproduced by this simple power law density model for the galaxy (SPL+ICM2) with abest fit power law index $\alpha = 1.54$ (corresponding toa beta model with $\beta = 0.47$ for $r >> r_c$ ), in reasonable agreement with the beta model results $\beta = 0.51$ , $r_c = 0.57$ kpc) for NGC 1404 by Paolillo"
+The IT» S(0) and S(1) lires are detected in both sets of observations (see Table 1) aud the ditfforeuc“OS in fluxes between the two setsare ~LO%..,The $_2$ S(0) and S(1) lines are detected in both sets of observations (see Table 1) and the differences in fluxes between the two sets are $\sim 10$.
+ This is well within the estimated total error of ~30%.. which is nostly due to uncertainties 1ithe flux calibration.," This is well within the estimated total error of $\sim 30$, which is mostly due to uncertainties in the flux calibration."
+ The $(3) aud $(5) lines are not detected down oa linüt ofνοT)«1015 Cresyes 1 cni (30)., The S(3) and S(5) lines are not detected down to a limit of $\sim (5-7)\times 10^{-15}$ erg $^{-1}$ $^{-2}$ $3\sigma$ ).
+ The spectra of the Πο S(0) and SOL) lines are displaved iu Figme 1 or revolution 668., The spectra of the $_{2}$ S(0) and S(1) lines are displayed in Figure 1 for revolution 668.
+ Since the turbuent and rotational velocities in «lisks are oulv a few kins 3 ves»ectivelv. oth lines are unresolved.," Since the turbulent and rotational velocities in disks are only a few km $^{-1}$, respectively, both lines are unresolved."
+ After sultraction of the continuum. the lines are fitted by Caussiaus wih widths fixed bv the instrumental resolutio rand the line fluxes are compued from 3s to τὸν TWIAL," After subtraction of the continuum, the lines are fitted by Gaussians with widths fixed by the instrumental resolution and the line fluxes are computed from $-3 \times$ to $+3 \times$ HWHM."
+ Iu the optically thin limit. theobserved line fluxes are directv rolated to the populations in the IT; J2 and 3 levels.," In the optically thin limit, theobserved line fluxes are directly related to the populations in the $_2$ $J$ =2 and 3 levels."
+ The derived Iutic temperature assume loca thermodynamic equilibriun (LTE) iS 110+10 K. where the error bar xcflects the tucertainty in tje fluxes.," The derived kinetic temperature assuming local thermodynamic equilibrium (LTE) is $110 \pm 10$ K, where the error bar reflects the uncertainty in the fluxes."
+ Although the 7=3 level has a factor of LO lower populatioi than the J=2 level in gas with a eniperature near LOO K. the radiative transition probability of the οί1) line. Asy=L5«110 1 Jas uuch larger than that of the S(0) transition. Aoy=2.9«10ts Ἐν," Although the $J$ =3 level has a factor of 40 lower population than the $J$ =2 level in gas with a temperature near 100 K, the radiative transition probability of the S(1) line, $A_{31}=4.8\times 10^{-10}$ $^{-1}$ , is much larger than that of the S(0) transition, $A_{20}=2.9\times 10^{-11}$ $^{-1}$."
+ Tu ackition. the spectral resolution at 17 gan ds somewhat higher than that at 28 jan. and the line-to-coutimmiun ratio is correspoucdiely larger.," In addition, the spectral resolution at 17 $\mu$ m is somewhat higher than that at 28 $\mu$ m, and the line-to-continuum ratio is correspondingly larger."
+ Both of these factors explain why the S(1) line is detectable as well., Both of these factors explain why the S(1) line is detectable as well.
+ The Imuüts ou the 839) aud S(5) lines inplv temperatures less than ~260 I& aud [50 Ik. respectively. neelecting anv correction for differential extinction.," The limits on the S(3) and S(5) lines imply temperatures less than $\sim 260$ K and $\sim 450$ K, respectively, neglecting any correction for differential extinction."
+ Siace the ISO beam is much larger than the size of the circumbinary disk. it is imporaut to check. whether anv of the observed Is emission may arise roni residual extended euvelope or clotd 1aaterial.," Since the ISO beam is much larger than the size of the circumbinary disk, it is important to check whether any of the observed $_2$ emission may arise from residual extended envelope or cloud material."
+" Il, lines 1) to S(9) are readily cletected toward embeded Ierde Ac stars with the ISO-SWS (AeBS. van deu Ancker et 11998)."," $_2$ lines up to S(9) are readily detected toward embedded Herbig Ae stars with the ISO-SWS (e.g., van den Ancker et 1998)."
+ In these cases. the emission is clomiuaed by the interaction of the voiug star with its surroundie envelope through shocks ane ultraviole photois.," In these cases, the emission is dominated by the interaction of the young star with its surrounding envelope through shocks and ultraviolet photons."
+ The typical Πω excitation te11]yeratures these regicons are Γι =500700 TS. 1iuchi arecr than the value found for GG Tau.," The typical $_2$ excitation temperatures for these regions are $T_{\rm exc}$ =500–700 K, much larger than the value found for GG Tau."
+" Deep ISO searches for II, S(O) aid 51) lines toward diffuse iux| traushiceit clous have been pertormed by Thi et ((1909", Deep ISO searches for the $_2$ S(0) and S(1) lines toward diffuse and translucent clouds have been performed by Thi et (1999b).
+ The lines are not detected down to 2x10H Cres 1  (20) iu gas with deusities of a few hundre to a few tvousane cin? exposed to the πουλια] interstellar radiation field., The lines are not detected down to $2\times 10^{-14}$ erg $^{-1}$ $^{-2}$ $2\sigma$ ) in gas with densities of a few hundred to a few thousand $^{-3}$ exposed to the normal interstellar radiation field.
+ The 0ο 10 emission arond (αζα Tau is less than 50 mls 30 ) (Skrutskieetal. 19933). more than an order of magnitude lower tha- found for diffuse clouds such as that toward ¢ Oph.," The $^{12}$ CO 1–0 emission around GG Tau is less than 50 mK $\sigma$ ) \cite{Skr}) ), more than an order of magnitude lower than found for diffuse clouds such as that toward $\zeta$ Oph."
+ The corresponding mass iu the SWS beams is estimated to be less than a few «10.1 AL. sienificautly lower than the mass derived from the IT»lines (see below).," The corresponding mass in the SWS beam is estimated to be less than a few $\times 10^{-4}$ $_{\odot}$ , significantly lower than the mass derived from the $_2$lines (see below)."
+" IKkeck images of GC Tau in A"" continu aud I emission iu the e-1 20 and e=2 51 S(1) lines at", Keck images of GG Tau in $K'$ continuum and $_2$ emission in the $v$ $\to$ 0 and $v$ $\to$ 1 S(1) lines at
+stronely indicative of the substantial advantage such a wide field instrument has when it comes to determining the extent of a population of transients svstems.,strongly indicative of the substantial advantage such a wide field instrument has when it comes to determining the extent of a population of transients systems.
+ Though parts of the SAIC have been regularly monitored by RAPE over a decade (Galache et al..," Though parts of the SMC have been regularly monitored by RXTE over a decade (Galache et al.,"
+. 2008). the smaller Ποια of view. combined with the collimator response makes it dillieult exactly to quantify the duty eveles of these ΛΙΝ systems.," 2008), the smaller field of view, combined with the collimator response makes it difficult exactly to quantify the duty cycles of these HMXB systems."
+ Other wide field instruments such as the Swift/D.NT simply to do not have the combined spatial ancl (ux sensitivity. of INTPECGRAL/IBIS., Other wide field instruments such as the Swift/BAT simply to do not have the combined spatial and flux sensitivity of INTEGRAL/IBIS.
+ Thus the conclusion of the INPECGRAL results presented here is that the total population of LIAINBs in the SAIC is probably much greater than the 6tFTU ο. so far catalogued. with consequent implications for the star formation rate in this galaxy.," Thus the conclusion of the INTEGRAL results presented here is that the total population of HMXBs in the SMC is probably much greater than the 60-70 systems so far catalogued, with consequent implications for the star formation rate in this galaxy."
+ Plus. there is probably. a significant population along the Bridge - that may decline in number (along with the age of the stellar population) as one moves [rom the SAIC to the LMC.," Plus, there is probably a significant population along the Bridge - that may decline in number (along with the age of the stellar population) as one moves from the SMC to the LMC."
+ Future IVPEGRAL observations should continue to enhance our understanding ofall parts of the Magellanic system., Future INTEGRAL observations should continue to enhance our understanding of all parts of the Magellanic system.
+ The OGLE project was partially supported by the Polish AINISW. grant N20303032/4275., The OGLE project was partially supported by the Polish MNiSW grant N20303032/4275.
+. LJ'E is supported by a AlavHower Scholarship from the University of Southampton., LJT is supported by a Mayflower Scholarship from the University of Southampton.
+The work is supported in part by the Once-Huncdred-‘Talent Proegran.. by NABRSE(G19990754). by NSEC(Νο.10043004). and by SED375.,"The work is supported in part by the One-Hundred-Talent Program, by NKBRSF(G19990754), by NSFC(No.10043004), and by SFB375."
+ We thank an anonymous referee for helpful comments., We thank an anonymous referee for helpful comments.
+"our data, a double Gaussian distribution is preferred to a single Gaussian distribution (y?=1.141 and y2=2.784, respectively).","our data, a double Gaussian distribution is preferred to a single Gaussian distribution $\chi^2_\nu = 1.141$ and $\chi^2_\nu =2.784$, respectively)."
+ Fitting a double Gaussian function to the data results in a blue peak at (V—R)o=0.46 mag and a red peak at (V—R)g=0.58 mag., Fitting a double Gaussian function to the data results in a blue peak at $(V-R)_0 = 0.46$ mag and a red peak at $(V-R)_0 = 0.58$ mag.
+" For a 13-Gyr population, these peaks correspond to metallicities of [Fe/H]~—2.2 dex and [Fe/H]«—0.4 dex, respectively (?).."," For a 13-Gyr population, these peaks correspond to metallicities of $\mathrm{[Fe/H]} \approx -2.2$ dex and $\mathrm{[Fe/H]} \approx -0.4$ dex, respectively \citep{2003MNRAS.344.1000B}."
+" These estimates are consistent with the values derived for metal-poor and GCs in a number of giant elliptical galaxies (?,andrefer- therein)..", These estimates are consistent with the values derived for metal-poor and metal-rich GCs in a number of giant elliptical galaxies \citep[][and references therein]{2006ARA&A..44..193B}.
+" Given this, we define for the following analyses (Sect. 4.2))"," Given this, we define for the following analyses (Sect. \ref{sec:kinematics}) )"
+" a blue (metal-poor) and a red (metal-rich) sub- with (V—R)o<0.5 mag and (V—R)o=0.5 mag, respectively."," a blue (metal-poor) and a red (metal-rich) sub-population with $(V-R)_0 < 0.5$ mag and $(V-R)_0 \geq 0.5$ mag, respectively."
+ The separating colour corresponds to a metallicity of [Fe/H]«—1.4 dex., The separating colour corresponds to a metallicity of $\mathrm{[Fe/H]} \approx -1.4$ dex.
+" With (V—R)o=0.60 mag, the brightest UCD in our sample (HUCD1) clearly belongs to the metal-rich sub-population and is located at the tip of the red GC sequence in the CMD."," With $(V-R)_0 = 0.60$ mag, the brightest UCD in our sample (HUCD1) clearly belongs to the metal-rich sub-population and is located at the tip of the red GC sequence in the CMD."
+" This confirms results of photometric studies, where the red GC sequence is found to extend to higher luminosities than the blue sequence (e.g.??).. "," This confirms results of photometric studies, where the red GC sequence is found to extend to higher luminosities than the blue sequence \citep[e.g.][]{2008ApJ...681.1233W, 2010ApJ...710.1672M}."
+"For 26 of the identified cluster GCs/UCDs, imaging is available in the HST archive."," For 26 of the identified cluster GCs/UCDs, imaging is available in the HST archive."
+" Two WFPC2 fields were observed in Cycle 6 with the PCI chip centred on NGC 3311 and NGC 3309, respectively."," Two WFPC2 fields were observed in Cycle 6 with the PC1 chip centred on NGC 3311 and NGC 3309, respectively."
+" The exposure times were 3700 s in F555W and 3800 s in F814W for the field centred on NGC 3311, and 4400 s in both filters for the NGC 3309 field (HST programme GO.06554.01-95A, PI: J.P. Brodie, see also ?))."," The exposure times were $3700$ s in $F555W$ and $3800$ s in $F814W$ for the field centred on NGC 3311, and $4400$ s in both filters for the NGC 3309 field (HST programme GO.06554.01-95A, PI: J.P. Brodie, see also \citealt{2000ApJ...543L..19B}) )."
+" The F555W and F814W images were bias, flat-field and bad pixel corrected."," The F555W and F814W images were bias, flat-field and bad pixel corrected."
+" The sub-pixel dithered imaging allowed us to eliminate cosmic-rays, hot pixels and improve the spatial resolution."," The sub-pixel dithered imaging allowed us to eliminate cosmic-rays, hot pixels and improve the spatial resolution."
+" To perform a PSF photometry on each individual WFPC2 detector, we created a grid of 100x PSFs with the software package (?).."," To perform a PSF photometry on each individual WFPC2 detector, we created a grid of $100\times100$ PSFs with the software package \citep{1995ASPC...77..349K}."
+ This library was used to create a spatially variable PSF model and perform PSF-fitting photometry with the task in IRAF., This library was used to create a spatially variable PSF model and perform PSF-fitting photometry with the task in IRAF.
+" Isolated stars were used to determine an aperture correction to 10 aperture diameter, which is the aperture used by ? to derive the most up to date CTE corrections and photometric zero points for each of the WFPC2 chips."," Isolated stars were used to determine an aperture correction to $1 \farcs 0$ aperture diameter, which is the aperture used by \citet{2009PASP..121..655D} to derive the most up to date CTE corrections and photometric zero points for each of the WFPC2 chips."
+" Finally, the VEGAMAG WFPC2 photometry was corrected for foreground extinction using ? dust maps towards the direction of both galaxies, and ? relations for Ry—3.1 to calculate the absorption at the effective wavelengths for the WFPC2 filters."," Finally, the VEGAMAG WFPC2 photometry was corrected for foreground extinction using \citet{1998ApJ...500..525S} dust maps towards the direction of both galaxies, and \citet{1989ApJ...345..245C} relations for $R_V=3.1$ to calculate the absorption at the effective wavelengths for the WFPC2 filters."
+ The following dereddening values were applied: Arsssw=0.261 mag and Arsiaw=0.156 mag., The following dereddening values were applied: $A_{F555W}=0.261$ mag and $A_{F814W}=0.156$ mag.
+" To transform the VEGAMAG WFPC2 magnitudes to the standard Johnson/Cousins magnitudes, we used the ? transformations with the coefficients in their table 4 for the (V—J) colour."," To transform the VEGAMAG WFPC2 magnitudes to the standard Johnson/Cousins magnitudes, we used the \citet{2009PASP..121..655D} transformations with the coefficients in their table 4 for the $(V-I)$ colour."
+" At a the adopted distance of 47.2 Mpc (m-M= mag), one PC1 and WF2,3,4 pixel corresponds to 11 and"," At a the adopted distance of $47.2$ Mpc $m-M = 33.37$ mag), one PC1 and WF2,3,4 pixel corresponds to 11 and"
+formation.,formation.
+ Iu this scenario. manv of the initially eas-rich. star forming galaxies iu the maim cluster would lave consumed much of their available eas in an earlier epoch of triggered star formation. explaining the lack of late-type galaxies in our current sample.," In this scenario, many of the initially gas-rich, star forming galaxies in the main cluster would have consumed much of their available gas in an earlier epoch of triggered star formation, explaining the lack of late-type galaxies in our current sample."
+ However. the relative lack of late-tvpe galaxies cannot explain why we do not see a signal across the shock front. especially considering that even among the late-tvpe population. we do nof detect a significant difference iu color in the post-shock versus pre-shock galaxies.," However, the relative lack of late-type galaxies cannot explain why we do not see a signal across the shock front, especially considering that even among the late-type population, we do not detect a significant difference in color in the post-shock versus pre-shock galaxies."
+" To trace specific star formation rate. we use IRAC color m, 5-1us."," To trace specific star formation rate, we use IRAC color $_{4.5}$ $_{8}$."
+ The Syn fux traces star formation via the PATI emission feature at rest wavelength 6.2521. which is redshifted into the Sida baud at the cluster redshift.," The $\micron$ flux traces star formation via the PAH emission feature at rest wavelength $\micron$, which is redshifted into the $\micron$ band at the cluster redshift."
+ The 1.5;12 fux can be taken as a proxy for stellar nass., The $\micron$ flux can be taken as a proxy for stellar mass.
+ Using spectroscopic redshifts. we confirmed 63 out of 200 TRAC sources (169 in direct path of shock frout | 31 above /helowthe shock front) as cluster members. aud found 12 IRAC sources that are interlopers.," Using spectroscopic redshifts, we confirmed 63 out of 200 IRAC sources (169 in direct path of shock front + 31 above/below the shock front) as cluster members, and found 42 IRAC sources that are interlopers."
+ Applving a magnitude cut at ups=17. we are left with 158 cluster candidates. of which 131 are in the direct path of the shock front.," Applying a magnitude cut at $_{4.5}=17$, we are left with 158 cluster candidates, of which 134 are in the direct path of the shock front."
+ We observe no sienificaut treucl i nij5-& ax a uction of distance from the shock frout., We observe no significant trend in $_{4.5}$ $_{8}$ as a function of distance from the shock front.
+ The KS-test reveals that the sample of galaxies behind the shock versus ahead of the shock front. ouly differs im color by ess than 0σ. for galaxies with various cuts in nij 5-1s color.," The KS-test reveals that the sample of galaxies behind the shock versus ahead of the shock front only differs in color by less than $\sigma$, for galaxies with various cuts in $_{4.5}$ $_{8}$ color."
+ This indicates that even in a dramatic mereer event ike the Bullet Cluster. the rain pressure induced frou a supersonic collision does not drastically trigger or quench current/recent star formation.," This indicates that even in a dramatic merger event like the Bullet Cluster, the ram pressure induced from a supersonic collision does not drastically trigger or quench current/recent star formation."
+ A possible explanation is hat the rami pressure can remove the gas reservoirs in the outer disk. while preserving star formation iu the central disk of a galaxy.," A possible explanation is that the ram pressure can remove the gas reservoirs in the outer disk, while preserving star formation in the central disk of a galaxy."
+ While ecometric effects; contamination roni iuterlopers. aud the lack of late-tvpe galaxies in ιο Bullet Cluster could contribute to diluting a signal. rev are unlikely to explain the complete lack of signal we observe.," While geometric effects, contamination from interlopers, and the lack of late-type galaxies in the Bullet Cluster could contribute to diluting a signal, they are unlikely to explain the complete lack of signal we observe."
+ Even among our small sample of late-tvpe ealaxies. we detect no change in color between the shock aud post-shock galaxies to within 1.50.," Even among our small sample of late-type galaxies, we detect no change in color between the pre-shock and post-shock galaxies to within $1.5\sigma$."
+ The authors would hike to thauk Rafael Barrena for sharing his redshift datawith us., The authors would like to thank Rafael Barrena for sharing his redshift datawith us.
+ We also acknuoxledge support for this work fro) NASA/Spitzer eraut 1319111., We also acknowledge support for this work from NASA/Spitzer grant 1319141.
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (A," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (AC," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (ACT," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (ACTS," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (ACTS-," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (ACTS-T," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (ACTS-T)," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (ACTS-T).," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+ MM was supported by NASA contract NÀSS- and Chandra eraut. GOS-9128X. Cliaistine Jones acknowledges support from Spitzer Contract 1265581. (IRAC).. (ACTS-T)..," MM was supported by NASA contract NAS8-39073 and Chandra grant GO8-9128X. Christine Jones acknowledges support from Spitzer Contract 1265584. , ,"
+Ouce the stellar wobble has been confirmed. it could still be due to another mechanisin than the pull by a companion.,"Once the stellar wobble has been confirmed, it could still be due to another mechanism than the pull by a companion."
+ As shown by Schucider (1999) aud coutirmed w Morais Correia (2011). if can be due to a distant jnary star with the same orbital period as for the plauct (false) candidate.," As shown by Schneider (1999) and confirmed by Morais Correia (2011), it can be due to a distant binary star with the same orbital period as for the planet (false) candidate."
+ But such events can Πίο only low nass planets on wide orbits. bevoud the present detection capability.," But such events can mimic only low mass planets on wide orbits, beyond the present detection capability."
+ It does not apply to objects currently present in the catalogue., It does not apply to objects currently present in the catalogue.
+ It nevertheless represents a dauegcr for utire discoveries with the Gaia astrometric space mission (Schneider Cabrera 2006)., It nevertheless represents a danger for future discoveries with the Gaia astrometric space mission (Schneider Cabrera 2006).
+ The final discrimination between this artefact aud a real companion will be provided by the detection of the candidate by direct spectro-imaging., The final discrimination between this artefact and a real companion will be provided by the detection of the candidate by direct spectro-imaging.
+ The same discussion holds for planets discovered by astrometry aud by timine. -, The same discussion holds for planets discovered by astrometry and by timing. -
+ The iain possible artefact is the preseuce of a backeround eclipsing binary (BEB) in the target Point Spread Function., The main possible artefact is the presence of a background eclipsing binary (BEB) in the target Point Spread Function.
+ A first test is then to check. if possible. that the transit is achromatic (we leave aside the discussion of self-buninous transiting planets) and that a BEB is not seen in high augular resolutiou Huases.," A first test is then to check, if possible, that the transit is achromatic (we leave aside the discussion of self-luminous transiting planets) and that a BEB is not seen in high angular resolution images."
+ A second step is then to measure the radial velocity variations due to the suspected trausiting object2005))., A second step is then to measure the radial velocity variations due to the suspected transiting object.
+ Ouce the trausit has been secured. it cau still mimic a Jupiter-zed planet if it is due to a brown dwarf. or a super-Earth if it is due to a white dwiuf.," Once the transit has been secured, it can still mimic a Jupiter-sized planet if it is due to a brown dwarf, or a super-Earth if it is due to a white dwarf."
+" is removed by RV measurements which give the companion Lass, -", is removed by RV measurements which give the companion mass. -
+ The main possible artefact is the confusion with a vackeround star., The main possible artefact is the confusion with a background star.
+ It can easily be removed bv verifviug hat the star aud the companion have a como proper notion., It can easily be removed by verifying that the star and the companion have a common proper motion.
+ Some authors nevertheless publish a planctary candidate το be confirmed”. prior to the check of common oxoper motion. when the probability of presence of a vackeround star in the tareet vicinity is “suffiicicuth” ow (e.g. Lagrange et al (," Some authors nevertheless publish a planetary candidate ""to be confirmed"", prior to the check of common proper motion, when the probability of presence of a background star in the target vicinity is ""sufficiently"" low (e.g. Lagrange et al. ("
+2008) for ) Pic b).,2008) for $\beta$ Pic b).
+" There is 10 commonly accepted value of the FAP for these cases and the decision between ~coufirmed” iud uucoufinued"" planet is arbitrary."," There is no commonly accepted value of the FAP for these cases and the decision between ""confirmed"" and ""unconfirmed"" planet is arbitrary."
+ Ilere again. there are two aspects: - Once the mass limit has been decided. au uucertainty ronuaius since the measured mass suffers from various inaccuracies.," Here again, there are two aspects: - Once the mass limit has been decided, an uncertainty remains since the measured mass suffers from various inaccuracies."
+ Tt introduces a fuzziness on what objects to include in the catalogue., It introduces a fuzziness on what objects to include in the catalogue.
+ For all type of methods there is zu dustrumental uncertainty., For all type of methods there is an instrumental uncertainty.
+" But iu case of planets detected by RV or bv timine here is an additional “sin/ uncertainty. meaning that he ouly observable is then Mj.sin/ istead of M, itself. due to the uuknown orbital iuclinatiou /."," But in case of planets detected by RV or by timing there is an additional $\sin i$ uncertainty"", meaning that the only observable is then $M_{pl}.\sin i$ instead of $M_{pl}$ itself, due to the unknown orbital inclination $i$."
+" The database ists all objects with M.sini less than 25 Mu,", The database lists all objects with $M.\sin i$ less than 25 $M_{Jup}$.
+ As it is well know. this uncertantv does not hok for transiting dlanets since / i8 then derived from the shape of the transit ight curve.," As it is well know, this uncertainty does not hold for transiting planets since $i$ is then derived from the shape of the transit light curve."
+ For planets detected by maging. constraints ou their mass cau be put thanks to multicolor photometry or spectroscopy and to atimosphere modelling.," For planets detected by imaging, constraints on their mass can be put thanks to multicolor photometry or spectroscopy and to atmosphere modelling."
+ In that case the unecrtainty is rather large (see for instance Neuhacuser et al 2005)., In that case the uncertainty is rather large (see for instance Neuhaeuser et al 2005).
+ The mass estimate them rests most often on the Daraffe ct al., The mass estimate then rests most often on the Baraffe et al.
+ model (2010) correlating the mass to the spectrum and age., model (2010) correlating the mass to the spectrum and age.
+ It is aportant to note that this model has not been tested vet for plauets with a lass kuown from radial velocity measurements. -, It is important to note that this model has not been tested yet for planets with a mass known from radial velocity measurements. -
+ For rigorous conipleteness one has to consider the possibility that less than 25 την compact objects are uot plaucts., For rigorous completeness one has to consider the possibility that less than 25 $M_{Jup}$ compact objects are not planets.
+ Preseutlv there is no known category of such objects., Presently there is no known category of such objects.
+ But. at least in the carly era of the first planet candidates when the abundance of exoplauets was colpletely unknown. oue could wouder if a new class of objects was not discovered instead.," But, at least in the early era of the first planet candidates when the abundance of exoplanets was completely unknown, one could wonder if a new class of objects was not discovered instead."
+ For instance oue could in principle invoke black holes or X-stur. analogous to neutron stars. where X is some new type of particle.," For instance one could in principle invoke black holes or ""X-stars"", analogous to neutron stars, where X is some new type of particle."
+ Planetary mass black holes have an evaporation time TDH=TpMpu/Mpnaa)!zWP? ��μ... (Yaug Chane 2009) and therefore survive evaporatiou over stellar lifetimes., Planetary mass black holes have an evaporation time $\tau _{BH} = \tau_{Planck}(M_{BH}/M_{Planck})^3 \approx 10^{63}(M_{BH}/M_{Jup})^3$ s (Yang Chang 2009) and therefore survive evaporation over stellar lifetimes.
+ But. having a (Scharzshild) radius of he the order of meters. they are made implausible by the observed radius derived from transit events.," But, having a (Scharzshild) radius of the the order of meters, they are made implausible by the observed radius derived from transit events."
+" N-stars would rave a mass Ay.=ALPο requiring a 30 neutron nass X particle for à 1M;,, object."," X-stars would have a mass $M_{X*} = M_{Planck}^3/M_X^2$, requiring a 30 neutron mass X particle for a $1 M_{Jup}$ object."
+ It is not excuded by he zoo of “hevoud the standard model” theories. but it would imply a discrete mass distribution with oulv a very ew dnass values (one per X-species).," It is not excluded by the zoo of ""beyond the standard model"" theories, but it would imply a discrete mass distribution with only a very few mass values (one per X-species)."
+ These perspectives are (or were) au müiportant issue eiven the philosophical significance of exoplauets aud that “extraordinary clainis require extraordinary proofs”.," These perspectives are (or were) an important issue given the philosophical significance of exoplanets and that ""extraordinary claims require extraordinary proofs""."
+ But of course today the case is Closed., But of course today the case is closed.
+ Iu couclusion. we have arbitrarily chosen a 25 ην | AA/imass limit with a | sigma imarein for AAS.," In conclusion, we have arbitrarily chosen a 25 $M_{Jup}$ + $\Delta M$ mass limit with a 1 sigma margin for $\Delta M$."
+ By having a generous ass limit. we allow the user to conipare casily planets and brown dwarts.," By having a ""generous"" mass limit, we allow the user to compare easily planets and brown dwarfs."
+ Here again a sharp ass hit. evenwith a l sigma niargin. is absurd since it excludes conipiuions with Af - AM = 25.1 AL yy).," Here again a sharp mass limit, evenwith a 1 sigma margin, is absurd since it excludes companions with $M$ - $\Delta M$ = 25.1 $M_{Jup}$ ."
+sense that older galaxies have more Salpeter-like IAIFs.,sense that older galaxies have more Salpeter-like IMFs.
+ As noted by Napolitanoetal.(2010).. this is the trend than is expected from our relatious.," As noted by \citet{napolitano}, this is the trend than is expected from our relations."
+ The cause of this discrepancy is unclear but may be related to how the samples are chosen., The cause of this discrepancy is unclear but may be related to how the samples are chosen.
+ The SLAC'S lenses are at higher redshifts than the Napolitanoetal.(2010) ealaxies aud are selected to have spectra that match old stellar templates: they would therefore ouly span ages vetween approximately 5 aud 10 Cr. for which the aec-IME treud is esseutially flat.," The SLACS lenses are at higher redshifts than the \citet{napolitano} galaxies and are selected to have spectra that match old stellar templates; they would therefore only span ages between approximately 5 and 10 Gyr, for which the age-IMF trend is essentially flat."
+ Likewise. Graves&Faber(2010) use trends between nass and mictallicity (Fe/II[. |Mg/II]. aud |Meg/Fo]) to investigate the role of the IME on the central dark natter fraction.," Likewise, \citet{graves} use trends between mass and metallicity ([Fe/H], [Mg/H], and [Mg/Fe]) to investigate the role of the IMF on the central dark matter fraction."
+ They show that the e-eubhanuceimeut of galaxies (c.g. the |Mg/Foe| ratio) is mass-depeudent and could therefore signal a more top-heavy IAIF in nore massive galaxies such that the relative number of corce-collapse supernovae increases du niore luassive ealaxies.," They show that the $\alpha$ -enhancement of galaxies (e.g., the [Mg/Fe] ratio) is mass-dependent and could therefore signal a more top-heavy IMF in more massive galaxies such that the relative number of core-collapse supernovae increases in more massive galaxies."
+ However. invokiug more supernovae vields to siper-abunudant metallicities compared to the data. aud another mechanisii must therefore be posited in order to remove these metals (Graves&Faber2010).," However, invoking more supernovae yields to super-abundant metallicities compared to the data, and another mechanism must therefore be posited in order to remove these metals \citep{graves}."
+". We have demonstrated. that. eiven our standard assumnptious. the IME of massive ETCs caunot be ""light. even in the presence of significant adiabatic contraction."," We have demonstrated that, given our standard assumptions, the IMF of massive ETGs cannot be “light”, even in the presence of significant adiabatic contraction."
+ Furthermore. the IME is oulv consisteut with being universal if the central dark matter profile is mass dependent. as in the AC models preseuted here.," Furthermore, the IMF is only consistent with being universal if the central dark matter profile is mass dependent, as in the AC models presented here."
+ It is clear that if one wants to preserve a light aud universal IME one las to abandon standard assuniptious about dark matter halos. like the NEW profile.," It is clear that if one wants to preserve a light and universal IMF one has to abandon standard assumptions about dark matter halos, like the NFW profile."
+ However. there are two inportant caveats.," However, there are two important caveats."
+ First. better constraints on the star-formation efficiency must be obtained frou the data in order to draw definitive conclusions about the role of a iass-dependent IME relative to CDM halo contraction.," First, better constraints on the star-formation efficiency must be obtained from the data in order to draw definitive conclusions about the role of a mass-dependent IMF relative to CDM halo contraction."
+ Second. although our conclusions are robust with respect to a wide variety of changes in our asstuuptious. we have not exhaustedafl possible families of theoretical models.," Second, although our conclusions are robust with respect to a wide variety of changes in our assumptions, we have not exhausted possible families of theoretical models."
+ We will address these caveats. quantity the relationship between the ceutral dark matter profile aud the IME in more detail aud quantity the consequences of our various assumptions iu a forthcoming paper (Augeretal.90100.inpreparation).," We will address these caveats, quantify the relationship between the central dark matter profile and the IMF in more detail, and quantify the consequences of our various assumptions in a forthcoming paper \citep[][in preparation]{haloimf}."
+ We thauk the referee for his/her helpful comuucuts. and. Aarou Dutton and Matteo Barnabe for their sugeestions.," We thank the referee for his/her helpful comments, and Aaron Dutton and Matteo Barnabè for their suggestions."
+ We acknowledge support from the NSF through CAREER award NSF-0612621. from the Packard Foundation through a Packard Fellowship to TT. and from NASA. through HIST exauts 11202. 107985. and 10191.," We acknowledge support from the NSF through CAREER award NSF-0642621, from the Packard Foundation through a Packard Fellowship to TT, and from NASA, through HST grants 11202, 10798, and 10494."
+ RG is supported by the Centre National des Etudes Spatiales., RG is supported by the Centre National des Etudes Spatiales.
+"(Ahrens,1954),, resulting in asymmetric error bars according to the skewness of the distribution.","\citep{ahrens54a}, resulting in asymmetric error bars according to the skewness of the distribution."
+ Note that consequently & is unitless., Note that consequently $\tilde \sigma$ is unitless.
+" With these definitions, the best-fit function has a median value of (x)g;=50.3 ccps and a multiplicative standard deviation of Gm=1.8."," With these definitions, the best-fit function has a median value of $\left< x\right>_\text{fit} =
+50.3$ cps and a multiplicative standard deviation of $\tilde{\sigma}_\text{fit} =
+1.8$."
+" Compared with that the measured count rates have a median flux of (x)meas= 48.2ccps and a multiplicative standard deviation of Gmeas=2.0, similar to the fit."," Compared with that the measured count rates have a median flux of $\left< x\right>_\text{meas} = 48.2$ cps and a multiplicative standard deviation of $\tilde \sigma_\text{meas}=2.0 $, similar to the fit."
+ The largest cause of source-independent variation that affects the data are background fluctuations., The largest cause of source-independent variation that affects the data are background fluctuations.
+" To obtain an estimate for the background noise level, we extracted a lightcurve in the kkeV energy band from a region in the sky without any known X-ray source (a= 9h16m 37s, 6=—44*07' "," To obtain an estimate for the background noise level, we extracted a lightcurve in the keV energy band from a region in the sky without any known X-ray source $\alpha = 9$ $\,16$ $\,37$ s, $\delta = -44\degr 07\arcmin 11\arcsec$ )."
+"The brightness distribution of this background is shown 11"").with a dotted line in Fig. 2..", The brightness distribution of this background is shown with a dotted line in Fig. \ref{fig:lcratehistall.ps}.
+" The median value of the background distribution is (x)yy,31.5 ccps, which is below of the source data."," The median value of the background distribution is $\left<x\right>_\text{bkg}\approx$ cps, which is below of the source data."
+" Note that the brightest data points of the background reach values of up to ccps, which is in the lower left flank of the source distribution."," Note that the brightest data points of the background reach values of up to cps, which is in the lower left flank of the source distribution."
+" They do not, however, change the overall shape and parameters of the distribution, so that a significant background influence can be ruled out."," They do not, however, change the overall shape and parameters of the distribution, so that a significant background influence can be ruled out."
+ We have seen in the previous section that deviations from a log-normal distribution are visible., We have seen in the previous section that deviations from a log-normal distribution are visible.
+" These deviations could be induced by averaging over all orbital phases, where possibly not all phases show the same statistical brightness variations."," These deviations could be induced by averaging over all orbital phases, where possibly not all phases show the same statistical brightness variations."
+" Systematic influences could be due to the complex and turbulent accretion geometry, consisting of a bow shock, an accretion wake, a possible photoionization wake or even a tidal stream (see,e.g.,Blondinetal.,1990,1991;Mauche2008).."," Systematic influences could be due to the complex and turbulent accretion geometry, consisting of a bow shock, an accretion wake, a possible photoionization wake or even a tidal stream \citep[see, e.g.,][]{blondin90a, blondin91a, mauche08a}."
+" An accretion wake in Vela X-1 was already postulated by Eadieetal.(1975),, showing up as absorption dips in the lightcurve of the Ariel V Sky Survey Experiment."," An accretion wake in Vela X-1 was already postulated by \citet{eadie75a}, showing up as absorption dips in the lightcurve of the Ariel V Sky Survey Experiment."
+" To analyze the orbital influences, we extracted orbital phase resolved histograms."," To analyze the orbital influences, we extracted orbital phase resolved histograms."
+" The data set comprises ~8 orbits overall, which allows us to extract histograms with Ao;=0.05."," The data set comprises $\sim$ 8 orbits overall, which allows us to extract histograms with $\Delta \phi_\text{orbit} = 0.05$."
+ We reduced the number of logarithmically spaced bins to 96 to collect more than 20 values in most bins to satisfy the requirements for y? statistics., We reduced the number of logarithmically spaced bins to 96 to collect more than 20 values in most bins to satisfy the requirements for $\chi^2$ statistics.
+" Figure 3 shows color-coded the probability of finding a datapoint in a countrate bin for a given histogram, equivalent to the height of the histogram in Fig. 2.."," Figure \ref{fig:phasres_histo.eps} shows color-coded the probability of finding a datapoint in a countrate bin for a given histogram, equivalent to the height of the histogram in Fig. \ref{fig:lcratehistall.ps}."
+" Additionally the average countrate in every bin is shown, together with its uncertainty."," Additionally the average countrate in every bin is shown, together with its uncertainty."
+" As the average countrate can be calculated from a large number of data points in the lightcurve, these uncertainties are rather small."," As the average countrate can be calculated from a large number of data points in the lightcurve, these uncertainties are rather small."
+ Significant variations from phase bin to phase bin can be seen in the figure., Significant variations from phase bin to phase bin can be seen in the figure.
+ Especially notable is a dip in the average countrate around phase ¢~0.3—0.4 and a sudden brightening to twice the value afterwards., Especially notable is a dip in the average countrate around phase $\phi\approx0.3$ $0.4$ and a sudden brightening to twice the value afterwards.
+ Goldsteinetal.(2004) analyzed data taken at Φου=0.25 and ¢oy=0.5 of one orbit and measured a stronger absorbed spectrum in the do=0.5 data., \citet{goldstein04a} analyzed data taken at $\phi_\text{orb}=0.25$ and $\phi_\text{orb}=0.5$ of one orbit and measured a stronger absorbed spectrum in the $\phi_\text{orb}=0.5$ data.
+ These authors explained the increased absorption by the accretion wake being in the light of sight during that episode., These authors explained the increased absorption by the accretion wake being in the light of sight during that episode.
+ The dip in Fig., The dip in Fig.
+ 3 occurs slightly before ¢=0.5., \ref{fig:phasres_histo.eps} occurs slightly before $\phi=0.5$ .
+" If an accretion wake is responsible for this feature, it seems to have shifted to earlier orbital phases in our data."," If an accretion wake is responsible for this feature, it seems to have shifted to earlier orbital phases in our data."
+" When analyzing the three data blocks (Table 1)) separately an orbital modulation is also visible, although with different strength in the different blocks."," When analyzing the three data blocks (Table \ref{tab:blocks}) ) separately an orbital modulation is also visible, although with different strength in the different blocks."
+ This indicates a strongly variable feature rather than an orbitally stable one and suggests that the variation seen in Fig., This indicates a strongly variable feature rather than an orbitally stable one and suggests that the variation seen in Fig.
+" 3 is biased by single orbits probably deviating strongly from the average orbital histogram, maybe due to a particularly active or quiet state of the system."," \ref{fig:phasres_histo.eps} is biased by single orbits probably deviating strongly from the average orbital histogram, maybe due to a particularly active or quiet state of the system."
+ The influence of different Ny values during the orbit can be investigated by analyzing data in softer X-rays., The influence of different $N_\text{H}$ values during the orbit can be investigated by analyzing data in softer X-rays.
+" Softer energy bands are far more affected by photoabsorption than the ISGRI band, because the cross-section of the photoabsorption is declining quickly with increasing energy as &«E~>."," Softer energy bands are far more affected by photoabsorption than the ISGRI band, because the cross-section of the photoabsorption is declining quickly with increasing energy as $\sigma \propto
+E^{-3}$."
+" A very good data set in the soft X-rays is provided by the All-Sky Monitor (ASM,Levineetal.,1996) aboard the “Rossi X-Ray Timing Explorer"" (RXTE))."," A very good data set in the soft X-rays is provided by the All-Sky Monitor \citep[ASM,][]{levine96a} aboard the “Rossi X-Ray Timing Explorer” )."
+" The ASM performs ssec dwells on an irregular basis onto different sources, including Vela X-1, and thus provides good statistics in the three energy bands between kkeV called A kkeV), B kkeV), and C kkeV) band."," The ASM performs sec dwells on an irregular basis onto different sources, including Vela X-1, and thus provides good statistics in the three energy bands between keV called A keV), B keV), and C keV) band."
+" We performed a similar analysis of the ASM data as of the ISGRI data, but binning the countrate directly to 256 bins instead of its logarithm."," We performed a similar analysis of the ASM data as of the ISGRI data, but binning the countrate directly to 256 bins instead of its logarithm."
+ The linear binning was necessary because especially in the eclipse many data points of the lightcurve are negative and cannot be presented ona logarithmic scale., The linear binning was necessary because especially in the eclipse many data points of the lightcurve are negative and cannot be presented ona logarithmic scale.
+ It has previously been shown that due to, It has previously been shown that due to
+are nearly free from scattered light and show a 1.6 - 2.1 times higher contrast than the untreated images. while the contrast of BPs increases by a factor of 1.8 - 2.8.,"are nearly free from scattered light and show a 1.6 - 2.1 times higher contrast than the untreated images, while the contrast of BPs increases by a factor of 1.8 - 2.8."
+ The stronger enhancement of the BP contrast is due to their smaller size than granules. which are dominantly responsible for the contrast of the whole image.," The stronger enhancement of the BP contrast is due to their smaller size than granules, which are dominantly responsible for the contrast of the whole image."
+As a supernova shock wave propagates through the progenitor star. it eventually emerges through the outer envelope region which has a low optical depth.,"As a supernova shock wave propagates through the progenitor star, it eventually emerges through the outer envelope region which has a low optical depth."
+" At this point the shock “breaks out"" of the star. initiating the first electromagnetic signal of the explosion."," At this point the shock “breaks out” of the star, initiating the first electromagnetic signal of the explosion."
+ The breakout results in a flash of radiation. as the internal energy deposited by the shock diffuses out of the shocked region on a timescale comparable to the dynamical time of the shock.," The breakout results in a flash of radiation, as the internal energy deposited by the shock diffuses out of the shocked region on a timescale comparable to the dynamical time of the shock."
+ Following breakout the ejected envelope expands so that the photoshpere recedes into the ejecta., Following breakout the ejected envelope expands so that the photoshpere recedes into the ejecta.
+ Any remaining internal energy. which has not been converted during adiabatic expansion. is then gradually radiated to power the supernova lighteurve.," Any remaining internal energy, which has not been converted during adiabatic expansion, is then gradually radiated to power the supernova lightcurve."
+ It has long been suspected that the shock breakout from a star would appear as an. X-ray flash (22)... followed by a UV/optical transient corresponding to emission from the outer layers of the ejecta (2)..," It has long been suspected that the shock breakout from a star would appear as an X-ray flash \citep{Colgate74,KC78}, followed by a UV/optical transient corresponding to emission from the outer layers of the ejecta \citep{Falk78}."
+ The interest in prompt signals from supernovae increased recently due to the new capabilities of modern searches for transients. such as and as well as the planned andWFIRST!.," The interest in prompt signals from supernovae increased recently due to the new capabilities of modern searches for transients, such as and as well as the planned and."
+. In fact. several supernovae lightcurves were observed early after the explosion both in the UV and the optical bands (222)..," In fact, several supernovae lightcurves were observed early after the explosion both in the UV and the optical bands \citep{Gezarial08,Gezarial10,Soderbergal08}."
+ In two cases. GRBO60218/SN2006a] (2) and XROOS109/SN2008D (?).. X-ray detectors on board the Swift satellite have captured a luminous X-ray outburst. which were later followed by observations of supernovae in the longer wavelengths.," In two cases, GRB060218/SN2006aj \citep{Campanaal06} and XRO08109/SN2008D \citep{Soderbergal08}, X-ray detectors on board the Swift satellite have captured a luminous X-ray outburst, which were later followed by observations of supernovae in the longer wavelengths."
+ Both events may very well have been the signature of shock breakout., Both events may very well have been the signature of shock breakout.
+ These new observations motivated detailed theoretical modeling of the early emission in supernovae., These new observations motivated detailed theoretical modeling of the early emission in supernovae.
+ The models focused on the early UV/optical lighteurve that follows breakout. on ime scales of order a day (222). for which the dynamics are somewhat simpler (see also 2) ," The models focused on the early UV/optical lightcurve that follows breakout, on time scales of order a day \citep{ChevFrans08,RW10,NS10}, for which the dynamics are somewhat simpler (see also \citet{Chevalier92}) )."
+Modeling the X-ray flash associated with breakout is more complicated. since it must include he properties and structure of the radiation-mediated shock (RMS) as it propagates through the sharply decreasing density of the outer envelope. where the shock width becomes non-negligible relative o the distance to the edge of the star.," Modeling the X-ray flash associated with breakout is more complicated, since it must include the properties and structure of the radiation-mediated shock (RMS) as it propagates through the sharply decreasing density of the outer envelope, where the shock width becomes non-negligible relative to the distance to the edge of the star."
+ ? have shown that if the shock velocity normalized by c. 3s20.07. free-free emission will be unable to produce a large enough number density of photons o establish a blackbody spectrum.," \citet{Katzal10} have shown that if the shock velocity normalized by $c$, $\beta_S\geq 0.07$, free-free emission will be unable to produce a large enough number density of photons to establish a blackbody spectrum."
+ As a result. the average photon energy in the shocked region will be significantly larger than the equilibrium blackbody photon temperature with the same energy density.," As a result, the average photon energy in the shocked region will be significantly larger than the equilibrium blackbody photon temperature with the same energy density."
+ Hence. any analysis of the breakout lightcurve based on the assumption of local thermodynamic equilibrium and a black body spectrum will not reproduce the photon flux and spectrum emitted from the shock region during breakout.," Hence, any analysis of the breakout lightcurve based on the assumption of local thermodynamic equilibrium and a black body spectrum will not reproduce the photon flux and spectrum emitted from the shock region during breakout."
+ The shock breakout lightcurve could be significantly influenced by any circumstellar material (CSM) surrounding the exploding star., The shock breakout lightcurve could be significantly influenced by any circumstellar material (CSM) surrounding the exploding star.
+ In fact. the aforementioned two candidates for an X-ray flash of a shock breakout origin are associated with a type Ibe supernova. in which the progenitor is believed to be a," In fact, the aforementioned two candidates for an X-ray flash of a shock breakout origin are associated with a type Ibc supernova, in which the progenitor is believed to be a"
+We confirm the accuracy of the measured. uncertainty on cach estimate of ALC(06;;) using the reduced? ol the WLSQ fit that determines 2.,We confirm the accuracy of the measured uncertainty on each estimate of $\asqzeta$ using the $\chi^2$ of the WLSQ fit that determines $\hat{A^2}$.
+ The reduced-4? of this fit is which has a value of 1.3 for the Verbiest ο al. (, The $\chi^2$ of this fit is which has a value of 1.3 for the Verbiest et al. (
+"2008. 2009) residuals. indicatingthat the uncertainty estimates c,2 are consistent with the rns variation of the estimates. 2ος.","2008, 2009) residuals, indicatingthat the uncertainty estimates $\sigma_{A_k^2}$ are consistent with the rms variation of the estimates $A_k^2$."
+ WeH obtain. an independent.. estimate. of the accuracy of the measured. errors. by making. use of he information contained in the imaginary part of the Cross power spectrum. which we denote ImagΑν]," We obtain an independent estimate of the accuracy of the measured errors by making use of the information contained in the imaginary part of the cross power spectrum, which we denote $\rmn{Imag}\left[X_{ij}(f)\right]$."
+" We calculate ImagBU(61,)] by evaluating equation (8)) with μπας.X;j(f)] in place of Realδν].", We calculate $\rmn{Imag}\left[\asqzeta\right]$ by evaluating equation \ref{eq:a2zeta}) ) with $\rmn{Imag}\left[X_{ij}(f)\right]$ in place of $\rmn{Real}\left[X_{ij}(f)\right]$.
+ We then process μπαςBU(45; in exactly the same wav as the real part is processed., We then process $\rmn{Imag}\left[\asqzeta\right]$ in exactly the same way as the real part is processed.
+ Sinceη correlation coellicients are real. we expect that Imag[lee(63; will contain no correlated signal.," Since correlation coefficients are real, we expect that $\rmn{Imag}\left[\asqzeta\right]$ will contain no correlated signal."
+ This means that we can calculateη the analogue of the reduceda-x2 using ImagBun(6;η: Similar to the v if the errors on ALC(6;;) are well-estimated then this quantity should be near unity.," This means that we can calculate the analogue of the $\chi^2$ using $\rmn{Imag}\left[\asqzeta\right]$: Similar to the $\chi^2$, if the errors on $\asqzeta$ are well-estimated then this quantity should be near unity."
+ For the Verbiest et al. (, For the Verbiest et al. (
+2008. 2009) residuals. we find OST. inclicatine that the errors are well-estimat,"2008, 2009) residuals, we find $\chi^2_{\rm r, im} = 0.87$ , indicating that the errors are well-estimated."
+" Although both xz and N24, show that the uncertaintiesecd. 042 are reliable on average. these uncertainties come from power spectral estimates so they are random. variables."," Although both $\chi^2_{\rm r}$ and $\chi^2_{\rm r, im}$ show that the uncertainties $\sigma_{A^2_k}$ are reliable on average, these uncertainties come from power spectral estimates so they are random variables."
+ We estimated the sensitivity. of 21? to variations in ay bw multiplying each o4» by à random factor. distributed as[ the square root of the product of two X7 random variables with two degrees of freedom.," We estimated the sensitivity of $\hat{A^2}$ to variations in $\sigma_{A_k^2}$ by multiplying each $\sigma_{A_k^2}$ by a random factor, distributed as the square root of the product of two $\chi^2$ random variables with two degrees of freedom."
+ This is the expected. distribution for cach e42., This is the expected distribution for each $\sigma_{A_k^2}$.
+ We found that 04 increased hy a factor of 1.6. incicating that the use of incorrect AARC(6;;) estimates cleeraces the sensitivity of the 412. measurement by only a factor of 1.6.," We found that $\sigma_{\hat{A^2}}$ increased by a factor of 1.6, indicating that the use of incorrect $\easqzeta$ estimates degrades the sensitivity of the $\hat{A^2}$ measurement by only a factor of 1.6."
+" However. the 21,4(8;;) awe not Gaussian: rather they come fromthe sum of two pairwise products of independent Gaussian variables and thus have a two-sicecl exponential distribution. which is rellected in Figure 3.."," However, the $\asqzeta$ are not Gaussian; rather they come fromthe sum of two pairwise products of independent Gaussian variables and thus have a two-sided exponential distribution which is reflected in Figure \ref{fig:distnull}. ."
+ This means that the maximunr likelihood estimator for ο is not a WLSQ estimator but a weighted least. absolute deviation (LAD) Lt (see.eg. ?). ," This means that the maximum likelihood estimator for $A^2$ is not a WLSQ estimator but a weighted least absolute deviation (LAD) fit \citep[see, e.g.,][]{2006Cox}. ."
+We tested. both weighted. and, We tested both weighted and
+intensities of our 3 families of structures (Fig.,intensities of our 3 families of structures (Fig.
+" 4, 6 and 7), we shall assume that thiourea derivatives (with compact and linear aromatics) and the aliphatic chains (Fig."," 4, 6 and 7), we shall assume that thiourea derivatives (with compact and linear aromatics) and the aliphatic chains (Fig."
+" 2, 3, 5, 8 and 9) are in the ratios 1:1:1."," 2, 3, 5, 8 and 9) are in the ratios 1:1:1."
+" Thus, in Fig.11, we concatenate the spectra of Fig."," Thus, in Fig.11, we concatenate the spectra of Fig."
+ 4 (red) and 6 (green) together with the spectrum of Fig., 4 (red) and 6 (green) together with the spectrum of Fig.
+ 10 (blue)., 10 (blue).
+" Because we have not considered enough individual structures in each family, the concatenated spectrum of all the selected structures is still not nearly continuous."," Because we have not considered enough individual structures in each family, the concatenated spectrum of all the selected structures is still not nearly continuous."
+ We therefore need some sort of interpolation and/or smoothing., We therefore need some sort of interpolation and/or smoothing.
+" Accordingly, the concatenated line spectrum was Fast Fourier Transform smoothed over 50 points, resulting in the continuous black line in Fig."," Accordingly, the concatenated line spectrum was Fast Fourier Transform smoothed over 50 points, resulting in the continuous black line in Fig."
+ 11., 11.
+" The 20-ym feature peaks at 20.7um; its apparent FWHM is 2.3 um. Since the integrated band intensity, J, is proportional to the absorbance, a, of the material (eq."," The $\mu$ m feature peaks at $\mu$ m; its apparent FWHM is 2.3 $\mu$ m. Since the integrated band intensity, $I$, is proportional to the absorbance, $\alpha$, of the material (eq."
+" 1), its radiative emission, P5, will be proportional to the product of / and the black body emission at the temperature, Tm, of the molecules."," 1), its radiative emission, $F_{\lambda}$, will be proportional to the product of $I$ and the black body emission at the temperature, $T_{m}$, of the molecules."
+" This will be multiplied by A,for comparison with the ÀF) spectra of Zhang et al. (2010).."," This will be multiplied by $\lambda$,for comparison with the $\lambda\,F_{\lambda}$ spectra of Zhang et al. \cite{zha}. ."
+" For the same reason, we adopt dust emissivities scaling like \~?."," For the same reason, we adopt dust emissivities scaling like $\lambda^{-2}$."
+"to the ""lens"" by about a factor of 10. which makes the “lens” almost invisible at 6 em.","to the “lens” by about a factor of 10, which makes the “lens” almost invisible at $\lambda$ 6 cm."
+ Several regions are located around £/=14075 within a 325x6:0 field (Fig. 9)., Several regions are located around $\ell = 140\fdg5$ within a $3\fdg5 \times 6\fdg0$ field (Fig. ).
+ Unfortunately. no information of the regions was given in the catalogue of?.," Unfortunately, no information of the regions was given in the catalogue of."
+ Three faint optically visible regions of the Lynds catalogue are: the semi-ring shaped LBN 676 (£=139257.b 2770)with a size of 028. LBN 677 (SH 2-202) (£=140:07.b.1264. size of 2:0). and the bar-like shaped LBN 679 (£=140777.b.—1.42. size of 2:0).," Three faint optically visible regions of the Lynds catalogue are: the semi-ring shaped LBN 676 $\ell = 139\fdg57, b = 2\fdg70$ )with a size of $0\fdg8$ , LBN 677 (SH 2-202) $\ell = 140\fdg07, b = 1\fdg64$, size of $2\fdg0$ ), and the bar-like shaped LBN 679 $\ell = 140\fdg77, b = -1.42\degr$, size of $2\fdg0$ )."
+" For morphology reasons we name the three regions the ""Drumstick"" in the following.", For morphology reasons we name the three regions the “Drumstick” in the following.
+ LBN 676 and LBN 679 are supposed to be at the same distance in the Perseus arm., LBN 676 and LBN 679 are supposed to be at the same distance in the Perseus arm.
+ They were already investigated in some detail by using DRAO Synthesis Telescope data at 408 MHz and discussed together with infrared and HI maps., They were already investigated in some detail by using DRAO Synthesis Telescope data at 408 MHz and discussed together with infrared and HI maps.
+ found that the elongated region LBN 679 coincides with a large spur located in the Perseus arm and pointed out that itis a thermal rather than a non-thermal feature as suggested by?., found that the elongated region LBN 679 coincides with a large spur located in the Perseus arm and pointed out that it is a thermal rather than a non-thermal feature as suggested by.
+ studied all three regions with I’ resolution using CGPS data at 1.4 GHz(?)., studied all three regions with $\arcmin$ resolution using CGPS data at 1.4 GHz.
+ The thermal character of LBN 679 was again confirmed and in addition they derived a thermal spectrum with B=—].85+0.1 for LBN 676. thus excluding a possible SNR identification considered by for this shell structure.," The thermal character of LBN 679 was again confirmed and in addition they derived a thermal spectrum with $\beta = -1.85\pm0.1$ for LBN 676, thus excluding a possible SNR identification considered by for this shell structure."
+ The 16ο em data agree with the thermal properties of all objects through acheck of their temperature spectral indices via TT-plots using Effelsberg 21 em data for comparison., The $\lambda$ 6 cm data agree with the thermal properties of all objects through a check of their temperature spectral indices via TT-plots using Effelsberg $\lambda$ 21 cm data for comparison.
+ All three LBNe appear to be depolarized at 26 em when large-scale polarized emission is added., All three LBNe appear to be depolarized at $\lambda$ 6 cm when large-scale polarized emission is added.
+ Estimates of the magnetic field strength from modeled RM;4 require the thermal radio continuum brightness temperature to find their EM., Estimates of the magnetic field strength from modeled $RM_{FS}$ require the thermal radio continuum brightness temperature to find their $EM$.
+" For a known distance the source size and 7, need to be caleulated in addition.", For a known distance the source size and $n_{e}$ need to be calculated in addition.
+ assumed that LBN 676 and LBN 679 are both at a distance of 3 kpe in the Perseus arm., assumed that LBN 676 and LBN 679 are both at a distance of 3 kpc in the Perseus arm.
+ However. the Perseus arm distance was revised by to be about 2 kpe by triangulation of W3OH. which is just about 7° apart in Galactic longitude.," However, the Perseus arm distance was revised by to be about 2 kpc by triangulation of W3OH, which is just about $\degr$ apart in Galactic longitude."
+ In the following we adopt this distance., In the following we adopt this distance.
+ It turns out that a ring average of the PA/PI differences for LBN 676 is difficult to perform because of confusion with LBN 677 emission., It turns out that a ring average of the PA/PI differences for LBN 676 is difficult to perform because of confusion with LBN 677 emission.
+ Thus we just take the data from the upper part., Thus we just take the data from the upper part.
+ The model fit gives the best result (Fig. 10), The model fit gives the best result (Fig. )
+ for the first five pixels as RM;= 280430 rad m., for the first five pixels as $RM_{FS} = 280\pm$ 30 rad $^{-2}$.
+ Foreground polarized emission comprises about 79% while f=0.80., Foreground polarized emission comprises about 79% while $\mathit{f} = 0.80$.
+ Likely most of the polarized emission originates in the local arm., Likely most of the polarized emission originates in the local arm.
+ The total intensity attributed to LBN 676 is about 50 mK Ty at 26 em., The total intensity attributed to LBN 676 is about 50 mK $\rm T_{B}$ at $\lambda$ 6 cm.
+ An apparent diameter of 0:8 equals to a path length of 28 pe for 2 kpe distance., An apparent diameter of $0\fdg8$ equals to a path length of 28 pc for 2 kpc distance.
+" We obtain 71, as about 3.3 em? and Bj s 3.7 HG. In the southern part of the 2° long filamentary LBN 679. we find large PA changes and also in its outskirts beyond."," We obtain $n_{e}$ as about 3.3 $^{-3}$ and $B_{\parallel}$ $\approx$ +3.7 $\mu$ G. In the southern part of the $\degr$ long filamentary LBN 679, we find large $PA$ changes and also in its outskirts beyond."
+ There is an inclined elongated PA structure. about 1? long. running from northeast to southwest.," There is an inclined elongated $PA$ structure, about $\degr$ long, running from northeast to southwest."
+ Model fitting (Fig. 11), Model fitting (Fig. )
+ is done for a 35° wide cone in south-western direction of the rim., is done for a $\degr$ wide cone in south-western direction of the rim.
+ Best fit for the first ten pixels gives RM;s=—1SS415 rad m7. a foreground polarization fraction of about 60% for f= 0.85.," Best fit for the first ten pixels gives $RM_{FS} = -155\pm$15 rad $^{-2}$, a foreground polarization fraction of about 60% for $\mathit{f} = 0.85$ ."
+" Fromthe 6 em brightness temperature of about 40 mK Ty and assuming a line-of-sight length of the source of 35 pe. we calculated By = —2.0 μΟ and 5, = 2.7 em."," Fromthe $\lambda$ 6 cm brightness temperature of about 40 mK $\rm T_{B}$ and assuming a line-of-sight length of the source of 35 pc, we calculated $B_{\parallel}$ = $-2.0~\mu$ G and $n_{e}$ = 2.7 $^{-3}$ ."
+ To the lower left, To the lower left
+"Astronomical surveys now gather data on huge numbers of astronomical objects: tvw 2 Micron ALL Sky Survey (24LASS: 7). the Sloan Digital Sky Survey (SDSS: 7)) and the UINIBCE Infrared. Deep Sky Survey (UIXIDSS: 2)) have all identified hundreds of millions ο ""distinct sources.","Astronomical surveys now gather data on huge numbers of astronomical objects: the 2 Micron All Sky Survey (2MASS; \citealt{Skrutskie_etal:2006}) ), the Sloan Digital Sky Survey (SDSS; \citealt{York_etal:2000}) ) and the UKIRT Infrared Deep Sky Survey (UKIDSS; \citealt{Lawrence:2006}) ) have all identified hundreds of millions of distinct sources."
+ Phe scale of these projects immediately necessitates an automated. approach to data analysis (although5 an intriguing5 alternative is The Galaxy Zoo project. described. by 2))., The scale of these projects immediately necessitates an automated approach to data analysis (although an intriguing alternative is The Galaxy Zoo project described by \citealt{Lintott_etal:2008}) ).
+. Considerable οσο has been put into developing algorithms which can decompose an im.age into a smooth background. and a catalogue of discree objects. the properties of which must be characterised as well.," Considerable effort has been put into developing algorithms which can decompose an image into a smooth background and a catalogue of discrete objects, the properties of which must be characterised as well."
+ Source positions. lluxes and shapes can all be estimated. reliably by using fairly simple moment-based. approaches(e.g... 2?)). but the separation of point-like stars from more extended: galaxies. generally requires at least some external astrophysical information be incluced.," Source positions, fluxes and shapes can all be estimated reliably by using fairly simple moment-based approaches, \citealt{Irwin:1985,Bertin_Arnouts:1996}) ), but the separation of point-like stars from more extended galaxies generally requires at least some external astrophysical information be included."
+ As such. the problem of cclassification is well suited to Bavesian methods in which the measurements of a given source are combined with prior knowledge of the astrophysical populations of which the source might be a member.," As such, the problem of classification is well suited to Bayesian methods in which the measurements of a given source are combined with prior knowledge of the astrophysical populations of which the source might be a member."
+ A practical formalism. for Bavesian classification is developed in this paper., A practical formalism for Bayesian classification is developed in this paper.
+ 1n the existing methods of Cclassification are reviewed. with particular emphasis on," In the existing methods of classification are reviewed, with particular emphasis on"
+analysis in Sect.3.3.. slightl,analysis in Sect. \ref{fund_param}.
+yIncluding overshoot in the BASTI models results in only different values; M/M..=1.50+0.10. R/R»=4.8440.35. age 1.520.2 Gyr. and loge=3.3340.07.," Including overshoot in the BASTI models results in only slightly different values; $M/M_\odot = 1.80\pm0.10$, $R/R_\odot = 4.84\pm0.35$, age $1.5\pm0.2$ Gyr, and $\log g = 3.33\pm0.07$."
+ In Table | we list the results using the standard BASTI isochrones without overshoot., In Table \ref{tab:results} we list the results using the standard BASTI isochrones without overshoot.
+ These two independent mass estimates are in rough agreement., These two independent mass estimates are in rough agreement.
+ A seismologically determinec mass estimate 1s potentially more accurate than isochrone fitting 1fone can measure the large separation Av or the position of maximum power vias (2).., A seismologically determined mass estimate is potentially more accurate than isochrone fitting ifone can measure the large separation $\Delta\nu$ or the position of maximum power $\nu_{\rm max}$ \citep{basu+2009}.
+" Since our data do not allow us to measure Av and the error on vy, is rather large. we adopt the mass estimate based on the SHOTGUN approach."," Since our data do not allow us to measure $\Delta\nu$ and the error on $\nu_{\rm max}$ is rather large, we adopt the mass estimate based on the SHOTGUN approach."
+ The rotational velocity of EK Eri is notoriously hard to measure due the the very slow rotation., The rotational velocity of EK Eri is notoriously hard to measure due the the very slow rotation.
+ The first accurate determination was done by ?. from R=120.000 spectra. from which they derived vsini=1.5€0.5 wwith an adopted value of the macroturbulent velocity S50 s. with a Fourier analysis of the line profile as a check of the validity of the adopted μωρο.," The first accurate determination was done by \citet{strassmeier+1999} from $R=120,000$ spectra, from which they derived $v\sin i = 1.5 \pm 0.5$ with an adopted value of the macroturbulent velocity $v_\mathrm{macro}=5.0$ , with a Fourier analysis of the line profile as a check of the validity of the adopted $v_\mathrm{macro}$."
+ In Paper I we derived vsini21.00.5 bby simple €profile fitting to 1D atmospheric model spectra. assuming a similar value of the macroturbulence.," In Paper I we derived $v\sin i=1.0\pm 0.5$ by simple profile fitting to 1D atmospheric model spectra, assuming a similar value of the macroturbulence."
+ As we will argue in Sec. 4.2..," As we will argue in Sec. \ref{rv-activity},"
+ we suggest that the rotational period of EK Ert could be twice the photometric period.i close to 617 d. We have derived 4.87+0.29 R. (see Sect. 3.2)).," we suggest that the rotational period of EK Eri could be twice the photometric period, i.e. close to 617 d. We have derived $4.87\pm0.29$ $_\odot$ (see Sect. \ref{seismology}) ),"
+ which ts in good agreement with the value ofR=4.78 R.. we derive using Eq., which is in good agreement with the value of $R = 4.78$ $_\odot$ we derive using Eq.
+ | from ?.., 1 from \citet{strassmeier+1999}.
+ Adopting Py=2Ppnor- this means that the radius of the star will be R/R=12.2vsin/. where we assumethat /z90°.," Adopting $P_\mathrm{rot} = 2 P_\mathrm{phot}$, this means that the radius of the star will be $R/R_\odot = 12.2\, v\sin i$, where we assumethat $i \simeq 90^\circ$."
+ With the derived radius of 4.87 Rz. this translates to an expectation for the measurement of vsini. namely vsin«0.40s!.," With the derived radius of $4.87$ $_\odot$, this translates to an expectation for the measurement of $v\sin i$, namely $v\sin i < 0.40$."
+". If on the other hand P4,=Pon. then vsin7 must be less than 0.80!.. as already noted by ?.. and this value can thus be regarded as a safe upper limit on vsini£."," If on the other hand $P_\mathrm{rot} = P_\mathrm{phot}$, then $v\sin i$ must be less than $0.80$, as already noted by \citet{strassmeier+1999}, and this value can thus be regarded as a safe upper limit on $v\sin i$."
+ Using the high-resolution CES data. and our newly derived stellar parameters. we have attempted to derive vsin? and μπάσο:," Using the high-resolution CES data, and our newly derived stellar parameters, we have attempted to derive $v\sin i$ and $v_\mathrm{macro}$."
+ For this analysis we chose 7 lines from the two CES spectra. chosen to be non-blended based on the lie list obtained from VALD (?) and visual evaluation of the line symmetry.," For this analysis we chose 7 lines from the two CES spectra, chosen to be non-blended based on the line list obtained from VALD \citep{kupka+1999} and visual evaluation of the line symmetry."
+" The synthetic spectrum 1s broadened with different values of vsin/ and v4, in a fine grid in the range 0—4 ffor both parameters. and the y is calculated."," The synthetic spectrum is broadened with different values of $v\sin i$ and $v_\mathrm{macro}$ in a fine grid in the range 0–4 for both parameters, and the $\chi^2$ is calculated."
+" At values of Viaao>+ tthe fits start to deviate significantly and we thus note that a value of viser=5 citeploriginallyfrom][as.à""typical""value]grav1992.— is inconsistent with our results."," At values of $v_\mathrm{macro} > 4$ the fits start to deviate significantly and we thus note that a value of $v_\mathrm{macro,RT} = 5$ \\citep[originally from][as a ``typical'' value]{gray1992} is inconsistent with our results."
+ There is a strong correlation between the broadening caused by rotation and by macroturbulence which is very difficult to separate even though their exact forms are slightly different., There is a strong correlation between the broadening caused by rotation and by macroturbulence which is very difficult to separate even though their exact forms are slightly different.
+ In. general we find a continuum of solutions. as illustrated in Fig. 7..," In general we find a continuum of solutions, as illustrated in Fig. \ref{fig:vsini},"
+ which shows the reduced y surface for one spectral line for different values of vsini and Όρων., which shows the reduced $\chi^2$ surface for one spectral line for different values of $v\sin i$ and $v_\mathrm{macro}$.
+ The v surfaces for the other lines look very similar., The $\chi^2$ surfaces for the other lines look very similar.
+ We find that lis so low that it cannot be determined reliably and only upper limits can be set., We find that is so low that it cannot be determined reliably and only upper limits can be set.
+ Note that we do not include Zeeman broadening. which in any case would contribute to ai even lower limit on vsinf.," Note that we do not include Zeeman broadening, which in any case would contribute to an even lower limit on $v\sin i$."
+ It is well known that stellar magnetic activity affects the shape of spectral lines and thereby the apparent RV (??).. which again affects the ability to detect planets by the Doppler technique.," It is well known that stellar magnetic activity affects the shape of spectral lines and thereby the apparent RV \citep[][]{gray1988,gray2005}, which again affects the ability to detect planets by the Doppler technique."
+ The best direct measure of activity available in the HARPS spectra are the emission cores of the calctum H and K lines. and for each spectrum we have derived the activity index Rux following the procedure of Paper I. It was shown by ? that the Bisector Inverse Slope (BIS) of the CCF is à good qualitative measure of the distortion of the spectral lines caused by activity. and we have calculated the BIS for all CCFs following the procedure of ?..," The best direct measure of activity available in the HARPS spectra are the emission cores of the calcium H and K lines, and for each spectrum we have derived the activity index $R_\mathrm{HK}$ following the procedure of Paper I. It was shown by \citet{queloz+2001}
+ that the Bisector Inverse Slope (BIS) of the CCF is a good qualitative measure of the distortion of the spectral lines caused by activity, and we have calculated the BIS for all CCFs following the procedure of \citet{dall+2006}."
+ The series of RV. BIS and aare shown in Fig.," The series of RV, BIS and are shown in Fig."
+ 8 along with the corresponding part of the photometric light curve., \ref{fig:rvbjd} along with the corresponding part of the photometric light curve.
+ In Figs., In Figs.
+ 9 and 10 we show the correlations of RV with activity index logRy« and BIS., \ref{fig:rhk} and \ref{fig:bis} we show the correlations of RV with activity index $\log R_\mathrm{HK}$ and BIS.
+ From Figs., From Figs.
+ 8 and 9.. it is clear that the large scale RV variations are very well correlated with the activity as measured by Ruk.," \ref{fig:tsAll} and \ref{fig:rhk}, , it is clear that the large scale RV variations are very well correlated with the activity as measured by ."
+. As the star gets dimmer. it gets redder (asnoted and the chromospheric activity gets higher. both facts pointing to cool spots with overlaid plages as the cause of the variation.," As the star gets dimmer, it gets redder \citep[as noted by][]{strassmeier+1999} and the chromospheric activity gets higher, both facts pointing to cool spots with overlaid plages as the cause of the variation."
+ The variation of RV seems to be correlated with, The variation of RV seems to be correlated with
+~LOO vr to several 1000 vr.,$\sim 100$ yr to several 1000 yr.
+ In spite of the heterogeneity of the data and (he uncertainties in individual cases. the figure shows that the ages of the jets and the ages of the tori are well correlated. and that the data points he not [ar from the line /;=/;.," In spite of the heterogeneity of the data and the uncertainties in individual cases, the figure shows that the ages of the jets and the ages of the tori are well correlated, and that the data points lie not far from the line $t_j=t_t$."
+ The implication is that the torus and jets in each object were ejected nearly simultaneously., The implication is that the torus and jets in each object were ejected nearly simultaneously.
+ It is also of interest to note that the ejection ages are roughly ordered according to the degree of evolution indicated by the spectral (wpe of the star and/or the development ol the nebula: i.e.. from voungest to oldest corresponds to AGB stars. through proto-PNe (typically B-ivpe spectra). to true PNe.," It is also of interest to note that the ejection ages are roughly ordered according to the degree of evolution indicated by the spectral type of the star and/or the development of the nebula: i.e., from youngest to oldest corresponds to AGB stars, through proto-PNe (typically B-type spectra), to true PNe."
+ This is consistent wil all ejections (except perhaps AI2-9 whose status is uncertain. see 822) happening in the final stages of the AGB phase. or shortly thereafter.," This is consistent with all ejections (except perhaps M2-9 whose status is uncertain, see 2) happening in the final stages of the AGB phase, or shortly thereafter."
+ The jets aud tori have similar ages. but the data points in Fie.," The jets and tori have similar ages, but the data points in Fig."
+ 1 lie svstematically below the line /;=/;., 1 lie systematically below the line $t_j = t_t$.
+ The elfect appears more pronounced Lor the vounger svstems in the log-log plot. but it is also present for the older svstems as well.," The effect appears more pronounced for the younger systems in the log-log plot, but it is also present for the older systems as well."
+" The data in a few cases are consistent with /;= within the uncertainties. but for all objects our best estimates of the time difference Al=/,—/; (column 9 in Table 1) are positive."," The data in a few cases are consistent with $t_j = t_t$ within the uncertainties, but for all objects our best estimates of the time difference $\Delta t = t_t-t_j$ (column 9 in Table 1) are positive."
+ This ordering has previously been noted in individual cases (e.g..Forveileetal.1993:IHIugeinsSánchezContrerasetal. 2004).. but the result [or the ensemble as a whole makes it much more concrete. especially in view of the fact that /; refers to the heads of the jets. and /; refers to the mean radius of the tori (or some similar measure) in most cases.," This ordering has previously been noted in individual cases \citep[e.g.,][]{for98,hug00,san04}, but the result for the ensemble as a whole makes it much more concrete, especially in view of the fact that $t_j$ refers to the heads of the jets, and $t_t$ refers to the mean radius of the tori (or some similar measure) in most cases."
+ In fact. given the diversity of the data used to compile the figure. the uniformity of the effect is striking.," In fact, given the diversity of the data used to compile the figure, the uniformity of the effect is striking."
+ There are (wo likely explanations of the observed effect that need to be considered. aud we discuss each in (turn.," There are two likely explanations of the observed effect that need to be considered, and we discuss each in turn."
+ The first is that the values of /; and /; ave good approximations to the actual travel times of the jets and tori (as previously assumed). but the jets and tori are not exactly simultaneous.," The first is that the values of $t_j$ and $t_t$ are good approximations to the actual travel times of the jets and tori (as previously assumed), but the jets and tori are not exactly simultaneous."
+ The observed effect is produced if the jets ave shehtly vounger ihan the tori. ie.. there is a small but svstematic delay in the launching of the jets relative to the tori.," The observed effect is produced if the jets are slightly younger than the tori, i.e., there is a small but systematic delay in the launching of the jets relative to the tori."
+ If so. this delay. or jet-lag. is given bv the value of A’.," If so, this delay, or jet-lag, is given by the value of $\Delta
+t$."
+ It might correspond. for exaniple. to a power-up time for the jets. or an accretion (ime scale (844.2.2).," It might correspond, for example, to a power-up time for the jets, or an accretion time scale 4.2.2)."
+ Inspection of Table 1 (column 9) shows that the largest value of A/ is for the oldest svstem (IxjPn 3). but the others are comparable. within the large uncertzinties.," Inspection of Table 1 (column 9) shows that the largest value of $\Delta t$ is for the oldest system (KjPn 8), but the others are comparable, within the large uncertainties."
+ The median jet-Iag is 300 vr., The median jet-lag is 300 yr.
+ The second explanation to consider is that (he jets and tori are ejected simultaneously. but the expansion times differ from Che actual travel Gaines.," The second explanation to consider is that the jets and tori are ejected simultaneously, but the expansion times differ from the actual travel times."
+ The expansion times of the massive molecular (ori are almost certainly good approximations to the travel times [or, The expansion times of the massive molecular tori are almost certainly good approximations to the travel times for
+uncertainties because of larger photometric scatter in the observations (by factors of 3-5 relative to the BV(RIJc photometry) in combination with the correction from infrared to optical reddening.,uncertainties because of larger photometric scatter in the observations (by factors of 3–5 relative to the $_C$ photometry) in combination with the correction from infrared to optical reddening.
+" They are nevertheless crucial for narrowing the range of potential color excesses for some of the sample Cepheids, such as IO Cas."," They are nevertheless crucial for narrowing the range of potential color excesses for some of the sample Cepheids, such as IO Cas."
+" A few objects in the sample appear to be Type II Cepheids, but that does not affect the results."," A few objects in the sample appear to be Type II Cepheids, but that does not affect the results."
+ In most cases the implied reddening near the Cepheid is fairly evident., In most cases the implied reddening near the Cepheid is fairly evident.
+" 'The results of the variable-extinction studies of the sample Cepheids are presented in Table 1,, which lists the inferred space reddening for each Cepheid derived from likely AF-type stars lying within different angular radii, from 2’ to 5’ distant, and the reddening inferred from 2MASS JHK, colors for stars within 5’ of the Cepheid."," The results of the variable-extinction studies of the sample Cepheids are presented in Table \ref{tab2}, which lists the inferred space reddening for each Cepheid derived from likely AF-type stars lying within different angular radii, from $2^\prime$ to $5^\prime$ distant, and the reddening inferred from 2MASS $_{\rm s}$ colors for stars within $5^\prime$ of the Cepheid."
+" The adopted color excess in each case, column (8), was the average for all AF-type stars lying within 5’ of the Cepheid and of comparable inferred distance, and the resulting reddening, which is equivalent to that for a star of spectral type BO, was converted in the last column of the table to a value appropriate for the derived intrinsic color of the Cepheid, again using the relationship of Fernie(1963)."," The adopted color excess in each case, column (8), was the average for all AF-type stars lying within $5^\prime$ of the Cepheid and of comparable inferred distance, and the resulting reddening, which is equivalent to that for a star of spectral type B0, was converted in the last column of the table to a value appropriate for the derived intrinsic color of the Cepheid, again using the relationship of \citet{fe63}."
+. The number of stars used to obtain the space reddening is listed in the second last column of Table 1.., The number of stars used to obtain the space reddening is listed in the second last column of Table \ref{tab2}. .
+" In rich fields that number is of order 710-20, whereas in more poorlypopulated regions or"," In rich fields that number is of order $\sim$ 10–20, whereas in more poorlypopulated regions or"
+indicates that the injection timescale should not exceed A/c (Chiaberge Ghisellini 1999).,indicates that the injection timescale should not exceed $R/c$ (Chiaberge Ghisellini 1999).
+ For these reasons we also explore here the situation where the injection lasts for a finite timescale. of order of the light crossing time.," For these reasons we also explore here the situation where the injection lasts for a finite timescale, of order of the light crossing time."
+ A physical scenario. where flares naturally occur and the injection timescale can be of order of the dynamical one ts that of internal shocks (e.g. Piran 1999; Ghisellint 1999a: Spada et al., A physical scenario where flares naturally occur and the injection timescale can be of order of the dynamical one is that of internal shocks (e.g. Piran 1999; Ghisellini 1999a; Spada et al.
+ 2001). in which the dissipation takes place during the collision of two shells of fluid moving at different speeds.," 2001), in which the dissipation takes place during the collision of two shells of fluid moving at different speeds."
+ Let us then adopt such scenario., Let us then adopt such scenario.
+ As in the steady state model just considered. the magnetic field is assumed to be homogeneous and tangled throughout the region.," As in the steady state model just considered, the magnetic field is assumed to be homogeneous and tangled throughout the region."
+ Note however that in the internal shock scenario the geometry of the emitting volume is a cylinder with, Note however that in the internal shock scenario the geometry of the emitting volume is a cylinder with
+"tidal eccentricity time rate of change due to tides is: where e is the eccentricity and k, is the tidal Love number of the secondary [1999).",tidal eccentricity time rate of change due to tides is: where $e$ is the eccentricity and $k_s$ is the tidal Love number of the secondary \citep{Murray1999}.
+. Tidal evolution can cause excitation or damping of the eccentricity depending on the system mass ratio and the tidal Love numbers of each component., Tidal evolution can cause excitation or damping of the eccentricity depending on the system mass ratio and the tidal Love numbers of each component.
+ There are two developed theories for the internal structure of asteroids: “monolith” and “rubble pile.’, There are two developed theories for the internal structure of asteroids: “monolith” and “rubble pile.”
+" Evidence that asteroids have a “rubble pile"" internal structure rather than a “monolithic” interior includes: the Hayabusa mission to Itokawa showing no obvious impact craters and the appearance of a structure made entirely from shattered fragments of different size scales (Fujiwaraetal|2006),, numerical modeling of collisions of asteroids 2001)., mass and volume measurements from the NEAR Shoemaker flyby of Mathilde and radar observations of 1999 KW, (OstroetaL]2006) showing mean densities that are lower than their constitutive elements, the rotational speed limit period (~2.2 hours) observed amongst asteroids with diameters larger than ~200 m, which corresponds to the critical disruption spin rate of a self-gravitating, strengthless body (Pravecetal.|2007),, and that rotational fission of strengthless bodies is responsible for the asteroid pair population (Pravecetal.|2010)."," Evidence that asteroids have a “rubble pile” internal structure rather than a “monolithic” interior includes: the Hayabusa mission to Itokawa showing no obvious impact craters and the appearance of a structure made entirely from shattered fragments of different size scales \citep{Fujiwara2006}, numerical modeling of collisions of asteroids \citep{Michel2001}, mass and volume measurements from the NEAR Shoemaker flyby of Mathilde \citep{Yeomans1997} and radar observations of 1999 $_4$ \citep{Ostro2006} showing mean densities that are lower than their constitutive elements, the rotational speed limit period $\sim 2.2$ hours) observed amongst asteroids with diameters larger than $\sim200$ m, which corresponds to the critical disruption spin rate of a self-gravitating, strengthless body \citep{Pravec2007a}, and that rotational fission of strengthless bodies is responsible for the asteroid pair population \citep{Pravec2010}."
+". All of this evidence suggests that a “rubble pile"" internal structure is a more realistic assessment than a “monolith” internal structure, however for completeness we will consider both theories."," All of this evidence suggests that a “rubble pile” internal structure is a more realistic assessment than a “monolith"" internal structure, however for completeness we will consider both theories."
+ The two theories become distinct when assessing the functional form of the tidal Love number., The two theories become distinct when assessing the functional form of the tidal Love number.
+" The dimensionless rigidity of the body jz, which can be thought of as the ratio of the fluid strain to the elastic strain, can be used to defined the tidal Love number: k=1.5/(14-i)©1.5/ji, this approximation is true when the fluid strain dominates the elastic strain (Gold (2009)."," The dimensionless rigidity of the body $\tilde{\mu}$, which can be thought of as the ratio of the fluid strain to the elastic strain, can be used to defined the tidal Love number: $k = 1.5 / ( 1 + \tilde{\mu} ) \approx 1.5 / \tilde{\mu}$, this approximation is true when the fluid strain dominates the elastic strain \citep{Goldreich2009}."
+". According to the canonical “monolith” theory the dimensionless rigidity has the form fi=19u/(2:03pR?), where µ is the rigidity of the body."," According to the canonical “monolith"" theory the dimensionless rigidity has the form $\tilde{\mu} =19 \mu / (2 \omega_d^2 \rho R^2)$, where $\mu$ is the rigidity of the body."
+" To first order, the tidal Love number for a “monolith” depends on the size of the body as ky;ος R?."," To first order, the tidal Love number for a “monolith” depends on the size of the body as $k_M \propto R^2$ ."
+Oue ofti6 priuarv motivations or studvius the oucr halos of galaxies is to study fie nature o lark mattex.,	One of the primary motivations for studying the outer halos of galaxies is to study the nature of dark matter.
+ The reasou the «mter halos of galaxies aro ΠΟΥaut for addressine the nature of dark matte rds that most galaxies are οilv «ark matteroninated a large distances frou their ceners., The reason the outer halos of galaxies are important for addressing the nature of dark matter is that most galaxies are only dark matter dominated at large distances from their centers.
+" Iu the inner rock1s of galaxie s(r3 R,). he observed stars aud gas vpically ma sen sigQvificaut contribution to the mass budeet."," In the inner regions of galaxies $r \ltsim R_{e}$ ), the observed stars and gas typically make a significant contribution to the mass budget."
+ Exactly what tha coutributic11 19 Or differeit galaxies renains a subject of active researc1 aud vigorns debato., Exactly what that contribution is for different galaxies remains a subject of active research and vigorous debate.
+" This vigorous debate is sviuptomaic of the problem of deterwining the «ark ntorY co1eut of the inner parts «ft ealaxies accurately,", This vigorous debate is symptomatic of the problem of determining the dark matter content of the inner parts of galaxies accurately.
+ Specifcalv. bec:USC m5 ()ISCLVE«0 barvons nake up a ron-neelieible fraction of he imas in the ner reelois of most. ealaxics. an accurate determination of he ddark iiater Duet)oerties iu these regions requires a verv accurate accountiues 6| the nass COLlributioi of these barvons.," Specifically, because the observed baryons make up a non-negligible fraction of the mass in the inner regions of most galaxies, an accurate determination of the dark matter properties in these regions requires a very accurate accounting of the mass contribution of these baryons."
+ Since it is a severe challenge to acΜονο such an accurate accounting of the barvouic mass in the inner regions of ealaxies. he dark matter properties of the inner regious of galaxies remain SOLLOWhat nucertain.," Since it is a severe challenge to achieve such an accurate accounting of the baryonic mass in the inner regions of galaxies, the dark matter properties of the inner regions of galaxies remain somewhat uncertain."
+ Moreover. im regions iu which barvous are a significant ractiou of he mass. the dissipation. star formation. aud feedback that can happen in the rvous nay affect the distribution of the dark matter.," Moreover, in regions in which baryons are a significant fraction of the mass, the dissipation, star formation, and feedback that can happen in the baryons may affect the distribution of the dark matter."
+ Although it is possible o study special classes of objects low surface bxiglituess gaaxies) for, Although it is possible to study special classes of objects low surface brightness galaxies) for
+"our method leads to better completeness. levels. at these lluxes (both for using £2)>0.5 and D,> 0.9).",our method leads to better completeness levels at these fluxes (both for using $\ps>0.5$ and $\ps>0.9$ ).
+ For Yc19 our method with £2>0.9. using the UINIDSS merged class labels and. using the UIXIDSS pipeline posteriors perform identically.," For $Y\simeq19$ our method with $\ps>0.9$, using the UKIDSS merged class labels and using the UKIDSS pipeline posteriors perform identically."
+ Dasing the selection. on the UINIDSS. Y-band class labels or on the posteriors from. our method. with D.0.5 leads to higher completeness but lower cllicicney (but our method. vields a much higher completeness than using the UIXIDSS Y-banc labels and. also a marginally arger ellicienev)., Basing the selection on the UKIDSS $Y$ -band class labels or on the posteriors from our method with $\ps>0.5$ leads to higher completeness but lower efficiency (but our method yields a much higher completeness than using the UKIDSS $Y$ -band labels and also a marginally larger efficiency).
+. Real dillerences can. however. be seen at Yc20: while our method with ἐν70.9 has a much lower completeness level than the UIXIDSS pipeline owed: methods. it also achieves a much higher elliciency.," Real differences can, however, be seen at $Y\simeq20$: while our method with $\ps>0.9$ has a much lower completeness level than the UKIDSS pipeline based methods, it also achieves a much higher efficiency."
+ If telescope time is limited and completeness not important. hen basing source selection on {3 can lead to a considerable reduction in wasted! observation time.," If telescope time is limited and completeness not important, then basing source selection on $\ps$ can lead to a considerable reduction in `wasted' observation time."
+ Using our method with [A0.5 leads to completeness ancl efficiency. levels more in line with the UIXIDSS pipeline based methods., Using our method with $\ps>0.5$ leads to completeness and efficiency levels more in line with the UKIDSS pipeline based methods.
+ We have developed a Bayesian formalism for Cclassification in optical or ssurvevs that combines the morphological properties of an object (as measured in multiple passhands) with prior knowledge of the star and ealaxy populations., We have developed a Bayesian formalism for classification in optical or surveys that combines the morphological properties of an object (as measured in multiple passbands) with prior knowledge of the star and galaxy populations.
+ AC fully Javesian approach must also include colour information for sell-consistencv: but. given the aim of combining morphological information correctly. a number of approximations were developed to maximize the influence of the morphological information.," A fully Bayesian approach must also include colour information for self-consistency; but, given the aim of combining morphological information correctly, a number of approximations were developed to maximize the influence of the morphological information."
+10 of the FOV7.,$10$ of the FOV.
+.. Tje statistics is reported in Fig. 2..," 	The statistics is reported in Fig. \ref{fig3},"
+ here the network classes are pointed out by arrows., 	here the network classes are pointed out 	by arrows.
+ Move in detail. the network-EN division performed by the routine has an equivaleit division in the classes.," More in detail, the network-IN division performed by the routine 	has an equivalent division in the classes."
+ Iudeed. the classes that are abundant in the network (1.6.. hose pointed out by he arrows) are scarce in the IN. andvice-versa?.," Indeed, 	the classes that are abundant in the network (i.e., those 	pointed out by the arrows) are scarce in the IN, 	and."
+. These results imply that the network js associated with particular Stokes V shapes that are iot found elsewhere., These results imply that the network is 	associated with particular Stokes $V$ shapes that are not found elsewhere.
+ In other words. a classical network identification aud the classification point out the saue reeions with two differcut methods: the first oue is uainlv based on the analysis of the spatial colereuce of strong Stokes V. sienals. whereas the second one does not ake tuto accouit the amplitude of Stokes V. (profiles are rormalizer before classification. see 3 A). aud refers just to profile srapes.," 	In other words, a classical network identification 	and the classification point out the same 	regions with two different methods: the first 	one is mainly based on the analysis of the spatial coherence 	of strong Stokes $V$ signals, whereas 	the second one does not take into account the amplitude 	of Stokes $V$ (profiles are normalized before classification, 	see \ref{Dataarch}) ), and refers just to profile shapes."
+" The ""uetwork &uuilv contaius he ollowius group of classes:0..3..4.. 6.. 7.. 12.. and (the ¢uster profiles of these classes are reported iu Fig."," 	The “network family” contains the following group of classes:, and 	(the cluster profiles of these classes are 	reported in Fig."
+ 1 with »urple error bars)., \ref{fig1} with purple error bars).
+ These are represcutative of about 28.6 of he total umber of analyzed pixels.,	 	These are representative of about $28.6$ 	of the total number of analyzed pixels.
+ By examining he shapes o the clusters in the network family one cau rotice tha hese profiles prescut the blue lobe arger han the1c red one. while the latter is more extened iu waveleugth (ts area is larger than the area of the blue obe).," 	By examining the shapes of the clusters in the network family 	one can notice that these profiles present 	the blue lobe larger than the red one, while 	the latter is more extended in wavelength (its area is larger than 	the area of the blue lobe)."
+ Such a shape is well known to be associated to jetwork regious (c.g..Steufloetal.198[:SanchezAlneidaetal. 1996).," 	Such a shape is well known to be associated to network regions \citep[e.g.,][]{Ste84,SanA96}."
+". Moreover. the amyplitieο of the Stosox V xofiles iu the two lines is verv siniluw in alb network classes: this highliehts the fact that asauration due o kG fields is preseut. as expected in network regions (sec.ο,ος,Socas-Naviuro&SduchezAlmeida2002:Steuflo 2010)."," 	Moreover, the amplitude of the Stokes $V$ profiles in the two lines is very similar 	in all network classes; this highlights the fact that a saturation due to 	kG fields is present, as expected in network regions \citep[see, e.g., ][]{SocNSanA02,Ste10}."
+. As explained above. each profile curereing frou κοιν with sufficient signal is associaed 0 à class axc. bv eroupius of classes through a ceradn criteria. to a anuly.," 	As explained above, each profile emerging from pixels 	with sufficient signal is associated to a class and, 	by grouping of classes through a certain criteria, 	to a family."
+ In Fie., In Fig.
+ 30 we identify. the position of a network xofile with a purple pixel., \ref{fig2} we identify 	the position of a network profile with a 	purple pixel.
+ The map shows how the retwork coutours hiehlieht laree purple patches., 	The map shows how the network contours highlight 	large purple patches.
+ Ou he other hand. clusters profiles in EN pixels present a mec varicty of asviunuetries," 	 	On the other hand, clusters profiles in IN pixels present a huge variety 	of asymmetries."
+ Ax au example.1.. he most abundant class in the IN. is represented by a cluster xofile cdoniinuatec bv the blue lobe.," As an example, the most abundant class in the IN, 	is represented by a cluster profile dominated by the blue lobe."
+ Maux. otlr Classes are represented by profiles dominated by a sinele oboe. eiher the blue one or the red one (here abwavs jegative becase of the condition inrposed in 3.13).," Many other classes are 	represented by profiles dominated by a single lobe, either the 	blue one or the red one (here always negative because of the condition 	imposed in \ref{Dataarch}) )."
+ The xofiles in these classes. represent approximately 3115€ of he anavzed pixels., The profiles in these classes 	represent approximately $34.4$ of the analyzed pixels.
+" Auioug these. 17.5 are the profiles comaed by he blue lobe. named the ""bluc-Iobe family. lo.1..5.. 10.. 15... aud (represented with blue error bars in Fie."," 	Among these, $17.5$ are the profiles 	dominated by the blue lobe, named the “blue-lobe family”, i.e., 	and (represented with blue error bars in Fig."
+ 1 and blue pixels in Fig. 31) ," \ref{fig1}
+ 	and blue pixels in Fig. \ref{fig2}) )"
+"and 9.6 by the red lobe. named the ""red-lobe family. ie.17..21.. 23.. and (represented wih red error bars in Fig."," 	and $9.6$ by the red lobe, named 	the “red-lobe family”, i.e., and 	(represented with red error bars in Fig."
+ 1 and red pixels iu Fig. 3)).," \ref{fig1}
+ 	and red pixels in Fig. \ref{fig2}) )."
+ Aloreover. a third family of extremely asvuunuetric profiles can be defined:25..26..30.. 31.. 33.. aud (rexeseuted with skv-blue error bars in Fie.," 	Moreover, a third family of extremely asymmetric profiles 	can be defined;, and (represented with sky-blue error bars in Fig."
+ 1 aud skav-bhiο pixels in Fig. 3))," \ref{fig1}
+ 	and sky-blue pixels in Fig. \ref{fig2}) )"
+" form the ""Q-like family”. which contaius 7.3> of the total umber of analyzed pixels."," form the $Q$ -like family”, 	which contains $7.3$ of the total number of analyzed pixels."
+ The classes still not considered are IN profiles preseutiic shapes that are uot clearly among the three families cchued above., 	 	The classes still not considered are IN profiles 	presenting shapes that are not clearly among the 	three families defined above.
+" In spite of this. many of these profiles sti] present evident asvnuunetries: the ""asvinunetrie funi“as formed by the2..8..11.. 13.. 14.. 16.. 18.. 22.. 27.. aud which are 29.6 of the analyzed profiles (represented with ereen error bars in Fig."," In spite of this, 	many of these profiles still present evident asymmetries; 	the “asymmetric family” is formed by the, and which are $29.6$ 	of the analyzed profiles (represented with green error bars in Fig."
+ 1 aud ereen pixels iu Fie. 3)).," \ref{fig1}
+ 	and green pixels in Fig. \ref{fig2}) )."
+ Finally. ouly two classes show alumost antisviunetrie profiles: these represent ouly 6.3 of the dataset (i.c.. and 20)) and are named the “antisvuunetric family” (represented with vellow error bars in Fie.," 	Finally, only two classes show almost antisymmetric 	profiles; these represent only $6.3$ of the dataset 	(i.e., and ) and are named the 	“antisymmetric family” (represented with yellow error bars in Fig."
+" and vellow pixels iu Fie. 3)),"," \ref{fig1}
+ 	and yellow pixels in Fig. \ref{fig2}) )."
+ AΟΠΟΥ iteresting result comes from the study of the position of the familics ou the FOV., 		 	Another interesting result comes from the study of the position 	of the families on the FOV.
+" This caibe done. for example. by studying t1ο LOS velocity (05,4) for the pixels. associate Lto each family."," This can be done, for example, by studying 	the LOS velocity $v_{los}$ ) for the pixels 	associated to each family."
+" Pixels with ej;«0 (Cus Z0) are placexb in erauules (utereranular lÓaues). τοι, upflow reesious (downflow regions)."," 	Pixels with $v_{los}<0$ $v_{los}>0$ ) are 	placed in granules (intergranular lanes), 	i.e., upflow regions (downflow regions)."
+ Fie., 	Fig.
+ |. contaius +tje result of such a study when calculating the statistics Of Cis fas derived w Stokes f. sce 3)) for the pixels oeοισος in the sis faiubes defined above.," \ref{fig4} contains the result of such a study 	when calculating the statistics of $v_{los}$ (as 	derived by Stokes $I$, see \ref{Analysis}) ) for 	the pixels included in the six families defined above."
+ The main result erived from this sπαν ds that each family is related to a certain clynanuics., 	The main result derived from this study is that 	each family is related to a certain dynamics.
+ This is very revealing because the classification xocedure docs not kuow anything about Cros because wavelength shits were removed when extracting Stokes !~ profiles in the waveleneth wiudows centered around 5okes £ cores 3.1)).,	 	This is very revealing because the classification 	procedure does not know anything about $v_{los}$ because 	wavelength shifts were removed when extracting 	Stokes $V$ profiles in the wavelength windows centered 	around Stokes $I$ cores \ref{Dataarch}) ).
+ From this one learns that the plasma cyanics. as derived frou Stokes I. je. the dvnamics mostly associated to the nou uaenetized plasma in the pixel. is iu some wav related he shape of Stokes V. profiles.," 	From this one learns that the plasma dynamics, as derived from Stokes $I$, 	i.e., the dynamics mostly associated to the non magnetized 	plasma in the pixel, is in some way related the shape of 	Stokes $V$ profiles."
+ ο1ο can elaborate a bit nore on this result., One can elaborate a bit more on this result.
+ The separation between the families oue lobe oulv (1.6... the bluc-lo)e aud the red-lobe) is evident.," 	The separation between the families 	with one lobe only (i.e., the blue-lobe and the red-lobe) is evident."
+ Ou the one haud. profiles dominated by the blue obe are mostly found iu upilow regious. Le. on top of he exanules.," 	On the one hand, profiles dominated by the blue lobe are mostly found in upflow regions, 	i.e., on top of the granules."
+ On. the other haud. xofiles dominated. bv he red lobe are mostly found in ¢ownflow regions. ic. lanes.," On 	the other hand, profiles dominated 	by the red lobe are mostly found in downflow regions, i.e., 	intergranular lanes."
+ The Q-like. network. aud asvinietric παλος are found in dowuflow regions.," 	The $Q$ -like, network, and asymmetric families 	are found in downflow regions."
+ Antisviunetric xofiles are found iu both upflow and cownflow pixels with almost the same probaημίν., Antisymmetric profiles 	are found in both upflow and downflow pixels with almost 	the same probability.
+ Overall. as expected. uost of the polarization signals reside in cdownuflows. lo. of the analyzed pixels.," 	Overall, as expected, most of the polarization signals reside 	in downflows, i.e., $68.3$ of the analyzed pixels."
+ Different families are also associated with different spatial extenta., 	 	Different families are also associated with different spatial extents.
+ When, 	When
+" From thedata onhyperon polarizations in pe collisions al pj, = 400GeV/c. we set ",the spin dependent parameter $C_{[ud]}^s$ in \ref{eqn:Pud00a}) ) is fixed from the data on $\Lambda$ polarization in $pBe$ 
+Eq. (17)),Eq. \ref{prode}) )
+ has a general formal structure which is verv similar (ο that of Eq. (1)).," has a general formal structure which is very similar to that of Eq. \ref{prod}) ),"
+ describing the electron-positron. pair production in a constant electric field., describing the electron-positron pair production in a constant electric field.
+ In. both cases the pair production rate is proportional to the intensity of the electric field. multiplied bv an integral over (he transverse momentum of the particles.," In both cases the pair production rate is proportional to the intensity of the electric field, multiplied by an integral over the transverse momentum of the particles."
+ However. since in the electrosphere E(z.T)—0 for large z. the particle production rate also naturally tends to zero aud outside (he electrosphere or in the regions with low electric field the particle production ceases.," However, since in the electrosphere $E(z,T)\rightarrow 0$ for large $z$, the particle production rate also naturally tends to zero and outside the electrosphere or in the regions with low electric field the particle production ceases."
+ The variation of the pair production rate in (he electrosphere. as a function of the distance z is represented for different values of the temperature in Fie.," The variation of the pair production rate in the electrosphere, as a function of the distance $z$ is represented for different values of the temperature in Fig."
+ 3., 3.
+ By introducing a new variable ¢=m?+k?/V(2.T). the pair production rate can also be represented as In order to find an approximate representation of (he pair production rate we power expand the integrand in Eq. (16)).," By introducing a new variable $\zeta
+=\sqrt{m_{e}^{2}+k^{2}}/V\left( z,T\right) $, the pair production rate can also be represented as In order to find an approximate representation of the pair production rate we power expand the integrand in Eq. \ref{sigma}) )."
+ In the first order in 7j we obtain Therefore in the first order of approximation σ is given bv By fixing the electron energy so that w=V(z.2) and taking into account the expression for the electric field in the electrosphere. given by Eq. (10)).," In the first order in $\eta $ we obtain Therefore in the first order of approximation $\sigma $ is given by By fixing the electron energy so that $\omega =V(z,T)$ and taking into account the expression for the electric field in the electrosphere, given by Eq. \ref{el}) ),"
+ it follows that the pair production rate can be written as, it follows that the electron-positron pair production rate can be written as
+from the model of Tavlor and Cordes (1993) are reproduced.,from the model of Taylor and Cordes (1993) are reproduced.
+ We see that the Norma spiral arm dominates (he radial range 0.5 (o 0.6 2. near (he tangent point. the Scutum-Crux arm al the near distance dominates radii 0.6 {ο 0.7 2... and Sagitarius-Carima dominates at about 0.85 H..," We see that the Norma spiral arm dominates the radial range 0.5 to 0.6 $R_{\odot}$ near the tangent point, the Scutum-Crux arm at the near distance dominates radii 0.6 to 0.7 $R_{\odot}$, and Sagitarius-Carina dominates at about 0.85 $R_{\odot}$."
+ There is a gap in the range 0.7 to 0.85 which is between arms., There is a gap in the range 0.7 to 0.85 which is between arms.
+ This appears to be reflected in a dip in & at about 0.8 22. on figure 9., This appears to be reflected in a dip in $\kappa$ at about 0.8 $R_{\odot}$ on figure 9.
+ Even in this dip. and at all other radii in the inner Galaxy. the opacity is larger than its solar circle value of about 5 km | |.," Even in this dip, and at all other radii in the inner Galaxy, the opacity is larger than its solar circle value of about 5 km $^{-1}$ $^{-1}$."
+ This narrow longitude range mav not be representative of (he entire inner galaxy. so we must wait for the full SGPS survey. area to be studied belore deciding conclusively that (he opacity increases as we &o to smaller Galactic radii.," This narrow longitude range may not be representative of the entire inner galaxy, so we must wait for the full SGPS survey area to be studied before deciding conclusively that the opacity increases as we go to smaller Galactic radii."
+ Iu the first quacrant. Garwood and Dickey (1989) found about the same result for the solar circle opacity. but. smaller. values for (he inner galaxy.," In the first quadrant, Garwood and Dickey (1989) found about the same result for the solar circle opacity, but smaller values for the inner galaxy."
+ However. recently lolpak et al. (," However, recently Kolpak et al. ("
+2002) have made more extensive observations with the VLA in the first quadrant. and thev find a trend similar to what we see on figure 9.,"2002) have made more extensive observations with the VLA in the first quadrant, and they find a trend similar to what we see on figure 9."
+ The difference may be due to the small numbers of lines of sight sampled in the Garwood and Dickey study; which may have preferentially saniplecl (he regions between the spiral arms. which appear here to show opacities of 5 km | ! or less. in contrast to the arm regions which can have <5>>20 km ! |.," The difference may be due to the small numbers of lines of sight sampled in the Garwood and Dickey study, which may have preferentially sampled the regions between the spiral arms, which appear here to show opacities of 5 km $^{-1}$ $^{-1}$ or less, in contrast to the arm regions which can have $<\kappa> > 20$ km $^{-1}$ $^{-1}$."
+ If the inerease in «&&> by a [actor of three to five with clecreasing Galactic radius from (he solar circle to the molecular ring region at about 0.5 22. is twpical of the disk overall. we may ask whether it is due simply to an increase in the surface densitv of gas.," If the increase in $<\kappa>$ by a factor of three to five with decreasing Galactic radius from the solar circle to the molecular ring region at about 0.5 $R_{\odot}$ is typical of the disk overall, we may ask whether it is due simply to an increase in the surface density of gas."
+" The surface density as traced by 21-cm emission surveys does not show an increase with decreasing I, in this longitude range (Burton 1983. fig."," The surface density as traced by 21-cm emission surveys does not show an increase with decreasing $_{gal}$ in this longitude range (Burton 1988, fig."
+ 1.15). but the molecular gas does.," 7.15), but the molecular gas does."
+ The strong variation of «&> tells us that the cool atomic gas. in contrast to the warm neutral medium that dominates the A21-em emissivity. is much more abundant in (he inner galaxy. than ad the solar circle.," The strong variation of $<\kappa>$ tells us that the cool atomic gas, in contrast to the warm neutral medium that dominates the $\lambda$ 21-cm emissivity, is much more abundant in the inner galaxy than at the solar circle."
+ This underlines the role of the CNM clouds as an intermediate population between structures in the WNM and the much denser and colder molecular clouds., This underlines the role of the CNM clouds as an intermediate population between structures in the WNM and the much denser and colder molecular clouds.
+" The integrals of the emissioni spectra are given on table 1. columns 8 ancl 9 in units of 107"" 7."," The integrals of the emission spectra are given on table 1, columns 8 and 9 in units of $^{20}$ $^{-2}$."
+ Column 8 gives the integral of the emission uncorrected [or absorption. i.e. where we integrate over (he full negative velocity range corresponding to the inner Galaxy.," Column 8 gives the integral of the emission uncorrected for absorption, i.e. where we integrate over the full negative velocity range corresponding to the inner Galaxy."
+ In column 9 we take the emission integral over the restricted velocity range AW corresponding to the absorption integral in column 7 (i.e. from zero to just short of the recombination line velocity)., In column 9 we take the emission integral over the restricted velocity range $\Delta V$ corresponding to the absorption integral in column 7 (i.e. from zero to just short of the recombination line velocity).
+ For (his velocity range we can correct the values of emission for absorption to get a better estimate of the (rue column density. assuming a one-phase medium (see next section)., For this velocity range we can correct the values of emission for absorption to get a better estimate of the true column density assuming a one-phase medium (see next section).
+ That correction is given by, That correction is given by
+sets are very noisy.,sets are very noisy.
+ Based on these data. Abe et ((2008)) found a rotation period of 7.6272€ hh. and a spin orientation of ey=33107. Bay=-20.07. indicating a prograde rotation.," Based on these data, Abe et \cite{abe08}) ) found a rotation period of $\pm$ h and a spin orientation of $\lambda_{\mathrm{ecl}}\,=\,331.0^{\circ}$, $\beta_{\mathrm{ecl}}\,=\,+20.0^{\circ}$, indicating a prograde rotation."
+ A more recent analysis by the same authors communication) resulted in à rotation model with slightly different values: ο=327.37. Bey=634.7. Pog hhours.," A more recent analysis by the same authors communication) resulted in a rotation model with slightly different values: $\lambda_{\mathrm{ecl}}\,=\,327.3^{\circ}$, $\beta_{\mathrm{ecl}}\,=\,+34.7^{\circ}$, $_{\mathrm{sid}}$ hours."
+ Both solutions were derived using the epoch and amplitude methods described by Magnusson (1986))., Both solutions were derived using the epoch and amplitude methods described by Magnusson \cite{magnusson86}) ).
+ These methods are reliable for trregularly shaped bodies and for sufficient lightcurve data covering various aspect angles., These methods are reliable for irregularly shaped bodies and for sufficient lightcurve data covering various aspect angles.
+ has a comparatively spherical shape and the available lightcurve data were from restricted directions., has a comparatively spherical shape and the available lightcurve data were from restricted directions.
+ Nevertheless. both solutions gave a reasonable match to the observed lighteurves. but it turned out that these solutions are not unique and other parameter sets with different rotation periods. spin-axis orientations and shape models could not be excluded (see Fig. 1)).," Nevertheless, both solutions gave a reasonable match to the observed lightcurves, but it turned out that these solutions are not unique and other parameter sets with different rotation periods, spin-axis orientations and shape models could not be excluded (see Fig. \ref{fig:lc}) )."
+ Since these values are crucial input parameters for our thermophysical model analysis. we repeated the search for possible shape and spin-vector solutions using the lghtcurve inversion method developed by Kaasalainen and Torppa (2001)).," Since these values are crucial input parameters for our thermophysical model analysis, we repeated the search for possible shape and spin-vector solutions using the lightcurve inversion method developed by Kaasalainen and Torppa \cite{kaasalainen01}) )."
+ Our goal was to derive a set of the most likely convex shape models that would fit all available lighteurves., Our goal was to derive a set of the most likely convex shape models that would fit all available lightcurves.
+ Because of the poor quality of some of the data. we were only able to determine the range of the sidereal rotation period to hhours.," Because of the poor quality of some of the data, we were only able to determine the range of the sidereal rotation period to hours."
+ Different shape models with pole orientations covering almost the entire celestial sphere (without any preference for pro- or retrograde solutions) fit the data equally well., Different shape models with pole orientations covering almost the entire celestial sphere (without any preference for pro- or retrograde solutions) fit the data equally well.
+ By scanning the period-pole parameter space. we derived 77 shape models that were physically acceptable (rotating around the shortest axis) and for which the y of the fit was no more than higher than the best-fit y.," By scanning the period-pole parameter space, we derived 77 shape models that were physically acceptable (rotating around the shortest axis) and for which the $\chi^2$ of the fit was no more than higher than the best-fit $\chi^2$."
+ These models corresponded to local minima in theparameter space., These models corresponded to local minima in theparameter space.
+ In addition. we included both of the original Abe et ((2008)) solutions and added another five with the pole fixed to the two Cli4. Bocr)-pairs mentioned above and rotation periods in the range given by Abe et ((2008)).," In addition, we included both of the original Abe et \cite{abe08}) ) solutions and added another five with the pole fixed to the two $\lambda_{\mathrm{ecl}}$ $\beta_{\mathrm{ecl}}$ )-pairs mentioned above and rotation periods in the range given by Abe et \cite{abe08}) )."
+ Only two of these models pass the y + limit on the basis of the visual lightcurves. the other five models have higher y7-values.," Only two of these models pass the $\chi^2$ + limit on the basis of the visual lightcurves, the other five models have higher $\chi^2$ -values."
+ Six of these shape models are elongated along the “Zz axis and therefore unphysical., Six of these shape models are elongated along the 'z' axis and therefore unphysical.
+ For all 84 models we performed a thermophysical model analysis for awide range of possible parameters (see Table 3))., For all 84 models we performed a thermophysical model analysis for awide range of possible parameters (see Table \ref{tbl:tpm_params}) ).
+ The mid-IR photometric data were already described in Hasegawa et ((2008))., The mid-IR photometric data were already described in Hasegawa et \cite{hasegawa08}) ).
+ The data set includes 15 N-band Subaru observations and two dedicated observations at 15 and um. We binned the single-epochSpitzer IRS data (Campins et citecampins09)) into 20 wavelength points (4 for band SL2. for SLI. 4 for LL2. 5 for LLI: see Fig. 5..," The data set includes 15 N-band Subaru observations and two dedicated observations at 15 and $\mu$ m. We binned the single-epoch IRS data (Campins et \\cite{campins09}) ) into 20 wavelength points (4 for band SL2, 7 for SL1, 4 for LL2, 5 for LL1; see Fig. \ref{fig:irs_obs_mod},"
+ bottom)., bottom).
+ The 20 wavelength points were chosen to give about equal weight to the two published data samples in terms of number of observations (17 in Hasegawa et citehasegawa08 and 20 for the Campins et citecampins09— sample).," The 20 wavelength points were chosen to give about equal weight to the two published data samples in terms of number of observations (17 in Hasegawa et \\cite{hasegawa08}
+ and 20 for the Campins et \\cite{campins09} sample)."
+ In this way the derived object properties are better connected to the entire data set. and they do not just match the measurements of one observing epoch.," In this way the derived object properties are better connected to the entire data set, and they do not just match the measurements of one observing epoch."
+ Each observation set also has a mixture of lower and higher quality data: The ground-based Subaru data are of lower quality than the data. while the quality of the IRS spectrum changes with wavelength.," Each observation set also has a mixture of lower and higher quality data: The ground-based Subaru data are of lower quality than the data, while the quality of the IRS spectrum changes with wavelength."
+ This again ensures that the final solutions are not biased towards à single measurement., This again ensures that the final solutions are not biased towards a single measurement.
+ All observations are listed in Table 2.., All observations are listed in Table \ref{tbl:tpm_obs}.
+ All 84 possible spin-vector and shape solutions from Sect., All 84 possible spin-vector and shape solutions from Sect.
+ have been used in combination with these thermal data., \ref{sec:sv} have been used in combination with these thermal data.
+ For our analysis we are using a thermophysical model (TPM) described by Lagerros (1996.. 1997.. 1998)) and Mülller Lagerros (1998)).," For our analysis we are using a thermophysical model (TPM) described by Lagerros \cite{lagerros96}, , \cite{lagerros97},, \cite{lagerros98}) ) and Mülller Lagerros \cite{mueller98}) )."
+ This TPM works with true, This TPM works with true
+region and the northern spiral arm (including complex C). are predominantly. composed. of lower-abundance star clusters.,"region and the northern spiral arm (including complex C), are predominantly composed of lower-abundance star clusters."
+ Similarly. the addition of NUR passbands appears to result in more solutions with lower extinction estimates (Figs.," Similarly, the addition of NIR passbands appears to result in more solutions with lower extinction estimates (Figs."
+ See and E)., \ref{met3310.fig}e e and f).
+ This is likely due to the better constraints on the extinction. from. optical-NIIS. compared to optical-only baselines., This is likely due to the better constraints on the extinction from optical-NIR compared to optical-only baselines.
+ However. whether this result. is. indeed significant also depends on the suitability of the starburst ealaxv-tvpe extinction. law adopted for our fitting routine (Calzetti ct al.," However, whether this result is indeed significant also depends on the suitability of the starburst galaxy-type extinction law adopted for our fitting routine (Calzetti et al."
+ 2000). of which the analysis is bevonc the scope of the present paper (but see Calzetti. Ixinnev. Storchi-Bergmann |1994]. for a. discussion).," 2000), of which the analysis is beyond the scope of the present paper (but see Calzetti, Kinney, Storchi-Bergmann [1994] for a discussion)."
+ Nevertheless. these elfects are negligible. or marginallv significant at worst. compared to the photometric uncertainties (0.10.—0.15 mag}.," Nevertheless, these effects are negligible, or marginally significant at worst, compared to the photometric uncertainties $\sim 0.10-0.15$ mag)."
+ We can conclude from this exercise that NGC 3310 underwent a significant burst of. cluster formation. some 3510° vr ago., We can conclude from this exercise that NGC 3310 underwent a significant burst of cluster formation some $3 \times 10^7$ yr ago.
+ The actual duration of the burst of cluster formation may have been shorter because uncertainties in the age determinations may have broacened the peak., The actual duration of the burst of cluster formation may have been shorter because uncertainties in the age determinations may have broadened the peak.
+ [t appears. therefore. that the peak of cluster formation in NGC 3310 coincides closely with the suspected galactic cannibalism or last tical interaction (van der Ixruit. 1976. BUS. POs4. P98. 896. 02). while the possibly more recent. marginally significant cluster formation at /10 vr can be interpreted as cluster formation associated. with the most recent (&LO Myr) enhanced star formation episode traced by the circumnuclear Ho-bright sources (P93. DOO).," It appears, therefore, that the peak of cluster formation in NGC 3310 coincides closely with the suspected galactic cannibalism or last tidal interaction (van der Kruit 1976, BH81, TG84, P93, S96, E02), while the possibly more recent, marginally significant cluster formation at $t
+\lesssim 10^7$ yr can be interpreted as cluster formation associated with the most recent $\lesssim 10$ Myr) enhanced star formation episode traced by the circumnuclear $\alpha$ -bright sources (P93, D00)."
+ The clusters older than ~105 vr are. smoothly distributed throughout the ealactic centre. roughly following the inner ring structure and the other concentrations of star clusters (sec Fig. 2)).," The clusters older than $\sim 10^7$ yr are smoothly distributed throughout the galactic centre, roughly following the inner ring structure and the other concentrations of star clusters (see Fig. \ref{clustercoords.fig}) )."
+ Llowever. the vounger clusters. with ages log(age/vr)<7.1 (1οι. the voungest two age bins in bie. 3)).," However, the younger clusters, with ages $\log( {\rm age/yr}) \le 7.1$ (i.e., the youngest two age bins in Fig. \ref{n3310ages.fig}) ),"
+ are. predominantly concentrated. in the Jumbo region (7: nomenclature from POS) and in the northern spiral arm. which also contains the bright star-forming regions C and E (see Fig. 2)).," are predominantly concentrated in the Jumbo region (“A”; nomenclature from P93) and in the northern spiral arm, which also contains the bright star-forming regions C and E (see Fig. \ref{clustercoords.fig}) )."
+ This is consistent with the very voung ages derived. for these regions. /10 Myr (Verlevich et al.," This is consistent with the very young ages derived for these regions, $t \lesssim 10$ Myr (Terlevich et al."
+ 1990. P93. DOO. E02). although we emphasize that both regions also contain clusters spanning the entire age range observed for the NGC 3310 star cluster svstenm.," 1990, P93, D00, E02), although we emphasize that both regions also contain clusters spanning the entire age range observed for the NGC 3310 star cluster system."
+ ligure Ὁ shows our best estimates of the clusters’ masses., Figure \ref{mass3310.fig} shows our best estimates of the clusters' masses.
+ These were obtained by scaling our mocel SEDs (for masses of 1.610ΑΠ. to the observed. SEDs for the appropriate combination of age. metallicity and (total) extinction.," These were obtained by scaling our model SEDs (for masses of $1.6 \times 10^9 M_\odot$ ) to the observed SEDs for the appropriate combination of age, metallicity and (total) extinction."
+ Our model SEDs are based on a Salpeter-type IME. consisting of stellar. massesa in the range 0.15.«mfAL.<(5070) (Schulz ct al., Our model SEDs are based on a Salpeter-type IMF consisting of stellar masses in the range $0.15 \le m/M_\odot \le (50-70)$ (Schulz et al.
+ 2002)., 2002).
+ The exact upper mass limit for our model SSPs depends on the metallicity and is determüned by the mass coverage of the Padova isochrones., The exact upper mass limit for our model SSPs depends on the metallicity and is determined by the mass coverage of the Padova isochrones.
+ The exact value for the upper mass limit is unimportant for the determination. of the total cluster κ..., The exact value for the upper mass limit is unimportant for the determination of the total cluster masses.
+ We realise. that recent determinations of the stellar IME deviate significantly [rom a Salpeter-tvpe EME at low, We realise that recent determinations of the stellar IMF deviate significantly from a Salpeter-type IMF at low
+This final result demonstrates that the conditional redshift posterior is essentially proportional to the intensity function of the Poissonian likelihood.,This final result demonstrates that the conditional redshift posterior is essentially proportional to the intensity function of the Poissonian likelihood.
+ It is also interesting to note. that the right hand side of equation (C173) does not depend on the galaxy coordinates (C173)," It is also interesting to note, that the right hand side of equation \ref{eq:cond_redshift_posterior_1d}) ) does not depend on the galaxy coordinates \ref{eq:cond_redshift_posterior_1d})"
+where is (he specific energv of a particle in the disk.,where $E^+$ is the specific energy of a particle in the disk.
+ When Jf=0. ie.. there is no magnetic coupling. equation (7)) and equation (8)) retum to the equations derived. by NovikovandThorne(1973).. and PageandThorne(1974).," When $H=0$, i.e., there is no magnetic coupling, equation \ref{con_ang}) ) and equation \ref{con_ener}) ) return to the equations derived by \citet{nov73}, and \citet{pag74}."
+". E is related to {, bv The flux of angular momentum (transfered from the black hole to the disk. 77. is defined bv the torque 77,5 produced by the black hole on the disk Comparing equation (10)) with equation (3)). we have if the magnetic field is smoothly distributed on the disk."," $E^+$ is related to $L^+$ by \citep{pag74}
+ The flux of angular momentum transfered from the black hole to the disk, $H$, is defined by the torque $T_{HD}$ produced by the black hole on the disk Comparing equation \ref{flux_ang}) ) with equation \ref{toq2}) ), we have if the magnetic field is smoothly distributed on the disk."
+ In a quasi-steady state. My is constant (hroughout (he disk (the conservation of mass).," In a quasi-steady state, $\dot{M}_D$ is constant throughout the disk (the conservation of mass)."
+ Then. equation (7)) aud equation (3)) can be solved lor F and g by using equation (9)).," Then, equation \ref{con_ang}) ) and equation \ref{con_ener}) ) can be solved for $F$ and $g$ by using equation \ref{el}) )."
+" The solutions are and where (he subscript ης denotes (he mareinally stable orbit which is assumed to be (he inner boundary of a thin Ixeplerian disk. g,,: is an integration constant. / is defined by The integration of / has been worked out bv PageandThorne(1974). and is given by their equation (151). we do not display it here since it is lengthy."," The solutions are and where the subscript $ms$ ” denotes the marginally stable orbit which is assumed to be the inner boundary of a thin Keplerian disk, $g_{ms}$ is an integration constant, $f$ is defined by The integration of $f$ has been worked out by \citet{pag74} and is given by their equation (15n), we do not display it here since it is lengthy."
+" It is easy to check that für«=ry)x=O and für"">rus)ςemRAApSEE LÀ.", It is easy to check that $f(r=r_{ms})=0$ and $f(r \gg r_{ms})\approx{3M_H\over 2r^2}$ .
+Iu chunk A the power spectrum las already: changed with respect to the earlier observations: the power at higher frequencies has disappeared. aud there is oulv power iu the rauge 0.0510 Tz.,"In chunk A the power spectrum has already changed with respect to the earlier observations: the power at higher frequencies has disappeared, and there is only power in the range 0.05–10 Hz."
+ Probably we are witucssing the onset ofthe 1 Wz modulation., Probably we are witnessing the onset of the 1 Hz modulation.
+ The average power spectrum shows a fractional amplitude of (2141)% rni in the 0.0510 Tz baud., The average power spectrum shows a fractional amplitude of $(21\pm 1)\%$ rms in the 0.05–10 Hz band.
+ In chunk Bowe find a 1 Tz QPO clearly present with a fractional amplitude of (GS+L1) rius., In chunk B we find a 1 Hz QPO clearly present with a fractional amplitude of $(68\pm4)\%$ rms.
+ We then caculated power spectra from 256 s data segments and determined the fractional rius amplitude in the 0.05-10 Iz baud. (, We then calculated power spectra from 256 s data segments and determined the fractional rms amplitude in the 0.05-10 Hz band. (
+see Fie. 8)).,see Fig. \ref{appears}) ).
+ Significant power is detected in three seeioeuts during the observations of chunk A. meaning that some low level activity is already preseut in the same frequency auge where the OPO will later appear.," Significant power is detected in three segments during the observations of chunk A, meaning that some low level activity is already present in the same frequency range where the QPO will later appear."
+ In chunk D power is nearly always detected: the ouset of the 1 Hz QPO inst have occurred. within the 5000s eap., In chunk B power is nearly always detected; the onset of the 1 Hz QPO must have occurred within the 5000 s gap.
+ At NLJD 53510.15 the amplitude shows au abrupt decrease from GO% ris down to <30% confidence level) on a time scale of 256 s. We also calculated the power in a umber of fequency bands in the rauge 1/128 to 256 Iz for all power spectra of chunk A. The power is always cousistent with zero at the three sigma level in all the frequency bands except for the 0.05-10 ITz baud in the three scements mentioned above., At MJD 53540.48 the amplitude shows an abrupt decrease from $\sim 60\%$ rms down to $<30\%$ confidence level) on a time scale of $256$ s. We also calculated the power in a number of frequency bands in the range 1/128 to 256 Hz for all power spectra of chunk A. The power is always consistent with zero at the three sigma level in all the frequency bands except for the 0.05-10 Hz band in the three segments mentioned above.
+ Tn Fig., In Fig.
+ 9 we show the cuerey dependence of the 1 Iz ΟΡΟ., \ref{energy} we show the energy dependence of the 1 Hz QPO.
+" The points in the figure refer to the observation (one per outburst) for which the fractional rims of the 1 Tz OPO in the 2-60 keV. euerev baud was the highest (110.5% riis in 2000. ObsId L0035-01-0L-01, total exposure 7 ks: τς rius in 2002. ObsId T0080-03-15-00. total exposure 2 ks: 125% ruis in 2005. ObsId 91115-01-02-05. total exposure 1.2 ks)."," The points in the figure refer to the observation (one per outburst) for which the fractional rms of the 1 Hz QPO in the 2-60 keV energy band was the highest $110.5\%$ rms in 2000, ObsId 40035-01-04-01, total exposure 7 ks; $117\%$ rms in 2002, ObsId 70080-03-15-00, total exposure 2 ks; $125\%$ rms in 2005, ObsId 91418-01-02-05, total exposure 1.2 ks)."
+ In these observations the characteristic frequencies of the 1 IIz QPO aud its, In these observations the characteristic frequencies of the 1 Hz QPO and its
+"of the photons is assumed to be E—8.7 keV. Therefore. m,22E,=11.4 keV. Since the size of a degenerate sterile neutrino halo with total mass M,<109M. and m,z17 keV is much smaller than 20 pc. the total energy [lux of the decaved photons within the field of view of Suzaku (solid angle Q) is given by where ry=8.5 kpe is the distance to the GC and P=ΛΙΓΟ/210! eres tis the total power emitted in Che sterile neutrino decays (the required power of the hard component of the hot gas nearQC).","of the photons is assumed to be $E_s=8.7$ keV. Therefore, $m_s \approx 2E_s=17.4$ keV. Since the size of a degenerate sterile neutrino halo with total mass $M_s \le 
+10^6M_{\odot}$ and $m_s \approx 17$ keV is much smaller than 20 pc, the total energy flux of the decayed photons within the field of view of $Suzaku$ (solid angle $\Omega$ ) is given by where $r_0=8.5$ kpc is the distance to the GC and $P=M_s \Gamma 
+c^2/2=10^{40}$ erg $^{-1}$ is the total power emitted in the sterile neutrino decays (the required power of the hard component of the hot gas nearGC)."
+ The excess energy [ιν observed by Suzaku is E!—(1120.6)0.8)xIO. ere em> !.," The excess energy flux observed by $Suzaku$ is $F_s'=(1.1 
+\pm 0.6) \times 10^{-5}E_s \approx (1.5 \pm 0.8) \times 10^{-13}$ erg $^{-2}$ $^{-1}$."
+ Therefore. only around of photons can escape from the GC. and most of the emitted photons are first absorbed by (he gas clouds inside (he sterile neutrino halo and nearby.," Therefore, only around of photons can escape from the GC, and most of the emitted photons are first absorbed by the gas clouds inside the sterile neutrino halo and nearby."
+ The optical depth 7 is given by where c; is (he ellective absorption cross section of 8.7 keV photons by different. gas components including cold molecular hydrogen gas(/ —I5). warm atomic gas (1-1. He) hot ionized gas (G—hot) and very hot ionized gas (G—vhoti). n; is the number densitv of the gas components.," The optical depth $\tau$ is given by where $\sigma_i$ is the effective absorption cross section of 8.7 keV photons by different gas components including cold molecular hydrogen gas$i=$ $_2$ ), warm atomic gas $i$ =H, He) hot ionized gas $i$ =hot) and very hot ionized gas $i$ =vhot), $n_i$ is the number density of the gas components."
+" Assuming all the gas components follow the same density profile near the GC. whieh ean be modelled by (Sehodeletal.2002) where r,—0.34 pc and iy; are the core radius andcore munber density for the;'- gas component."," Assuming all the gas components follow the same density profile near the GC, which can be modelled by \citep{Schodel}
+ where $r_c=0.34$ pc and $n_{0,i}$ are the core radius andcore number density for the$i^{\rm th}$ -type gas component."
+" For HI molecular aud atomic gases. the average number densities within 450 pe are T4 aE and 0.9 aκ respectively (Ferriereetal. 2007).. which corresponds {ο No,=12x10* 7 and nij;=1.5x10 *."," For H molecular and atomic gases, the average number densities within 450 pc are 74 $^{-3}$ and 0.9 $^{-3}$ respectively \citep{Ferriere}, , which corresponds to $n_{0,H_2}=1.2 
+\times 10^7$ $^{-3}$ and $n_{0,H}=1.5 \times 10^5$ $^{-3}$."
+ Since hydrogen and helium constitute and by mass in the interstellar medium respectively (Ferriere 2001).. we have Notte&OLXοὗog)=2.5x105 7.," Since hydrogen and helium constitute and by mass in the interstellar medium respectively \citep{Ferriere2}, , we have $n_{0,He} \approx 0.1 \times 
+(2n_{0,H2}+n_{0,H})=2.5 \times 10^6$ $^{-3}$."
+ The hot ionic (10?—10° IX) and very hot ionic (10°—10 Ix) gas components constitute44-54% and0.3%—0.5% by mass of all interstellar gas respectively (Ferriereοἱal. 2007).. which correspond to rope=(1—1.2)x109) a. and Novy=(7—12)x10! 7.," The hot ionic $10^5-10^6$ K) and very hot ionic $10^7-10^8$ K) gas components constitute$-$ and $-$ by mass of all interstellar gas respectively \citep{Ferriere}, , which correspond to $n_{0, \rm hot}=(1-1.2) \times 10^6$ $^{-3}$ and $n_{0, \rm vhot}=(7-12) \times 10^4$ $^{-3}$ ."
+" The effective photoionization cross sections of Ils. II. He. hot metal ionsand very hot metal ions by &.7keV photons are oy,=2.2x107° cnr. eq=l.lx1075 οι”. ay.=3.8x107 απ. ong=(1.6—1.9)x1077 ene? and Tout=(1.5—1.6)x107! em? respectively"," The effective photoionization cross sections of $_2$ , H, He, hot metal ionsand very hot metal ions by 8.7keV photons are $\sigma_{\rm H_2}=2.2 \times 10^{-26}$ $^2$, $ 
+\sigma_{\rm H}=1.1 \times 10^{-26}$ $^2$ , $\sigma_{\rm He}=3.8 \times 10^{-25}$ $^2$ , $\sigma_{\rm 
+hot}=(1.6-1.9) \times 10^{-24}$ $^2$ and $\sigma_{\rm 
+vhot}=(1.5-1.6) \times 10^{-24}$ $^2$ respectively"
+are related to the integrated power spectrum of the data.,are related to the integrated power spectrum of the data.
+ Each simulated light curve was constructed as follows., Each simulated light curve was constructed as follows.
+ We first generated a time series with 10 s time resolution that is longer bv up to a factor of 2 (han an actual lieht curve., We first generated a time series with 10 s time resolution that is longer by up to a factor of 2 than an actual light curve.
+ For example. to simulate light curves for Ton S180 a time series of 524.288 points was required.," For example, to simulate light curves for Ton S180 a time series of 524,288 points was required."
+ A light curve was then drawn with random starting time trom Chis time series. (hen resampled and binned exactly as the Iheht curve.," A light curve was then drawn with random starting time from this time series, then resampled and binned exactly as the light curve."
+ Ho was then rescaled by adding the average count rate so that Poisson noise could be added by replacing each point by one drawn from a Poisson distribution havine (he same number of counts., It was then rescaled by adding the average count rate so that Poisson noise could be added by replacing each point by one drawn from a Poisson distribution having the same number of counts.
+ Periodograms for each of these 100 light curves were conputec., Periodograms for each of these 100 light curves were computed.
+ These periodograms were rebinned logarithmically. and the average and standard deviation ab each frequency were computed.," These periodograms were rebinned logarithmically, and the average and standard deviation at each frequency were computed."
+ They were then compared using 47 with the power spectrum [rom the observed light curve., They were then compared using $\chi^2$ with the similarly-rebinned power spectrum from the observed light curve.
+ The minimunm in 47 was found as a [unction of a and total variance., The minimum in $\chi^2$ was found as a function of $\alpha$ and total variance.
+ Mininnun 4? were obtained for a=1.55. 1.3. and 1.3 for4711.. Ton S180. and III 0419577. respectively.," Minimum $\chi^2$ were obtained for $\alpha = 1.55$, 1.3, and 1.3 for, Ton S180, and 1H 0419–577, respectively."
+ We note that the minima are generally [airlv. broad. so the range in a and σ must contribute to uncertainty in the final significance of the candidate signals.," We note that the minima are generally fairly broad, so the range in $\alpha$ and $\sigma$ must contribute to uncertainty in the final significance of the candidate signals."
+ Finally. 1.000 simulated light curves were generated for each object using its best-fittec values of a and σ of the lisht curve. their periodograms were generated. ancl (he resulting exponential probability distribution of power at each Irequeney. was calculated.," Finally, 1,000 simulated light curves were generated for each object using its best-fitted values of $\alpha$ and $\sigma$ of the light curve, their periodograms were generated, and the resulting exponential probability distribution of power at each frequency was calculated."
+ Upper limits on the simulated: power al various confidence levels are graphed in comparison with the actual periodograms in Figures 7- 9.., Upper limits on the simulated power at various confidence levels are graphed in comparison with the actual periodograms in Figures \ref{ps1}- \ref{ps3}.
+ For example. the confidence limit for Ton S180 is below the observed peak at 2.08 davs. which means that fewer than of the simulations had as much power al 2.08 clavs as is observed.," For example, the confidence limit for Ton S180 is below the observed peak at 2.08 days, which means that fewer than of the simulations had as much power at 2.08 days as is observed."
+" In Figure δ., one can also see features in (he simulations of Ton S130 near a period of 1 day ancl its harmonic (hat are caused by eaps in light curve where bad points in the observation were removed: (his process has no effect on the significance of the detected signal at 2.08 days."," In Figure \ref{ps2}, one can also see features in the simulations of Ton S180 near a period of 1 day and its harmonic that are caused by gaps in light curve where bad points in the observation were removed; this process has no effect on the significance of the detected signal at 2.08 days."
+ We note that the wav simulated periodograms are created for Chis application constitutes a conservative test in (hat we distribute of the observed variance in the light curve among the red-noise components. whereas some of the observed. variance could be due to the periodic signal that we are evaluating. if it is real.," We note that the way simulated periodograms are created for this application constitutes a conservative test in that we distribute of the observed variance in the light curve among the red-noise components, whereas some of the observed variance could be due to the periodic signal that we are evaluating, if it is real."
+" The formal. single-trial probability [or chance occurrence of the 2.08 day. signal in Ton S180 is 1.6x10.1,"," The formal, single-trial probability for chance occurrence of the 2.08 day signal in Ton S180 is $1.6 \times 10^{-4}$."
+ However. one must also take into account the number of independent frequencies searched ancl (he possibility that a false signal might arise al any one of these.," However, one must also take into account the number of independent frequencies searched and the possibility that a false signal might arise at any one of these."
+ In the analvsis of Ton 8180. we searched 128 independent frequencies [rom 3.5x10* + to 4.6105 I which is half the Nvquist frequency for a sampling interval of one satellite orbit.," In the analysis of Ton S180, we searched 128 independent frequencies from $3.5 \times 10^{-7}$ $^{-1}$ to $4.6 \times 10^{-5}$ $^{-1}$, which is half the Nyquist frequency for a sampling interval of one satellite orbit."
+" While there is no rigorous wav ol extending the sinele-trial significance to other frequencies that have different mean power levels. we can approximate the ""global significance’ of the detection in a manner similar to Benllochetal.(2001).. bv multiplving the single-trial"," While there is no rigorous way of extending the single-trial significance to other frequencies that have different mean power levels, we can approximate the “global significance” of the detection in a manner similar to \cite{bwe01}, by multiplying the single-trial"
+fraction increases with increasing primary spectral type (for reviews see e.g.. Sterzik Durisen 2004: Kohler et al..,"fraction increases with increasing primary spectral type (for reviews see e.g., Sterzik Durisen 2004; Kohler et al.,"
+ 2006: Lada. 2006: Bouy et al..," 2006; Lada, 2006; Bouy et al.,"
+ 2006: Bate. 2009).," 2006; Bate, 2009)."
+ Direct observations of massive stars reveal that most of them are born às binary members (e.g.. Kobulnicky Fryer. 2007; Sana et al.," Direct observations of massive stars reveal that most of them are born as binary members (e.g., Kobulnicky Fryer, 2007; Sana et al."
+ 2008; Mason et al., 2008; Mason et al.
+ 2009 )., 2009 ).
+ By using the adaptive optics imaging surveys of Shatsky Tokovinin (2002) of the association Sco OB2. Kouwenhoven et al. (," By using the adaptive optics imaging surveys of Shatsky Tokovinin (2002) of the association Sco OB2, Kouwenhoven et al. ("
+2007) concluded that the primordial intermediate mass binary frequency of the aggregate was close to100%.. are intermedate mass binaries.,"2007) concluded that the primordial intermediate mass binary frequency of the aggregate was close to, are intermedate mass binaries."
+ Duquennoy Mayor (1991) find a Solar-type star multiplicity fraction of 0.58. the M-dwarf binary frequency = 0.42 (Fisher Marcy (1992) whereas probably no more than 20-25 of the very low mass objects (masses between 0.003 M. and 0.1 M.) are binary components (for a review. see e.g.. Burgasser et al..," Duquennoy Mayor (1991) find a Solar-type star multiplicity fraction of 0.58, the M-dwarf binary frequency = 0.42 (Fisher Marcy (1992) whereas probably no more than 20-25 of the very low mass objects (masses between 0.003 $M_{\odot}$ and 0.1 $M_{\odot}$ ) are binary components (for a review, see e.g., Burgasser et al.,"
+ 2007)., 2007).
+ It is interesting to remark that the hydrodynamical simulations of star cluster formation of Bate et al. (, It is interesting to remark that the hydrodynamical simulations of star cluster formation of Bate et al. (
+2003) and Bate (2009) closely reproduce the varying observed binary frequency discussed above.,2003) and Bate (2009) closely reproduce the varying observed binary frequency discussed above.
+ An indirect indication that the intermediate mass close binary frequency is high comes from population synthesis simulations., An indirect indication that the intermediate mass close binary frequency is high comes from population synthesis simulations.
+ The latter reveal that in order to explain the observed supernova type Ia rate in galaxies. the intermediate mass close binary frequency must be very large. corresponding to the observations of Kouwenhoven et al.," The latter reveal that in order to explain the observed supernova type Ia rate in galaxies, the intermediate mass close binary frequency must be very large, corresponding to the observations of Kouwenhoven et al."
+ mentioned above (e.g.. Han Podsiadlowski. 2004; Yungelson Livio. 2000: Ruiter et al..," mentioned above (e.g., Han Podsiadlowski, 2004; Yungelson Livio, 2000; Ruiter et al.,"
+ 2009; Mennekens et al., 2009; Mennekens et al.
+ 2010)., 2010).
+ The question then naturally arises: was the initial binary frequency of the first generation of stars in. the Globular Clusters similar as discussed above?, The question then naturally arises: was the initial binary frequency of the first generation of stars in the Globular Clusters similar as discussed above?
+ The observations yield much smaller values: Albrow et al. (, The observations yield much smaller values: Albrow et al. (
+2001) conclude that the binary fraction in 47 Tue ts smaller than whereas Hubble Space Telescope observations of NGC 6397 yield a binary frequency x (Davis et al..,"2001) conclude that the binary fraction in 47 Tuc is smaller than whereas Hubble Space Telescope observations of NGC 6397 yield a binary frequency $\le$ (Davis et al.,"
+ 2008)., 2008).
+ A recent study of Milone et al. (, A recent study of Milone et al. (
+2011) confirms the low binary fraction in a large sample of Globular Clusters.,2011) confirms the low binary fraction in a large sample of Globular Clusters.
+ However. the age of these clusters is very large (up to 10 Gyr and larger). and the low mass binary frequency may have changed considerably compared to the initial one due to the dynamical interaction of stars and binaries in dense clusters as shown by Hurley et al. (," However, the age of these clusters is very large (up to 10 Gyr and larger), and the low mass binary frequency may have changed considerably compared to the initial one due to the dynamical interaction of stars and binaries in dense clusters as shown by Hurley et al. ("
+2007). Fregeau et al. (,"2007), Fregeau et al. ("
+2009). Chatterjee et al. (,"2009), Chatterjee et al. ("
+2010). Vesperini et al. (,"2010), Vesperini et al. ("
+2010).,2010).
+ In the present paper we will focus on the effects of massive and intermediate mass stars and binaries on the chemical evolution of Globular Clusters and we will assume that initially the Globular Cluster population is similar as the field population (similar binary frequencies. similar binary period and mass ratio distributions).," In the present paper we will focus on the effects of massive and intermediate mass stars and binaries on the chemical evolution of Globular Clusters and we will assume that initially the Globular Cluster population is similar as the field population (similar binary frequencies, similar binary period and mass ratio distributions)."
+ Since massive and intermediate mass stars have an evolutionary lifetime < 300 Myr. the latter four studies illustrate that it is not very likely that this population is much affected by stellar dynamics.," Since massive and intermediate mass stars have an evolutionary lifetime $\le$ 300 Myr, the latter four studies illustrate that it is not very likely that this population is much affected by stellar dynamics."
+ De Mink et al. (, De Mink et al. (
+2009) propose non-conservative evolution of massive close binaries as the source of abundance anomalies in Globular Clusters.,2009) propose non-conservative evolution of massive close binaries as the source of abundance anomalies in Globular Clusters.
+ To illustrate their thesis. they calculate one massive binary evolutionary model. a 20 M. (primary) + 15 M. (secondary) system with an initial period of 12 days.," To illustrate their thesis, they calculate one massive binary evolutionary model, a 20 $M_{\odot}$ (primary) + 15 $M_{\odot}$ (secondary) system with an initial period of 12 days."
+ During the Roche lobe overflow of the primary (the mass loser) the secondary (the mass gainer) spins up and reaches the critical rotation velocity., During the Roche lobe overflow of the primary (the mass loser) the secondary (the mass gainer) spins up and reaches the critical rotation velocity.
+ At that moment the authors assume that the mass that is lost by the loser leaves the binary taking the specific orbital angular momentum of the gainer with it., At that moment the authors assume that the mass that is lost by the loser leaves the binary taking the specific orbital angular momentum of the gainer with it.
+ However. the critical remarks made above when the WFRMS scenario Was discussed also apply here. re. rotation may not be responsible for large stellar mass loss and since one or both components of the massive binary will explode as a SN one may wonder if the lost matter will not be ejected from the Globular Cluster. swept up by the SN shell.," However, the critical remarks made above when the WFRMS scenario was discussed also apply here, i.e. rotation may not be responsible for large stellar mass loss and since one or both components of the massive binary will explode as a SN one may wonder if the lost matter will not be ejected from the Globular Cluster, swept up by the SN shell."
+ Here. we propose an alternative binary scenario.," Here, we propose an alternative binary scenario."
+ Most of the intermediate mass close binaries (MCB) do not experience a SN explosion whereas many lose a large amount of mass at small velocity due to Roche lobe overflow and/or the common envelope process., Most of the intermediate mass close binaries (IMCB) do not experience a SN explosion whereas many lose a large amount of mass at small velocity due to Roche lobe overflow and/or the common envelope process.
+ Part of this lost mass resembles pristine gas. part of it is helium enriched but not affected by the hot CNO burning processes.," Part of this lost mass resembles pristine gas, part of it is helium enriched but not affected by the hot CNO burning processes."
+ We will demonstrate in the present paper that the mass lost by a significant population of intermediate mass close binaries in combination. with the single star AGB pollution scenario may explain most of the chemical properties of a second generation of stars in Globular Clusters., We will demonstrate in the present paper that the mass lost by a significant population of intermediate mass close binaries in combination with the single star AGB pollution scenario may explain most of the chemical properties of a second generation of stars in Globular Clusters.
+ Calculating the evolution of a population of intermediate mass close binaries relies on the evolution of these systems., Calculating the evolution of a population of intermediate mass close binaries relies on the evolution of these systems.
+ The latter is mainly determined by three initial. parameters: the mass of the primary. the mass ratio and the binary period.," The latter is mainly determined by three initial parameters: the mass of the primary, the mass ratio and the binary period."
+ To explain some specific binaries. also the initial eccentricity may be important but for overall population synthesis related to the topic of the present paper (section 3) it is a second order parameter (see also Hurley et al..," To explain some specific binaries, also the initial eccentricity may be important but for overall population synthesis related to the topic of the present paper (section 3) it is a second order parameter (see also Hurley et al.,"
+ 2002)., 2002).
+ The evolution of IMCBs has been extensively studied since the early 60s up to the eighties by many authors (for reviews. see Paczynsski 1971: Plavee 1973: Van den Heuvel 1976: Thomas 1977) mainly to understand Algols. cataclysmie variables. double WD binaries and the link with Type Ia supernova. symbiotic stars. low mass X-ray binaries.," The evolution of IMCBs has been extensively studied since the early 60s up to the eighties by many authors (for reviews, see Paczyńsski 1971; Plavec 1973; Van den Heuvel 1976; Thomas 1977) mainly to understand Algols, cataclysmic variables, double WD binaries and the link with Type Ia supernova, symbiotic stars, low mass X-ray binaries."
+ De Greve and Vanbeveren (1980) compared the properties of a set of 151 theoretical binary mass exchange computations with the observed properties of Algol binaries in order to restrict the parameter space of the physics of the Roche lobe overflow., De Greve and Vanbeveren (1980) compared the properties of a set of 151 theoretical binary mass exchange computations with the observed properties of Algol binaries in order to restrict the parameter space of the physics of the Roche lobe overflow.
+ As computer power increased. much larger sets of evolutionary computations were published. e.g. Iben and Tutukov 1985. 1987; De Loore and Vanbeveren 1995; Vanbeveren et al.," As computer power increased, much larger sets of evolutionary computations were published, e.g. Iben and Tutukov 1985, 1987; De Loore and Vanbeveren 1995; Vanbeveren et al."
+ Γο0δο; Nelson and Eggleton 2001)., 1998c; Nelson and Eggleton 2001).
+ These sets resulted ina general evolutionary scenario (see for example Hurley et al..," These sets resulted in a general evolutionary scenario (see for example Hurley et al.,"
+ 2002 for a general overview). a scenario that allows us to study population synthesis of intermediate mass close binaries.," 2002 for a general overview), a scenario that allows us to study population synthesis of intermediate mass close binaries."
+ We discuss briefly this scenario., We discuss briefly this scenario.
+ The evolution of a star in a binary differs from. the evolution of that star when it is single mainly due to the Roche lobe overflow (RLOF) process., The evolution of a star in a binary differs from the evolution of that star when it is single mainly due to the Roche lobe overflow (RLOF) process.
+ Following the classification of Kippenhahn and Weigert (1967) and Lauterborn (1970) we distinguish three main phases of RLOF. which correspond to the three major expansion phases during stellar evolution: case A where the RLOF takes place during the core hydrogen," Following the classification of Kippenhahn and Weigert (1967) and Lauterborn (1970) we distinguish three main phases of RLOF, which correspond to the three major expansion phases during stellar evolution: case A where the RLOF takes place during the core hydrogen"
+such a dusty object. it would likely also be strongly starforming.,"such a dusty object, it would likely also be strongly starforming."
+ We would then expect to see Ha emission in our spectrum over most of the likely redshift range. but the signal-to-noise in our spectrum is about a factor of two worse than our other spectra.," We would then expect to see $\alpha$ emission in our spectrum over most of the likely redshift range, but the signal-to-noise in our spectrum is about a factor of two worse than our other spectra."
+ Considering also our poor constraint on the bulge-to-dise fit for HR14. we cannot conclusively decide if this is a starforming object or a more quiescent galaxy at a high redshift (we note that Pozzetti Mannucci classify ΠΕΙ as an “elliptical”. although within the errors of the photometry it does not lie unambiguously in this region of their colour-colour diagram).," Considering also our poor constraint on the bulge-to-disc fit for HR14, we cannot conclusively decide if this is a starforming object or a more quiescent galaxy at a high redshift (we note that Pozzetti Mannucci classify HR14 as an “elliptical”, although within the errors of the photometry it does not lie unambiguously in this region of their colour–colour diagram)."
+ Our central observational conclusion is that in this small sample. one of our targets is unambigously starforming. despite ifs morphological appearance as a well-formed elliptical and the clear indication from its near-IR colours that the bulk of its stellar population is not dust-enshrouded.," Our central observational conclusion is that in this small sample, one of our targets is unambigously starforming, despite its morphological appearance as a well-formed elliptical and the clear indication from its near-IR colours that the bulk of its stellar population is not dust-enshrouded."
+ With our data we cannot determine whether THO9 is a undergoing a very short-lived episode of star formation. or whether is it undergoing a process wherein the old stellar population is being augmented by a longer episode of star formation. perhaps by a substantial increase in its cold gas reservoir via a merger.," With our data we cannot determine whether TH09 is a undergoing a very short-lived episode of star formation, or whether is it undergoing a process wherein the old stellar population is being augmented by a longer episode of star formation, perhaps by a substantial increase in its cold gas reservoir via a merger."
+ However. we stress that the presence of an object such as THO9 in such a small observation sample is indicative that such objects must represent a non-negligible fraction of the ERO population.," However, we stress that the presence of an object such as TH09 in such a small observation sample is indicative that such objects must represent a non-negligible fraction of the ERO population."
+ If this is indeed that case. it would provide support for a picture in which ellipticals frequently do not form in a single monolithie star-formation.," If this is indeed that case, it would provide support for a picture in which ellipticals frequently do not form in a single monolithic star-formation."
+ To determine precisely the rate of this continuing growth in the stellar mass of the elliptical galaxy population. it will be essential to determine spectroscopically which of the EROS are truly quiescent. which host an AGN. and which are undergoing star formation.," To determine precisely the rate of this continuing growth in the stellar mass of the elliptical galaxy population, it will be essential to determine spectroscopically which of the EROs are truly quiescent, which host an AGN, and which are undergoing star formation."
+ Such observations are ideally suited to Subaru Telescope's second generation instrument FMOS (Maihara et 22000: Kimura et 22003). a wide-field. tibre-fed. OH-suppression spectrograph which will be able to simultaneously take up to 400 spectra across a 30-areminute field over the crucial 0.9—1.8 rrange.," Such observations are ideally suited to Subaru Telescope's second generation instrument FMOS (Maihara et 2000; Kimura et 2003), a wide-field, fibre-fed, OH-suppression spectrograph which will be able to simultaneously take up to 400 spectra across a 30-arcminute field over the crucial 0.9–1.8 range."
+ We have presented new OHS/CISCO near-IR spectroscopy and an analysis of new and existing HST and ground-based imaging of a small sample of EROs., We have presented new OHS/CISCO near-IR spectroscopy and an analysis of new and existing HST and ground-based imaging of a small sample of EROs.
+ Our conclusions are as follows: We warmly thank the referee for helpful comments on our first draft. and Katherine Inskip and Steve Rawlings for useful discussions.," Our conclusions are as follows: We warmly thank the referee for helpful comments on our first draft, and Katherine Inskip and Steve Rawlings for useful discussions."
+ This work is based in part on data collected at Subaru Telescope. which is operated by the National Astronomical Observatory of Japan: and in part on observations made with the NASA/ESA Hubble Space Telescope. obtained from the data archive at the Space Telescope Science institute.," This work is based in part on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan; and in part on observations made with the NASA/ESA Hubble Space Telescope, obtained from the data archive at the Space Telescope Science institute."
+ STScI is operated by the Association of Universities for research in. Astronomy. Inc.. under NASA contract NAS 5-26555.," STScI is operated by the Association of Universities for research in Astronomy, Inc., under NASA contract NAS 5-26555."
+. GC. acknowledges support from PPARC grants KKZA/014 RG35306 and KKZA/0I RG32761., GC acknowledges support from PPARC grants KKZA/014 RG35306 and KKZA/01 RG32761.
+ CS acknowledges a PPARC Advanced Fellowship., CS acknowledges a PPARC Advanced Fellowship.
+ RCB acknowledges a PPARC Ph.D. research studentship., RCB acknowledges a PPARC Ph.D. research studentship.
+ It is a pleasure to express our gratitude and respect to the indigenous people of Hawaii. on whose sacred mountain Mauna Kea our infrared observations were made.," It is a pleasure to express our gratitude and respect to the indigenous people of Hawai`i, on whose sacred mountain Mauna Kea our infrared observations were made."
+ὃν making use of the quantities defined in Section 2 and of eq. (5)).,"By making use of the quantities defined in Section 2 and of eq. \ref{eq:I}) ),"
+ the clustering term in eq. (6)), the clustering term in eq. \ref{eq:ctheta}) )
+ takes the form: where is the ellective volume emissivity. (see Volfolatti et al., takes the form: where is the effective volume emissivity (see Toffolatti et al.
+ 1998)., 1998).
+ Following the prescriptions of Section 2. we have then evaluated C'(6) in eq.07)) separately for the three cases of low-. intermecdiatc- ancl high-mass objects. by plugging in eq. (8))," Following the prescriptions of Section 2, we have then evaluated $C(\theta)$ in \ref{eq:cth}) ) separately for the three cases of low-, intermediate- and high-mass objects, by plugging in eq. \ref{eq:jeff}) )"
+ the appropriate expressions for the luminosity. function., the appropriate expressions for the luminosity function.
+ Note that in this case. the quantity DPMai.z) should. indeed. be read as 5(Adin.Mais2). where Mia is the maximum halo mass corresponding to the maximum visible bulee mass (i.e. upper limit for the mass locked into stars whose values are given in Section 2).," Note that in this case, the quantity $b^2(M_{\rm min},z)$ should indeed be read as $b^2(M_{\rm min},M_{\rm max},z)$, where $M_{\rm max}$ is the maximum halo mass corresponding to the maximum visible bulge mass (i.e. upper limit for the mass locked into stars whose values are given in Section 2)."
+ This corresponds in eq.63)) to a replacement in the upper limit of the integrals of oc with Mya, This corresponds in \ref{eq:b}) ) to a replacement in the upper limit of the integrals of $\infty$ with $M_{\rm max}$.
+ Figure 4 shows predictions for the correlation signal due o intensity Hluctuations (in units ergs? em+ 2 7) in the case of SCUBA galaxies., Figure \ref{fig:lowb} shows predictions for the correlation signal due to intensity fluctuations (in units $^2$ $^{-4}$ $^{-2}$ $^{-2}$ ) in the case of SCUBA galaxies.
+ Dotted. dashed. ancl solid ines are respectively for spheroids of MZ=q.1077 and 107A4...," Dotted, dashed and solid lines are respectively for spheroids of $M_{sph}^{min}=10^9,\;10^{10}$ and $10^{11}\; M_{\odot}$."
+ In cach plot the two solid. curves correspond o different detection limits. the 100 απν case represented ον the higher curve and the LO mv case by the lower one.," In each plot the two solid curves correspond to different detection limits, the 100 mJy case represented by the higher curve and the 10 mJy case by the lower one."
+ ote that. due to the steepness of the bright end of source counts. the results are independent of Sy if Syz100 mv.," Note that, due to the steepness of the bright end of source counts, the results are independent of $S_d$ if $S_d\ge 100$ mJy."
+" tesults are shown for Mia;/AMa=10 (left-hand. panel) and Aya,Maus=100 (right-hand panel)."," Results are shown for $M_{\rm halo}/M_{\rm sph}=10$ (left-hand panel) and $M_{\rm
+halo}/M_{\rm sph}=100$ (right-hand panel)."
+ Note once again 1o steep drop of all the curves around 6217 caused by the vbsence of nearby sources of this kind and the very small 'ontribution eiven by the low-mass objects. mostly due to wir small bias function and their modest. contribution to 10 elfective volume emissivitv.," Note once again the steep drop of all the curves around $\theta\simeq 1^\circ$ caused by the absence of nearby sources of this kind and the very small contribution given by the low-mass objects, mostly due to their small bias function and their modest contribution to the effective volume emissivity."
+ In order to evaluate the total contribution of the justering to intensity [fluctuations one has to add. up all 1e values of C'(8) in eq. (7)), In order to evaluate the total contribution of the clustering to intensity fluctuations one has to add up all the values of $C(\theta)$ in eq. \ref{eq:cth}) )
+ obtained for the different mass intervals and also to take into account the cross-correlation erms between objects of dilferent masses. so that the final expression for C'(@) is given by where the indexes ἐ] stand for hieh. intermediate and low masses.," obtained for the different mass intervals and also to take into account the cross-correlation terms between objects of different masses, so that the final expression for $C(\theta)$ is given by where the indexes $i$ $j$ stand for high, intermediate and low masses."
+ The angular power spectrum of the intensity Duetuations eqs. (7)) (9))], The angular power spectrum of the intensity fluctuations [eqs. \ref{eq:cth}) \ref{eq:ctot}) )]
+" ean be obtained starting from the well known expression for the expansion of temperature I[uctuations in spherical harmonics Η the Luctuations are a stationary process it can be shown (Pechles 1993) that the previous expression does not depend on the index m (i.c. on the ó angle). so that we can write: (see ""Tollolatti et al."," can be obtained starting from the well known expression for the expansion of temperature fluctuations in spherical harmonics If the fluctuations are a stationary process it can be shown (Peebles 1993) that the previous expression does not depend on the index $m$ (i.e. on the $\phi$ angle), so that we can write: with (see Toffolatti et al."
+ 1998). where Z£(cos8) are the which relates intensity and temperature Ductuations 26 Ix. Mather et αἱ.," 1998), where $P_l({\rm cos}\theta)$ are the Legendre polynomials, and $\delta T/T$ is expressed as which relates intensity and temperature fluctuations (T=2.726 K, Mather et al."
+ 1994)., 1994).
+" Figure 5. shows the predicted: values for the quantity dl,=VIT|D€,/2z (in units of IX) at 353 CGllIz (850jum) - the central frequeney of one of the channels of the Ligh Frequeney Instrument (HET) of the ESA’s Planck mission - as a function of the multipole / up to /21000."," Figure \ref{fig:dT} shows the predicted values for the quantity $\delta T_l=\sqrt{l(l+1)C_l/2\pi}$ (in units of K) at 353 GHz $850\,\mu$ m) - the central frequency of one of the channels of the High Frequency Instrument (HFI) of the ESA's Planck mission - as a function of the multipole $l$ up to $l=1000$."
+ Results are plotted for two values of the source detection limit (5=100 and 10 mv: note that the source detection limit for this Planck channel is likelv to be 100m.Jv) and the usual two values of ΑμαMin.," Results are plotted for two values of the source detection limit $S_d=100$ and 10 mJy; note that the source detection limit for this Planck channel is likely to be $>100\,$ mJy) and the usual two values of $M_{\rm halo}/M_{\rm
+sph}$."
+ Also shown. lor comparison. is the power spectrum of |rimary (CAIB) anisotropies (solid line) predicted. by a standard Cold Dark Matter mocel for a ACDAL CX= 0.7. O4=0.3. fy= 0.7). computed with the CMDEAST code developed by Seljak Zaldarriaga (1996).," Also shown, for comparison, is the power spectrum of primary (CMB) anisotropies (solid line) predicted by a standard Cold Dark Matter model for a $\Lambda$ CDM $\Lambda=0.7$ , $\Omega_0=0.3$, $h_0=0.7$ ), computed with the CMBFAST code developed by Seljak Zaldarriaga (1996)."
+ At this frequeney our model. predicts. Ductuations of amplitude due to clustering comparable to. (and. possibly even larger than) those obtained for primary CMD anisotropies at /zz50., At this frequency our model predicts fluctuations of amplitude due to clustering comparable to (and possibly even larger than) those obtained for primary CMB anisotropies at $l\simgt 50$.
+ As illustrated by Fig. 4...," As illustrated by Fig. \ref{fig:lowb},"
+ most of the clustering signal comes from massive galaxies with fluxes S1 mJv. which lic at substantial redshilts (see Figs. 1))," most of the clustering signal comes from massive galaxies with fluxes $S\simgt 1$ mJy, which lie at substantial redshifts (see Fig. \ref{fig:N_z}) )"
+ and are therefore highly biased. tracers of. the uncderlving mass distribution., and are therefore highly biased tracers of the underlying mass distribution.
+ Also the stronely negative |x-correction increases their contribution to the cllective volume emissivity eq. (8))], Also the strongly negative K-correction increases their contribution to the effective volume emissivity [eq. \ref{eq:jeff}) )]
+ and. therefore to the We therefore expect. this elect το give ao substantial contribution to the power spectrum extracted e.g. from the BOOAERanG map (de DBernardis et al., and therefore to the We therefore expect this effect to give a substantial contribution to the power spectrum extracted e.g. from the BOOMERanG map (de Bernardis et al.
+ 2000) at 400 Cillz., 2000) at 400 GHz.
+ For /900 its contribution may be comparable to that due to inhomogeneous emission by theinterstellar dust., For $l \simgt 200$ its contribution may be comparable to that due to inhomogeneous emission by theinterstellar dust.
+ , 
+ 1. Introduction Blazars areamoung the most variableobjects known. Fluctuatious occur ou allobserved timescales.,attributed to synchrotron emission from high-energy particles accelerated in regions of shocked gas in the jet.
+ auc shiftsof several magnitudes iuoptical briglituess havebeen seenover the spau ofjustmouths (see.e.g.. Metcall and. Magliocchetti (2006))).The mechanistus large of variability behavior. et (2008): Raiteri carefully etal.," These high-energy particles can also reprocess synchrotron radiation from the low-energy peak, thermal photons from the accretion disk, or other external photons via Compton upscattering to form the high-energy peak."
+ (2008))). Characteristicrevealing fluctuationrange timescales measured for(e.g.specilieChatterjee ," Hadron synchrotron radiation or decay may also contribute to the high-energy peak (see, e.g., \cite{mucke01}) )."
+blazarsrange fromhours tovears (e.g.Heidt and Wae- uer(1996): Ciprini al.(2003): Ixartaltepe aud Balouek(2007 ))). Traditionally. blazar. variability has beenstudied through intensive monitoring ofindivid," Blazars are often divided into subcategories of BL Lac objects and flat spectrum radio quasars (FSRQs), based in part on the absence or presence, respectively, of broad optical emission lines superimposed on the characteristic jet continuum emission."
+ual jects. Withnew wide-field surveysit ispossible totakeacomplementary approach by studyiug [ew, These broad lines are accretion disk emission reprocessed by energetic gas that is close to the system's central black hole.
+ measurementsof many blazars. ILtle ΡΟ class iscompletely homogeneous. the eusemble opertieswill be representative ofeach object. Iftheclass isheterogeneous.then the 'operties sta," These emission lines are visible in FSRQs but are most prominent in type I quasars, which are AGN aligned somewhat off-axis such that the observer views flux predominantly from the accretion disk."
+tisticallythis (e.g. describiugVar sample.Berket al. (2001):: Bauervariability etal. (2009)))type quasarsancl compared with detailed individual objects (e.g.Collier Petersou.," Reviews of the typical constituent parts of an AGN and the orientation-based theory of AGN unification linking quasars and blazars can be found in, for example, \cite{urry} and \cite{peterson}."
+ (2001)))in order to uuderstaux jetterthe of variabilityaud the of physical atworkinthese systems.," FSRQs and BL Lacs are believed to differ not just in their emission lines, but also in the structure and power of their jets."
+ Observationsof individualquasar ΑΝ suggest that rangeblazar variability processesexhibits much moreciversity of," Radio galaxies, or AGN with misaligned jets, are divided into two basic categories: Fanaroff-Riley types I and II (FR I and FR II) \cite{fr74}) )."
+ amplitude aud timescale than quasar variability. Nolarge ensemble studyof blazar variability," FR I jets generally have lower energies and are spatially limited, with the power concentrated close to the center of the galaxy."
+ has vet. been published to study this effect. however. due tothe lackof repeated measurements large enough," FR II jets are more energetic and extend much further out, with the most luminous areas in lobes far from the center of the galaxy."
+ areas ofsky. Wide-field surveys such asPalomar-QUEST areable to observe," FSRQs and BL Lacs are believed to belong to the same population as FR IIs and FR Is, respectively."
+ large uumbers of blazars and.begiu to s," FSRQs are more luminous than BL Lacs, as FR IIs are compared to FR Is."
+tudy thecollective variability properties of theclass. and ) examine tlie extentof heterogeneity within the sample. The characterization ofblazar variability as seenby several measurements over awide-field su," Furthermore, the relative number counts of FSRQs with respect to FR IIs, and BL Lacs with respect to FR Is, are consistent with the blazars being those radio galaxies that are by chance aligned to within $\sim 15^{\circ}$ of the line of sight \cite{urry}) )."
+rveyis useful ]xuowledge discovery of physicsblazarsis part[or better relatively uud, This blazar/radio galaxy unification model implies that FSRQs have higher energy jets than BL Lacs.
+erstaudiugof the objects’ behavior. Furthermore. all blazars currentlyuuportantstudied have been selected bytheirflux," Because of the morphological, energetic, and spectral differences between FSRQs and BL Lacs, we study the two types separately."
+ ratios intle racio and/or X-raybands.nearly possibly contributingto significant selection ellects intheensemble., Distinctions between the optical variability of the groups can improve our understanding of the fundamental differences between these subclasses.
+ typically exhibitingstroug radio aud X-ray. emission. blazarsare, The dramatic variability of blazar flux has several possible causes.
+ also someof the mostoptically variable objects in the sky. By studying the appearanceof," For example, the intermittent nature of shocks in the jet produces highly variable emission."
+ blazars asa class in sparsely sampled variability surveyswecau perhaps facilitatetlie discoveryof π, Much work has been done to model shock dissipation mechanisms in order to understand blazars' variability and energy distributions.
+αν morein currentsurveyslike Palomar-QUEST. aswell as imminent surveys suchas Pan-STARRS. Blazarsare currently uncerstoocto be activegalacticnuclei (," However, current models are not well constrained; even a single model framework can yield many possible variability characteristics (see, e.g., \cite{li00}) )."
+A&CGN) with jets.where thejet axis is aligned. aloug the observer'slineof sight.Their f[Iuxis therefore dominated by beamed jet euission the electromagnetic spectrum. The broadbaud," It is also uncertain how the sources of different wavelength emission are physically related to each other, for example if the radio, optical, and/or X-ray flux are generated in the same jet location."
+spectral clistributious(SEDs) ol blazars exhibit two broad peaks. at lower energies in frequencyenerey from radioto X-rays) audat higher energies (from X-raysto TeV (rangingrays).The low-enereypeak iscommonl," Multiwavelength monitoring of several blazars has revealed some cases in which radio, optical, and X-ray flares appear to have a common source; other cases show one band's emission to be unaffected by another's fluctuations"
+spectral clistributious(SEDs) ol blazars exhibit two broad peaks. at lower energies in frequencyenerey from radioto X-rays) audat higher energies (from X-raysto TeV (rangingrays).The low-enereypeak iscommonly," Multiwavelength monitoring of several blazars has revealed some cases in which radio, optical, and X-ray flares appear to have a common source; other cases show one band's emission to be unaffected by another's fluctuations"
+computationally expeusive aud an alternative methocl. used w Liou&Zooς(1999). and Martin&Malhotra(2¢002 ).. is to siniulate a simmaller utmber o‘particles aud 'egularly record their positions. which cai then be sumined to produce «istributio naps.,"computationally expensive and an alternative method, used by \citet{lz99} and \citet{mm02}, is to simulate a smaller number of particles and regularly record their positions, which can then be summed to produce distribution maps."
+ À similatiou is terminated Hler a fixed time «o» when all particles have beer destroyec| through οἱer ejection via close eucounters with platets or accretion onto the central :star., A simulation is terminated after a fixed time or when all particles have been destroyed through either ejection via close encounters with planets or accretion onto the central star.
+ In this way. loig Tis'ecd particles (trapped -—1 resonances) contribute more to the Lina dist clist‘ibution.," In this way, long lived particles (trapped in resonances) contribute more to the final dust distribution."
+ If. we assurje tUe dust distribution is ergocdie aud that the total dust mass is coistaui (i.e. je dust. productior rate equals the loss rate). us approach shoulc give an accurate re)‘exeuation of the actual stea«ventivue’ disk structure.," If we assume the dust distribution is ergodic and that the total dust mass is constant (i.e. the dust production rate equals the loss rate), this approach should give an accurate representation of the actual `steady-state' disk structure."
+ Iu previous studies (e.g.Moro-Mari&Malhora2002 ile pa‘ticle locatious have been ‘ausformecd iuto a 'eference frame whic1 roales wlh the oLerOst jxanet. but this approach is uot appropriate [or planets with ecceUric orbits.," In previous studies \citep[e.g.][]{mm02} the particle locations have been transformed into a reference frame which rotates with the outermost planet, but this approach is not appropriate for planets with eccentric orbits."
+ For this reasol.. we lustead eusure that the j»article recording takes place witl the plale in specilic orbial pisses.," For this reason, we instead ensure that the particle recording takes place with the planet in specific orbital phases."
+" Dno this way a number of ""uapshots of the systeln are taken wil1 he planet iu varyviug locations. allowing the effect. of dlanetary plasc o be sldied (Wilierοἱal.2002:Ixuclner&Holman2003)."," In this way a number of `snapshots' of the system are taken with the planet in varying locations, allowing the effect of planetary phase to be studied \citep{whk02,kh03}."
+. For example. many sinulations utiised a pkwet with orbita period D2OO veaIs ald a particle recording time of 1000 years. resulting in paricle distribuious for three differe1 planetaΝ plases.," For example, many simulations utilised a planet with orbital period 300 years and a particle recording time of 1000 years, resulting in particle distributions for three different planetary phases."
+ Lt should be note hat there are soie LimialLous t¢» this »ocedure of simulatiug a debris disk using a small number «) “test particles., It should be noted that there are some limitations to this procedure of simulating a debris disk using a small number of test particles.
+ As ciscised by Moro-Marin&Malhotra(2002).. becaIse both the probability of being trapped in a particuar resolance ald the time spent ina 'esolianice depeud sensitively ou 1itial coucitious. using a ini(ος uumber o ‘particles may over-exaggerate the lunportauce of a yarticar resouauce.," As discussed by \citet{mm02}, because both the probability of being trapped in a particular resonance and the time spent in a resonance depend sensitively on initial conditions, using a limited number of particles may over-exaggerate the importance of a particular resonance."
+ We have ried to overcome tus problem by usiug a1 order of maguituce more particles than previous simula101s (e.g.QuilleL&Thorudike2002:Mo‘o-Nlartin&Malhotra2002) in igh resolution simulatious (see following sections).," We have tried to overcome this problem by using an order of magnitude more particles than previous simulations \citep[e.g.][]{qt02,mm02} in high resolution simulations (see following sections)."
+ An aclclitlona problem. with this procedure is tiat all of the test particles are released fi'om thelr parent body at the sale ime. with a specific panetary phase.," An additional problem with this procedure is that all of the test particles are released from their parent body at the same time, with a specific planetary phase."
+ In reality. dust particles would be released. contiulously. at ando planetary phases.," In reality, dust particles would be released continuously, at random planetary phases."
+ Tus effect. however. is uot likely to |ave a significant inipact. since the est particles are typically reeased a long way from the planet αμα are distributed andomly around he star.," This effect, however, is not likely to have a significant impact, since the test particles are typically released a long way from the planet and are distributed randomly around the star."
+ Setting a fixed mΧΙΙur sinulatiou time may also be problematic., Setting a fixed maximum simulation time may also be problematic.
+ For tie stall particle uunber approach to be acerate. the integration should be continued until all test. particles have een destroyed.," For the small particle number approach to be accurate, the integration should be continued until all test particles have been destroyed."
+ The terminaion of a simulation at a predetermined time means tlat some particles uay still be active at the ed of the simulation., The termination of a simulation at a predetermined time means that some particles may still be active at the end of the simulation.
+ One could argue. however. tha extremely particles are unlikely to exist due to destruction via particle collisious.," One could argue, however, that extremely long-lived particles are unlikely to exist due to destruction via particle collisions."
+ 5since RMVS3 is a collisionless code. the destriction of particles in this way is not accounted for.," Since RMVS3 is a collisionless code, the destruction of particles in this way is not accounted for."
+ We investigate hese effects in more detail in Seetiou [.1.., We investigate these effects in more detail in Section \ref{subsec:collisionalproc}. .
+" In order to obtain estimates for the dust temperatures in these systems, we also fit a modified blackbody spectrum to the 100um and 160um emission: f,«[1—exp(—t,)]B,, where B, is the Planck function, and r, is the frequency-dependent optical depth, v/vo?."," In order to obtain estimates for the dust temperatures in these systems, we also fit a modified blackbody spectrum to the $\mu$ m and $\mu$ m emission: $f_\nu \propto [ 1 -
+\exp(- \tau_\nu) ] B_\nu$, where $B_\nu$ is the Planck function, and $\tau_\nu$ is the frequency-dependent optical depth, ${\nu/\nu_0}^\beta$."
+" We assume fj=1.5 as the emissivity power law index, and as the frequency at which the emission becomes optically thick, but note that the integrated luminosity depends very weakly on the values assumed for 6€[0.7—2] and vo, while the best fit temperature varies at most by 20%."," We assume $\beta=1.5$ as the emissivity power law index, and as the frequency at which the emission becomes optically thick, but note that the integrated luminosity depends very weakly on the values assumed for $\beta \in
+[0.7-2]$ and $\nu_0$, while the best fit temperature varies at most by $20\%$."
+" The BCG is much brighter than the rest of the galaxies in the cluster, with fluxes of 255.6+3.9 mJy (160um) and 297.8+3.1 mJy (100um) and an estimated [πι consistent with previous Spitzer/MIPS observations of Abell 1835 (9).."," The BCG is much brighter than the rest of the galaxies in the cluster, with fluxes of $255.6\pm3.9$ mJy $\mu$ m) and $\pm$ 3.1 mJy$\mu$ m) and an estimated $L_{\rm TIR}$ consistent with previous Spitzer/MIPS observations of Abell 1835 \citep{Egami2006}."
+" We derive star formation rates from the combined H, (from our spectroscopic measurements) and Ly; emission using the following relation from ?:: SFR[Kroupa](Meyr!)=5.5x10?[L(Ha@)ps+0.002Ly](ergs s!)."," We derive star formation rates from the combined $_{\alpha}$ (from our spectroscopic measurements) and $_{\rm TIR}$ emission using the following relation from \citet{Kennicutt2009}: ${\rm SFR [Kroupa]}
+(M_\odot \,yr^{-1}) = 5.5 \times 10^{-42} [L(H\alpha)_{obs} +
+0.002\,L_{\rm TIR}] \ ({\rm ergs\,s^{-1}})$ ."
+ This yields SFRs in the range: Moyr! for the general cluster population.," This yields SFRs in the range: $M_\odot \, yr^{-1}$ for the general cluster population."
+ The BCG’s high luminosity and special location at the center of a massive cool core separate it from the rest of the cluster population and we therefore exclude it from the subsequent analysis., The BCG's high luminosity and special location at the center of a massive cool core separate it from the rest of the cluster population and we therefore exclude it from the subsequent analysis.
+" We find that, of all galaxies with far-infrared emission detected byHerschel,, only 2 have X-ray counterparts down to Ly~4x10+ ergs sl, which is below the classical 10? ergs s! threshold for AGN definition."," We find that, of all galaxies with far-infrared emission detected by, only 2 have X-ray counterparts down to $_X \approx 4\times10^{41}$ ergs $^{-1}$, which is below the classical $10^{42}$ ergs $^{-1}$ threshold for AGN definition."
+" Only four galaxies (including the couple of X-ray point sources) in the sample are spectroscopically classified AGN according to their position in the classic BPT diagram (?),, defined by the emission line flux ratios of [NII] 6587/H,, versus [OIII] 5007/Hg."," Only four galaxies (including the couple of X-ray point sources) in the sample are spectroscopically classified AGN according to their position in the classic BPT diagram \citep{Baldwin1981}, defined by the emission line flux ratios of [NII] $_\alpha$ versus [OIII] $_\beta$."
+" We thus conclude that the bulk of the infrared luminous population is being heated by star formation and not an active nucleus, and we exclude the 4 likely AGN from our analysis."," We thus conclude that the bulk of the infrared luminous population is being heated by star formation and not an active nucleus, and we exclude the 4 likely AGN from our analysis."
+ It is apparent from Fig., It is apparent from Fig.
+" 1 that, apart from the BCG, there are few luminous infrared galaxies in the central regions of the cluster."," 1 that, apart from the BCG, there are few luminous infrared galaxies in the central regions of the cluster."
+" In fact, the radial surface density profiles in Fig."," In fact, the radial surface density profiles in Fig."
+" 3 show that most of the infrared luminosity is concentrated at large (r 2.5Mpc) radial distances, which is not the case for the K-band luminosity density which, as an approximate tracer of stellar mass, declines smoothly with radius."," 3 show that most of the infrared luminosity is concentrated at large $r \sim 2.5
+$ Mpc) radial distances, which is not the case for the K-band luminosity density which, as an approximate tracer of stellar mass, declines smoothly with radius."
+" Upon closer inspection, however, it is clear that this overdensity at the virial radius does not occur isotropically around the cluster, but rather peaks along specific directions."," Upon closer inspection, however, it is clear that this overdensity at the virial radius does not occur isotropically around the cluster, but rather peaks along specific directions."
+ The velocity distribution of K-band selected cluster galaxies (black line in Fig., The velocity distribution of K-band selected cluster galaxies (black line in Fig.
+" 4) is distinctly non-Gaussian, and shows signs of significant substructure."," 4) is distinctly non-Gaussian, and shows signs of significant substructure."
+" The most luminous infrared sources, located toward the south of the cluster, are dynamically segregated from the rest of the cluster population, clustering around an elongated structure in the foreground (6v~2000 km/s) of the cluster, which is also traced out in the K-band luminosity density maps (see 1)."," The most luminous infrared sources, located toward the south of the cluster, are dynamically segregated from the rest of the cluster population, clustering around an elongated structure in the foreground $\delta v \sim 2000$ km/s) of the cluster, which is also traced out in the K-band luminosity density maps (see )."
+" A gradient is observed in the velocities projected along the axis of elongation, where the relative velocity with respect to the cluster center progressively decreases as the inner region of the cluster is approached."," A gradient is observed in the velocities projected along the axis of elongation, where the relative velocity with respect to the cluster center progressively decreases as the inner region of the cluster is approached."
+" Figure 5 shows Subaru i’-band cutouts of the most luminous infrared galaxies, in decreasing order of Ly."," Figure 5 shows Subaru $i'$ -band cutouts of the most luminous infrared galaxies, in decreasing order of $L_{\rm TIR}$."
+" Many galaxies appear to have undergone recent interactions - ~ half of them have close companions, and half present distorted morphologies indicative of recent merging."," Many galaxies appear to have undergone recent interactions - $\sim$ half of them have close companions, and half present distorted morphologies indicative of recent merging."
+" We can attempt to quantify the importance of mergers in our small sample by comparing the fraction, f., of sources with close"," We can attempt to quantify the importance of mergers in our small sample by comparing the fraction, $f_c$ , of sources with close"
+every usable night from UTC 2000 August 13 (ILJD 2451769.5) through UTC 2000 October 14 (UD 2451831.5). UTC 2001 May. 20 (HLJD. 2452049.5) through UTC 2001 November 3 (ILJD 2452216.5). and UTC 2002 June 9 (ILJD 2452434.5) through UTC 2002 September 25 (ILJD 2452542.5).,"every usable night from UTC 2000 August 13 (HJD 2451769.5) through UTC 2000 October 14 (HJD 2451831.5), UTC 2001 May 20 (HJD 2452049.5) through UTC 2001 November 3 (HJD 2452216.5), and UTC 2002 June 9 (HJD 2452434.5) through UTC 2002 September 25 (HJD 2452542.5)."
+ Each night. a set of seven slightly offset images were obtained and then combined and trimmed to form a final nightly image.," Each night, a set of seven slightly offset images were obtained and then combined and trimmed to form a final nightly image."
+ The //-band images consist of 30 5 exposures. aud it took about four minutes to obtain the group of seven images:the A-band images consist of 10 s exposures and took about. (vo minutes to obtain.," The $H$ -band images consist of 30 s exposures, and it took about four minutes to obtain the group of seven images;the $K$ -band images consist of 10 s exposures and took about two minutes to obtain."
+ The final 512x pixel images. corresponding to a field of view of 112x areseconds. are approximately centered on the Galactic center.," The final $512 \times 512$ pixel images, corresponding to a field of view of $112 \times 112$ arcseconds, are approximately centered on the Galactic center."
+ After image quality cuts were made. there are a lolal of 144 Z£-band and 137 A-band images.," After image quality cuts were made, there are a total of 144 $H$ -band and 137 $K$ -band images."
+ The seeing ranges [rom 0793 to 1793 half maximum: in general. (he. //-band images are of higher quality (Uvpical seeing ~ 173) than the A-band (with tvpical seeing ~1745).," The seeing ranges from $0\farcs 93$ to $1\farcs 93$ full-width half maximum; in general, the $H$ -band images are of higher quality (typical seeing $\sim 1\farcs 3$ ) than the $K$ -band (with typical seeing $\sim 1\farcs 45$ )."
+ Because the field is crowded. we reduced the data using the ISIS difference image analvsis package CAlard2000;Hartmanοἱal.2004).," Because the field is crowded, we reduced the data using the ISIS difference image analysis package \citep{alard00,hartman04}."
+.. This analysis revealed a sign error in (he original DePovetal.(2004) reduction: LIRSIGSW is clearly an eclipsing binary.," This analysis revealed a sign error in the original \citet{depoy04}
+ reduction; IRS16SW is clearly an eclipsing binary."
+ Because IRS165W is subject to significant blendingthere are roughly half à dozen sources in the Ottet catalog within about one arcsecond of IRSIGSWwe calibrated our light-curves with the Ottetal.(1999). data as presented by Martins, Because IRS16SW is subject to significant blending—there are roughly half a dozen sources in the \citet{ott99} catalog within about one arcsecond of IRS16SW—we calibrated our light-curves with the \citet{ott99} data as presented by \citet{martins06}.
+etal.(2006).. The {ραπ photometry gives a mean magnitude of 0.2 mag higher (han the Ottetal.(1999) mean. and ineludes (wo more seasons of data. Ottetal.," The \citet{martins06} $K$ -band photometry gives a mean magnitude of 0.2 mag higher than the \citet{ott99} mean, and includes two more seasons of data. \citeauthor{ott99}"
+ [found a variability amplitude of 0.55 mag: using DAOPhot photometry to seale the ISIS fluxes. we find an amplituce of ~0.35 mag in both J and Jv (seeHartmanetal.2001. D)..," found a variability amplitude of 0.55 mag; using DAOPhot photometry to scale the ISIS fluxes, we find an amplitude of $\sim 0.35$ mag in both $H$ and $K$ \citep[see][Appendix B]{hartman04}."
+ This substantial amplitucle difference is indicative of sienificant blending in our dala., This substantial amplitude difference is indicative of significant blending in our data.
+ We used (he period of 19.45 days reported by DePoyetal.(2004) ancl to scale our A-band light-curve to have the same mean magnitude ancl amplitude as the Martinsetal. ddata., We used the period of 19.45 days reported by \citet{depoy04} and \citet{martins06} to scale our $K$ -band light-curve to have the same mean magnitude and amplitude as the \citeauthor{martins06} data.
+ There were 110 nights for which data were obtained in both the //- and A-bancds. providing contemporaneous measurements of the //2A color.," There were 110 nights for which data were obtained in both the $H$ - and $K$ -bands, providing contemporaneous measurements of the $H-K$ color."
+ Using the DAOPhot photometry. we find a constant {1—A color with an rms of 0.05 mag: this color does not vary. with time. phase. I/-. or A-band magnitude. as shown in Figure 1..," Using the DAOPhot photometry, we find a constant $H-K$ color with an rms of 0.05 mag; this color does not vary with time, phase, $H$ -, or $K$ -band magnitude, as shown in Figure \ref{fig:color}. ."
+ Lacking properly calibrated -band data. we scaled the similarly blended. -band data to have the same amplitude and mean magnitude as the A-band data.," Lacking properly calibrated $H$ -band data, we scaled the similarly blended $H$ -band data to have the same amplitude and mean magnitude as the $K$ -band data."
+ These scaled light-curves form the basis of our analysis: Table 1. eives the final scaled time-series photometry., These scaled light-curves form the basis of our analysis; Table \ref{tbl:data} gives the final scaled time-series photometry.
+We plot the IFMR for our sample in Figure 1..,We plot the IFMR for our sample in Figure \ref{fig:ifmr}.
+" The observed, final masses and their calculated progenitor masses are represented by the gray points."," The observed, final masses and their calculated progenitor masses are represented by the gray points."
+ The weighted mean ppair for each cluster is in black., The weighted mean pair for each cluster is in black.
+" When compared with the literature from which the initial data came, our final masses are within the errors of earlier measurements."," When compared with the literature from which the initial data came, our final masses are within the errors of earlier measurements."
+" Typically, the final masses of a given cluster are consistent across different IFMR studies despite coming from various theory groups (e.g.Williamsetal.2009a)."," Typically, the final masses of a given cluster are consistent across different IFMR studies despite coming from various theory groups \citep[\eg][]{Williams09}."
+". On the other hand, the cluster parameters that dictate the calculated initial masses are less well constrained: cluster age. in particular, is notoriously difficult to determine."," On the other hand, the cluster parameters that dictate the calculated initial masses are less well constrained; cluster age, in particular, is notoriously difficult to determine."
+" The uncertainties in the inputs to MS isochrone fitting: namely composition, reddening, and distance, all contribute to uncertainty in the cluster age and hence, progenitor masses."," The uncertainties in the inputs to MS isochrone fitting: namely composition, reddening, and distance, all contribute to uncertainty in the cluster age and hence, progenitor masses."
+" Progenitor masses calculated in our analysis rather than in SO9 still show good agreement with those of other studies assuming the same cluster ages.e.g... in M37 our weighted mean M;23.2150.11M.. for Td72635634312:3Myrvr and andFFerrarioettal.(2005).. (2005). whowh use 7,,,2520z:80520480Mvr. find Mj=3.2740.12M..Mj=3."," Progenitor masses calculated in our analysis rather than in S09 still show good agreement with those of other studies assuming the same cluster ages, in M37 our weighted mean $=3.21\pm0.11 \msol$ for $=563\pm31\ Myr$ and \citet{Ferrario05}, , who use $=520\pm80\ Myr$, find $=3.27\pm0.12 \msol$."
+"27+0.12M Irrespective of the assumption of convective overshooting, the average fractional uncertainty in our sample for M;((M;)) 1s 3%(7%)."," Irrespective of the assumption of convective overshooting, the average fractional uncertainty in our sample for ) is $3\% (7\%)$."
+ The red line in Figure 1. connects the anticipated core mass at the first thermal pulse ΙΡ) for cach cluster's weighted mean initial mass., The red line in Figure \ref{fig:ifmr} connects the anticipated core mass at the first thermal pulse ) for each cluster's weighted mean initial mass.
+ We determine geiven παπά aaccording to Pietrinfernietal.(2004) (see Section ??))., We determine given and according to \citet{Pietrinferni04} (see Section \ref{sec:m1tp}) ).
+ The core mass at the first thermal pulse is a monotonically increasing function of initial mass for a given metallicity., The core mass at the first thermal pulse is a monotonically increasing function of initial mass for a given metallicity.
+ The clusters in our sample span 0.32 dex in[Fe/H]., The clusters in our sample span $0.32$ dex in.
+. The differing compositions of the two clusters with the lowest progenitor mass explains the small dip in the red line at M;-2M..., The differing compositions of the two clusters with the lowest progenitor mass explains the small dip in the red line at $\sim 2 M_{\odot}$.
+ The relationship between the core mass at the first thermal pulse and the progenitor mass is a robust theoretical prediction: comparison between the models used here and those of Girardietal.(2000). yield identical results to within 2%., The relationship between the core mass at the first thermal pulse and the progenitor mass is a robust theoretical prediction; comparison between the models used here and those of \citet{Girardi00} yield identical results to within $2\%$.
+" In Figure I., all the cluster mean final masses are larger than the predicted core mass at the first thermal pulse."," In Figure \ref{fig:ifmr}, all the cluster mean final masses are larger than the predicted core mass at the first thermal pulse."
+" Given nominal assumptions regarding stellar evolution, current cluster observations demand that the core grows substantially during the TP-AGB phase: thus, TP-AGB stars must contribute sienificantly to cluster luminosity."," Given nominal assumptions regarding stellar evolution, current cluster observations demand that the core grows substantially during the TP-AGB phase; thus, TP-AGB stars must contribute significantly to cluster luminosity."
+" While this idea has been gaining traction in the literature (c.g.Marastonetal.2006)., we show, using the IFMR and accepted stellar interior theory, thatTP-AGB stars must be considered a large luminosity source in population synthesis models."," While this idea has been gaining traction in the literature \citep[\eg][]{Maraston06}, we show, using the IFMR and accepted stellar interior theory, thatTP-AGB stars must be considered a large luminosity source in population synthesis models."
+gas and the dynamical implications of supernovae explosions in particular.,gas and the dynamical implications of supernovae explosions in particular.
+ In this instance. it could well be that. during most of the active phase of the SMBH. the dense gas is no longer present in sufficient quantity on the spatial scale required by our model.," In this instance, it could well be that, during most of the active phase of the SMBH, the dense gas is no longer present in sufficient quantity on the spatial scale required by our model."
+ Furthermore. it is doubtful that the structure of the molecular gas were dominated. in such a phase. even if present. by a rotationally supported configuration out to several 100 pc.," Furthermore, it is doubtful that the structure of the molecular gas were dominated, in such a phase, even if present, by a rotationally supported configuration out to several 100 pc."
+ Only direct observations 1n molecular lines. with the adequate angular resolution. will àswer these questions.," Only direct observations in molecular lines, with the adequate angular resolution, will answer these questions."
+ To reach simple and sharp conclusions. we therefore consider two instances.," To reach simple and sharp conclusions, we therefore consider two instances."
+ In Sects., In Sects.
+ 5 and 6. in order to emphasise only the consequences of what was described in Sect.," 5 and 6, in order to emphasise $\it only$ the consequences of what was described in Sect."
+ 2. we shall posit that. for all AGN no matter the morphological type of the host galaxies. the model applies equally well and a single representative distribution of€ will be used.," 2, we shall posit that, for all AGN no matter the morphological type of the host galaxies, the model applies equally well and a single representative distribution of $\Sigma$ will be used."
+ In Sect., In Sect.
+" 7. we posit that the above applies only to the late type galaxies. while in the active early type galaxies the ""covering factor"" is basically unknown."," 7, we posit that the above applies only to the late type galaxies, while in the active early type galaxies the “covering factor” is basically unknown."
+ The latter instance requires that F7;(Mj) be kept separate from Fr(Mgj)., The latter instance requires that $F_L(M_{BH})$ be kept separate from $F_E(M_{BH})$.
+" The fraction of absorbed AGN. ων in a given luminosity bin (Lp, < Lp € Los) can be computed using the equation where""hern οι−=TotaLE (tym peng) and; utr,=LOGSLis Gl€ D. f(x is the fractional distribution in the quantity 3. introduced to take properly into account the strong dependence of C(X) on low values of the argument."," The fraction of absorbed AGN, $f_{abs}$, in a given luminosity bin $_{b1}$ $<$ $_{b}$ $<$ $_{b2}$ ) can be computed using the equation where $\lambda_{1}=\frac{L_{b1}}{L_{Edd}(M_{BH,i})}$ $\lambda_{1} \geq \lambda_{peak}$ ) and $\lambda_{2}=\frac{L_{b2}}{L_{Edd}(M_{BH,i})}$ $\lambda_{2} \leq$ 1), $f(\Sigma_{j})$ is the fractional distribution in the quantity $\Sigma$, introduced to take properly into account the strong dependence of $C(\Sigma)$ on low values of the argument."
+" Inour application we used mass bins 0.3 dex wide. and the results are presented in bins of Ly, that are 0.5 dex wide."," Inour application we used mass bins 0.3 dex wide, and the results are presented in bins of $L_b$ that are 0.5 dex wide."
+A normalization of the function Φ. as explained in Sect.,"A normalization of the function $\Phi$, as explained in Sect."
+ 2. will depend on the assumption we make on the quantity q(Mpj) given by Eq.(6)).," 2, will depend on the assumption we make on the quantity $q(M_{BH})$ given by \ref{fraction}) )."
+ To predict fupLp) with Eq. (07).," To predict $f_{abs}(L_b)$ with Eq. \ref{formulafabs}) ),"
+ we only consider the active SMBH with ων)., we only consider the active SMBH with $\lambda > \lambda_{peak}(M_{BH})$.
+ When the relationshps (3) and (5) are combined. itturns out that. at one particular value of the redshift. 220.35. Upeae(OMp)xMan very closely.," When the relationshps (3) and (5) are combined, itturns out that, at one particular value of the redshift, $z$ =0.35, $\lambda_{peak}(M_{BH})\propto M_{BH}^{-1}$ very closely."
+ Therefore. all masses in the range 107— 10° M. to which at 220.35 both (3) and (5) apply. contribute above the same value of Ly. equal to 10?ergs!," Therefore, all masses in the range $^{7.5}$ – $^9$ $M_{\odot}$, to which at $z$ =0.35 both (3) and (5) apply, contribute above the same value of $L_b$, equal to $^{45} erg s^{-1}$."
+ In other words. the active BH in this range of mass with 4/εως contributes only below 10?eres7!.," In other words, the active $BH$ in this range of mass with $\lambda < \lambda_{peak}$ contributes only below $^{45} erg s^{-1}$."
+ Notably. this is the Lege of Mpy = 10 Mo.," Notably, this is the $L_{edd}$ of $M_{BH}$ = $^7$ $M_{\odot}$."
+" Thus. to obtain a “complete” estimate of f,,, 1n the interval of logl, = 45-45.5. only the contribution of the masses in the range 10-107? M. needs to be added. which we did by resorting to a small extrapolation of the curve labelled 107 Af; in Fig."," Thus, to obtain a “complete” estimate of $f_{abs}$ in the interval of $L_b$ = 45–45.5, only the contribution of the masses in the range $^7$ $^{7.5}$ $M_{\odot}$ needs to be added, which we did by resorting to a small extrapolation of the curve labelled $^7$ $M_{\odot}$ in Fig."
+ | beyond its formal limit at z20.3., 1 beyond its formal limit at $z$ =0.3.
+ This convenient situation will be exploited to make some clear point. and to predict {ο at à given z to be used in Sect.," This convenient situation will be exploited to make some clear point, and to predict $f_{abs}$ at a given $z$ to be used in Sect."
+ 6 for comparison with the predictions at higher and lower values of the redshift., 6 for comparison with the predictions at higher and lower values of the redshift.
+ As anticipated in Sects., As anticipated in Sects.
+" 3 and 4.we first adopt the ""shape"" of the local FCMj). Fig. 3.."," 3 and 4,we first adopt the “shape” of the local $F(M_{BH})$, Fig. \ref{BHMF_Marconi},"
+" as representative of the ensamble of active SMBH in both late and early type galaxies. and assume that the ""model"" applies equally well to the two types of hosts with the same distribution of X."," as representative of the ensamble of active SMBH in both late and early type galaxies, and assume that the “model” applies equally well to the two types of hosts with the same distribution of $\Sigma$."
+ Concerning the quantity q (Eq. (60).," Concerning the quantity $q$ (Eq. \ref{fraction}) )),"
+ we start with the assumption that it 1s equal to a constant. namely independent of My.," we start with the assumption that it is equal to a constant, namely independent of $M_{BH}$."
+ It is perhaps worth emphasising that the value of this constant is irrelevant for obtainmg the fraction fi». while its precise knowledge would be necessary if we wanted to predict the number. for instance per unit volume. of the two types of AGN.," It is perhaps worth emphasising that the value of this constant is irrelevant for obtaining the fraction $f_{abs}$ , while its precise knowledge would be necessary if we wanted to predict the number, for instance per unit volume, of the two types of AGN."
+ The result is given as a histogram in Fig. 6.., The result is given as a histogram in Fig. \ref{Abs_Lbol_sigma}.
+ The first aspect deserving a comment is the sensitivity of the result to the choice concerning Y., The first aspect deserving a comment is the sensitivity of the result to the choice concerning $\Sigma$.
+" The values of f,p.(2,) 1n the individual bins of the histogram have been computed using the ""observed"". local. distribution of the 44 values of X referred to in Sect."," The values of $f_{abs}(L_b)$ in the individual bins of the histogram have been computed using the “observed”, local, distribution of the 44 values of $\Sigma$ referred to in Sect."
+ 4.and the bars represent the confidence intervals calculated on the basis of the limited number of objects in that distribution.," 4,and the bars represent the confidence intervals calculated on the basis of the limited number of objects in that distribution."
+ For the sake of neatness. in the following similar histograms. the bars will no longer be reproduced.," For the sake of neatness, in the following similar histograms, the bars will no longer be reproduced."
+" For comparison. in the same figure fp, is also given for two choices of a “single” representative value of X. namely the median (150 M pe77) and the mean (270 M. pec) of the same distribution."," For comparison, in the same figure $f_{abs}$ is also given for two choices of a “single” representative value of $\Sigma$, namely the median (150 $M_{\odot}$ $^{-2}$ ) and the mean (270 $M_{\odot}$ $^{-2}$ ) of the same distribution."
+ In both cases. especially in the second. the prediction is significantly higher in all bins.," In both cases, especially in the second, the prediction is significantly higher in all bins."
+ This follows from the fact that (see Fig. 4)).," This follows from the fact that (see Fig. \ref{Fabs_sigma}) ),"
+ when a substantial fraction of objects span the interval from 0 to about 300 4: pc7. the covering factor is extremely sensitive to the way the values ofX are distributed.," when a substantial fraction of objects span the interval from 0 to about 300 $M_{\odot}$ $^{-2}$, the covering factor is extremely sensitive to the way the values of $\Sigma$ are distributed."
+ In this respect.the prediction of firs -xbbtained adopting X270 M. pc (as it can be judged from F9) ," In this respect,the prediction of $f_{abs}$ obtained adopting $\Sigma$ =270 $M_{\odot}$ $^{-2}$ (as it can be judged from Fig. \ref{Fabs_sigma}) )"
+is rather close to the ceiling attainable by pushing the distribution toward much higher values of X than adopted here: a point on which we return in Sect., is rather close to the ceiling attainable by pushing the distribution toward much higher values of $\Sigma$ than adopted here: a point on which we return in Sect.
+ 7., 7.
+" It is important to note that. without the ""physical"" limit of unity for οἱ. the result would be different."," It is important to note that, without the “physical” limit of unity for $\lambda$, the result would be different."
+ For instance. if no cut-off at t=1 m Eq. (1))," For instance, if no cut-off at $\lambda$ =1 in Eq. \ref{lorentz}) )"
+" is imposed. the computation of fup, results i$ approximately a constant value. irrespective of luminosity."," is imposed, the computation of $f_{abs}$ results is approximately a constant value, irrespective of luminosity."
+" This occurs because a single mass range. 10/- 10? M... turns out to dominate at all luminosities. and the value of fj», to be close to the one pertinent to these mass values."," This occurs because a single mass range, $^7$ $^{7.3}$ $M_{\odot}$, turns out to dominate at all luminosities, and the value of $f_{abs}$ to be close to the one pertinent to these mass values."
+ Moreover. in this instance. we have no information to adopt from NT2007 in order to estimate the contribution of the masses below 10 M. a very serious cause of uncertainties. due to the relatively large number of SMBH in the mass range 105-107 M...," Moreover, in this instance, we have no information to adopt from NT2007 in order to estimate the contribution of the masses below $^{7}$ $M_{\odot}$, a very serious cause of uncertainties, due to the relatively large number of SMBH in the mass range $^6$ $^{7}$ $M_{\odot}$ ."
+ Because the relationship between ως and Mgg remains qualitatively the same at the redshift values dealt with in the following section. the same type of result. with the same type of uncertainties. is also obtained at theother To illustrate the consequence of taking into account the downsizing effect mentioned in Sect.," Because the relationship between $\lambda_{peak}$ and $M_{BH}$ remains qualitatively the same at the redshift values dealt with in the following section, the same type of result, with the same type of uncertainties, is also obtained at theother To illustrate the consequence of taking into account the downsizing effect mentioned in Sect."
+ 3. we only need to change the leading assumption on q. namely adopt a proper dependence of g on Mj.," 3, we only need to change the leading assumption on $q$ , namely adopt a proper dependence of $q$ on $M_{BH}$ ."
+ The shape of the broad-line AGN mass function in the local universe. as obtained by Greene Ho (2007) above 10 M... suggests a dependence ofthe type qxGOd»pg/107) 79. ," The shape of the broad-line AGN mass function in the local universe, as obtained by Greene Ho (2007) above $^7$ $M_{\odot}$ suggests a dependence ofthe type $q \propto (M_{BH}/10^{7})^{-0.5}$ ."
+To show that the knowledge of the exact difference between F and F is not very important for this range of mass. we assumed an even steeper correction às anexercise.∣↽ namely⋅ DEMgx My.," To show that the knowledge of the exact difference between $F^*$ and $F$ is not very important for this range of mass, we assumed an even steeper correction as anexercise, namely $q \propto M_{BH}^{-1}$ ."
+ The result is shown in Fig. 7.., The result is shown in Fig. \ref{Abs_z035}. .
+ The, The
+explanation of the observed O-C diagram. namely a continually changing period or two constant periods.,"explanation of the observed O-C diagram, namely a continually changing period or two constant periods."
+ After the year 1998 Trigilioetal.(2008.2011) observed another increase in the period of radio pulses of AP=1.12 s with respect to the period determined by Pyperetal.(1998).," After the year 1998 \citet{trigi,trigi11} observed another increase in the period of radio pulses of $\Delta P=1.12$ s with respect to the period determined by \citet{pyper98}."
+. However. the 2010 measurements indicate some period decrease.," However, the 2010 measurements indicate some period decrease."
+ The rate of the deceleration can be evaluated using the spin-down time. r. defined as r=P/P. where P(r) is the instantaneous rotation period at the time ¢ and P is the mean rate of the rotational deceleration.," The rate of the deceleration can be evaluated using the spin-down time, $\tau$, defined as $\tau=P/\bar{\dot{P}}$, where $P(t)$ is the instantaneous rotation period at the time $t$ and $\bar{\dot{P}}$ is the mean rate of the rotational deceleration."
+ The paradox of iis according to Mikuláseketal.(2011) that its spin down time. τ-6xIO? years. is more than two orders of magnitude shorter than the estimated age of the star — 9x10’ yr Bagnulo20006 ).," The paradox of is according to \citet{unstead} that its spin down time, $ \tau \sim 6\times 10^5$ years, is more than two orders of magnitude shorter than the estimated age of the star – $9\times10^7$ yr \citep{kobacu}."
+". V90] Ori = HD 37776 is à young hot star IIV) residing in the emission nebula IC 432. with a global. extraordinarily strong (B,=20 kG). and unusually complex magnetic field (Thompson&Landstreet1985:Kochukhovetal.2011)."," V901 Ori = HD 37776 is a young hot star IV) residing in the emission nebula IC 432, with a global, extraordinarily strong $(B_s\approx20$ kG), and unusually complex magnetic field \citep{thomla,ko901}."
+. The observed moderate light variations are caused by the spots of overabundant silicon and helium (Krtiékaetal.2007)., The observed moderate light variations are caused by the spots of overabundant silicon and helium \citep{krt901}.
+. Three decades of precise photometric and spectroscopic monitoring have revealed a continuous rotational deceleration (Mikuldseketal.2008).. increasing the period of about 115387 by a remarkable 18 s!," Three decades of precise photometric and spectroscopic monitoring have revealed a continuous rotational deceleration \citep{mik901}, increasing the period of about $1\fd5387$ by a remarkable 18 s!"
+ Ruling out (a) the light-time effect in a binary star. (b) the precession of the star's rotational axis. and (ο) evolutionary effects as possible causes of the period change. we interpreted the deceleration in terms of the rotational braking of the outer stellar layers as caused by the angular momentum loss in the stellar magnetosphere.," Ruling out (a) the light-time effect in a binary star, (b) the precession of the star's rotational axis, and (c) evolutionary effects as possible causes of the period change, we interpreted the deceleration in terms of the rotational braking of the outer stellar layers as caused by the angular momentum loss in the stellar magnetosphere."
+ However. this cannot explain the discrepancy between the spin-down time. 7=2.5x10° yr. and the stars age of one million years or older (Mikuláseketal.2008.2010.2011:Kochukhov&Bagnulo 2006).," However, this cannot explain the discrepancy between the spin-down time, $\tau = 2.5 \times 10^5$ yr, and the star's age of one million years or older \citep{mik901,mik7355,unstead,kobacu}."
+. The interpretation of the rotation period evolution by simple angular motion loss was also questioned by the negative value of the second derivative of the period: P=-29(3)x107d'. which indicated that braking could change into acceleration (Mikuláseketal. 2008).," The interpretation of the rotation period evolution by simple angular motion loss was also questioned by the negative value of the second derivative of the period: $\ddot{P}=-29(13)\times
+10^{-13}\,\mathrm{d}^{-1}$, which indicated that braking could change into acceleration \citep{mik901}."
+. Nevertheless. the significance of this conclusion was fairly low.," Nevertheless, the significance of this conclusion was fairly low."
+ For the period analyses of wwe obtained some new photometric or spectroscopic data. specifically: 1) JK derived 210 individual spectrophotometric magnitudes in 21 bands (201 - 299 nm)by processing 10 IUE spectra. in. 1979 January - March (fordetailsseeKrticékaetal.20115: 22 GWH obtained 374 precise BY measurements at Fairborn Observatory Arizona. USA. in 2010 February - 2011 3) JL obtained 38 V measurements using his own microtelescope at the SAAO. South Africa and Brno. CR. in 2010 April - June: 4) JJ obtained 251 vb observations at Suhora Observatory. Poland. in 2011 March: and JJ+JL obtained 536 UBV measurements using the 0.5m reflector of the SAAO. South Africa. in 2011 May.," For the period analyses of we obtained some new photometric or spectroscopic data, specifically: 1) JK derived 210 individual spectrophotometric magnitudes in 21 bands (201 - 299 nm)by processing 10 IUE spectra, in 1979 January - March \citep[for details
+see][]{krtcu}; 2) GWH obtained 374 precise $\mathit{BV}$ measurements at Fairborn Observatory Arizona, USA, in 2010 February - 2011 3) JL obtained 38 $V$ measurements using his own microtelescope at the SAAO, South Africa and Brno, CR, in 2010 April - June 4) JJ obtained 251 $\mathit{vb}$ observations at Suhora Observatory, Poland, in 2011 March; and JJ+JL obtained 536 $\textit{UBV}$ measurements using the 0.5m reflector of the SAAO, South Africa, in 2011 May."
+ We collected a total of 8965 individual measurements of oobtained between 1949 and 2011 (62 years or 43662 revolutions of the star) including. 8270. photometric and spectrophotometric measurements in photometric bands from 200 to 753 nm às well as spectroscopic. spectropolarmmetric and radiometric observations.," We collected a total of 8965 individual measurements of obtained between 1949 and 2011 (62 years or 43662 revolutions of the star) including 8270 photometric and spectrophotometric measurements in photometric bands from 200 to 753 nm as well as spectroscopic, spectropolarimetric and radiometric observations."
+ The diverse observational data originated from 39 sources (see the list in Appendix Β)., The diverse observational data originated from 39 sources (see the list in Appendix B).
+ The ddata listed in. MikulaSeketal.(2008) were enhanced with 98 precise BV measurements obtained by GWH at Fairborn Observatory Arizona. USA from September 2009 to February 2011.," The data listed in \citet{mik901} were enhanced with 98 precise $BV$ measurements obtained by GWH at Fairborn Observatory Arizona, USA from September 2009 to February 2011."
+ We used a total of 2611 individual measurements obtained from 1976 to 2011 from 14 sources., We used a total of 2611 individual measurements obtained from 1976 to 2011 from 14 sources.
+ In addition to 2409 photometric measurements. the data include 202 measurements of equivalent widths of eleven selectedi lines.," In addition to 2409 photometric measurements, the data include 202 measurements of equivalent widths of eleven selected lines."
+" The method of data processing. outlined in the Appendix A. assumes that phase curves of all monitored quantities are constant. while the ""instant period? of changes P(t) is variable in (me ¢."," The method of data processing, outlined in the Appendix A, assumes that phase curves of all monitored quantities are constant, while the `instant period' of changes $P(t)$ is variable in time $t$."
+ It is advantageous to introduce a “phase function? (r.b). which is a continuous monotonic function of time 1.," It is advantageous to introduce a `phase function' $\vartheta(t,\mathbf{b})$, which is a continuous monotonic function of time $t$."
+ The fractional part of it corresponds to the common phase ¢. the integer part is the so-called epoch (E) and b is a set of gp free parameters describing a model of the period development.," The fractional part of it corresponds to the common phase $\varphi$, the integer part is the so-called epoch $(E)$ and $\mathbf{b}$ is a set of $g_{\mathrm{b}}$ free parameters describing a model of the period development."
+ Quantities P(r) and (r.b) are bound by the simple equality P(t)= 1/8. where the dot refers to the first time derivative," Quantities $P(t)$ and $\vartheta(t,\mathbf{b})$ are bound by the simple equality $P(t)=1/\dot{\vartheta}$ , where the dot refers to the first time derivative"
+We use near-infrared spectroscopy obtained with the imaging spectrograph SINFONI and the echelle spectrograph CRIRES (both on the ESO Very Large Telescope) to quantify the dissipation of kinetic energy in the molecular gas on scales of few 10s to 100s of pe.,We use near-infrared spectroscopy obtained with the imaging spectrograph SINFONI and the echelle spectrograph CRIRES (both on the ESO Very Large Telescope) to quantify the dissipation of kinetic energy in the molecular gas on scales of few 10s to 100s of pc.
+" To this end we observed the region around ""knot 5 of ?.. the brightest H» peak m the overlap region."," To this end we observed the region around “knot 5” of \citet{brandl09}, the brightest $_2$ peak in the overlap region."
+ This knot. which is near the massive and IR-bright cluster No.28 of ?.. coincides with the super-giant molecular cloud SGMC 2 observed by ? in CO.," This knot, which is near the massive and IR-bright cluster No.28 of \citet{whitmore10}, coincides with the super-giant molecular cloud SGMC 2 observed by \citet{wilson00} in CO."
+ SGMC 2 is one of the most massive molecular clouds in the overlap region. with a mass of 4x105M...," SGMC 2 is one of the most massive molecular clouds in the overlap region, with a mass of $4\tten{8}\, $."
+ Our pointing also coincides with knot K2a in the Herschel dust imaging of ?.., Our pointing also coincides with knot K2a in the Herschel dust imaging of \citet{klaas10}.
+ Thus. overall. by selecting an area with particularly bright line emission of warm molecular hydrogen. we observed one of the main sites of molecular gas and star formation in the Antennae overlap region.," Thus, overall, by selecting an area with particularly bright line emission of warm molecular hydrogen, we observed one of the main sites of molecular gas and star formation in the Antennae overlap region."
+ The paper is organized as follows., The paper is organized as follows.
+ The observations anc data reduction are described in Sect., The observations and data reduction are described in Sect.
+ 2., 2.
+ In Sect., In Sect.
+ 3 we focus oi the structure and kinematics of the warm geas., 3 we focus on the structure and kinematics of the warm gas.
+ We identify two components of eemission: diffuse emission extended over the 600 pe wide field of view of SINFONL and a compact source.," We identify two components of emission: diffuse emission extended over the 600 pc wide field of view of SINFONI, and a compact source."
+ We show that. for the compact and diffuse extended emission. the eemission is powered by shocks rather than UV heating. and estimate the bolometric lluminosity of each component combining Spitzer. anc SINFONI observations (Sect.," We show that, for the compact and diffuse extended emission, the emission is powered by shocks rather than UV heating, and estimate the bolometric luminosity of each component combining Spitzer and SINFONI observations (Sect."
+ 4)., 4).
+ In Sect., In Sect.
+ 5 we propose ai interpretation of the observations which relates the H» emissior to the formation of bound clouds through gas aceretion., 5 we propose an interpretation of the observations which relates the $_2$ emission to the formation of bound clouds through gas accretion.
+ [i Sect., In Sect.
+ 6 we discuss the observations and our interpretation within à broader astrophysical context. highlighting the role of turbulence for the formation of super star clusters and the regulation of the star formation efficiency in galaxy mergers.," 6 we discuss the observations and our interpretation within a broader astrophysical context, highlighting the role of turbulence for the formation of super star clusters and the regulation of the star formation efficiency in galaxy mergers."
+" Section 7 gathers the main conclusions,", Section 7 gathers the main conclusions.
+ Our analysis is based on observations carried out with the near-infrared integral-field spectrograph SINFONI (??) on the ESO Very Large Telescope.," Our analysis is based on observations carried out with the near-infrared integral-field spectrograph SINFONI \citep{eisenhauer03,bonnet04} on the ESO Very Large Telescope."
+ Data were obtained in service mode under Program ID 383.B-0789 (PI Nesvadba)., Data were obtained in service mode under Program ID 383.B-0789 (PI Nesvadba).
+" SINFONI is an image slicer with a field of view of 8x8"" aand a pixel scale of 250 mas in the seeing-limited mode.", SINFONI is an image slicer with a field of view of $\times$ and a pixel scale of 250 mas in the seeing-limited mode.
+ The spectral resolution is R=4000 att=2.2 pm.We observed three regions in the Antennae galaxy merger in the K-band. one near each nucleus and one region in the overlap region.," The spectral resolution is ${\rm R=4000}$ at $\lambda=2.2~\mu$ m.We observed three regions in the Antennae galaxy merger in the $K$ -band, one near each nucleus and one region in the overlap region."
+ These pointings were selected based on their bright pure-rotational H» line emission detected with Spitzer/IRS., These pointings were selected based on their bright pure-rotational $_2$ line emission detected with Spitzer/IRS.
+ This paper presents an analysis of molecular emission of the pointing in the overlap region., This paper presents an analysis of molecular emission of the pointing in the overlap region.
+ All data were taken in June and July 2009 under good and stable atmospheric conditions., All data were taken in June and July 2009 under good and stable atmospheric conditions.
+ We obtained a total of 1800 seconds of exposure time split into individual exposures of 600 seconds., We obtained a total of 1800 seconds of exposure time split into individual exposures of 600 seconds.
+ We adopted a dither pattern where one sky frame was taken in-between two object frames to allow for an accurate subtraction of the night sky., We adopted a dither pattern where one sky frame was taken in-between two object frames to allow for an accurate subtraction of the night sky.
+ Data reduction was done with the standard IRAF tools to reduce longslit spectroscopy (?).. which we extended with a set of SINFONI-specifie IDL routines.," Data reduction was done with the standard IRAF tools to reduce longslit spectroscopy \citep{tody93}, which we extended with a set of SINFONI-specific IDL routines."
+ For details see. e.g.. ??..," For details see, e.g., \citet{nesvadba07, nesvadba08}."
+ The telluric correction was done using bright (K~6-8 mag) stars observed at a similar air mass as our target., The telluric correction was done using bright $K\sim6-8$ mag) stars observed at a similar air mass as our target.
+ We used the light profiles of these stars to measure the PSF of our data. FWHM=077x06.," We used the light profiles of these stars to measure the PSF of our data, ${\rm FWHM=0\farcs7\times0\farcs6}$."
+ Comparison with high-resolution HST/NICMOS imaging of ?. shows that our PSF estimate Is robust., Comparison with high-resolution HST/NICMOS imaging of \citet{whitmore10} shows that our PSF estimate is robust.
+ To complement ourSINFONI observations we also obtained a high-resolution longslit spectrum with CRIRES on the VLT (program ID 386.B-0942: PI Herrera) and near-infrared imaging of the Antennae through the H» 1-0 SCI) and the ffilters with WIRCAM at the Canada-France Hawaii Telescope (CFHT program ID O9AF98; PI Nesvadba)., To complement ourSINFONI observations we also obtained a high-resolution longslit spectrum with CRIRES on the VLT (program ID 386.B-0942; PI Herrera) and near-infrared imaging of the Antennae through the $_2$ $-$ 0 S(1) and the filters with WIRCAM at the Canada-France Hawaii Telescope (CFHT program ID 09AF98; PI Nesvadba).
+ With CRIRES we obtained a total of 5 hrs of integration time of the H» 1-0 S(I) line at 2.12 gm in the rest-frame for one pointing centered on the compact molecular source (Sect. 3.3)), With CRIRES we obtained a total of 5 hrs of integration time of the $_2$ $-$ 0 S(1) line at 2.12 $\mu$ m in the rest-frame for one pointing centered on the compact molecular source (Sect. \ref{sec:cs}) )
+ and at position angle PA=67°., and at position angle ${\rm PA=67^{\circ}}$.
+ Atmospheric conditions were good and stable., Atmospheric conditions were good and stable.
+ We nodded our target along the slit to ensure a robust sky subtraction., We nodded our target along the slit to ensure a robust sky subtraction.
+ Data were dark-subtracted. flat-fielded. and sky subtracted.," Data were dark-subtracted, flat-fielded, and sky subtracted."
+ The wavelength calibration was done on telluric OH lines straddling the wavelength of the H» 1-0 SCI) line., The wavelength calibration was done on telluric OH lines straddling the wavelength of the $_2$ $-$ 0 S(1) line.
+ Using a wwide slit. we reached a spectral resolution of 6 km s!. a factor - 20 higher than the effective resolution of our SINFONI data.," Using a wide slit, we reached a spectral resolution of 6 km $^{-1}$, a factor $\sim$ 20 higher than the effective resolution of our SINFONI data."
+ This allowed us to analyze the line profile of the compact molecular source at a spectral resolution which ts comparable to that reached by ? for CO., This allowed us to analyze the line profile of the compact molecular source at a spectral resolution which is comparable to that reached by \citet{wilson00} for CO.
+ We determined the astrometry of the near-IR CFHT images through eross-correlation with the 2MASS K-band image of NGC4038/39., We determined the astrometry of the near-IR CFHT images through cross-correlation with the 2MASS $K$ -band image of NGC4038/39.
+ We give a short description of these data in Sect. ??.., We give a short description of these data in Sect. \ref{sec:ec}.
+ They have also been used to improve the astrometric accuracy of the SINFONI observations (Sect. ??))., They have also been used to improve the astrometric accuracy of the SINFONI observations (Sect. \ref{sec:ali}) ).
+ We complement our near-IR spectro-imaging data with CO(1-0) observations obtained with the Caltech Millimeter Array., We complement our near-IR spectro-imaging data with $-$ 0) observations obtained with the Caltech Millimeter Array.
+ C. Wilson. generously shared her zeroth and first-moment maps of NGC4038/4039 with us., C. Wilson generously shared her zeroth and first-moment maps of NGC4038/4039 with us.
+ For a full discussion of these observations. see ??..," For a full discussion of these observations, see \citet{wilson00, wilson03}."
+ Our SINFONI pointing in the overlap region coincides with their super-giant molecular cloud 2 (SGMC 2)., Our SINFONI pointing in the overlap region coincides with their super-giant molecular cloud 2 (SGMC 2).
+ Within the spatial resolution of the CO data. the peak of SGMC 2 coincides with the warm H» peak #55 in ?..," Within the spatial resolution of the CO data, the peak of SGMC 2 coincides with the warm $_2$ peak 5 in \citet{brandl09}."
+ In our analysis we adopt the mass. size and line width of SGMC 2 listed in Table | of ?..," In our analysis we adopt the mass, size and line width of SGMC 2 listed in Table 1 of \citet{wilson00}."
+ In addition. we use published results from Spitzer/IRS mid- spectroscopy by ?..," In addition, we use published results from Spitzer/IRS mid-infrared spectroscopy by \citet{brandl09}. ."
+ Our pointing corresponds to their, Our pointing corresponds to their
+to calibrate the reddeuing-free indices and correcting he transformation of theoretical isochrone libraries iuto he WEC3 photonetre svsteni.,to calibrate the reddening-free indices and correcting the transformation of theoretical isochrone libraries into the WFC3 photometric system.
+ With relatively well measured values of metallicity aud reddening (E(B-Vj=0.0+£0.02.Salarisetal.2007:ILuris 2010).. the distance vecoles the crucial iuput piuiuueter to determine the age of 17 Tuc.," With relatively well measured values of metallicity and reddening \citep[E(B~-~V)~$=0.04\pm0.02$, ][]{salaris07,harris10}, the distance becomes the crucial input parameter to determine the age of 47 Tuc."
+ A consideraje effort has gone iuto deterwining the distance to this cluster., A considerable effort has gone into determining the distance to this cluster.
+ We list recent work in Table 1.., We list recent work in Table \ref{tab:distance}.
+ There las been quite a variety of ioethodoogics used to determine the distance of this cluster. iucudiue main sequence fitting (IIesseretal.1987:Cratton1997:Reid1998:Car-2002:xCirattonetal.2003:Berebusch&Stetson 2009).. the tip of the red eiaut branch (Ferraroetal.Bonoctal. 2008).. horizoual brauch fitting (IIesseretraroetal.1999:Salariset:d. 2007).. RR Lyrae (Stormetal.1991:Do10 2008). €uster stellar duematics (MeLaughliuetal.2006).. eclipsing binaries (sahizivetal.2007:Thompson 2OL0).. and white dwart (WD) cooling sequence ftti1ο CZoccalietal.2001).," There has been quite a variety of methodologies used to determine the distance of this cluster, including main sequence fitting \citep{hesser87,gratton97,reid98,carretta00,grundahl02,percival02,gratton03,bergbusch09}, the tip of the red giant branch \citep{ferraro00,bono08}, , horizontal branch fitting \citep{hesser87,kaluzny98,salaris98,ferraro99,salaris07}, RR Lyrae \citep{storm94,bono08}, cluster stellar kinematics \citep{mclaughlin06}, eclipsing binaries \citep{kaluzny07,thompson10}, and white dwarf (WD) cooling sequence fitting \citep{zoccali01}."
+. This set of mieasuremioenuts spans unreddenued distance moduli between 13.02—13.57 for I: Τιc Which has important iuplicatious for the derived age., This set of measurements spans unreddened distance moduli between $13.02-13.57$ for 47 Tuc which has important implications for the derived age.
+ Exanuuinege only the nian sequence turnoff huulnosity for a cluster at the extremes m distance moduli. this vields a difference iu age of ~6Cvrtt.. altering any inferred history of the formation of the Πιν Wavy that one night derive from 17 Tuc.," Examining only the main sequence turnoff luminosity for a cluster at the extremes in distance moduli, this yields a difference in age of $\sim6$, altering any inferred history of the formation of the Milky Way that one might derive from 47 Tuc."
+ Our goal is to determine the distauce to 17 Tue ‘ollowing the example of Zoccalictal.(2001). by using he clusters WD population., Our goal is to determine the distance to 47 Tuc following the example of \cite{zoccali01} by using the cluster's WD population.
+" Zoccalictal.(2001) astructed a WD sequeice using local calibrators witli ueasured trigonometric parallaxes and magütudes to upare with the WD cooling sequence of lí Tuc. vieldiug au unreddcned ¢istance modulus of {u-Al,- 13.15x0.11 (see Tab 13."," \cite{zoccali01} constructed a WD sequence using local calibrators with measured trigonometric parallaxes and magnitudes to compare with the WD cooling sequence of 47 Tuc, yielding an unreddened distance modulus of (m - $_{\rm{o}}=13.15\pm0.14$ (see Table \ref{tab:distance}) )."
+ Hore. rather than use local ibrators. we deteriunue the distauce by conruius our liensred photometry to the hydrogen atinosphere (DÀ) WD models of Tremblayetal.(2011) aud the lel atinosphere. (DB) models of Berecronoetal. (2011)..," Here, rather than use local calibrators, we determine the distance by comparing our measured photometry to the hydrogen atmosphere (DA) WD models of \cite{tremblay11} and the helium atmosphere (DB) models of \cite{bergeron11}. ."
+ Previous spectroscopic studies have found solely DÀ WDs in elobular clusters within the temperaure range of our saniple. 1ichiding the elobular clusters ALL etal. 20093... NCC 6752 aud NCC 6397 (Alochler 2001)... as well as iu the open cluster NGC 2099 (Ixalirai 2005).," Previous spectroscopic studies have found solely DA WDs in globular clusters within the temperature range of our sample, including the globular clusters M4 \citep{davis09}, NGC 6752 and NGC 6397 \citep{moehler04}, as well as in the open cluster NGC 2099 \citep{kalirai05}."
+. Wule this has not been tested in LF Tuc to date. we make he assumption that our sample cousists of DÀ WDs. but we also compare our results to that obtained consiering DB models aud. a mixture of DA aud DD WD.," While this has not been tested in 47 Tuc to date, we make the assumption that our sample consists of DA WDs, but we also compare our results to that obtained considering DB models and a mixture of DA and DB WDs."
+ WDs within elobular clusters are faint objects ii crowded fields. making their ideutification quite challenging.," WDs within globular clusters are faint objects in crowded fields, making their identification quite challenging."
+— Iowevor. with— carefully plauned observatious with the IIST. combined with thorough data reduction and analysis. it becomes possible to obtain high quality multi-baud plotometry ofWDs in large sampleswithin one cluster system (seetherecent 2008)..," However, with carefully planned observations with the HST, combined with thorough data reduction and analysis, it becomes possible to obtain high quality multi-band photometry ofWDs in large sampleswithin one cluster system \citep[see the recent observations of WDs in M4 and NGC 6397 from][]{richer02,hansen02,hansen07,richer08}. ."
+Cosmology has developed dramatically in recent years: from being restricted to the realms of philosophy. our observational abilities have advanced it to the point where we may obtain precise evidence with which to shape our models and understanding.,"Cosmology has developed dramatically in recent years; from being restricted to the realms of philosophy, our observational abilities have advanced it to the point where we may obtain precise evidence with which to shape our models and understanding."
+ In this age of precision cosmology. the fine tuning of surveys can dramatically improve their performance.," In this age of precision cosmology, the fine tuning of surveys can dramatically improve their performance."
+ This requires us to think in terms of designer surveys rather than using a build-and-point approach., This requires us to think in terms of designer surveys rather than using a build-and-point approach.
+ For any given problem in cosmology (we use that of dark energy hereafter) many surveys of varving capabilities will be proposed. and a combination will make it through the conceptual stages to see the light of night.," For any given problem in cosmology (we use that of dark energy hereafter) many surveys of varying capabilities will be proposed, and a combination will make it through the conceptual stages to see the light of night."
+ For example. there are several stages detined by the Dark Energy Task Force (DETF) to classify dark energy surveys: stage IT surveys are complete. tthe Sloan Digital Sky Survey (SDSS): several stage HT surveys are now taking data. TThe Baryon Oscillation Spectroscopic Survey (BOSS) and WiggleZ. with others at the manufacturing stage. TThe Dark Energy Survey (DES) and the Hobby—Eberly Telescope Dark Energy EXperiment (HETDEX): and Stage IV. surveys are still in the design phase BBigBOSS. the Square Kilometre Array (SKA). the Wide-Field Infrared Survey Telescope (WFIRST). and the recently-approved Euclid satellite mission.," For example, there are several stages defined by the Dark Energy Task Force (DETF) to classify dark energy surveys; stage II surveys are complete, the Sloan Digital Sky Survey (SDSS); several stage III surveys are now taking data, The Baryon Oscillation Spectroscopic Survey (BOSS) and WiggleZ, with others at the manufacturing stage, The Dark Energy Survey (DES) and the Hobby–Eberly Telescope Dark Energy EXperiment (HETDEX); and Stage IV surveys are still in the design phase BigBOSS, the Square Kilometre Array (SKA), the Wide-Field Infrared Survey Telescope (WFIRST), and the recently-approved Euclid satellite mission."
+ When considering the large investments of time. money and expertise involved in these projects. it is imperative that designers identify the survey configuration that maximises the science return.," When considering the large investments of time, money and expertise involved in these projects, it is imperative that designers identify the survey configuration that maximises the science return."
+ Given the number of surveys all targeting the same goal. it is also important that they identify the appropriate niche: doing so maximises the overall science return from the combined effort of all relevant surveys.," Given the number of surveys all targeting the same goal, it is also important that they identify the appropriate niche; doing so maximises the overall science return from the combined effort of all relevant surveys."
+ Moreover. naive optimization can be wasteful unless there is an absolute scale of performance that can be used to determine survey configurations that are good enough for a given task. especially when dealing with time or cost restrictions. or with multiple probes within a survey.," Moreover, naive optimization can be wasteful unless there is an absolute scale of performance that can be used to determine survey configurations that are good enough for a given task, especially when dealing with time or cost restrictions, or with multiple probes within a survey."
+ The importance of optimization for dark energy surveys was first stressed by Bassett (2005) and Bassett. Parkinson Nichol (2005b).," The importance of optimization for dark energy surveys was first stressed by Bassett (2005) and Bassett, Parkinson Nichol (2005b)."
+ The concept of optimization is universal to design. regardless of the product in hand: a scalar rating or Figure of Merit (FoM) is defined. and the configuration of product variables that optimizes this number is identified.," The concept of optimization is universal to design, regardless of the product in hand; a scalar rating or Figure of Merit (FoM) is defined, and the configuration of product variables that optimizes this number is identified."
+ When it comes to designing a survey's observational parameters it is usual to exploit Monte Carlo Markov Chain (MCMC) methods to vary things like survey time. area. exposure time. and redshift range to identify the extreme of a FoM. In developing their roadmap for the future. the Dark Energy Task. Foree detined à FoM for comparing proposed dark energy The FoM is based on parameter estimation. quantifying the errors measured on the ACDAL values of the dark energy equation of state.," When it comes to designing a survey's observational parameters it is usual to exploit Monte Carlo Markov Chain (MCMC) methods to vary things like survey time, area, exposure time, and redshift range to identify the extreme of a FoM. In developing their roadmap for the future, the Dark Energy Task Force defined a FoM for comparing proposed dark energy The FoM is based on parameter estimation, quantifying the errors measured on the $\Lambda\rm CDM$ values of the dark energy equation of state."
+ This has subsequently become the standard for, This has subsequently become the standard for
+phases approximately settles to the following values: (1) 15 kms. (11) 30 km/s and (11) 75 km/s. What is very striking in the plot of the velocity dispersions. is the high velocities ( 500 km/s) attained by the hot. low density component of the gas.,"phases approximately settles to the following values: (I) 15 km/s, (II) 30 km/s and (III) 75 km/s. What is very striking in the plot of the velocity dispersions, is the high velocities $\sim$ 500 km/s) attained by the hot, low density component of the gas."
+ The densest structures which provide the raw material for star formation. collide and break apart. but are also subject to stripping via hvdrodyvnamical and thermal instabilities when this hot. low density material Dows rapiclv vast them.," The densest structures which provide the raw material for star formation, collide and break apart, but are also subject to stripping via hydrodynamical and thermal instabilities when this hot, low density material flows rapidly past them."
+ The picture of a violent. interstellar mecdium? (MeCray Snow 1979) emerges., The picture of a “violent interstellar medium” (McCray Snow 1979) emerges.
+ Reearcding the evolution of the thermal state of the gas. his is well portrayed in phase diagrams of the gas (bottom. row of figure 4)) which show the distribution of the mass raction of the gas as a function of its temperature and density.," Regarding the evolution of the thermal state of the gas, this is well portrayed in phase diagrams of the gas (bottom row of figure \ref{thermalevolution_4}) ) which show the distribution of the mass fraction of the gas as a function of its temperature and density."
+ Given our initial conditions of uniform density and emperature. if we were to plot a phase diagram of the gas at time /—0 Myr. all the gas would occupy a single point.," Given our initial conditions of uniform density and temperature, if we were to plot a phase diagram of the gas at time $t = 0$ Myr, all the gas would occupy a single point."
+ Decause the initial temperature (107 IX) of the gas coincides with the peak of the cooling curve. by 9. Myr. (first. panel of bottom row of figure 4)) the majority of the gas quickly raciativelv cools to an approximately isothermal state at a temperature corresponding to the minimum of the cooling curve. Le. 310 WK. As we instantaneously imprint a spectrum of velocity perturbations at the beginning of the simulation. the gas acquires a range of density values anc therefore has a spread in densities by this time.," Because the initial temperature $10^5$ K) of the gas coincides with the peak of the cooling curve, by 9 Myr (first panel of bottom row of figure \ref{thermalevolution_4}) ) the majority of the gas quickly radiatively cools to an approximately isothermal state at a temperature corresponding to the minimum of the cooling curve, i.e. 310 K. As we instantaneously imprint a spectrum of velocity perturbations at the beginning of the simulation, the gas acquires a range of density values and therefore has a spread in densities by this time."
+ Fhereafter. with the injection of hot gas into the medium. a tail of low density. hot gas appears.," Thereafter, with the injection of hot gas into the medium, a tail of low density, hot gas appears."
+" However as gas with temperatures LO” [Ix « Tox do10"" Wy (phase LED) is thermally unstable. it gradually vanishes from the medium. dividing the gas into two parts in the phase diagram."," However as gas with temperatures $10^5$ K $<$ T $<$ $4 \times 10^6$ K (phase III) is thermally unstable, it gradually vanishes from the medium, dividing the gas into two parts in the phase diagram."
+" Ehe majority of the coldest CE  300 Ix) gas dilfers by approximately a one order of magnitude pressure jump from gas with T 2 5 10"" Ix. Finally the pressure of both the hot and cold gas changes with time.", The majority of the coldest (T $\sim$ 300 K) gas differs by approximately a one order of magnitude pressure jump from gas with T $\geq$ 5 $\times$ $10^5$ K. Finally the pressure of both the hot and cold gas changes with time.
+ Lt rises as more and more hot gas fills the simulation volume. a situation that would probably be dillerent if hot gas were allowed to escape the box.," It rises as more and more hot gas fills the simulation volume, a situation that would probably be different if hot gas were allowed to escape the box."
+ Although complex. pictures. of the gas density. and temperature distribution in à two-dimensional slice through the simulation volume. do not capture the intricacy of the three-dimensional structure.," Although complex, pictures of the gas density and temperature distribution in a two-dimensional slice through the simulation volume, do not capture the intricacy of the three-dimensional structure."
+ In an attempt to display. this structure. in figure 5. we plot isodensity surfaces of the gas [or p—10.7.1. and 50 fem? at 50 Mrs.," In an attempt to display this structure, in figure \ref{gas3d} we plot isodensity surfaces of the gas for $\rho = 10^{-3}, 1, 10,$ and 50 ${\mathrm {cm}^3}$ at 50 Myrs."
+ Lt is clear [rom these figures that the hot. low density component fills most of the volume. while the densest regions fill the smallest fraction of the space. and are scattered throughout the box.," It is clear from these figures that the hot, low density component fills most of the volume, while the densest regions fill the smallest fraction of the space, and are scattered throughout the box."
+ A three-dimensional rendering of the stellar density at the same time instant (fig. 6)).," A three-dimensional rendering of the stellar density at the same time instant (fig. \ref{star3d}) ),"
+ reveals traces of the imprint of the high density gas. distribution and encouragingly bears some qualitative resemblance to the distribution of Ia emission in disk galaxies (c.g. NGC 4631. Wang et al.," reveals traces of the imprint of the high density gas distribution and encouragingly bears some qualitative resemblance to the distribution of $\alpha$ emission in disk galaxies (e.g. NGC 4631, Wang et al."
+ 2001)., 2001).
+ In an effort to assess what determines star formation rates. we systematically examine how different physical processes change the structure and the energeties of the interstellar mecium.," In an effort to assess what determines star formation rates, we systematically examine how different physical processes change the structure and the energetics of the interstellar medium."
+ Phe sequence of runs listed in table 1 are designed to isolate the effects of successively more. complicated physical processes., The sequence of runs listed in table \ref{sims} are designed to isolate the effects of successively more complicated physical processes.
+ A plot. comparing the star formation rates for this sequence of runs (figure 7)) invites us to study what keeps star formation at a minimum and alternatively what is necessary to drive it to high values., A plot comparing the star formation rates for this sequence of runs (figure \ref{sfrs}) ) invites us to study what keeps star formation at a minimum and alternatively what is necessary to drive it to high values.
+ Resolution ellects immediately. manifest. themselves in figure 7.., Resolution effects immediately manifest themselves in figure \ref{sfrs}.
+ The 64 and 1287 runs start. [rom the same initial conditions., The $64^3$ and $128^3$ runs start from the same initial conditions.
+ Preceecing star formation. feedback is non-existent. but gravity plavs a larger role in the 64% run where a grid cell of equivalent. density to that in the 128% grid is 8 times more massive.," Preceeding star formation, feedback is non-existent, but self-gravity plays a larger role in the $64^3$ run where a grid cell of equivalent density to that in the $128^3$ grid is 8 times more massive."
+ Therefore in the 645 run with only sellgravity (run C2). the SER rises more rapidly at earlier times than for the comparable run performed on the 128° grid (run B2).," Therefore in the $64^3$ run with only self-gravity (run C2), the SFR rises more rapidly at earlier times than for the comparable run performed on the $128^3$ grid (run B2)."
+ Once [cedback occurs. a mechanism supplementary to turbulence exists for creating high density contrasts which are stronger in the higher resolution runs.," Once feedback occurs, a mechanism supplementary to turbulence exists for creating high density contrasts which are stronger in the higher resolution runs."
+ Εις causes higher peaks of SER. in the 128? runs with feedback. (runs D3 and. 4) as compared to the equivalent 64% runs (C3 and. C4)., This causes higher peaks of SFR in the $128^3$ runs with feedback (runs B3 and B4) as compared to the equivalent $64^3$ runs (C3 and C4).
+ On the other hand. feedback also creates an extra source of pressure to fight self-eravity which explains why the C2 run leads to higher SEIs at earlier times than the €3 and C4 runs.," On the other hand, feedback also creates an extra source of pressure to fight self-gravity which explains why the C2 run leads to higher SFRs at earlier times than the C3 and C4 runs."
+ What, What
+of γε the whole [ range is seemingly allowed.,of $\g_e$ the whole $\G$ range is seemingly allowed.
+" However, this only happens at γε«62,34,10 for a=1.,0.5,0 respectively, and therefore conflicts with the soft X-ray observations."," However, this only happens at $\g_e<62,34,10$ for $\alpha = 1.,0.5,0$ respectively, and therefore conflicts with the soft X-ray observations."
+ In cases when Yy>1 most of the electrons energy is emitted as very high energy (TeV) gamma-rays., In cases when $Y_H>1$ most of the electrons energy is emitted as very high energy (TeV) gamma-rays.
+" When the scattering is in the KN regime, that is if Eq."," When the scattering is in the KN regime, that is if Eq."
+" 3 holds, the scattered photon, that has an energy of almost y.m.c? (in the fluid’s rest frame) and can therefore produce a pair when it encounters a typical low energy y-ray photon with energy hv,/T (in this frame)."," \ref{fKN} holds, the scattered photon, that has an energy of almost $\gamma_e m_e c^2 $ (in the fluid's rest frame) and can therefore produce a pair when it encounters a typical low energy $\gamma$ -ray photon with energy $h \nu_\g/\G$ (in this frame)."
+" More specifically, for a head on collision between a photon with energy hvy=£Pmec?/4. and an electron with a Lorentz factor γε, the energies of the electron and the photon after the collision are: and The resulting photon has now enough energy to collide with a photon with energy hv, "," More specifically, for a head on collision between a photon with energy $h\nu_\g = \xi \G m_e c^2 / \g_e $ and an electron with a Lorentz factor $\g_e$, the energies of the electron and the photon after the collision are: and The resulting photon has now enough energy to collide with a photon with energy $h \nu_\g$ "
+The instrument properties have changed over the 11 vears of operation in orbit because of at least three effects.,The instrument properties have changed over the 14 years of operation in orbit because of at least three effects.
+ First. the gain of SSC 1 (of SSCs 1-3) increased at about per vear for the first 13 vears or so because of a leak in the couuter entrance window (the rate of Increase appears to have moderated in vear 11).," First, the gain of SSC 1 (of SSCs 1-3) increased at about per year for the first 13 years or so because of a leak in the counter entrance window (the rate of increase appears to have moderated in year 14)."
+ This has resulted. first. in progressive reductions in the effective sensitivity of SSC Liu the 5-12 keV band aud. receutly. ο sole extent iu the 3-5 keV baud because in the last ew vears the pulse heights of higher euergv eveuts are often saturated aud thus unusable.," This has resulted, first, in progressive reductions in the effective sensitivity of SSC 1 in the 5-12 keV band and, recently, to some extent in the 3-5 keV band because in the last few years the pulse heights of higher energy events are often saturated and thus unusable."
+ The second of the effects 1 present because he one-dimensional positiou-seunsiug capability. of the xoportional counters is based on the charge-division echuique and the resistivity of carbou-coated quartz-Bber anodes., The second of the effects is present because the one-dimensional position-sensing capability of the proportional counters is based on the charge-division technique and the resistivity of carbon-coated quartz-fiber anodes.
+ The process of eveut detection iu the counters results in the redistribution or removal of the carbon., The process of event detection in the counters results in the redistribution or removal of the carbon.
+ The resistivity of cach anode thus evolves over iue ecnerally becoming larger aud less uniformi. aud so he functions that relate charge ratio to position have changed over the course of tle mission.," The resistivity of each anode thus evolves over time generally becoming larger and less uniform, and so the functions that relate charge ratio to position have changed over the course of the mission."
+ Third. in the last few wears there is often a high vackeround of events that have pulse heights that sugecst hat they are due to the detection of Al Na photous.," Third, in the last few years there is often a high background of events that have pulse heights that suggest that they are due to the detection of Al $\alpha$ photons."
+ We do not understand in detail why this is happening but we speculate that it may be due to leakage of ultraviolet radiation through the Al-coated Ἱναρίοι thermal shields hat are inuuecdiately above the coded mask plates., We do not understand in detail why this is happening but we speculate that it may be due to leakage of ultraviolet radiation through the Al-coated Kapton thermal shields that are immediately above the coded mask plates.
+ UV radiation would produce plotoclectrous at the counter xlv surfaces that happen to be. bv desigu of the iustrmucut. at potentials of approximately LakV. The xiotoelectrous would be accelerated toward aud would ut. e.g. the aluniuun collimators that are near the eround potential.," UV radiation would produce photoelectrons at the counter body surfaces that happen to be, by design of the instrument, at potentials of approximately $-1.8$ kV. The photoelectrons would be accelerated toward and would hit, e.g., the aluminum collimators that are near the ground potential."
+ Some fraction of the Al Ite photons hat would be produced by this process would then be detected., Some fraction of the Al $\alpha$ photons that would be produced by this process would then be detected.
+ Both the eun chauges aud charge ratio to position calibration changes are accounted for in detail in the detector model used iu our cocded-aperture analysis., Both the gain changes and charge ratio to position calibration changes are accounted for in detail in the detector model used in our coded-aperture analysis.
+ However. the calibrations are imperfect aud the derived source intensities are therefore somewhat affected hy systematic errors on both short aud long time scales.," However, the calibrations are imperfect and the derived source intensities are therefore somewhat affected by systematic errors on both short and long time scales."
+ The coded aperture analysis is also designed to eliminate the effects of backgrounds that are moth: and. slowly-varving with regard to position iu the detector: this is largely but not cutirely effective., The coded aperture analysis is also designed to eliminate the effects of backgrounds that are smooth and slowly-varying with regard to position in the detector; this is largely but not entirely effective.
+ Searches for periodicitics where large ummbers of Lnüieasnrenuents are involved have ecnerally proceeded by averaging the measured intensities in equally-spacecd time bins. Fourier transtormune the resulting time series producing a power density spectrmu by takine the absolute square of the complex amplitude for cach frequency. aud scarcling for peaks in the power spectrma," Searches for periodicities where large numbers of measurements are involved have generally proceeded by averaging the measured intensities in equally-spaced time bins, Fourier transforming the resulting time series, producing a power density spectrum by taking the absolute square of the complex amplitude for each frequency, and searching for peaks in the power spectrum."
+ Our search kweely follows this outline but we have implemented a few siguificant niodifications as described below., Our search largely follows this outline but we have implemented a few significant modifications as described below.
+ Iu some searches for periodicities. a Lomb-Scarele periodogram (Scarele1982:: sce also Pressetal. 1992)) is used iu place of a Fourier traustori.. but. when large ununnboers of 10ea8ureimmoents aud time bius are being analyzed. the differences between the two are not miportaut except at extremely low frequencies.," In some searches for periodicities, a Lomb-Scargle periodogram \citealt{scargle82}; see also \citealt{pressetal92}) ) is used in place of a Fourier transform, but, when large numbers of measurements and time bins are being analyzed, the differences between the two are not important except at extremely low frequencies."
+ For the present analyses. the conventional discrete Fourier transform is adequate so we do not make use of the Lonib- variation.," For the present analyses, the conventional discrete Fourier transform is adequate so we do not make use of the Lomb-Scargle variation."
+ We enplov two unusual strategies with the goal 6:dquprovius the search sensitivity., We employ two unusual strategies with the goal of improving the search sensitivity.
+ The first strategy involves the use of weights such as the reciprocals of the variances m Fourier aud other types of analysis since the individual ASM measurements have a wide range of associated uncertainties., The first strategy involves the use of weights such as the reciprocals of the variances in Fourier and other types of analysis since the individual ASM measurements have a wide range of associated uncertainties.
+ Since in the preseut instance the auplitudes iu the Fourier transforms may be regarded as statistical averages. they may be better determined bv the use of appropriate weights.," Since in the present instance the amplitudes in the Fourier transforms may be regarded as statistical averages, they may be better determined by the use of appropriate weights."
+ This is discussed briefiv in Corbetetal.(2007a} aud at exeater leneth iu Corbetetal.(2007))., This is discussed briefly in \citet{corbetal07} and at greater length in \citet{corbmt07}.
+. These two reports also discuss the use of weights which involve the intrinsic source variability. but we do uot use such weights in the work reported herem.," These two reports also discuss the use of weights which involve the intrinsic source variability, but we do not use such weights in the work reported herein."
+ The second strateey stems from the fact that the observations of the source are obtained with a low duty evcle. ie. the window fiction is sparse but nevertheless luehly structured.," The second strategy stems from the fact that the observations of the source are obtained with a low duty cycle, i.e., the window function is sparse but nevertheless highly structured."
+ The properties of the window fiction. in combination with the presence of slow variations of the source intensity. act to hinder the detection of varlatious on short time scales.," The properties of the window function, in combination with the presence of slow variations of the source intensity, act to hinder the detection of variations on short time scales."
+ The window function power density spectrum (PDS) has substantial power at hieh frequencies. ce. 1l evele  and 1 cycle per spacecraft orbit (~95 müuute period).," The window function power density spectrum (PDS) has substantial power at high frequencies, e.g., 1 cycle $^{-1}$ and 1 cycle per spacecraft orbit $\sim95$ minute period)."
+ Since the data nav be regarded as the product of a (hypothetical) continuous set of source intensity measurements with the window function. a Fourier transform of tle data is cquivalent to the convolution of a transform of a continuous set of intensity measurements with the window fuuction transtorm.," Since the data may be regarded as the product of a (hypothetical) continuous set of source intensity measurements with the window function, a Fourier transform of the data is equivalent to the convolution of a transform of a continuous set of intensity measurements with the window function transform."
+ The high frequency structure in the window function transform acts to spread power at low frequencies in the source intensity to high frequeucies in the calculated transform (or. equivalently the PDS).," The high frequency structure in the window function transform acts to spread power at low frequencies in the source intensity to high frequencies in the calculated transform (or, equivalently, the PDS)."
+ This effectively raises the noise level at high frequencics., This effectively raises the noise level at high frequencies.
+ Thus. in most of our analyses. a smoothlied version of the helt curve is subtracted from the binned data prior to the Fourier analysis.," Thus, in most of our analyses, a smoothed version of the light curve is subtracted from the binned data prior to the Fourier analysis."
+ The details of the analvsis of cach ASM dwell-by-dwell helt curve are as follows., The details of the analysis of each ASM dwell-by-dwell light curve are as follows.
+ The time of each iuteusitv lcasurement was adjusted to approximately remove the effects of the motion of the Earth around the Solar System baryceuter aud the adjusted times were used to assign the measurements to equally-spaced time bins of duration 5 minutes (0.0031722 d) or more., The time of each intensity measurement was adjusted to approximately remove the effects of the motion of the Earth around the Solar System barycenter and the adjusted times were used to assign the measurements to equally-spaced time bins of duration 5 minutes (0.0034722 d) or more.
+ Weights are used in averaging those measurements that fall iuto a eiven tine bin., Weights are used in averaging those measurements that fall into a given time bin.
+ The »; measurements that fall iuto the time bin which is labelled bv the iudex / are denoted with iudex ; that runs from 1. to »;., The $n_i$ measurements that fall into the time bin which is labelled by the index $i$ are denoted with index $j$ that runs from 1 to $n_i$.
+ We denote the jtl lüueasured source iuteusitv i a time biu by s; aud its associated 1-sgnia uncertainty by σι., We denote the $j$ th measured source intensity in a time bin by $s_j$ and its associated 1-sigma uncertainty by $\sigma_j$.
+ The averaging in this. stage uses weights. wj;=41σ]be “where Avy is. a selected weieliting index (sce values in Table 1))., The averaging in this stage uses weights $w_j = 1/\sigma_j^{k_A}$ where $k_A$ is a selected weighting index (see values in Table \ref{tbl:mthds}) ).
+ The Intensity in bin / is then For convenience we denote the denominator of this expression by IV; so that, The weighted-average intensity in bin $i$ is then For convenience we denote the denominator of this expression by $W_i$ so that
+ssunples we can derive some information on the spatial distribution of the stellar populations.,samples we can derive some information on the spatial distribution of the stellar populations.
+ This Προς dealing with smaller samples and thus a deterioration in nuuber statistics., This implies dealing with smaller samples and thus a deterioration in number statistics.
+ Diagrams derived from the current velocitics are qualitatively the same for all the samples., Diagrams derived from the current velocities are qualitatively the same for all the samples.
+ A closer inspection of the Toomme diagram shows that a majority of our objects are stars with thin disk orbit characteristics and only a iumor fraction is of a halo like tvpe., A closer inspection of the Toomre diagram shows that a majority of our objects are stars with thin disk orbit characteristics and only a minor fraction is of a halo like type.
+ Disk stars are iucluded in all the samples in a large enough quantity for aualvses., Disk stars are included in all the samples in a large enough quantity for analyses.
+ In the C. V-diagraun. the bimodality of the data cau still be seen. even in the ssunple.," In the $U,V$ -diagram, the bimodality of the data can still be seen, even in the sample."
+ This is due to the fact that the separation is observed for disk stars. and they are adequatolv represented in all the percentage sets.," This is due to the fact that the separation is observed for disk stars, and they are adequately represented in all the percentage sets."
+ Iu contrast. ouly three of the IPS stars have a relative parallax error below.," In contrast, only three of the HPS stars have a relative parallax error below."
+. This low umber implies that with these we caunot make a mieaningful analyses for the halo., This low number implies that with these we cannot make a meaningful analyses for the halo.
+ For AG the uuniber of IIPS stars is twelve.," For $\leq$, the number of HPS stars is twelve."
+ The orbital data thus coufirms the similarity of the conrpositioun of the error-selected samples., The orbital data thus confirms the similarity of the composition of the error-selected samples.
+ For the disk. stars. the histograms are of simular shape aud IPS stars of the acd ssaluple are also present with eccentricities larger than 0.1.," For the disk stars, the histograms are of similar shape and HPS stars of the and sample are also present with eccentricities larger than 0.4."
+ We conclude that no kinematical preference cau be found in the closer samples., We conclude that no kinematical preference can be found in the closer samples.
+ Even though we have very low unnuibers of stars for a D-distribution analysis in the sample (see rofstartab)). we give rough estimates ou the scale heights.," Even though we have very low numbers of stars for a $z$ -distribution analysis in the sample (see \\ref{startab}) ), we give rough estimates on the scale heights."
+ They are 0.12 kpe for the thin disk. 0.31 kpe for the thick disk and 5 spe for the halo.," They are 0.12 kpc for the thin disk, 0.34 kpc for the thick disk and 8 kpc for the halo."
+ Scale heights of the BHuple are 0.12 kpc. 0.16 kpc. aud. 9 kpe. respectively.," Scale heights of the sample are 0.12 kpc, 0.46 kpc, and 9 kpc, respectively."
+ These results are not different iu an essential wav from those of 1t startal.., These results are not different in an essential way from those of \\ref{startab}.
+ Another. perliaps critical aspect of our ROB selection frou ddata is the heavily populated location of the window border at the RC brauch.," Another, perhaps critical aspect of our RHB selection from data is the heavily populated location of the window border at the RG branch."
+ By excluding more of the critical red stars. we can check for influences of elauts and sec if our above results still hold.," By excluding more of the critical red stars, we can check for influences of giants and see if our above results still hold."
+ A new separation was adopted. by shiftiug the line towards brighter bby 0.5 mae (see dashed line in πο MDsel)y," A new separation was adopted, by shifting the line towards brighter by 0.5 mag (see dashed line in \\ref{CMDsel}) )."
+ We thereby excluded almost half of the full ssuuple eudiug up with some 1912 sstars., We thereby excluded almost half of the full $\leq$ sample ending up with some 1942 stars.
+ From. these stars. τοι. hack τοῦ lueasurenieuts doctunented iu BB2k (1327 stars in he full sample. see refstartal)).," From these stars, 731 had 789 measurements documented in BB2k (1327 stars in the full sample, see \\ref{startab}) )."
+ With this new aud restricted sample we noted several thines., With this new and restricted sample we noted several things.
+ The velocity listogram shows that some of the stars lageine behind the disk rotation arc now missing., The velocity histogram shows that some of the stars lagging behind the disk rotation are now missing.
+ The histogram of thick disk velocity stars now consists of several peaks., The histogram of thick disk velocity stars now consists of several peaks.
+ Stars with thin disk velocitics arc also less numerous. but the peak is still very simular in shape.," Stars with thin disk velocities are also less numerous, but the peak is still very similar in shape."
+" Looking at the Toonmre diagram allows the ideutification of the ""iuissing stars.", Looking at the Toomre diagram allows the identification of the “missing” stars.
+ We found that only 15 of the TIPS stars are stil iucludec in the restricted sample. but the relative wmuber of IIPS to disk stars stavs the same compared to or the base-line sample).," We found that only 15 of the HPS stars are still included in the restricted sample, but the relative number of HPS to disk stars stays the same compared to for the base-line sample)."
+ In addition. a reduction of star nunubers in the range of 100 οςep€150 wwas noticed.," In addition, a reduction of star numbers in the range of 100 $\leq v_{\perp}\leq$ 150 was noticed."
+ The abseuce of these objects explains new eaps in the cee-listogram (at eec=0.35) and ©.Eu- diagram (at O=175 Ly," The absence of these objects explains new gaps in the $ecc$ -histogram (at $ecc=0.35$ ) and $\Theta,E_{\rm kin}$ -diagram (at $\Theta=175$ )."
+ Tn contrast. the Bottlinger diagram aud the listoeram of normalized τοπίο are qualitatively the same as with the base-line sample.," In contrast, the Bottlinger diagram and the histogram of normalized $z$ -extent are qualitatively the same as with the base-line sample."
+" Thus. we derived the scale heights for this restriced suuple via a \o-fit frou the +-probability distribution results for the disks in feb,=0.2 kpe aud 0.16 kpe."," Thus, we derived the scale heights for this restriced sample via a $\chi^2$ -fit from the $z$ -probability distribution results for the disks in $h_{\rm thin}=0.12$ kpc and $h_{\rm thick}=0.46$ kpc."
+" For the halo we get Mya,=1.2 kpc.", For the halo we get $h_{\rm halo}=4.2$ kpc.
+ Iu the selection window there is a potential contamination of stars that do not beloug to the IIB. such as massive core heliuumi burners or RG stars (see rofRIIDRC )).," In the selection window there is a potential contamination of stars that do not belong to the HB, such as massive core helium burners or RG stars (see \\ref{RHBRC}) )."
+ Core Heli burners are more massive than RIB stars and evolve faster., Core Helium burners are more massive than RHB stars and evolve faster.
+ The probability to fiud these inside our window is small., The probability to find these inside our window is small.
+ RG stars of solar auetallicitv found within ou selection window are vounger than 1.25 bilhon wears (Yousci-Yale isoclaoues by 2))., RG stars of solar metallicity found within our selection window are younger than 1.25 billion years (Yonsei-Yale isochrones by \citealt{YY}) ).
+ Asstmine a constant Star Formation Rate (SFR) and an observed ROB ouly of solar metallicityv. we would expect that an ROB of 0.5 de.cwouldlieroughlg(.25 maytothcbluc," Assuming a constant Star Formation Rate (SFR) and an observed RGB only of solar metallicity, we would expect that an RGB of $=-0.5$ dex would lie roughly 0.25 mag to the blue."
+ RG Starsofthatimnetallic. Valeccolutionarytracksby?))\. sothismetal RGB wouldinthe€ AL Dbeattwo thirdsoftheobserved RGDB'sdeusity.," RG Stars of that metallicity develop faster by some (derived from Yonsei-Yale evolutionary tracks by \citealt{Yi03}) ), so this metal-poor RGB would in the CMD be at two-thirds of the observed RGB's density."
+Howecer thedensitul," However, the density bluewards of the observed RGB is much lower than that."
+luewe," Had the SFR in the past been smaller, then the contamination of our sample is smaller, too."
+ards CCMD.," If the SFR had been larger in the past, we would see many more RGs of bluer colour than actually are seen in the CMD."
+" Althisshowsthatacontaminationofour RII Bsample ifpr RC poorstars,"," All this shows that a contamination of our RHB sample (if present) comes from young stars and not from old, metal-poor RG stars."
+ We check for this contamination of our RUB sample by separating the selection window in two parts., We check for this contamination of our RHB sample by separating the selection window in two parts.
+ The Blic Sample is defined as having Wy<0.85 mae. while the other stars are assigned to the Red Sample (κος refC MDzel)).," The Blue Sample is defined as having $\leq0.85$ mag, while the other stars are assigned to the Red Sample (see \\ref{CMDsel}) )."
+ Stars that belong to the bluer sample are certain to be RIIDBs., Stars that belong to the bluer sample are certain to be RHBs.
+ These 99 stars include 6 IIPS stars., These 99 stars include 6 HPS stars.
+ The Red Sample contains 1228 stars. inclicding 19 TIPS stars.," The Red Sample contains 1228 stars, including 19 HPS stars."
+ The analvses of the two samples sugeest that there are uo sigmificant differences in kincmatic behaviour., The analyses of the two samples suggest that there are no significant differences in kinematic behaviour.
+ Tle Bluc Sample does have a comparatively higher fraction of halo stars. meludiug the two stars that have orbits reaching out to z>10 kpe.," The Blue Sample does have a comparatively higher fraction of halo stars, including the two stars that have orbits reaching out to $z>10$ kpc."
+ Scale heights of the Blue Sample are, Scale heights of the Blue Sample are
+LIL). its effective gravity could still be sieutficautly lower thaithe gravity of a maiu-sequeuce star (Oke1977:Murdinetal. 1980)..,"III), its effective gravity could still be significantly lower than the gravity of a main-sequence star \citep{oke1977,murdin1980}. ."
+ Shalibaz.Baucyopadhvay&Charles(1999) used ouly dwarl stars or their template spectra. so their results could not cistiugiish gravity depencent changes in the À2.29 jun PCO(2.0) line sreneth.," \citet{shahbaz1999} used only dwarf stars for their template spectra, so their results could not distinguish gravity dependent changes in the $\lambda2.29\ \mu$ m $^{12}$ CO(2,0) line strength."
+ Secoud. observations of cataclysmic variables have shown that in some systems. CO absoption is weaker relative to the streneths of the atomic absorption lines han expected for their dorOr star spectral types (Harrison.Szkody&Jolusou1999).," Second, observations of cataclysmic variables have shown that in some systems, $^{12}$ CO absorption is weaker relative to the strengths of the atomic absorption lines than expected for their donor star spectral types \citep{harrison1999}."
+. In addition. he ratio of the !12CO baiehead equivalent widtl to those of the atomic lines can vary with time {Ramseyeretal.1993).," In addition, the ratio of the $^{12}$ CO bandhead equivalent width to those of the atomic lines can vary with time \citep{ramseyer1993}."
+. Tlie reason for these abiormalities in the CO absorption liue streugtlis is rot yet. uuderstood., The reason for these abnormalities in the CO absorption line strengths is not yet understood.
+ Τlus. sine the streneth of tve A2.29jun 1I?C0(2.0) bancdheacl to measure the contribution of the disx to the iufrared {tix gives tiireliable results.," Thus, using the strength of the $\lambda2.29\ \mu$ m $^{12}$ CO(2,0) bandhead to measure the contribution of the disk to the infrared flux gives unreliable results."
+ The measurements by Charles(1999) do not settle tie question of how much flux the cdisk contributes to the light curve of &0620-00 at infrared waveletigthis., The measurements by \citet{shahbaz1999} do not settle the question of how much flux the disk contributes to the light curve of A0620-00 at infrared wavelengths.
+ The orbital inclination could. therefore. be significantly higher than the lower limit set by the ellipsoklal variations (2387) aud tlie mass of the compact star much less than the upper limit (M4<13.6 )).," The orbital inclination could, therefore, be significantly higher than the lower limit set by the ellipsoidal variations $i > 38\arcdeg$ ) and the mass of the compact star much less than the upper limit $M_{1} < 13.6$ )."
+ l., 1.
+ The infrared light curves of the Jack hole binary A0620-00 show an asymmetric. moculation with extra enission in the peak at phase ὁ = 0.75.," The infrared light curves of the black hole binary A0620–00 show an asymmetric, double-humped modulation with extra emission in the peak at phase $\phi$ = 0.75."
+ There were 1ο gross Changes in the morphology of the light curves ove τα oue year period from 1995 December to 1996 December., There were no gross changes in the morphology of the light curves over a one year period from 1995 December to 1996 December.
+ The mean inlrared. colors aud the shape of the ligit curve are also the same as οerved iu 1990 January., The mean infrared colors and the shape of the light curve are also the same as observed in 1990 January.
+ There were no sharp dips iu tle light curve nor reversals of the asyuunery between te two lumps as were seen in observations circa 1986 — 1989., There were no sharp dips in the light curve nor reversals of the asymmetry between the two humps as were seen in observations circa 1986 – 1989.
+ 2., 2.
+Tie light curves are consistent with ellipsoidal variatious [rom the Ix star pus beamec ῃx [rom a spot on the accretion disk., The light curves are consistent with ellipsoidal variations from the K star plus beamed flux from a spot on the accretion disk.
+ A precessing disk is au iulikely source of light «Mirve modulation curing the observation period., A precessing disk is an unlikely source of light curve modulation during the observation period.
+ Star spots are also ruled ou 1uless the spot is fixed in location d size on the Ix star surface., Star spots are also ruled out unless the spot is fixed in location and size on the K star surface.
+ 3., 3.
+ Based on fits to the lower peak in the light curve (yelween ó = 0 aud 0.51). the ΤΙ inclination in 0620-00 is 7>387.," Based on fits to the lower peak in the light curve (between $\phi$ = 0 and 0.51), the minimum inclination in A0620–00 is $i \geq
+38\arcdeg$."
+" From the abseuce of eclipse features in the ligib curves. moclels including an accretion disk of radius dy=0.5Rp, give an ipper limit to the iuclilation oli<75°"," From the absence of eclipse features in the light curves, models including an accretion disk of radius $R_{d} = 0.5\ R_{L_{1}}$ give an upper limit to the inclination of $i \leq 75\arcdeg$."
+" The mass of the compact star in &0620-00 is 3.3<AL,<13.6AL.", The mass of the compact star in A0620–00 is $3.3 < M_{1} < 13.6$.
+.. |., 4.
+ Ellipsoidal variatious provide ouly a lower limit to the inclination of noi-eclipsiug binary systems when the contribution of the accretion disk is unkιο., Ellipsoidal variations provide only a lower limit to the inclination of non-eclipsing binary systems when the contribution of the accretion disk is unknown.
+ A previous attetupt to determine the relative contribution of the accretion disk and Ix star to the near-infrared flux ised the strength of the A2.29 jum 1I?C0(2.0)bandhbead.," A previous attempt to determine the relative contribution of the accretion disk and K star to the near-infrared flux used the strength of the $\lambda2.29\ \mu$ m $^{12}$ CO(2,0)bandhead."
+ This method is not 'eliable., This method is not reliable.
+Cyenus-A is locally the most powerful Fanarolf-Hüley (iFRIL) radio galaxy and therefore it is also the best studied in ternis of spatial resolution.,Cygnus-A is locally the most powerful Fanaroff-Riley (FRII) radio galaxy and therefore it is also the best studied in terms of spatial resolution.
+ Ht has even. had. an entire workshop devoted to it and. global properties of the object are well reviewed! in these references (Carilli Barthel 1996. Carilli Harris L996. ancl Carilli ct 1998).," It has even had an entire workshop devoted to it and global properties of the object are well reviewed in these references (Carilli Barthel 1996, Carilli Harris 1996, and Carilli et 1998)."
+ Cyenus-A has featured) prominently in the quest for the unification of powerful radio galaxies and quasars. being cast as the classical case of a quasar in the plane of the sky. (Antonnuci Mailer 1985. Barthel 1989).," Cygnus-A has featured prominently in the quest for the unification of powerful radio galaxies and quasars, being cast as the classical case of a quasar in the plane of the sky (Antonnuci Miller 1985, Barthel 1989)."
+ In terms of the svnchrotron emission. at. low radio frequencies the two giant lobes domünate the emission. (llarerave Ryle 1974). but at higher frequencies. the hotspots. or working surfaces in the lobes. become more prominent along with the galaxy core.," In terms of the synchrotron emission, at low radio frequencies the two giant lobes dominate the emission (Hargrave Ryle 1974), but at higher frequencies, the hotspots, or working surfaces in the lobes, become more prominent along with the galaxy core."
+ The southern and northern hotspots are respectively H op D. ⋅↱≻∪⋜⋯∠⋀∪⇂↓⋅∪⊔↓∣⇂⊓⋅≼∙∪↓⋅∢, The southern and northern hotspots are respectively $50''$ and $70''$ from the core.
+⋅⋡⇂⊳∖↓⊔⋏∙≟⋜↧∐⊔⋡∣⋡⇂⋖⋅≼∙∪⊔≱∖↿⋜⋯∣∪⇂ ⊥ ↳∖↓↓≻≼⇍↓⋜⋯∠⇂⋜↧↓⋅∢⊾∠⇂⊳∖↓↥⊲↓∐∪⇂⋅∶∶≱⋅∪⋅↱≻≼⊔↿∖∺↥⋯∙↳↿∪⊔ ∢⋅↿⋜↧↓⋡↓≤⋗≤⊔⊐⋡≺⊲, Using a Hubble constant of $^{-1}$ $^{-1}$ and a redshift of $z = 0.0562$ (Stockton et al.
+∙∖⇁⋏∙≟⊔⊔≱∖−⇀∖∐⋖⋅⊳∖⋜↧⇂⋜↧∠∐⊳∖↿⋜⋯⊓⊾∪⇂⋅⇉⇉⊤↳∖∟∖↓," 1994), Cygnus-A lies at a distance of Mpc."
+↓⋊⋱ Deep VLA images reveal the thin jet which transports energv from the AGN core to the radio lobes (Perley. Dreher. Cowan 1984).," Deep VLA images reveal the thin jet which transports energy from the AGN core to the radio lobes (Perley, Dreher, Cowan 1984)."
+ VLBI observations by Linfield (1984) and other authors. and most recently by. Ixrichbaunm et ((1998). have extended: this picture at. sub-milliaresecond resolution.," VLBI observations by Linfield (1984) and other authors, and most recently by Krichbaum et (1998), have extended this picture at sub-milliarcsecond resolution."
+ The fact that the electron svnchrotron lifetime in the hotspots is less than the light travel time from the central core (Llarerave Ryle 1974) means that electron re- acceleration must take place in the lobes. and the hotspots ave believed. to be the working surfaces at which this re- acceleration takes place.," The fact that the electron synchrotron lifetime in the hotspots is less than the light travel time from the central core (Hargrave Ryle 1974) means that electron re- acceleration must take place in the lobes, and the hotspots are believed to be the working surfaces at which this re- acceleration takes place."
+ Indeed. the prominence of hotspots ab the outer edges of the lobes is one of the identifving features of ΙΕ radio sources (Fanaroll Riley 1974).," Indeed, the prominence of hotspots at the outer edges of the lobes is one of the identifying features of FRII radio sources (Fanaroff Riley 1974)."
+ The precise mechanism for the electron acceleration is uncertain and determination. of the svnchrotron. spectral index is important as it can rule out some potential mechanisms., The precise mechanism for the electron acceleration is uncertain and determination of the synchrotron spectral index is important as it can rule out some potential mechanisms.
+ For à steady injection model (Bell 1978) a steepening of 0.5 in the index is expected at à certain frequency as the electrons lose energy. more rapidly than is being supplied by the acceleration. process., For a steady injection model (Bell 1978) a steepening of 0.5 in the index is expected at a certain frequency as the electrons lose energy more rapidly than is being supplied by the acceleration process.
+ At some higher frequency. the electron spectrum. will cut-olf as the high energy. electrons rapidlv lose energy and become depleted: this. produces the final turnover ancl steep downturn of the svnchrotron emission.," At some higher frequency, the electron spectrum will cut-off as the high energy electrons rapidly lose energy and become depleted; this produces the final turnover and steep downturn of the synchrotron emission."
+ The emission from the hot-spots (denoted by A and D in the convention derived. from. VLA maps. where A ds the north-west and. D is the south-cast hot-spot) has been well observed at a number of radio frequencies (Muxlow," The emission from the hot-spots (denoted by A and D in the convention derived from VLA maps, where A is the north-west and D is the south-east hot-spot) has been well observed at a number of radio frequencies (Muxlow"
+We consider that the negative density fluctuations evolve within a medium of density pi=poll|z)*. where po=Pinon|fuo Che subseript w stands for the physical parameters of the Universe).,"We consider that the negative density fluctuations evolve within a medium of density $\rho_{\rm t}=\rho_0(1+z)^3$, where $\rho_0=\rho_{\rm bu0}+\rho_{\rm du0}$ (the subscript `u' stands for the physical parameters of the Universe)."
+ Thus. in our model we take into account the inlluence of the external medium on the evolution of the underdense region. and therefore as external boundary condition we have the pressure of the Universe. namely lt is also worth stressing that in many studies of structure cmation the influence of the external medium on the garucture under study is not taken into consideration.," Thus, in our model we take into account the influence of the external medium on the evolution of the underdense region, and therefore as external boundary condition we have the pressure of the Universe, namely It is also worth stressing that in many studies of structure formation the influence of the external medium on the structure under study is not taken into consideration."
+ In order to calculate the temperature of the matter of the Universe we consider in the cooling function only the contributions due to the Compton cooling-heating and the cooling due to the recombination processes., In order to calculate the temperature of the matter of the Universe we consider in the cooling function only the contributions due to the Compton cooling-heating and the cooling due to the recombination processes.
+ The her physical processes considered. for the evolution of 10 barvonic density perturbation are not relevant for the evolution of the matter temperature of the Universe., The other physical processes considered for the evolution of the baryonic density perturbation are not relevant for the evolution of the matter temperature of the Universe.
+ The energy equation thus reads Similarly. the collisional ionization is not important for the matter of the Universe.," The energy equation thus reads Similarly, the collisional ionization is not important for the matter of the Universe."
+ Phus. we use the study of Peebles (LOGS). where the degree of ionization is given by where now Equations (19) and (20) are. respectively. the equations of conservation of thermal energy and ionization balance. with Zin the temperature of the matter of the Universe and rna its ionization fraction.," Thus, we use the study of Peebles (1968), where the degree of ionization is given by where now Equations (19) and (20) are, respectively, the equations of conservation of thermal energy and ionization balance, with $T_{\rm bu}$ the temperature of the matter of the Universe and $x_{\rm bu}$ its ionization fraction."
+ The barvonic ambient density. Lo. the barvonic density of the Universe. is given by where z is the redshift and Yo is the present cosmic background radiation temperature.," The baryonic ambient density, i.e., the baryonic density of the Universe, is given by where $z$ is the redshift and $T_0$ is the present cosmic background radiation temperature."
+" The equation for the time evolution of the radiation temperature and. therefore. the time evolution of the Hubble parameter is given by where 4,=Oi|Qa is the density parameter of the matter."," The equation for the time evolution of the radiation temperature and, therefore, the time evolution of the Hubble parameter is given by where $\Omega_{\rm m}=\Omega_{\rm b}+\Omega_{\rm d}$ is the density parameter of the matter."
+ To solve the above set of equations. a hvdrodynamical Lagrangian numerical code has been written. which basically follows the method developed. by Ryehtmver Morton (1967).," To solve the above set of equations, a hydrodynamical Lagrangian numerical code has been written, which basically follows the method developed by Rychtmyer Morton (1967)."
+ Lt is important to stress that the total amount of barvonic matter and non-barvonic dark matter are taken to be constant throughout the calculations., It is important to stress that the total amount of baryonic matter and non-baryonic dark matter are taken to be constant throughout the calculations.
+ We use in our models 9000 shells for the barvonie matter and 9000 shells for the dark matter., We use in our models 9000 shells for the baryonic matter and 9000 shells for the dark matter.
+ Models. with 45000 shells for each component showed no considerable dillerence in the results.," Models with $45\, 000$ shells for each component showed no considerable difference in the results."
+ We also test our code. with good agreement. against he one-zone code of de Araujo Όρμος (1991) for he collapse of barvonic anc non-barvonic dark matter Population LLL objects (positive density perturbation).," We also test our code, with good agreement, against the one-zone code of de Araujo Opher (1991) for the collapse of baryonic and non-baryonic dark matter Population III objects (positive density perturbation)."
+ Also. as an additional test we reproduce the results of de Araujo Opher (1993) for voids formecl from negative density »rturbations. where a one-zone model was usec.," Also, as an additional test we reproduce the results of de Araujo Opher (1993) for voids formed from negative density perturbations, where a one-zone model was used."
+" The orogram conserves mass and total οποίον of the svstem to at least one part in 10"".", The program conserves mass and total energy of the system to at least one part in $10^6$.
+ In the present paper we use dillerent power spectra of density perturbations to study their role on the scale of voids formed. and to see if cosmological parameters could be constrained., In the present paper we use different power spectra of density perturbations to study their role on the scale of voids formed and to see if cosmological parameters could be constrained.
+ First we take a typical CDAL power spectrum of perturbations given by (Pacdmanabhan 1993) where ο is the spectrum amplitude. 2=L37(057).àMpc. (€=90(0h7)?“Alpe? and D=1.0(ON7)“Alpe7.," First we take a typical CDM power spectrum of perturbations given by (Padmanabhan 1993) where $A$ is the spectrum amplitude, $B=1.7(\Omega
+h^2)^{-1}{\rm Mpc}$, $C=9.0(\Omega h^2)^{-3/2}{\rm Mpc}^{3/2}$ and $D=1.0 (\Omega h^2)^{-2}{\rm Mpc}^{-2}$."
+ [n the above equation the power index is related to à by the relation n=(42a)., In the above equation the power index is related to $\alpha$ by the relation $n=(4-2\alpha)$.
+ The variance of Ductuations on scale ru reads We analvsed two cdillorent. window functions WArg): Gaussian and top-hat filters.," The variance of fluctuations on scale $r_0$ reads We analysed two different window functions $W(k\,r_0)$: Gaussian and top-hat filters."
+ Our results are very. weakly dependent on VW(GEru) so. for convenience and to compare our results with previous studies. we present here the results using à top-hat window function.," Our results are very weakly dependent on $W(k\,r_0)$ so, for convenience and to compare our results with previous studies, we present here the results using a top-hat window function."
+ Then. we have with the volume of the top-hat window given by The mass lluctuation in a radius £2 reads The input parameters that describe. the cosmological scenarios are taken from a recent study performed by Con (1998a.b). who analvsed mass ancl autocorrelation functions of rich clusters of galaxies from. linear. density," Then, we have with the volume of the top-hat window given by The mass fluctuation in a radius $R$ reads The input parameters that describe the cosmological scenarios are taken from a recent study performed by Cen (1998a,b), who analysed mass and autocorrelation functions of rich clusters of galaxies from linear density"
+This paper is organized as follows.,This paper is organized as follows.
+ Section 2 describes the data reduction., Section 2 describes the data reduction.
+ In Sect., In Sect.
+ 3 and 4 we describe the temporal and spectral analysis and in Sect., 3 and 4 we describe the temporal and spectral analysis and in Sect.
+ 5 we discuss our results., 5 we discuss our results.
+" The raw BAT survey data of the first 54 months of the mission were retrieved from the HEASARC public and processed with a dedicated software (?),, thatarchive)| performs screening, mosaicking and source detection on data and produces spectra and light curves for any given sky position."," The raw BAT survey data of the first 54 months of the mission were retrieved from the HEASARC public and processed with a dedicated software \citep{segreto10}, that performs screening, mosaicking and source detection on data and produces spectra and light curves for any given sky position."
+" JJ05007—7047, IGRJJ13186—-6257 and JJ17354—3255 were detected in the all-sky map with a significance maximized in the 15—50 keV band of 16.1, 10.0, and 18.9 standard deviations, respectively."," $-$ 7047, $-$ 6257 and $-$ 3255 were detected in the all-sky map with a significance maximized in the 15–50 keV band of 16.1, 10.0, and 18.9 standard deviations, respectively."
+ Light curves were extracted in the 15-150 keV energy range with the maximum available time resolution (300 s) and corrected to the solar system barycentre (SSB) by using the taskEARTH2SUN., Light curves were extracted in the 15–150 keV energy range with the maximum available time resolution $\sim 300$ s) and corrected to the solar system barycentre (SSB) by using the task.
+" The 15-150 keV spectra were obtained extracting the source fluxes from the sky maps in sixteen energy bands, and were analyzed using an appropriately rebinned CALDB response matrix."," The 15–150 keV spectra were obtained extracting the source fluxes from the sky maps in sixteen energy bands, and were analyzed using an appropriately rebinned CALDB response matrix."
+ observed JJ05007—7047 between 2010-10-25 and 2010-10-28 (ObsId 31846); IGRJJ13186—6257 was observed on 2008-05-01 and on 2010-05-25 (ObsId 37071); JJ17354—3255 was observed on 2008-03-11 and 2009- (ObsId 37054)., observed $-$ 7047 between 2010-10-25 and 2010-10-28 (ObsId 31846); $-$ 6257 was observed on 2008-05-01 and on 2010-05-25 (ObsId 37071); $-$ 3255 was observed on 2008-03-11 and 2009-04-17 (ObsId 37054).
+ The sources were all observed in Photon Counting mode., The sources were all observed in Photon Counting mode.
+ Source and background spectra and the relevant ancillary files were created and retrieved from the product on-line builder maintained by the data center at the University of Leicester (?).., Source and background spectra and the relevant ancillary files were created and retrieved from the product on-line builder maintained by the data center at the University of Leicester \citep{evans09}.
+" The spectra were rebinned to have at least 20 counts per bin, and energy channels were grouped to obtain a S/N ratio 7 3 per channel."," The spectra were rebinned to have at least 20 counts per bin, and energy channels were grouped to obtain a S/N ratio $\geq$ 3 per channel."
+ This choice allows us theuse of the X? statistics., This choice allows us theuse of the $\chi^2$ statistics.
+" In the following, we report errors at stated otherwise."," In the following, we report errors at stated otherwise."
+ Table [I] reports the log of the and observations used for the present analysis., Table \ref{log} reports the log of the and observations used for the present analysis.
+ We analyzed the long term light curve to search for intensity modulations by applying a folding technique and searching in the dd period range., We analyzed the long term light curve to search for intensity modulations by applying a folding technique and searching in the d period range.
+" The period resolution is given by P?/(N ATgar), where N=16 is the number of trial profile phase bins and ATgar is the data span length (?).."," The period resolution is given by $^{2}/(N \,\Delta$ $_{\rm BAT})$, where $N=16$ is the number of trial profile phase bins and $\Delta$ $_{\rm BAT}$ is the data span length \citep{buccheri85}."
+" The average rate in each phase bin was evaluated by weighting the light curve rates by the inverse square of the corresponding statistical error where R is the average rate in the j-th phase bin of the trial profile, r; are the rate of the light curve whose phase fall into the j-th phase bin and er; are the corresponding statistical errors."," The average rate in each phase bin was evaluated by weighting the light curve rates by the inverse square of the corresponding statistical error where $R_j$ is the average rate in the j-th phase bin of the trial profile, $r_i$ are the rate of the light curve whose phase fall into the j-th phase bin and $er_i$ are the corresponding statistical errors."
+ The error on R; is (Jz1lei).," The error on $R_j$ is $\left( \sqrt{\sum{1/er_i^2}}
+\right)^{-1}$."
+ The weighting procedure was adopted to deal with the large span of er; and it is justified because the BAT data are background dominated., The weighting procedure was adopted to deal with the large span of $er_i$ and it is justified because the BAT data are background dominated.
+" The significance of a feature P, in the resulting periodogram cannot be evaluated using the Xx? statistics because, if the source is variable on time scales comparable to those under investigation, the average X? is far from the value expected for white noise (N—1)."," The significance of a feature $_o$ in the resulting periodogram cannot be evaluated using the $\chi^2$ statistics because, if the source is variable on time scales comparable to those under investigation, the average $\chi^2$ is far from the value expected for white noise $(N-1)$."
+" Therefore, the significance of P, is evaluated as follows (??):: 1."," Therefore, the significance of $_o$ is evaluated as follows \citep{laparola10,cusumano10}: : 1."
+" We fit each periodogram with a second order polynomial and subtract the trend to the y? distribution, obtaining a new Xe 2."," We fit each periodogram with a second order polynomial and subtract the trend to the $\chi^2$ distribution, obtaining a new $\chi^2_C$ 2."
+" We build the histogram of the new Xe distribution excluding the interval around P, and those around any features deriving from it (e.g. periods multiple of Ρο). 3.", We build the histogram of the new $\chi^2_C$ distribution excluding the interval around $_o$ and those around any features deriving from it (e.g. periods multiple of $_o$ 3.
+ We fit the tail of the resulting distribution with an exponential function and evaluate the integral of the best-fit function beyond Xe(Po)., We fit the tail of the resulting distribution with an exponential function and evaluate the integral of the best-fit function beyond $\chi^2_C(\rm P_o)$.
+ This integral yields the number of chance occurrences due to noise., This integral yields the number of chance occurrences due to noise.
+ Figure [μα shows the periodogram obtained for JJO5007—7047., Figure \ref{period1}a a shows the periodogram obtained for $-$ 7047.
+" We find significant evidence for the presence of a periodicity (y?~210) at a period of Po=30.77+ 0.03dd, where the period uncertainty is the period resolution at P,."," We find significant evidence for the presence of a periodicity $\chi^2\sim 210$ ) at a period of $P_{0}=30.77\pm0.03$ d, where the period uncertainty is the period resolution at $_o$."
+" The other two significant features in the periodogram correspond to 2P, and 3P,.", The other two significant features in the periodogram correspond to $_o$ and $_o$.
+" The value of Xe (i.e. x? after subtracting the periodogram trend) at P, is 177.", The value of $\chi^2_C$ (i.e. $\chi^2$ after subtracting the periodogram trend) at $_o$ is 177.
+" Figure [ipb shows the Xe distribution excluding the values around P, and its multiples.", Figure \ref{period1}b b shows the $\chi^2_C$ distribution excluding the values around $_o$ and its multiples.
+" We then fitted the tail Xe>15) of the resulting distribution with an exponential function and evaluated the integral of the best-fit function beyond P,.", We then fitted the tail $\chi^2_C >15$ ) of the resulting distribution with an exponential function and evaluated the integral of the best-fit function beyond $_o$.
+" The number of chance occurrences to have a value of 177 or higher is 3.9x10:13, that corresponds to ~7 standard deviations in Gaussian statistics."," The number of chance occurrences to have a value of 177 or higher is $3.9\times 10^{-12}$, that corresponds to $\sim 7$ standard deviations in Gaussian statistics."
+ The pulsed profile (Fig., The pulsed profile (Fig.
+" [ῆρο) folded at P, with Tepoch=54159.753 MJD, shows a roughly sinusoidal modulation with a minimum value consistent with zero intensity."," \ref{period1}c c) folded at $_o$ with $_{\rm epoch}= 54159.753$ MJD, shows a roughly sinusoidal modulation with a minimum value consistent with zero intensity."
+" The centroid of the minimum, evaluated by fitting the data around the dip with a Gaussian model is at phase 0.23."," The centroid of the minimum, evaluated by fitting the data around the dip with a Gaussian model is at phase 0.23."
+" The uncertainty depends on the rough binning of the folded light curve and we conservatively assume an error corresponding to one bin phase; the dip phase, therefore, corresponds to MJD (54166.8+1.0) nxP, MJD."," The uncertainty depends on the rough binning of the folded light curve and we conservatively assume an error corresponding to one bin phase; the dip phase, therefore, corresponds to MJD $(54166.8\pm 1.0)\pm n \times $ $_o$ MJD."
+ The phase coverage of the light curve bins shows that this dip is not due to an accidental under-sampling of the light curve at this phase., The phase coverage of the light curve bins shows that this dip is not due to an accidental under-sampling of the light curve at this phase.
+ We self-consistently checked this aspect also for the dips of JJ17354—3255 and JJ13186—6257., We self-consistently checked this aspect also for the dips of $-$ 3255 and $-$ 6257.
+" The three Swift-XRT observations were performed at the orbital phase intervals of 0.369—0.380, 0.428—0.432 and 0.469— respectively."," The three -XRT observations were performed at the orbital phase intervals of 0.369–0.380, 0.428–0.432 and 0.469--0.482, respectively."
+The 0.2-10 keV light curve shows a persistent emission with a variability of a factor ~5 and an average count rate of 0.058+0.004 count s~!.,The 0.2–10 keV light curve shows a persistent emission with a variability of a factor $\sim5$ and an average count rate of $\pm$ 0.004 count $^{-1}$ .
+" The source events selected in the three observations are 77, 37 and 509,"," The source events selected in the three observations are 77, 37 and 509,"
+Choudhuri. Schüsssler. Dikpati (1995) that an equatorward propagation is still possible in (his situation. if (he (ime scale of meridional circulation is shorter than the time scale ol diffusion between the lavers of a and velocity shear.,"Choudhuri, Schüsssler, Dikpati (1995) that an equatorward propagation is still possible in this situation, if the time scale of meridional circulation is shorter than the time scale of diffusion between the layers of $\alpha$ and velocity shear."
+ This opens up the possibility of buikling models of the solar cdvnamo in which we have a positive a near (he surface and a positive gradient of differential rotation at the bottom of the convection zone., This opens up the possibility of building models of the solar dynamo in which we have a positive $\alpha$ near the surface and a positive gradient of differential rotation at the bottom of the convection zone.
+ The meridional circulation has to play à very crucial role in such models in ensuring the desired behaviour., The meridional circulation has to play a very crucial role in such models in ensuring the desired behaviour.
+ While using Durnev8 double ring method. the signs of the magnetic field in the (wo rings have to be chosen such (hat there is a correspondence with the positive a situation.," While using Durney's double ring method, the signs of the magnetic field in the two rings have to be chosen such that there is a correspondence with the positive $\alpha$ situation."
+ With the double ring method also. we found that the diamo wave at the bottom of the convection zone propagates equatorward only when there is a strong meridional flow and propagates poleward when this flow is switched oll.," With the double ring method also, we found that the dynamo wave at the bottom of the convection zone propagates equatorward only when there is a strong meridional flow and propagates poleward when this flow is switched off."
+ In 822 we discuss the details of our model., In 2 we discuss the details of our model.
+ Then we go on to present our main results in 833., Then we go on to present our main results in 3.
+ Our conelusions are summarised in 844., Our conclusions are summarised in 4.
+ We assume axisvmmeltry in all our caleulations., We assume axisymmetry in all our calculations.
+" The magnetic and velocity fields can be written as where D and d respectively represent the toroidal aud poloidal components of the magnetic field: O is the angular velocity. and v,=(0,6,+yey is the meridional circulation."," The magnetic and velocity fields can be written as where $B$ and $A$ respectively represent the toroidal and poloidal components of the magnetic field; $\Omega$ is the angular velocity, and $\vf_p = v_r {\bf e}_r + v_{\theta} {\bf e}_{\theta}
+$ is the meridional circulation."
+ We substitute equations (1) and (2) in the induction equation, We substitute equations (1) and (2) in the induction equation
+The properties of galaxies in the Universe are determined bv the behaviour of both the dark matter and the baryonic material from which they are made.,The properties of galaxies in the Universe are determined by the behaviour of both the dark matter and the baryonic material from which they are made.
+ The dynamics of the cdark matter. which are determined by gravity alone. are now reasonably well understood.," The dynamics of the dark matter, which are determined by gravity alone, are now reasonably well understood."
+ N-body simulations (e.g. Davisetal. 1985)) provide an accurate description of the evolution of structure into the highlv non-linear regime where dark matter halos form (see. for example 19999).," N-body simulations (e.g. \pcite{defw}) ) provide an accurate description of the evolution of structure into the highly non-linear regime where dark matter halos form (see, for example )."
+ Xnalvtically. the Press-Schechter theory (Press&Schechter1974). predicts to within 50% the distribution of halo masses found in N-bocky simulations for a specified cosmology. whilst theoretically motivated fitting functions do even better (Sheth&Tormen.1999:Sheth.Mo 2000).," Analytically, the Press-Schechter theory \cite{PS74} predicts to within $\sim 50\%$ the distribution of halo masses found in N-body simulations for a specified cosmology, whilst theoretically motivated fitting functions do even better \cite{st,smt,arj00}."
+. Extensions ofthis theory (Bondοἱal.1991:BowerLacey&Cole1993) predict. with reasonable accuracy. the hierarchical build-up of halos through the mergers of smaller progenitors (sec. for example. Lacey&Cole1994:Somervilleetal. 2000)).," Extensions of this theory \cite{BKCE,rgb91,lc93}
+ predict, with reasonable accuracy, the hierarchical build-up of halos through the mergers of smaller progenitors (see, for example, \pcite{lc94,rsetal98}) )."
+ The, The
+one.,one.
+ We find later. in our simulation results. that the inclination appears (to be the least important parameter for predicüng which objects will leave the scattered disk on gigavear timescales. suggesting that higher or lower inclinations do not much affect an SDOs orbital stability.," We find later, in our simulation results, that the inclination appears to be the least important parameter for predicting which objects will leave the scattered disk on gigayear timescales, suggesting that higher or lower inclinations do not much affect an SDO's orbital stability."
+ The debiasimg procedures described in Sections ?? and ?? were used to generate initial conditions for a simulated population of 1849 test particles., The debiasing procedures described in Sections \ref{ss:ae} and \ref{ss:i} were used to generate initial conditions for a simulated population of 1849 test particles.
+ The clones of individual observed objects were given tlie same eccentricitv. longitude of ascending node. and argument of perihelion as the original object.," The clones of individual observed objects were given the same eccentricity, longitude of ascending node, and argument of perihelion as the original object."
+ The semi-major axes of all the clones for a given observed object were evenly spread in à interval around the observed. value: this spread is similar to the observational uncertainties in this orbitalparameter?., The semi-major axes of all the clones for a given observed object were evenly spread in a interval around the observed value; this spread is similar to the observational uncertainties in this orbital.
+. The distributions of a. e. and q for the observed population and the simulated population are shown in Figure 1..," The distributions of $a$ , $e$, and $q$ for the observed population and the simulated population are shown in Figure \ref{f:dist}."
+ The inclinations were randomly selected from the best fit inclination distribution. and tlie mean anomaly lor each test particle was randomly selected from 0 to 23.," The inclinations were randomly selected from the best fit inclination distribution, and the mean anomaly for each test particle was randomly selected from 0 to $\pi$."
+ Ou debiased population can be compared to the debiased population of ? who modeled the scattered disk and Centaur populations., Our debiased population can be compared to the debiased population of \citet{disisto07} who modeled the scattered disk and Centaur populations.
+ The authors employ a similar debiasing procedure (ο our own for the semi-major axis distribution. with the addition that they superimpose a power law distribution of semi-major axes aller applying the correction factor for heliocentric distance.," The authors employ a similar debiasing procedure to our own for the semi-major axis distribution, with the addition that they superimpose a power law distribution of semi-major axes after applying the correction factor for heliocentric distance."
+ Their result is a semimajor axis distribution (hat is peaked in the 40—TO AU region. whereas oursstarts at 50 AU and is relatively flat in the 50-90 AU range.," Their result is a semimajor axis distribution that is peaked in the $40-70$ AU region, whereas oursstarts at $50$ AU and is relatively flat in the $50-90$ AU range."
+ While the classical disk does seem to follow a power law in heliocentric distance. (here are indications that the scattered populations heliocentric distance distribution is much more flat in the range of40—GO AU (?)..," While the classical disk does seem to follow a power law in heliocentric distance, there are indications that the scattered population's heliocentric distance distribution is much more flat in the range of$40-60$ AU \citep{kavelaars08}. ."
+ ?. also ealeulate a debiased inclination distribution:, \citet{disisto07} also calculate a debiased inclination distribution;
+SZ (Iuffeubereer&Seljak2001) (we included their estimate for the IR source power spectiuui at 115 Cz iu Figure 8)).,SZ \citep{Huffenberger:2004gm} (we included their estimate for the IR source power spectrum at 145 GHz in Figure \ref{Cl_contribs}) ).
+ Ou the other hand the assumption of perfect cleaning of the thermal SZ effect is safer because we understand its frequeney dependence well., On the other hand the assumption of perfect cleaning of the thermal SZ effect is safer because we understand its frequency dependence well.
+ To compare the expertueutal coustraints with the power spectra extracted from our simulations on degree patches on the sky. we bin the errors iuto bands of width Al=360.," To compare the experimental constraints with the power spectra extracted from our simulations on degree patches on the sky, we bin the errors into bands of width $\Delta l=360$."
+ The predicted measurcients errors are plotted in Figure 9.. together with power spectra generated from tje “cleaned” maps of primordial CMD and kinetic SZ ος1ubined.," The predicted measurements errors are plotted in Figure \ref{distinguish-pow}, together with power spectra generated from the “cleaned” maps of primordial CMB and kinetic SZ combined."
+ Since the kinetic SZ appears almost featureless on the scales accessible to the next eenecration of exeliments. relonizatiou models can be distinguished by heir average amplitude on these scales.," Since the kinetic SZ appears almost featureless on the scales accessible to the next generation of experiments, reionization models can be distinguished by their average amplitude on these scales."
+ The error bars in the relevant band-powers of ACT are colmbined using 024=(Soa;2ji, The error bars in the relevant band-powers of ACT are combined using $\sigma^2_{tot}=({\sum_i \sigma^{-2}_i})^{-1}$.
+ Tt follows that the overall aurplitude cau be measured with an accuracy of gap=Ομ (we assumed a 14 calibration uncertaintv).," It follows that the overall amplitude can be measured with an accuracy of $\sigma_{\Delta
+T}=0.011 \mu K$ (we assumed a $1\%$ calibration uncertainty)."
+ Given that the plateau of the extended patchy model (Model C) lies at AT~2.107448. while Model A has an average amplitude on these scales of ATz2.0761. the two will be distinguishable at the 300 level with ACT. if we ignore contamination ly volt sources and by the thermal SZ. which is uot a realistic asuniptià.," Given that the plateau of the extended patchy model (Model C) lies at $\Delta T \simeq 2.407 \mu K$, while Model A has an average amplitude on these scales of $\Delta T \simeq 2.076 \mu K$, the two will be distinguishable at the $30 \sigma$ level with ACT, if we ignore contamination by point sources and by the thermal SZ, which is not a realistic assumption."
+ This is the same for the South Pole Telescope. which as better angular resolution. but will ook less deep (it js a much larecr survey area).," This is the same for the South Pole Telescope, which has better angular resolution, but will look less deep (it has a much larger survey area)."
+ It is also of interest to ask whether additional constraints on relouization scenarios could be achieved * usns non-Cuusan statistics. in particular the four mt function (tle skewness arisine from a," It is also of interest to ask whether additional constraints on reionization scenarios could be achieved by using non-Gaussian statistics, in particular the four point function (the skewness arising from a"
+shown in the micelle panel of figure 1..,shown in the middle panel of figure \ref{allFig}.
+" This wind is comprised of the 6 trajectories defined in Table 2 of Paper 1, Ii these outflows. neutrinos convert [ree protons into neutrons."," This wind is comprised of the 6 trajectories defined in Table 2 of Paper I. In these outflows, neutrinos convert free protons into neutrons."
+ Capture of these neutrons by (n.p) reactions on long-lived nuclei along the path of the rp-process accelerate (he progress to nuclei as heavy as Pd.," Capture of these neutrons by (n,p) reactions on long-lived nuclei along the path of the $rp$ -process accelerate the progress to nuclei as heavy as Pd."
+ Final conditions lor these wind trajectories are provided in table 2.., Final conditions for these wind trajectories are provided in table \ref{5512570_final}.
+ We defer a discussion of the nuclear flows and the potential for production of heavier elements to section 3.., We defer a discussion of the nuclear flows and the potential for production of heavier elements to section \ref{broaderStudy}.
+ For comparison. we also show. in ligure 2.. the integrated results (summed over the same six wind trajectories) of caleulations that do include neutrino capture on protons.," For comparison, we also show, in figure \ref{allFigNoNeutrinos}, the integrated results (summed over the same six wind trajectories) of calculations that do include neutrino capture on protons."
+ Apart from the neglect of neutrinos the trajectories studied here are identical to those described in (he middle panel of figure 1.., Apart from the neglect of neutrinos the trajectories studied here are identical to those described in the middle panel of figure \ref{allFig}.
+ The differences are dramatic., The differences are dramatic.
+ When neutrino captures are ignored nucleosynthesis stops some 40 mass units lower than when thev are included., When neutrino captures are ignored nucleosynthesis stops some 40 mass units lower than when they are included.
+ The calculations in (he previous section are quite uncertain owing to both an uncertainty in (he supernova explosion model and to our procedure for extrapolating expansion alter the explosion calculation has stopped., The calculations in the previous section are quite uncertain owing to both an uncertainty in the supernova explosion model and to our procedure for extrapolating expansion after the explosion calculation has stopped.
+ Modest alterations are reasonable to the asvinptotic wind velocity. the electron fraction of the wind. ancl neutrino capture rates in (he outflow.," Modest alterations are reasonable to the asymptotic wind velocity, the electron fraction of the wind, and neutrino capture rates in the outflow."
+ More extreme changes (hat might reflect the influence of novel physical processes operating in theearly wind are explored in section 3.., More extreme changes that might reflect the influence of novel physical processes operating in theearly wind are explored in section \ref{broaderStudy}. .
+observed 2-10keV X-ray luminosity (as in Fig. 5)),observed 2-10keV X-ray luminosity (as in Fig. \ref{fig:lutz}) )
+ will help us avoid model assumptions about AGN structure and the AGN environment (see refsec:discuss below)., will help us avoid model assumptions about AGN structure and the AGN environment (see \\ref{sec:discuss} below).
+ The observed IR band emission from AGN should consist of some fraction of reprocessed radiation from the AGN typically peaking in the mid-IR (e.g. Efstathiouetal. (2000))) and some fraction of emission from star formation typically peaking in the far-IR (e.g. Efstathiou&Rowan- (2008))) plus intrinsic IR from the SED.," The observed IR band emission from AGN should consist of some fraction of reprocessed radiation from the AGN typically peaking in the mid-IR (e.g. \citet{b25}) ) and some fraction of emission from star formation typically peaking in the far-IR (e.g. \citet{b26,b31,b29}) ) plus intrinsic IR from the SED."
+" Therefore, by comparing the observed mid- and far-IR luminosities in group 1 and group 2 AGN, we may be able to distinguish dominant components of IR in different AGN."," Therefore, by comparing the observed mid- and far-IR luminosities in group 1 and group 2 AGN, we may be able to distinguish dominant components of IR in different AGN."
+" Figure 6 shows the mean far-IR (100um) luminosity plotted against the mean mid-IR (12um) luminosity for the 225 AGN in our sample without jets, for which both mid-IR and far-IR, data were available."," Figure \ref{fig:ir12v100} shows the mean far-IR $\micron$ ) luminosity plotted against the mean mid-IR $\micron$ ) luminosity for the 225 AGN in our sample without jets, for which both mid-IR and far-IR data were available."
+" 115 AGN are group 1 (black circles), 110 are group 2 (red crosses)."," 115 AGN are group 1 (black circles), 110 are group 2 (red crosses)."
+" Also plotted are lines denoting constant luminosity ratios of 10,1,0.1 respectively."," Also plotted are lines denoting constant luminosity ratios of $10,1,0.1$ respectively."
+ The populations of group 1 and group 2 AGN are very well separated by the luminosity ratio Liooum/Li2um=1., The populations of group 1 and group 2 AGN are very well separated by the luminosity ratio $L_{100\micron}/L_{12\micron}=1$.
+ Only 22/110 (~ 20%) of group 2 have Liooum/Lizum 1., Only 22/110 $\sim 20\%$ ) of group 2 have $L_{100\micron}/L_{12\micron}<1$ .
+ The three outlier group 2 AGN in this plot near the Liooum/Lizum=0.1 line are nearby LINERs that could be cross-classified., The three outlier group 2 AGN in this plot near the $L_{100\micron}/L_{12\micron}=0.1$ line are nearby LINERs that could be cross-classified.
+ From Fig., From Fig.
+ 6 we can say that the Liooum/ ratios for group 2 AGN are on average an order of magnitudeLizum larger than those for group 1 AGN., \ref{fig:ir12v100} we can say that the $L_{100\micron}/L_{12\micron}$ ratios for group 2 AGN are on average an order of magnitude larger than those for group 1 AGN.
+ The dispersion in both groups is also similar (roughly an order of magnitude)., The dispersion in both groups is also similar (roughly an order of magnitude).
+" The ratio also appears to neatly separate out the ’archetypal’ Liooum/Li»umGroup 2 AGN, in order of most ’cold’ (100um) dust to least ’cold’ dust."," The $L_{100\micron}/L_{12\micron}$ ratio also appears to neatly separate out the 'archetypal' Group 2 AGN, in order of most 'cold' $\micron$ ) dust to least 'cold' dust."
+" Evidently the starburst Arp 220 (open star) has by far the highest value of L1ooum/Li2um (as we should expect e.g. Efstathiouetal.(2000);Mas-Hesse (2008))) followed by the Compton-thick AGN seen in transmission (NGC 4945- open square), Circinus (open diamond) and NGC 1068 (open triangle) respectively."," Evidently the starburst Arp 220 (open star) has by far the highest value of $L_{100\micron}/L_{12\micron}$ (as we should expect e.g. \citet{b25,b29}) ) followed by the Compton-thick AGN seen in transmission (NGC 4945- open square), Circinus (open diamond) and NGC 1068 (open triangle) respectively."
+ Fig., Fig.
+ 3 demonstrates that the luminosity ratio (Rip/.) of the 245 AGN in our sample is not a scatterplot., \ref{fig:irx12} demonstrates that the luminosity ratio $R_{ir/x}$ ) of the 245 AGN in our sample is not a scatterplot.
+" So, are the unoccupied regions of the RH;,;, plot genuinely depopulated, providing a model-independent constraint on AGN structure, or is there a selection effect at work?"," So, are the unoccupied regions of the $R_{ir/x}$ plot genuinely depopulated, providing a model-independent constraint on AGN structure, or is there a selection effect at work?"
+" First, flux limits from all-sky surveys in both the IR band and the X-ray band allow us to establish which regions of Rip, are probably genuinely unoccupied by AGN."," First, flux limits from all-sky surveys in both the IR band and the X-ray band allow us to establish which regions of $R_{ir/x}$ are probably genuinely unoccupied by AGN."
+ Fig., Fig.
+ 7 is as Fig., \ref{fig:fluxlimits} is as Fig.
+" 3 except superimposed are flux detectability thresholds from (solid line) the ROSAT All-sky survey (2x10P? ergs cm?s! from Voges (1993)) and the IRAS 12um complete sample (ce 0.3Jy from Sandersetal. (2003))), corresponding to luminosity distances of 10Mpc,100Mpc and 1000Mpc."," \ref{fig:irx12} except superimposed are flux detectability thresholds from (solid line) the ROSAT All-sky survey $2 
+\times 10^{-13}$ ergs $\rm{cm}^{-2} \rm{s}^{-1}$ from \citet{b95}) ) and the IRAS $\micron$ complete sample $\sim 0.3$ Jy from \cite{b60}) ), corresponding to luminosity distances of 10Mpc,100Mpc and 1000Mpc."
+ The ROSAT survey was carried out in soft X-rays (0.1-2.4keV) and will undercount heavily absorbed AGN (ie Group 2 AGN) in this band., The ROSAT survey was carried out in soft X-rays (0.1-2.4keV) and will undercount heavily absorbed AGN (ie Group 2 AGN) in this band.
+" So, since Group 1 AGN are not very heavily absorbed in the soft X-ray band and since 115/116 of the Group 1 AGN in our sample are detectable in the 2-10keV band at ROSAT flux levels, the ROSAT limits are a reasonable proxy for X-ray flux limits for Group 1 AGN."," So, since Group 1 AGN are not very heavily absorbed in the soft X-ray band and since 115/116 of the Group 1 AGN in our sample are detectable in the 2-10keV band at ROSAT flux levels, the ROSAT limits are a reasonable proxy for X-ray flux limits for Group 1 AGN."
+" This is the best that we can do since surveys overlapping the 2-10keV band, are neither all-sky nor complete (Revnitsevetal.2004;Tueller 2008).."," This is the best that we can do since surveys overlapping the 2-10keV band, are neither all-sky nor complete \citep{b98,b37}. ."
+We shall assume that the GRB-huninosity funetion is redshift independent. i.e. without cosmological evolution of the nature of ils progenitors.,"We shall assume that the GRB-luminosity function is redshift independent, i.e., without cosmological evolution of the nature of its progenitors."
+ We take a log-normal probability density for the luminosity shape-Iunction. with mean j£ and width σ given by are ," We take a log-normal probability density for the luminosity shape-function, with mean $\mu$ and width $\sigma$ given by where $L$ is normalized with respect to $^{-2}$ $^{-1}$."
+This notation means (hat the estimated parameters can be either (122.3.4). (123.3.2). (125.2.8). or (126.2.6). but not (122.2.6) for instance.," The optimized parameters of our model are This notation means that the estimated parameters can be either (122,3.4), (123,3.2), (125,2.8), or (126,2.6), but not (122,2.6) for instance."
+ Our results are compatible to the expectations of Sethi et Dhargavi (2001) who derive a log-normal Iuminosity ΠΙΟΠΟ with ιτ129 and στ2 from a different. flux-Hmited sample., Our results are compatible to the expectations of Sethi et Bhargavi (2001) who derive a log-normal luminosity function with $\mu = 129$ and $\sigma = 2$ from a different flux-limited sample.
+ They mentioned that because of selection effects. the inferred average luminosity over-estimates (he true mean Iuminosityv by a factor of 2 or 3. and the variance is within of the true variance.," They mentioned that because of selection effects, the inferred average luminosity over-estimates the true mean luminosity by a factor of 2 or 3, and the variance is within of the true variance."
+ The observed and predicted. redshift distribution. based on the SER. are shown in Fie.," The observed and predicted redshift distribution, based on the SFR, are shown in Fig."
+ lin case of optimal parameters (7))., 1 in case of optimal parameters \ref{EQN_PAR}) ).
+ The results indicate a good fit to the data. suggesting that selection effects are aclequaly modeled.," The results indicate a good fit to the data, suggesting that selection effects are adequaly modeled."
+ The fraction ofdetectable GRBs as a Iunction of redshift shows a the steep decrease as the luminosity Chreshold increases. making high-redshift GRBs less likely to be detected.," The fraction of GRBs as a function of redshift shows a the steep decrease as the luminosity threshold increases, making high-redshift GRBs less likely to be detected."
+ A satisfactory test [or our model is provided by comparing the predicted (his distribution of GRBs above the sensitivity threshold with the peak flux distribution of a sample of G7 GRBs observed with IPN (Fig., A satisfactory test for our model is provided by comparing the predicted flux distribution of GRBs above the sensitivity threshold with the peak flux distribution of a sample of 67 GRBs observed with IPN (Fig.
+ 3)., 3).
+ The fluxes derived from our luminosity function in keV have been extrapolated to the IPN range of 25-100 keV. assuming an E? enerey spectrum and using (he formula given in Áppendix D of Sethi et Bhareavi (2001).," The fluxes derived from our luminosity function in 50-300 keV have been extrapolated to the IPN range of 25-100 keV, assuming an $\mathrm{E}^{-2}$ energy spectrum and using the formula given in Appendix B of Sethi et Bhargavi (2001)."
+ The conversion [actor from ergem7s! to photonem7s! has been taken to be 0.87x103. and the sensitivity threshold equal to 5 photonem7s! (Ibuley. private communication).," The conversion factor from $\mathrm{erg}\,\mathrm{cm}^{-2}\,\mathrm{s}^{-1}$ to $\mathrm{photon}\,\mathrm{cm}^{-2}\,\mathrm{s}^{-1}$ has been taken to be $0.87 \times 10^{-7}$, and the sensitivity threshold equal to 5 $\mathrm{photon}\,\mathrm{cm}^{-2}\,\mathrm{s}^{-1}$ (Hurley, private communication)."
+" A quantitative result provided by our model is the ratio L/f, of the tine GRB-event rate (what would be seen in case of a zero-fIux limit in the detector) and the observed GRB-event", A quantitative result provided by our model is the ratio $1/f_r$ of the true GRB-event rate (what would be seen in case of a zero-flux limit in the detector) and the observed GRB-event
+"The elobal HI enmüssiou profile of DDO210. obtained fro 11"".x data cube. is shown in Fie. Ll.","The global HI emission profile of DDO210, obtained from $''\times37''$ data cube, is shown in Fig. \ref{fig:HI_spec}."
+ A Cassia fit to the profile gives a central velocity (helioceutric) of 139.5€2.0|., A Gaussian fit to the profile gives a central velocity (heliocentric) of $-139.5 \pm 2.0$.
+ The integrated flux is 12.1c1.2 Jwl, The integrated flux is $12.1\pm1.2$ Jy.
+a These are duo good agreement with the values of 11042.0 τιaud 11.551.2 Ivkuns lobtained from sinele dish observations (IIuchtiuneier Richter 1986)., These are in good agreement with the values of $-140 \pm 2.0 $ and $11.5\pm1.2$ Jy obtained from single dish observations (Huchtmeier Richter 1986).
+ A eood agreement between the CAIRT flux aud the single dish flux shows that no flux was nüssed because of the missing short spacings in the interferometric observations., A good agreement between the GMRT flux and the single dish flux shows that no flux was missed because of the missing short spacings in the interferometric observations.
+ The velocity width at 50 level of peak emission (AVS ) is found to be 19.141. which again is in excellent agreement with the Αο value of 19.0 from the single dish observations.," The velocity width at 50 level of peak emission $\Delta V_{50}$ ) is found to be $19.1 \pm 1$, which again is in excellent agreement with the $\Delta V_{50}$ value of 19.0 from the single dish observations."
+ The IIT mass obtained from the integrated profile (taking the distance to the galaxy to be 1.0 Mpc) is 2840.3«&10937... aud the AyΡΕ ratio is found to be ~1.0 in solar units.," The HI mass obtained from the integrated profile (taking the distance to the galaxy to be 1.0 Mpc) is $2.8\pm0.3 \times{10}^{6} M_\odot$, and the $M_{\rm{HI}}/L_{\rm{B}}$ ratio is found to be $\sim 1.0$ in solar units."
+ Fie., Fig.
+" 2 shows the iuteerated ITI cussion fou DDO210 at L1""«37"" resolution. overlaved on the digitised kv survey (DSS) image."," \ref{fig:ov} shows the integrated HI emission from DDO210 at $''\times37''$ resolution, overlayed on the digitised sky survey (DSS) image."
+ The UD isodeusitv contours are elongated in easterü aud southern half of the galaxy i.e. the density is cuhanced in these directions., The HI isodensity contours are elongated in eastern and southern half of the galaxy i.e. the density is enhanced in these directions.
+ These deusity chhancements are highlighted in Fie., These density enhancements are highlighted in Fig.
+" Lo which shows the integrated ITE emission at hieh resolution (12""« 11"").", \ref{fig:ov_20} which shows the integrated HI emission at high resolution $''\times11''$ ).
+ Note that he faint extended chussion seeu in the low resolution ap is resolved out in this image., Note that the faint extended emission seen in the low resolution map is resolved out in this image.
+ The optical oenuüsson in the ealaxy shows two conrponeuts; a central bright compact component and an outer faint extended component: both elongated iu the east-west direction (Lee ct al.," The optical emission in the galaxy shows two components, a central bright compact component and an outer faint extended component; both elongated in the east-west direction (Lee et al."
+ 1999)., 1999).
+ However. the ceuters for the two conmponeuts do not overlap: the bright component les more castwards than the ceuter of extended componcut.," However, the centers for the two components do not overlap; the bright component lies more eastwards than the center of extended component."
+ From ai deep €-MD diagram of DDO?10. Tolstoyetal.(2000) found a unuunber of faint stars older than several Cr in the ealaxy.," From a deep C-M diagram of DDO210, \cite{tolstoy00} found a number of faint stars older than several Gyr in the galaxy."
+ vanZee(2000) estimated an inclination of 68 degrees for DDO210 from her observatious. Leeetal.(1999)," \cite{vanzee00} estimated an inclination of 68 degrees for DDO210 from her observations. \cite{lee99},"
+.. based onu deeper observations. estimated. the inclination of the aut extended coimpoueut to be 60 degrees. but noted iat the preseuce of several bright foreground stars 11akes Us estimate of iuclination very uncertain.," based on deeper observations, estimated the inclination of the faint extended component to be $\sim$ 60 degrees, but noted that the presence of several bright foreground stars makes this estimate of inclination very uncertain."
+ From a visual iuspectiou of the overlay (Fig. 2)).," From a visual inspection of the overlay (Fig. \ref{fig:ov}) ),"
+ the stern TT density cuhaucements seem to correlate with 160 optical ciuissiou in the galaxy., the eastern HI density enhancements seem to correlate with the optical emission in the galaxy.
+ As can be secu the melt component of the optical enmuissiou is not centered 3 the HI eiission. but is instead offse to the cast.," As can be seen the bright component of the optical emission is not centered on the HI emission, but is instead offset to the east."
+ Ou i6 other haud. no optical ciission is seen along the southern III deusitv enhancement.," On the other hand, no optical emission is seen along the southern HI density enhancement."
+ As a further check or optical enüssion associated with the southern III density euhlaucenieut. we computed. uuuber counts of stars iu that region using the deep optical image obtained w Tolstoyctal.(2000).," As a further check for optical emission associated with the southern HI density enhancement, we computed number counts of stars in that region using the deep optical image obtained by \cite{tolstoy00}."
+ No significant increasei in the star counts from the backerouud was detected., No significant increase in the star counts from the background was detected.
+ As seen in the Fig. 2..," As seen in the Fig. \ref{fig:ov},"
+ the faint extended: cussion in the lower resolution III distribution is less distorted by the presence of the IIT density culaucemeuts., the faint extended emission in the lower resolution HI distribution is less distorted by the presence of the HI density enhancements.
+" ence. an estimate of the morphological ceuter. position angele (PA) aud iuclinatiou of the ealaxy was obtained by fittingo elliptical aunuli to the ""o.31"" vosolution III distribution."," Hence, an estimate of the morphological center, position angle (PA) and inclination of the galaxy was obtained by fitting elliptical annuli to the $''\times37''$ resolution HI distribution."
+ The estimated IIT morphological ceuter lies close to the center ofthe faint extended optical component (but as noted above is offset from the ceuter of the xieh optical enüssion)., The estimated HI morphological center lies close to the center of the faint extended optical component (but as noted above is offset from the center of the bright optical emission).
+ The variation of the fitted norphological PA with the ealactoceutric radius is οἴνοιι in Fie. 5.., The variation of the fitted morphological PA with the galactocentric radius is given in Fig. \ref{fig:pa}.
+ The inclination estimated from the outer IHI contours (assuming the intrinsic shape of the IIT disk to ο circular) is 2747.0 degrees., The inclination estimated from the outer HI contours (assuming the intrinsic shape of the HI disk to be circular) is $\pm$ 7.0 degrees.
+ The inclination im the ΠΙΟ: regions of the galaxwv could not be coustraimec reliably due to the significant distortion in the IIT contours., The inclination in the inner regions of the galaxy could not be constrained reliably due to the significant distortion in the HI contours.
+ For he same reason. ellipse fitting to the higher resolution IIT data (where the smmooth emission is resolved out) is not reliable.," For the same reason, ellipse fitting to the higher resolution HI data (where the smooth emission is resolved out) is not reliable."
+ The inclination derived from the IIT distribution is significantly ciffercut from the optical inclination of the ealaxy., The inclination derived from the HI distribution is significantly different from the optical inclination of the galaxy.
+ It is likely that the optical eiissiou docs not arise from au axi-svaiunetric disk. but instead that the stars are concentrated in a patchy clongated region in DDO210.," It is likely that the optical emission does not arise from an axi-symmetric disk, but instead that the stars are concentrated in a patchy elongated region in DDO210."
+ Ileuce. in the absence of anv other reliable estimate. the value of inclination obtained from the outer III contours was assigned to the whole galaxy.," Hence, in the absence of any other reliable estimate, the value of inclination obtained from the outer HI contours was assigned to the whole galaxy."
+ The deprojected IIT radial surface density profiles were theu obtained by averaging over elliptical auuuli iu the plane of the galaxy., The deprojected HI radial surface density profiles were then obtained by averaging over elliptical annuli in the plane of the galaxy.
+" The profiles derived from the LI!&377 aud 29%ς23"" resolution Π distributions are eiven in Fig. 3..", The profiles derived from the $''\times37''$ and $''\times23''$ resolution HI distributions are given in Fig. \ref{fig:smd}.
+ As can be seen. the two distributions are in reasonable agreement.," As can be seen, the two distributions are in reasonable agreement."
+" The flux inteeral estimated from the 29""423"" resolution TT moment 0 map was fouud to be —11 smaller than hat estimated frou the E17&37"" resolutiou HT distribution.", The flux integral estimated from the $''\times23''$ resolution HI moment 0 map was found to be $\sim$ smaller than that estimated from the $''\times37''$ resolution HI distribution.
+ For the vest of the analysis. we will use the H7<37” III profile.," For the rest of the analysis, we will use the $''\times37''$ HI profile."
+ This profile is well represented by a gaussian. 1.6. We Lave," This profile is well represented by a gaussian, i.e. we have"
+the 1D extracted quartz lamp spectra. so we made a master flat for each spectral order for each observing night bv averaging the individually normalized 1D quartz spectra.,"the 1D extracted quartz lamp spectra, so we made a master flat for each spectral order for each observing night by averaging the individually normalized 1D quartz spectra."
+ For the red orders of interest here there is large aumplitude fringing in the ARCES spectra. which is parüallv removed from the stellar spectra by dividing bx the normalized master flats.," For the red orders of interest here there is large amplitude fringing in the ARCES spectra, which is partially removed from the stellar spectra by dividing by the normalized master flats."
+ There are slight (1 pixel) offsets between the positions of (he spectral orders in (he quartz ancl stellar frames. resulting in incomplete fringe removal.," There are slight $\sim$ 1 pixel) offsets between the positions of the spectral orders in the quartz and stellar frames, resulting in incomplete fringe removal."
+ Alodeling telluric H5O absorption in our A star spectra requires high-resolution telluric templates., Modeling telluric $_{2}$ O absorption in our A star spectra requires high-resolution telluric templates.
+ In ?. an observed atmospheric (rausmission spectrum. derived [rom solar observations bv ?.. was scaled to fit telluric absorption features due to CHI; around 2300 nm.," In \citet{blake2010} an observed atmospheric transmission spectrum, derived from solar observations by \citet{livingston1991}, was scaled to fit telluric absorption features due to $_{4}$ around 2300 nm."
+ We found (his approach to be problematic for absorption by Π.Ο between 500 nm and 1000 nm. so instead we created theoretical iransmission spectra.," We found this approach to be problematic for absorption by $_{2}$ O between 800 nm and 1000 nm, so instead we created theoretical transmission spectra."
+ Closely following 7— we developed a custom line-bv-Hine radiative transler code to calculate template spectra that contain only molecular absorption (no scattering. no line mixing. no IIO. conünuum absorption) given an atmospheric model and molecular transition data from the HITILAN 2008 database (?)..," Closely following \citet{gordley1994} we developed a custom line-by-line radiative transfer code to calculate template spectra that contain only molecular absorption (no scattering, no line mixing, no $_{2}$ O continuum absorption) given an atmospheric model and molecular transition data from the HITRAN 2008 database \citep{rothman2009}."
+" Following the Beer-Lambert law for extinction through a medium by a single transition of a single absorber. (he transmission. Z,. of the Earth's atmosphere at [requeney. v is given by whereo() is the lrequency-dependent cross section aud wis (e mass path where wv is (he path leneth through the atmosphere. 7? is pressure. { is temperature. his Bollziman’s constant. ancl q is the Volume Mixing Ration (VMIR). which is defined as the fraction of the particles in a volume that are the given constituent."," Following the Beer-Lambert law for extinction through a medium by a single transition of a single absorber, the transmission, $T_{\nu}$, of the Earth's atmosphere at frequency $\nu$ is given by where$\sigma(\nu)$ is the frequency-dependent cross section and $u$ is the mass path where $x$ is the path length through the atmosphere, $P$ is pressure, $T$ is temperature, $k$ is Boltzman's constant, and $q$ is the Volume Mixing Ration (VMR), which is defined as the fraction of the particles in a volume that are the given constituent."
+ Earth's atusosphere contains many gasses and (he properties of the atmosphere. including 2 aud T. change with altitude and time.," Earth's atmosphere contains many gasses and the properties of the atmosphere, including $P$ and $T$, change with altitude and time."
+ The transmission through the atmosphere can be calculated by breaking (he atmosphere up into a number of levels in the vertical direction. z. within which the bulk properties are assumed to be constant.," The transmission through the atmosphere can be calculated by breaking the atmosphere up into a number of levels in the vertical direction, $z$, within which the bulk properties are assumed to be constant."
+ The total (transmission due to a number of different absorbing gasses. g. each with a number of absorption lines. 7. can then be calculated The shapes and strengths of the absorption lines depend on the cross section. o(r).," The total transmission due to a number of different absorbing gasses, $g$, each with a number of absorption lines, $i$, can then be calculated The shapes and strengths of the absorption lines depend on the cross section, $\sigma({\nu})$ ."
+ Absorption lines in Earth's atmosphere have intrinsie shapes both due to the bulk properties, Absorption lines in Earth's atmosphere have intrinsic shapes both due to the bulk properties
+density profile which are slightly different.,density profile which are slightly different.
+" In FARGO, the disk density at the position of the inner planet is slightly higher compared with GENESIS while it is slightly lower at the position of the outer For calculations in which Eq."," In FARGO, the disk density at the position of the inner planet is slightly higher compared with GENESIS while it is slightly lower at the position of the outer For calculations in which Eq."
+" 3 is solved at each timestep, the time evolution of the period ratio for four realizations with y=6x10? and y=1.3x107* is displayed in Fig. 11.."," \ref{eq:damp} is solved at each timestep, the time evolution of the period ratio for four realizations with $\gamma=6\times 10^{-5}$ and $\gamma=1.3\times 10^{-4}$ is displayed in Fig. \ref{damp}."
+" In that case, all the realizations performed with y=6x10? resulted in the formation of the 3:2 resonance whereas for y=13x1075, two of the four realizations resulted in capture in that resonance by the end of the run."," In that case, all the realizations performed with $\gamma=6\times 10^{-5}$ resulted in the formation of the 3:2 resonance whereas for $\gamma=1.3\times 10^{-4}$, two of the four realizations resulted in capture in that resonance by the end of the run."
+" For y=1.3x107+, the time evolution of the resonant angles associated with the 3:2 resonance is displayed in Fig. 12.."," For $\gamma=1.3\times 10^{-4}$, the time evolution of the resonant angles associated with the 3:2 resonance is displayed in Fig. \ref{angleswithdamping}."
+" Compared with previous runs in which the surface density profile was altered by turbulence, we see that capture in resonance tends to occur later in runs where a roughly constant surface density profile is maintained."," Compared with previous runs in which the surface density profile was altered by turbulence, we see that capture in resonance tends to occur later in runs where a roughly constant surface density profile is maintained."
+" This occurs because the disk density at the position of the inner planet tends to be higher in runs where the initial disk surface density profile is restored, leading to a slower differential migration between the two planets."," This occurs because the disk density at the position of the inner planet tends to be higher in runs where the initial disk surface density profile is restored, leading to a slower differential migration between the two planets."
+" For other models, repeating the previously decribed procedure leads to analytical estimates of y?""=1.2x104 for model G2 and y?""~9.3x10? for model G4."," For other models, repeating the previously decribed procedure leads to analytical estimates of $\gamma_c^{ana}=1.2\times 10^{-4}$ for model G2 and $\gamma_c^{ana} \sim 9.3\times 10^{-5}$ for model G4."
+" Given that the simulations performed with GENESIS indicate that 1.3x107*<y,«1.9x107 and 6x10?«y,<1.31074 for models G2 and G4 respectively, we see that again the previous analytical estimates compare reasonably well with the results of our simulations."," Given that the simulations performed with GENESIS indicate that $1.3\times 10^{-4}\le\gamma_c<1.9\times 10^{-4}$ and $6\times 10^{-5}\le\gamma_c<1.3\times 10^{-4}$ for models G2 and G4 respectively, we see that again the previous analytical estimates compare reasonably well with the results of our simulations."
+" For systems with g=1/2, the results of the simulations indicate that the 3:2 resonance can be maintained only in cases where the disk is close to being inviscid."," For systems with $q=1/2$, the results of the simulations indicate that the 3:2 resonance can be maintained only in cases where the disk is close to being inviscid."
+ Fig., Fig.
+" 13. shows the results for modelG5 in which m;=1.6 Mg and m,=33 Mg and for a single realization of the different values of y we considered.", \ref{run5} shows the results for modelG5 in which $m_i=1.6$ $M_\oplus$ and $m_o=3.3$ $M_\oplus$ and for a single realization of the different values of $\gamma$ we considered.
+" Moving from left to right and from top to bottom the panels display the time evolution of the planets’ semi-major axes, eccentricity of the inner planet, eccentricity of the outer one, and period ratio."," Moving from left to right and from top to bottom the panels display the time evolution of the planets' semi-major axes, eccentricity of the inner planet, eccentricity of the outer one, and period ratio."
+" The evolution of semi-major axes shows strong similarities with that of model Gl, with the migration rates of the planets observed to decrease with increasing the value for y, as discussed in Sect. 4.1.1.."," The evolution of semi-major axes shows strong similarities with that of model G1, with the migration rates of the planets observed to decrease with increasing the value for $\gamma$, as discussed in Sect. \ref{sec:orbit}."
+" Because of stochastic density fluctuations, the planets eccentricities are highly variables quantities and a clear trend of higher eccentricities for higher values of y is again observed at the beginning of the simulations."," Because of stochastic density fluctuations, the planets eccentricities are highly variables quantities and a clear trend of higher eccentricities for higher values of $\gamma$ is again observed at the beginning of the simulations."
+" For y=0, the time evolution of the period ratio shows that trapping in the 3:2 resonance occurs at {~10° orbits."," For $\gamma=0$, the time evolution of the period ratio shows that trapping in the 3:2 resonance occurs at $t\sim 10^3$ orbits."
+" This resonant interaction causes eccentricity growth of both planets with the eccentricities of the inner and outer planets reaching peak values of e;~0.035 and e,~0.02 respectively, although convergence is not fully established at the end of the simulation."," This resonant interaction causes eccentricity growth of both planets with the eccentricities of the inner and outer planets reaching peak values of $e_i\sim 0.035$ and $e_o\sim 0.02$ respectively, although convergence is not fully established at the end of the simulation."
+" Comparing Figs 3 and 13,, we see that the period ratio oscillates with greater amplitude in model GS,"," Comparing Figs \ref{run1} and \ref{run5}, , we see that the period ratio oscillates with greater amplitude in model G5,"
+range.,.
+ Indeed. is known to be a K2 giant and. hence. the spectrum of this source will have a weak SiO band.," Indeed, is known to be a K2 giant and, hence, the spectrum of this source will have a weak SiO band."
+ Careful inspection of the 9.7 jim features in Fig., Careful inspection of the 9.7 $\mu$ m features in Fig.
+ 4. shows that the photospheric gas phase SiO band is indeed present in. all but one of the molecular sightlines as à blue shoulder between 7.5 and about 8.5 yam. Onlyforabsent. since T Is Zero at about 8 uim. where the SiO band peaks.," \ref{fig:10feats2} shows that the photospheric gas phase SiO band is indeed present in all but one of the molecular sightlines as a blue shoulder between 7.5 and about 8.5 $\mu$ m. Onlyfor, since $\tau$ is zero at about 8 $\mu$ m, where the SiO band peaks."
+ So for this source. the observed difference with the Galactic feature be explained by the presence of a gas phase SiO band in the intrinsic spectrum of the background star.," So for this source, the observed difference with the Galactic feature be explained by the presence of a gas phase SiO band in the intrinsic spectrum of the background star."
+ However. Fig.," However, Fig."
+ 12. shows that the difference in the shape of the 9.7 um silicate features. 1n most cases.cannot be explained only by SiO. This is because gas phase SiO normally peaks at about 8 iim. whereas Fig.," \ref{fig:mol_galcent} shows that the difference in the shape of the 9.7 $\mu$ m silicate features, in most cases, be explained only by SiO. This is because gas phase SiO normally peaks at about 8 $\mu$ m, whereas Fig."
+ 12. shows features that peak at about ϱ-95 ium. except for J163346.2—242753.," \ref{fig:mol_galcent} shows features that peak at about 9–9.5 $\mu$ m, except for $-$ 242753."
+" Since the spectral type of is known we were able to extract the photospherie gas phase SiO band for this source using a model atmosphere spectrum calculated by (Decinetal..1997,2004). based on MARCS models (Gustafssonetal..1975.2008) (see Fig. 13.."," Since the spectral type of is known we were able to extract the photospheric gas phase SiO band for this source using a model atmosphere spectrum calculated by \citep{Decin97,2004ApJS..154..408D} based on MARCS models \citep{Gustafsson75,2008A&A...486..951G} (see Fig. \ref{fig:K2III},"
+ upper panel)., upper panel).
+ The photospheric gas phase SiO band in terms of optical depth was extracted from this model spectrum by taking a power-law continuum from 7.4 to 12.3 jm and using equation I (see Fig. 13..," The photospheric gas phase SiO band in terms of optical depth was extracted from this model spectrum by taking a power-law continuum from 7.4 to 12.3 $\mu$ m and using equation 1 (see Fig. \ref{fig:K2III},"
+ lower panel)., lower panel).
+ Subsequently. we subtracted this optical depth profile from the observed 9.7 yam feature.," Subsequently, we subtracted this optical depth profile from the observed 9.7 $\mu$ m feature."
+ The correction never exceeds 0.03 in τ., The correction never exceeds 0.03 in $\tau$.
+ The result is shown in Fig. 14.., The result is shown in Fig. \ref{fig:testElias13}.
+ Correcting for SiO in this spectrum is not sufficient to explain the difference in the shape of the 9.7 jim silicate feature between and the Galactic Centre., Correcting for SiO in this spectrum is not sufficient to explain the difference in the shape of the 9.7 $\mu$ m silicate feature between and the Galactic .
+ Different types of ices could affect the shape of the 9.7 jum silicate feature in molecular sightlines, Different types of ices could affect the shape of the 9.7 $\mu$ m silicate feature in molecular sightlines
+z=2 to z=0.,$z=2$ to $z=0$.
+" From this plot we can see that early on the stars are centrally concentrated, but a significant core starts developing at a redshift z=1."," From this plot we can see that early on the stars are centrally concentrated, but a significant core starts developing at a redshift z=1."
+" The core seems to be in place by a redshift z=0.6, or about 3 Gyrs, which is a cluster crossing timescale, but about 100 crossing times at the half light radius of 10 kpc."," The core seems to be in place by a redshift z=0.6, or about 3 Gyrs, which is a cluster crossing timescale, but about 100 crossing times at the half light radius of 10 kpc."
+" The core in the AGN-ON BCG doesn’t form naturally, the AGN-OFF simulation has no such effect."," The core in the AGN-ON BCG doesn't form naturally, the AGN-OFF simulation has no such effect."
+ Something has caused stars to move away from the inner region and this requires a significant amount of energy., Something has caused stars to move away from the inner region and this requires a significant amount of energy.
+shear are computed in Section 4. the contribution of ealaxy-ealaxy lensing (o cosmic shear is computed in Section 5. and a discussion of cosmic variance in relation to the field size in presented in Section 6.,"shear are computed in Section 4, the contribution of galaxy-galaxy lensing to cosmic shear is computed in Section 5, and a discussion of cosmic variance in relation to the field size in presented in Section 6."
+ The conclusions are summarized in Section T., The conclusions are summarized in Section 7.
+ To investigate (he frequency ancl effects of multiple weak deflections in galaxv-galaxy lensing. Monte Carlo simulations are constructed.," To investigate the frequency and effects of multiple weak deflections in galaxy-galaxy lensing, Monte Carlo simulations are constructed."
+ The lens galaxies in the Monte Carlo simulations are relatively bright galaxies (hat are contained within a circle of radius 4 arcmin. centered on the Lhibble Deep Field North (IIDF-N) (Williams et 11996).," The lens galaxies in the Monte Carlo simulations are relatively bright galaxies that are contained within a circle of radius 4 arcmin, centered on the Hubble Deep Field North (HDF-N) (Williams et 1996)."
+ This region of sky was (he subject of a deep redshift survey (Cohen et 11996: Steidel et 11996: Lowenthal et 11997: Phillips et 11997: Cohen οἱ 22000) as well as an extensive multicolor photometric investigation (llogg et 22000)., This region of sky was the subject of a deep redshift survey (Cohen et 1996; Steidel et 1996; Lowenthal et 1997; Phillips et 1997; Cohen et 2000) as well as an extensive multicolor photometric investigation (Hogg et 2000).
+ As a result. both the redshifts (Cohen et 22000. Tables 24 and 2B) and rest-frame blue huninosities. Lj; (Cohen 2001. Table 1). of 590 galaxies in this region of the skv are known.," As a result, both the redshifts (Cohen et 2000, Tables 2A and 2B) and rest-frame blue luminosities, $L_B$ (Cohen 2001, Table 1), of 590 galaxies in this region of the sky are known."
+ For the simulations. then. the locations of the lenses in redshift space are known verv accurately. and the relative strengths of the different lenses can be inferred quite well from their relative huminosities.," For the simulations, then, the locations of the lenses in redshift space are known very accurately, and the relative strengths of the different lenses can be inferred quite well from their relative luminosities."
+ Therelore. it is possible to make detailed theoretical predictions for the weak galaxv-galaxy. lensing shear field that should be expected in this region of sky.," Therefore, it is possible to make detailed theoretical predictions for the weak galaxy-galaxy lensing shear field that should be expected in this region of sky."
+" For simplicity. the dark matter halos of the lens galaxies are taken {ο have a mass density given bv where c, is the velocity dispersion of the halo. G is Newtons constant. and s is a characteristic halo radius (see e.g.. BBS: Hudson et 11998: Fischer et 22000: IHIoekstra et 22004)."," For simplicity, the dark matter halos of the lens galaxies are taken to have a mass density given by where $\sigma_v$ is the velocity dispersion of the halo, $G$ is Newton's constant, and $s$ is a characteristic halo radius (see e.g., BBS; Hudson et 1998; Fischer et 2000; Hoekstra et 2004)."
+ It is then convenient to scale the depths of the potential wells of lens galaxies with differing Iuminosities. Lj. according to a Tullv-Fisher or Faber-Jackson (vpe of relation where σ is the velocity dispersion of the halo of a lens galaxy with rest-frame blue huninosity L5.," It is then convenient to scale the depths of the potential wells of lens galaxies with differing luminosities, $L_B$, according to a Tully-Fisher or Faber-Jackson type of relation where $\sigma_v^\ast$ is the velocity dispersion of the halo of a lens galaxy with rest-frame blue luminosity $L_B^\ast$."
+ Again for simplicity. it is assumed that the mass-to-light ratio of a galaxy is constant independent of its laminosity.," Again for simplicity, it is assumed that the mass-to-light ratio of a galaxy is constant independent of its luminosity."
+" Therefore. (he characteristic radii of the halos of galaxies with LyσεLj scale with the radii of the halos of L5, galaxies as"," Therefore, the characteristic radii of the halos of galaxies with $L_B \ne L_B^\ast$ scale with the radii of the halos of $L_B^\ast$ galaxies as"
+ introduced in the context of plasma physics hy Morrison(1980) defined by where 3.1 stands for the functional. derivative of sl. which is the Elinear part of lf|8f].Af].,"– introduced in the context of plasma physics by \citet{M80} – defined by where $\frac{\delta A}{\delta f}$ stands for the functional derivative of $A$, which is the linear part of $A\left[ f+\delta f\right] - A\left[ f\right]$."
+ One can easily obtain the Lamiltonian formulation of CBP system where £f] is any functional of f., One can easily obtain the Hamiltonian formulation of CBP system where $F\left[ f\right]$ is any functional of $f$.
+ The stability of equilibrium states Is a very old. problem of theoretical stellar dynamics. and a large variety of methocls has been used to tackle it.," The stability of equilibrium states is a very old problem of theoretical stellar dynamics, and a large variety of methods has been used to tackle it."
+ A clear consensus was found about the global stabilityof isotropic spherical svstems with distribution function fy(42) monotonically decreasing: after the pioneering works hy Antonov(1961)... linear. stability was obtained usingenergv methods after a long series of papers by WKanerup&Svenct(1985). (and. references within) or see also Perez&Aly(1996). for a comparison of the dilferent.results: using direct nornial mode techniques. complicated proofs were also obtained (Fridmann&Polv-achenko 1984: Palmer 1994)).," A clear consensus was found about the global stabilityof isotropic spherical systems with distribution function $f_{0} \left( E \right) $ monotonically decreasing: after the pioneering works by \citet{Antonov}, linear stability was obtained usingenergy methods after a long series of papers by \citet{KS} (and references within) or see also \citet{Perezaly1} for a comparison of the differentresults; using direct normal mode techniques, complicated proofs were also obtained \citealt{FetP}; \citealt{Palmer}) )."
+ Non-linear stabilitv of such spherical isotropic svstenis was also proven for some specific models (seoRein2002.andreferenceswithin).., Non-linear stability of such spherical isotropic systems was also proven for some specific models \citep[see][and references within]{Rein}.
+ The stability of anisotropic spherical svstemis is à more difficult problem., The stability of anisotropic spherical systems is a more difficult problem.
+ Ehe distribution function depends both on the one-particle energy Z£ and on the squared onc particle angular momentum L7:—ρα(ρα)., The distribution function depends both on the one-particle energy $E$ and on the squared one particle angular momentum $L^{2}:=\mathbf{p}^{2}\mathbf{q}^{2} - (\mathbf{p}\cdot\mathbf{q})^{2}$.
+ The most general result in this context was obtained by Perez&Alv(1996) and concerns linear stability for the restricted. case of preserving perturbations. which includes radially svnunietric ones.," The most general result in this context was obtained by \citet{Perezaly1} and concerns linear stability for the restricted case of preserving perturbations, which includes radially symmetric ones."
+" Non-linear stability is assured for. some classes of eeneralized. polvtropes for which fü(E.L)=E'L?"" with adapted values of & and. p (see Rein(2002) and references within)."," Non-linear stability is assured for some classes of generalized polytropes for which $f_{0}\left( E,L^{2}\right) = E^{k}L^{2p}$ with adapted values of $k$ and $p$ (see \citet{Rein} and references within)."
+ Some verv technical approaches using normal modes claim linear instability for anisotropic svstenis. composed only of radial orbits (sec Eridmann&Polvachenko 1984: Palmer 1994)): this is known as the racial orbit. instability., Some very technical approaches using normal modes claim linear instability for anisotropic systems composed only of radial orbits (see \citealt{FetP}; ; \citealt{Palmer}) ): this is known as the radial orbit instability.
+ See Merritt&Aguilar(1985) for one of the first relevant numerical approaches. Perezctal.(1996). for an intermediate position or Barnesetal.(2009). and references within for the most recent. situation of this problem.," See \citet{merrit et aguilar} for one of the first relevant numerical approaches, \citet{perez et al} for an intermediate position or \citet{ROI Moderne2} and references within for the most recent situation of this problem."
+ A complete historical account of radial orbit instability is given in. Maréchal&Perez(2009)., A complete historical account of radial orbit instability is given in \citet{vlasovia}.
+. In the next. section. we present some Κον features of this process.," In the next section, we present some key features of this process."
+ Radial orbit instability (hereafter. ROL) appears in self-eravitating system dynamics with the pionnering works of Antonov(1973) ancl Hénon(1973).., Radial orbit instability (hereafter ROI) appears in self-gravitating system dynamics with the pionnering works of \citet{Antonov73} and \citet{henon}.
+ A decade later Polvachenko&Shukhman(1981) propose a stability criterion based on the ratio of radial over tangential kinetic energies. when it is to small the svstem must leave its spherical svmimetey and form a bar.," A decade later \citet{polya} propose a stability criterion based on the ratio of radial over tangential kinetic energies, when it is to small the system must leave its spherical symmetry and form a bar."
+ This work is criticized by Palmer&Papaloizou(1987). which suggest. using normal modes techniques. that ROL can occur for arbitrary small values of the Russian ratio. provided the distribution function of the svstem is unbounded for orbits with zero angular momentum.," This work is criticized by \citet{palmer87} which suggest, using normal modes techniques, that ROI can occur for arbitrary small values of the Russian ratio, provided the distribution function of the system is unbounded for orbits with zero angular momentum."
+ This paper is also the first one to propose a relevant. physical mechanism to understand ROL based on resonant. trapped. orbits: we note that this mechanism needs a coupling between orbits., This paper is also the first one to propose a relevant physical mechanism to understand ROI based on resonant trapped orbits; we note that this mechanism needs a coupling between orbits.
+" Several factors show that a radial svstem needs a ""seed? from which ROL can appear.", Several factors show that a radial system needs a “seed” from which ROI can appear.
+ This is developed in detail in Rov&Perez(2004)., This is developed in detail in \citet{roy}.
+. The general idea is that there has to be a near-equilibrium state. so that coupling between orbits has the time to develop. and the instability to grow.," The general idea is that there has to be a near-equilibrium state, so that coupling between orbits has the time to develop, and the instability to grow."
+ In AlaeMillanetal.(1999)... density. profiles in a power law are considered. which means that the core has the time to stabilize before outer zones collapse. causing ROL to appear: it also shows that adding clumps tends to accelerate the process.," In \citet{MCMillan}, density profiles in a power law are considered, which means that the core has the time to stabilize before outer zones collapse, causing ROI to appear; it also shows that adding clumps tends to accelerate the process."
+ A counterexample can be found in Trenti&Bertin(2006).. which shows that homogeneous spherical haloes do not undergo ROL as the svstem tends to isotropy before reaching equilibrium.," A counterexample can be found in \citet{supprROI}, which shows that homogeneous spherical haloes do not undergo ROI, as the system tends to isotropy before reaching equilibrium."
+ For this reason. our study. will focus on ROL emerging from equilibrium states.," For this reason, our study will focus on ROI emerging from equilibrium states."
+ Although ROL is a natural candidate. to produce triaxialitv which can occur in self-eravitating svstenmis. it was noted Ixatz1991)) that this spatial counterpart of ROL could disapear during the merging process of the ealaxy formation.," Although ROI is a natural candidate to produce triaxiality which can occur in self-gravitating systems, it was noted \citealt{Nkatz}) ) that this spatial counterpart of ROI could disapear during the merging process of the galaxy formation."
+ Llowever. Llussetal.(1999) and. more recently MacMillanetal.(1999) have shown that ROL is a fundamental initial process which shapes the phase space of the galaxy progenitor and allows it to get the eood final mass density. profile.," However, \citet{huss} and more recently \citet{MCMillan} have shown that ROI is a fundamental initial process which shapes the phase space of the galaxy progenitor and allows it to get the good final mass density profile."
+ It is therefore important to understand fully the nature of ROL., It is therefore important to understand fully the nature of ROI.
+ In this context. the objective of this paper is twofold.," In this context, the objective of this paper is twofold."
+ On the first hand. in section 2.. we present a general moethod for investigate instability of self-gravitating svstems.," On the first hand, in section \ref{section-method}, we present a general method for investigate instability of self-gravitating systems."
+ This approach couples a general mathematical result. by Blochetal.(1994) which generalizes Lyapunov theory. ancl the svimplectic approach of the stability problem of CBP system (sce Bartolomew1971: Ixandrup. 1990: and Perez&Aly 1906)) it must be noted that Ixandrup(1991). has already used. this technique for non spherical svstems without the complete mathematical background.," This approach couples a general mathematical result by \citet{BKMR} which generalizes Lyapunov theory, and the symplectic approach of the stability problem of CBP system (see \citealt{Bartho}; \citealt{Kandrup1}; and \citealt{Perezaly1}) ) – it must be noted that \citet{Kandrup2} has already used this technique for non spherical systems without the complete mathematical background."
+ On the second hand. in section 3. we apply this method to obtain a direct energy. proof of the racial orbit instability when the system can clissipate energy.," On the second hand, in section \ref{section-instab} we apply this method to obtain a direct energy proof of the radial orbit instability when the system can dissipate energy."
+ Consider the first-order variation of an equilibrium Ho," Consider the first-order variation of an equilibrium $f_{0}\rightarrow
+f_{0}+f^{(1)}$ ."
+tt is well-known (see Bartolomew (CI971).. Kanelrup and Perez&Aly (1996))) that there exists a phase space function gCD. D. such that ‘This function is called a of theperturbation.," It is well-known (see \citet{Bartho}, , \citet{Kandrup1} and \citet{Perezaly1}) ) that there exists a phase space function $g\left(\mathbf{\Gamma},t\right)$ , such that This function is called a of theperturbation."
+". WrittenNN in. this. form.] fo""! albeis the largest class of physical perturbations which can be considered as acting on fy."," Written in this form, $f^{(1)}$ is the largest class of physical perturbations which can be considered as acting on $f_{0}$ ."
+ In other words. fli is a deformation of fy and then," In other words, $f^{(1)}$ is a deformation of $f_{0}$ and then"
+sources on the final sample (actually the highest redshift LBG) have SNR2 (see Table 1)).,sources on the final sample (actually the highest redshift LBG) have $\ga$ 6 (see Table \ref{Table}) ).
+ The final 0.7«€z6«1.6 HerMES LBG sample contains 12 objects and the 1.6<z«2.8 sample contains 1 object.," The final $0.7\,{<}\,z\,{<}\,1.6$ HerMES LBG sample contains 12 objects and the $1.6\,{<}\,z\,{<}\,2.8$ sample contains 1 object."
+" Although it is difficult to determine what fraction of z ~2 LBGs should be detected in the FIR, we can use the 24 um data to estimate how many z~1 LBGs are expected to be detected at 250 jum. From the present sample, the mean ratio S259/S24=54+11 for the SPIRE-detected LBGs and so, the present wm detection limit of 7.6mmJy correspond to 140pJy at 24,um. At this level, and based on the um catalog, we would expect 20 z~1 LBGs to be detected with SPIRE at jum. We find 12 objects, i.e., of the expectation."," Although it is difficult to determine what fraction of $z\sim$ 2 LBGs should be detected in the FIR, we can use the $24\,\mu$ m data to estimate how many $z\sim1$ LBGs are expected to be detected at $\,\mu$ m. From the present sample, the mean ratio $S_{\rm 250}/S_{24} = 54 \pm 11$ for the SPIRE-detected LBGs and so, the present $\,\mu$ m detection limit of mJy correspond to $\,\mu$ Jy at $\,\mu$ m. At this level, and based on the $\,\mu$ m catalog, we would expect 20 $z\sim1$ LBGs to be detected with SPIRE at $\,\mu$ m. We find 12 objects, i.e., of the expectation."
+ This is consistent with the completeness at this flux level in Roseboometal. (2010).., This is consistent with the completeness at this flux level in \citet{roseboom10}. .
+Inspired by the work of Hui&Zhang(2002).. here we suggest a new formula to model the halo spin-spin correlations: where a nonlinear correlation parameter 2.) is introduced to measure the strength of the nonlinear tidal effect on the halo spin-spin correlations.,"Inspired by the work of \citet{hui-zha02}, here we suggest a new formula to model the halo spin-spin correlations: where a nonlinear correlation parameter $\varepsilon_{\rm nl}$ is introduced to measure the strength of the nonlinear tidal effect on the halo spin-spin correlations."
+" In (his model the linear correlation parameter 15 notated as a, since its value could be different between the (wo cases that the nonlinear effect is ignored as in equation (4)) ancl taken into account as in (8)).", In this model the linear correlation parameter is notated as $a_{\rm l}$ since its value could be different between the two cases that the nonlinear effect is ignored as in equation \ref{eqn:1st}) ) and taken into account as in \ref{eqn:2nd}) ).
+ Equation (3)) implies that if the nonlinear contribution exists. (hen the halo spin-spin correlations must exist on larger scale than predicted by the linear model (eq.[4]]) since ib scales linearly with the densitv correlation function in the nonlinear model.," Equation \ref{eqn:2nd}) ) implies that if the nonlinear contribution exists, then the halo spin-spin correlations must exist on larger scale than predicted by the linear model \ref{eqn:1st}] ]) since it scales linearly with the density correlation function in the nonlinear model."
+ That is. the nonlinear tidal effect. would generate large-scale correlations of the halo spin axes.," That is, the nonlinear tidal effect would generate large-scale correlations of the halo spin axes."
+ OI course. (he values of the nonlinear correlation parameter oy) as well as the linear correlation parameter have to be determined empirically.," Of course, the values of the nonlinear correlation parameter $\varepsilon_{\rm nl}$ as well as the linear correlation parameter have to be determined empirically."
+ The cross-correlation. gc(r) can be also modeled in accordance with (3)) as It is worth mentioning that the value of ἐμ for the cross correlation je(r) would not be necessarily lower than its value for the auto-correlation (r) in equation (3)). unlike the case of (he linear tidal torque model in 82.2.," The cross-correlation, $\eta_{C}(r)$ can be also modeled in accordance with \ref{eqn:2nd}) ) as It is worth mentioning that the value of $\epsilon_{\rm nl}$ for the cross correlation $\eta_{C}(r)$ would not be necessarily lower than its value for the auto-correlation $\eta(r)$ in equation \ref{eqn:2nd}) ), unlike the case of the linear tidal torque model in 2.2."
+ In the nonlinear regime where (he non-Gaussianity in (he density field erows. the occurrence of the halo merging is quite frequent.," In the nonlinear regime where the non-Gaussianity in the density field grows, the occurrence of the halo merging is quite frequent."
+ This occurrence of the halo merging plavs a role of transferring the orbital angular momentum generated by the external tidal field into the spin angular momentum of a merged halo., This occurrence of the halo merging plays a role of transferring the orbital angular momentum generated by the external tidal field into the spin angular momentum of a merged halo.
+ In other words. the orbital angular momentum of a svstem composed of small individual halos become the spin angular momentum of a large halo lormed through merging of the small halos.," In other words, the orbital angular momentum of a system composed of small individual halos become the spin angular momentum of a large halo formed through merging of the small halos."
+" Given that the orbital angular momentum of the svstem before merging is generated bv the external tidal fieldl on larger scale (Vilvilskaοἱ,al.2002).. the transfer of the orbital angular momentum into the spin angular momentum creates eross-correlations of the tidal fields between different scales. which in turn generates large-scale cross-correlations ge: in the nonlinear regime."," Given that the orbital angular momentum of the system before merging is generated by the external tidal field on larger scale \citep{vit-etal02}, the transfer of the orbital angular momentum into the spin angular momentum creates cross-correlations of the tidal fields between different scales, which in turn generates large-scale cross-correlations $\eta_{C}$ in the nonlinear regime."
+We note (he spectrum below (he break is not exactly a power-law allhough resembles one Closely. therefore to compare with observations. which have normally poor energy resolution AEeOLE (e.g.2).. we fitted our simulated spectra with simple power-law fits.,"We note the spectrum below the break is not exactly a power-law although resembles one closely, therefore to compare with observations, which have normally poor energy resolution $\Delta E\sim 0.1 E$ \citep[e.g.][]{Lin_etal95}, we fitted our simulated spectra with simple power-law fits."
+" The spectral index below the break energy 0, 15 alwavs smaller than the spectral index above the break energy Opjgh and correlates (Figure 2)) with dpi, 1n a remarkably similar manner as observed bv ?..", The spectral index below the break energy $\delta_{low}$ is always smaller than the spectral index above the break energy $\delta_{high}$ and correlates (Figure \ref{fig:2}) ) with $\delta_{high}$ in a remarkably similar manner as observed by \citet{Krucker_etal09}.
+ The range of ὃς appear in a rather narrow range between 2 ancl 2.5 for a wide range ol injected spectral indices between 3 and 5 (Figure 2))., The range of $\delta_{low}$ appear in a rather narrow range between 2 and 2.5 for a wide range of injected spectral indices between 3 and 5 (Figure \ref{fig:2}) ).
+" The actual value of 05, is also dependent on the background plasma density ancl will be different should the heliospheric density model change.", The actual value of $\delta_{low}$ is also dependent on the background plasma density and will be different should the heliospheric density model change.
+ The break energy range for all simulations is between + keV ancl 30 keV (Figure 2)). with the exact break energy being dependent on the initial spectral index of the beam. 957. and the initial density of the beam.," The break energy range for all simulations is between $4$ keV and $80$ keV (Figure \ref{fig:2}) ), with the exact break energy being dependent on the initial spectral index of the beam, $\delta_{high}$, and the initial density of the beam."
+ As the initial density of the beam increases. the break of spectral index occurs at higher energies.," As the initial density of the beam increases, the break of spectral index occurs at higher energies."
+ Indeed. the larger number of injected electrons. the faster the generation of plasma waves proceeds ancl hence (he stronger (he interaction between electrons ancl plasma waves.," Indeed, the larger number of injected electrons, the faster the generation of plasma waves proceeds and hence the stronger the interaction between electrons and plasma waves."
+ This also explains the dependence of break energy {ο (he injected spectral index. with lower spectral indices having a larger population of higher energv electrons and hence having higher break energies.," This also explains the dependence of break energy to the injected spectral index, with lower spectral indices having a larger population of higher energy electrons and hence having higher break energies."
+ The fInence al the break energy correlates to the break energy itself (Figure 2)) with hieher break energies corresponding to lower [Iuence magnitudes., The fluence at the break energy correlates to the break energy itself (Figure \ref{fig:2}) ) with higher break energies corresponding to lower fluence magnitudes.
+ The particles arriving at the Earth (1.2 AU) show near time-ol-Iight dispersion (Figure 3)). often observed by satellites in impulsive solar electron events.," The particles arriving at the Earth (1.2 AU) show near time-of-flight dispersion (Figure \ref{fig:3}) ), often observed by satellites in impulsive solar electron events."
+ Assuming scatter-[ree propagation. ie. without any interaction wilh plasma. one can produce (he apparent injection profile at the Sun. as ??. did for observations.," Assuming scatter-free propagation, i.e. without any interaction with plasma, one can produce the apparent injection profile at the Sun, as \citet{Krucker_etal1999,Wang_etal06} did for observations."
+ Under scatter-Iree assumption. the for energetic electrons from (he Sun {ο the spacecralt is simply /4=L/e(E). where L—1.2 AU is the distance and ve(£) is the speed of electrons for various energies.," Under scatter-free assumption, the time-of-flight for energetic electrons from the Sun to the spacecraft is simply $t_A=L/v(E)$, where $L=1.2$ AU is the distance and $v(E)$ is the speed of electrons for various energies."
+" The time of arrival is (hen simply /4=/;,;+L/((E). where /;,; is the apparent injection time."," The time of arrival is then simply $t_A=t_{inj}+L/v(E)$, where $t_{inj}$ is the apparent injection time."
+ These apparent injection profiles with background added have been calculated from the simulated [Iuxes al 1.2 AU (see Figure 4))., These apparent injection profiles with background added have been calculated from the simulated fluxes at 1.2 AU (see Figure \ref{fig:4}) ).
+" If the electrons propagate scatter-Iree they would require 10—20 minutes earlier onset of injection /;,; of low (3—12 keV) energy electron injection and a delaved. maximum of the injection to explain observations.", If the electrons propagate scatter-free they would require $10-20$ minutes earlier onset of injection $t_{inj}$ of low $3-12$ keV) energy electron injection and a delayed maximum of the injection to explain observations.
+ An identical simulationwas ran with an electron beam not interacting with the backeround plasma (scatter free propagation) and the results are compared in Figure 4.., An identical simulationwas ran with an electron beam not interacting with the background plasma (scatter free propagation) and the results are compared in Figure \ref{fig:4}. .
+ Comparing the calculated injection, Comparing the calculated injection
+cosmic far-infrared background.,cosmic far-infrared background.
+" The and aims to survey 550 and detect >300,000 galaxies."," The and aims to survey 550 $^2$ and detect $> 300{,}000$ galaxies."
+" A new methodology for selecting gravitational lenses using wide-area submillimeter (sub-mm) surveys (Blain 1996, Perrotta et al."," A new methodology for selecting gravitational lenses using wide-area submillimeter (sub-mm) surveys (Blain 1996, Perrotta et al."
+ 2002; Negrello et al., 2002; Negrello et al.
+ 2007) has been tested for the first time during Herschel’s science demonstration phase (SDP; Negrello et al., 2007) has been tested for the first time during 's science demonstration phase (SDP; Negrello et al.
+" 2010, N10 hereafter)."," 2010, N10 hereafter)."
+" Candidates are selected using a flux cut of 100mJy at jum, a limit based on the steep number counts slope in the sub-mm (Negrello et al."," Candidates are selected using a flux cut of $\,$ mJy at $\,\mu$ m, a limit based on the steep number counts slope in the sub-mm (Negrello et al."
+ , 2007).
+Above this flux limit only gravitationally lensed objects 2007).and easily identifiable ‘contaminants’ remain e.g. blazars and local spiral galaxies., Above this flux limit only gravitationally lensed objects and easily identifiable `contaminants' remain e.g. blazars and local spiral galaxies.
+" Initial results show that this selection method has ~100% efficiency and should deliver a sample of hundreds of new gravitational lenses in planned wide areaHerschel surveys, probing galaxies with intrinsic fluxes below theHerschel confusion limit."," Initial results show that this selection method has $\sim$ efficiency and should deliver a sample of hundreds of new gravitational lenses in planned wide area surveys, probing galaxies with intrinsic fluxes below the confusion limit."
+ Five lens candidates in the SDP H-ATLAS data et al., Five lens candidates in the SDP H-ATLAS data (Pascale et al.
+ 2010; Rigby et al., 2010; Rigby et al.
+ 2010; N10) were confirmed(Pascale with spectroscopic redshifts obtained via the detection of carbon monoxide (CO) emission lines (Lupu et al., 2010; N10) were confirmed with spectroscopic redshifts obtained via the detection of carbon monoxide (CO) emission lines (Lupu et al.
+ 2010; Frayer et al., 2010; Frayer et al.
+ 2011) of the background galaxies and optical spectra (Negrello et al., 2011) of the background galaxies and optical spectra (Negrello et al.
+ 2010) of the lens galaxies., 2010) of the lens galaxies.
+" In this Letter we study two of these gravitational lenses, ATLAS J090311.6+003906 (SDP.81) and H-ATLAS J091305.0-005343H- (SDP.130), which are submillimeter galaxy at redshifts of z—3.04 and z—2.63 being lensed by (SMG)ellipticals at z—0.30 and z=0.22, respectively."," In this Letter we study two of these gravitational lenses, H-ATLAS J090311.6+003906 (SDP.81) and H-ATLAS J091305.0-005343 (SDP.130), which are submillimeter galaxy (SMG) at redshifts of $z = 3.04$ and $z = 2.63$ being lensed by ellipticals at $z = 0.30$ and $z = 0.22$, respectively."
+" Submillimeter Array (SMA) imaging reveals the sub-mm morphology, consistent with a lensing event, with multiple peaks distributed around the position of the foreground elliptical galaxy (N10)."," Submillimeter Array (SMA) imaging reveals the sub-mm morphology, consistent with a lensing event, with multiple peaks distributed around the position of the foreground elliptical galaxy (N10)."
+" While no lensed background images were detectable in optical imaging, the spectral energy distribution models that flux from the background sources(SED) should be suggestdetectable above ~10yuJy at near-IR wavelengths, from 2 to 5 jum "," While no lensed background images were detectable in optical imaging, the spectral energy distribution (SED) models suggest that flux from the background sources should be detectable above $\sim$ $\,\mu$ Jy at near-IR wavelengths, from 2 to 5 $\,\mu$ m (N10)."
+"At these wavelengths, the emission from the foreground(N10). lenses, at z~0.25, and the background SMGs becomes comparable."," At these wavelengths, the emission from the foreground lenses, at $z\sim0.25$, and the background SMGs becomes comparable."
+" In this Letter we present Spitzer/IRAC imaging, light profile models, photometry and physical characteristics for SDP.81 and SDP.130 derived from SED fitting."," In this Letter we present /IRAC imaging, light profile models, photometry and physical characteristics for SDP.81 and SDP.130 derived from SED fitting."
+" In the next section, we present a summary of theSpitzer data, while in Section 3 we discuss modeling of the light profile of the foreground lenses."," In the next section, we present a summary of the data, while in Section 3 we discuss modeling of the light profile of the foreground lenses."
+ We perform SED modeling of background SMGs and present results related to the properties of these two sources in Section 4., We perform SED modeling of background SMGs and present results related to the properties of these two sources in Section 4.
+ Throughout the Letter we assume flat ACDM cosmology with Qu —0.3 and Hp=70 kms~! Mpc!.," Throughout the Letter we assume flat $\Lambda$ CDM cosmology with $\Omega_{\rm M}\,$ $\,0.3$ and $H_0\,=\,70$ $\,$ $^{-1}\,$ $^{-1}$."
+" The IR imaging data are part ofSpitzer program 548 (PI: CCooray), released for analysis on 2010 July 30."," The IR imaging data are part of program 548 (PI: Cooray), released for analysis on 2010 July 30."
+" For this program, Infrared Array Camera (IRAC; Fazio et al."," For this program, Infrared Array Camera (IRAC; Fazio et al."
+ images of SDP.81 and SDP.130 were taken at 3.6m 2004)(Channel 1) and m (Channel 2).," 2004) images of SDP.81 and SDP.130 were taken at $\,\mu$ m (Channel 1) and $\,\mu$ m (Channel 2)."
+" Both lensing systems were imaged with a 36-position dither pattern and an exposure time of 30 s for each frame to achieve an effective total exposure time of 1080 s. The rms depth reached is 0.3 and 0.4 wJy in Channels 1 and 2, respectively."," Both lensing systems were imaged with a 36-position dither pattern and an exposure time of 30 s for each frame to achieve an effective total exposure time of 1080 s. The rms depth reached is 0.3 and 0.4 $\mu$ Jy in Channels 1 and 2, respectively."
+" Corrected basic calibrated data pre-processed by theSpitzer Science Center (SSC), using the standard pipeline version S18.18.0, were spatially aligned, resampled, and combined into a mosaic image using version 18.3.1 of the SSC's MOPEX software suite (Makovoz & Marleau 2005)."," Corrected basic calibrated data pre-processed by the Science Center (SSC), using the standard pipeline version S18.18.0, were spatially aligned, resampled, and combined into a mosaic image using version 18.3.1 of the SSC's MOPEX software suite (Makovoz $\&$ Marleau 2005)."
+" The IRAC mosaics have a resampled pixel size of 0.6"" and angular resolution of 2 to 2.5"".", The IRAC mosaics have a resampled pixel size of $''$ and angular resolution of 2 to $''$.
+" For the work presented here, we also use the IRAC point spread function (PSF; version 2010 April) file as provided by the SSC5"," For the work presented here, we also use the IRAC point spread function (PSF; version 2010 April) file as provided by the SSC."
+"6, We also make use of optical images of SDP.81 and SDP.130 that were acquired on 2010 March 10 using the dual-arm Low Resolution Imaging Spectrometer (LRIS; Oke et al.", We also make use of optical images of SDP.81 and SDP.130 that were acquired on 2010 March 10 using the dual-arm Low Resolution Imaging Spectrometer (LRIS; Oke et al.
+" 1995, McCarthy et al."," 1995, McCarthy et al."
+ 1998) on the Keck I, 1998) on the Keck I
+Tf >=1. the set of equatious (1)) are called the Manakoy equations.,"If $\gamma=1$, the set of equations \ref{eq: manakov0}) ) are called the Manakov equations."
+ The Manakov equations are well-studied aud solutions can be obtained in several wavs |17.Ls.19.20.21.22.23].," The Manakov equations are well-studied and solutions can be obtained in several ways \cite{Manakov,Zakharov,Radh,Shches,Zen1,Zen2,Zen3}."
+. It has to be remarked that the transformation (3)) cads to special solutions of Eq.(1)). iu which the effect of the birefriugence appears in the phase factor.," It has to be remarked that the transformation \ref{trans}) ) leads to special solutions of \ref{eq: manakov}) ), in which the effect of the birefringence appears in the phase factor."
+ Since under experimental conditions usually the optical power is measured [us]? y. the transformation (3)) leads to special solutious of Eq.(1)). which cau not explain the decay of optical pulses due to birefringence (nce the birefiugence is iu he phase).," Since under experimental conditions usually the optical power is measured $|u_{i}|^{2}$ ), the transformation \ref{trans}) ) leads to special solutions of \ref{eq: manakov}) ), which can not explain the decay of optical pulses due to birefringence (since the birefringence is in the phase)."
+ Ou the other hand in Rof.|16]. solutions are reported which show that Eq.(1)) has modes which propagate with a different velocity., On the other hand in \cite{Kath} solutions are reported which show that \ref{eq: manakov0}) ) has modes which propagate with a different velocity.
+ A general iiethod to solve the equations (1)) has never becu given |11].., A general method to solve the equations \ref{eq: manakov}) ) has never been given \cite{Hasegawa}.
+ There are more approaches possible which can be used to find solutions of Eq.(1))., There are more approaches possible which can be used to find solutions of \ref{eq: manakov}) ).
+ Although IHirotas method has con proved to be successful for many types of equations. we approach the problem iu this paper usine a method developed by one of us (see 211) to find special solutions of Eq.(1)).," Although Hirota's method has been proved to be successful for many types of equations, we approach the problem in this paper using a method developed by one of us (see \cite{Dorren}) ) to find special solutions of \ref{eq: manakov}) )."
+ These solutious have the property that the solitonic part of the modes propagate with a different velocity., These solutions have the property that the solitonic part of the modes propagate with a different velocity.
+ The existence of such a solution might be inportaut or explaining the decay of solitonic pulses in birefringent optical fibers., The existence of such a solution might be important for explaining the decay of solitonic pulses in birefringent optical fibers.
+ For linear birefriugence it is well known that he decay of initially localized pulses can be explained by the fact that both modes propagate with a differcut velocity along the birefringence axes |0].., For linear birefringence it is well known that the decay of initially localized pulses can be explained by the fact that both modes propagate with a different velocity along the birefringence axes \cite{Agrawal}.
+ The results of this paper suggests that for solitonic pulses a similar behavior takes lace., The results of this paper suggests that for solitonic pulses a similar behavior takes place.
+ Iu Sect., In Sect.
+IL. we apply concepts derived in earlier publications to find special solutions of the birefringence problem.,"II, we apply concepts derived in earlier publications to find special solutions of the birefringence problem."
+ We derive solutions consisting of of two unperturbed solitons propagating with a differeut velocity along the principal ürefiugence axes., We derive solutions consisting of of two unperturbed solitons propagating with a different velocity along the principal birefringence axes.
+ Moreover the solutious contain interactions terms between the pure solitouic solutions., Moreover the solutions contain interactions terms between the pure solitonic solutions.
+ It is concluded that two effects play a role in the decay of solitouic pulses., It is concluded that two effects play a role in the decay of solitonic pulses.
+ The effect is simular as iu linear birefrineciuce and is due to the unequal propagation velocity of the modes., The effect is similar as in linear birefringence and is due to the unequal propagation velocity of the modes.
+ The secoud effect. which is absent iu linear birefringcuce is the nonlinear solitou-soliton interaction between the modes.," The second effect, which is absent in linear birefringence is the nonlinear soliton-soliton interaction between the modes."
+ The paper is coucluded with a discussion., The paper is concluded with a discussion.
+ For solving Eq.(1)). we try a starting poiut the following free wave solutions: If we substitute Eq.(5)) iuto Eq.(1)). we obtain: Iu Eq.(6)). it is assumed that A; and w; are both complex (the bars indicate the complex conjugates).," For solving \ref{eq: manakov}) ), we try a starting point the following free wave solutions: If we substitute \ref{freewave}) ) into \ref{eq: manakov}) ), we obtain: In \ref{f_check}) ), it is assumed that $k_{i}$ and $\omega_{i}$ are both complex (the bars indicate the complex conjugates)."
+ We require that the following dispersion relationships are valid: We can conclude that for nonzero coefficients A and B. the triy-solutiou (5)) does not satisfy Eq.(1)) and hence the birefringence equations have uo free wave solutions.," We require that the following dispersion relationships are valid: We can conclude that for nonzero coefficients $A$ and $B$, the try-solution \ref{freewave}) ) does not satisfy \ref{eq: manakov}) ) and hence the birefringence equations have no free wave solutions."
+ Before deriving special solutions of Eq.(1)) the following sinplificatious are miposed to simplify the computations which will follow:, Before deriving special solutions of \ref{eq: manakov}) ) the following simplifications are imposed to simplify the computations which will follow:
+"The sealing relations between some global properties of ellipticals (Es). known as the ""Fundamental Plane’ (FP: Dressler et al.","The scaling relations between some global properties of ellipticals (Es), known as the `Fundamental Plane' (FP; Dressler et al."
+ 1987. Djorgovski Davis 1987). is clearly a manifestation of a great ‘regularity’ in the properties of Es.," 1987, Djorgovski Davis 1987), is clearly a manifestation of a great `regularity' in the properties of Es."
+ Although a proof of its ‘universality’ is still lacking. the FP has been widely applied to various problems: from the study of peculiar motions (Lynden-Bell et al.," Although a proof of its `universality' is still lacking, the FP has been widely applied to various problems: from the study of peculiar motions (Lynden-Bell et al."
+ 1988. Jorgensen et al.," 1988, rgensen et al."
+ 1996). to the study of the evolution of M/L ratios (van Dokkum Franx 1996. Kelson et al.," 1996), to the study of the evolution of M/L ratios (van Dokkum Franx 1996, Kelson et al."
+ 1997). and measurements of gy (Bender et 11998).," 1997), and measurements of $q_0$ (Bender et 1998)."
+ The ‘regularity’ of the E’s properties. shown by the FP. must be connected to their gross structural characteristics and to their formation and evolutionary mechanisms (Djorgovski et al.," The `regularity' of the E's properties, shown by the FP, must be connected to their gross structural characteristics and to their formation and evolutionary mechanisms (Djorgovski et al."
+ 1995 and references therein)., 1995 and references therein).
+ To understand this connection is however a very hard task. especially in view of the problems posed by the ‘thinness’ of the FP itself (Ciotti et al.," To understand this connection is however a very hard task, especially in view of the problems posed by the `thinness' of the FP itself (Ciotti et al."
+ 1996)., 1996).
+ Apart from the obvious similarities. it is a matter of fact that Es do not constitute a homogeneous family of galaxies: Es not only have substantially different photometric and kinematical properties. but they also host a variety of sub-components. like inner disks. rottatting cores. outer rings. dust lanes. ete. (," Apart from the obvious similarities, it is a matter of fact that Es do not constitute a homogeneous family of galaxies: Es not only have substantially different photometric and kinematical properties, but they also host a variety of sub-components, like inner disks, ting cores, outer rings, dust lanes, etc. ("
+Hlingworth 1983. Kormendy Djorgovski 1989. Caon et al.,"Illingworth 1983, Kormendy Djorgovski 1989, Caon et al."
+ 1993. Michard Marchal 1994. Capaecioli Longo 1994. Scorza Bender 1996. Graham et 11996. Graham Colless 1997. Busarello et al.," 1993, Michard Marchal 1994, Capaccioli Longo 1994, Scorza Bender 1996, Graham et 1996, Graham Colless 1997, Busarello et al."
+ 1997. Prugniel Simien 1997. Carollo et al.," 1997, Prugniel Simien 1997, Carollo et al."
+ 1997. and references therein).," 1997, and references therein)."
+ The possibility that many Es have disks that escape photometric detection (Simien de Vaucouleurs 1986. Capaccioli 1987. Rix White 1990. Saglia et al.," The possibility that many Es have disks that escape photometric detection (Simien de Vaucouleurs 1986, Capaccioli 1987, Rix White 1990, Saglia et al."
+ 1997). and the dichotomy between the two families of boxy and disky Es lead to the possibility of a new classification scheme. in which galaxies formerly belonging to the same Hubble type are separated according to tsophotal shape. and in which a continuous sequence exists between Es and SOs (Kormendy Bender 1996).," 1997), and the dichotomy between the two families of boxy and disky Es lead to the possibility of a new classification scheme, in which galaxies formerly belonging to the same Hubble type are separated according to isophotal shape, and in which a continuous sequence exists between Es and S0s (Kormendy Bender 1996)."
+ Since the above scenario is manifestly at odds with the ‘regularity’ of the FP. the question arises to what extent the ‘global’ quantities entering into the construction of the FP are representative of the real status of the galaxies.," Since the above scenario is manifestly at odds with the `regularity' of the FP, the question arises to what extent the `global' quantities entering into the construction of the FP are representative of the real status of the galaxies."
+ In this respect. an important question is whether the central velocity dispersion συ is fully representative of the dynamical status of a galaxy.," In this respect, an important question is whether the central velocity dispersion $\sigma_0$ is fully representative of the dynamical status of a galaxy."
+ There i$ much evidence that this is not the case: due to differences in velocity dispersion profiles and anisotropy. masses derived from oy can be wrong by a factor of ten (Tonry 1983. Richstone Tremaine 1986. Mathews 1988); in presence of central. kinematically distinct sub-compponnents. the value of συ is likely to be more connected to the sub-compponnent than to the galaxy as a whole: Es are dynamically mmolloggous systems: the slope of the velocity dispersion profile is correlated to a) in such a way that there is a systematic departure of the kinetic energy from a simple scaling of oy (Graham Colless 1997. Busarello et al.," There is much evidence that this is not the case: due to differences in velocity dispersion profiles and anisotropy, masses derived from $\sigma_0$ can be wrong by a factor of ten (Tonry 1983, Richstone Tremaine 1986, Mathews 1988); in presence of central, kinematically distinct nents, the value of $\sigma_0$ is likely to be more connected to the nent than to the galaxy as a whole; Es are dynamically gous systems: the slope of the velocity dispersion profile is correlated to $\sigma_0$ in such a way that there is a systematic departure of the kinetic energy from a simple scaling of $\sigma_0$ (Graham Colless 1997, Busarello et al."
+ 1997); finally. the rotation. although not dominant in Es. plays a non-negligible role (Busarello et al.," 1997); finally, the rotation, although not dominant in Es, plays a non-negligible role (Busarello et al."
+" 1992, Prugniel"," 1992, Prugniel"
+Images of the eluster were obtained during two separate observing runs. in. 1985 and. 1993. with telescopes located on Cerro Tololo and Mt. Palomar respectively (see Table 1).,"Images of the cluster were obtained during two separate observing runs, in 1985 and 1993, with telescopes located on Cerro Tololo and Mt. Palomar respectively (see Table 1)."
+ The total field covered is shown in Fig. 1.., The total field covered is shown in Fig. \ref{fig-map}.
+ The first run was at the CPLO Blanco tra telescope using the prime focus CCD camera with CA2., The first run was at the CTIO Blanco 4m telescope using the prime focus CCD camera with RCA2.
+ The chip had. good «quantum elliciency even in the near-UV/blue. but a Large read-out noise (00 rms)," The chip had good quantum efficiency even in the near-UV/blue, but a large read-out noise (90 $^-$ rms)."
+ Using the doublet. corrector. this CCD at prime focus vielded a spatial sampling of 0.58 arcesec per pixel.," Using the doublet corrector, this CCD at prime focus yielded a spatial sampling of 0.58 arcsec per pixel."
+ Alrough this particular observing run consisted. of several nights. the cluster. was observed for. only two of these nights. both of which were photometric.," Although this particular observing run consisted of several nights, the cluster was observed for only two of these nights, both of which were photometric."
+ minimum of 12 stancard stars. as established by Graham (1982) and by Landolt (1983). were observed several times cach night and through a range of airmass (1 to 2).," A minimum of 12 standard stars, as established by Graham (1982) and by Landolt (1983), were observed several times each night and through a range of airmass (1 to 2)."
+ The colours of the observed standard stars spanned from —-0.25/0D-V 11.35. encompassinglhewholerangeofcoloursfoundintheclustermembers.," The colours of the observed standard stars spanned from $-$ 0.25 to 1.35, encompassing the whole range of colours found in the cluster members."
+"T, "," The cluster observations were generally performed near the meridian, and never with an airmass more than 1.2."
+The Glters used. for the observations (UDGWVIBL) consisted of an unusual mixture of passbancl definitions., The filters used for the observations (UBGVRI) consisted of an unusual mixture of passband definitions.
+ The passbands of B anc Vo were. defined by. standard combinations of Schott coloured-filter glass., The passbands of B and V were defined by standard combinations of Schott coloured-filter glass.
+ Phe passbancds of Roane EL were defined. by interference filters. common to CTIO/IKPNO at that time.," The passbands of R and I were defined by interference filters, common to CTIO/KPNO at that time."
+ The U and CG filters were those of Hamilton (unpublished)., The U and G filters were those of Hamilton (unpublished).
+" The U filter used in this case was à combination of Schott UCG5 with a liquid copper sulfate ο(ο,", The U filter used in this case was a combination of Schott UG5 with a liquid copper sulfate cut-off.
+ Schott. UG5 is slightly redder and broader than the more standard UCG2 commonly used at CIO since 19585., Schott UG5 is slightly redder and broader than the more standard UG2 commonly used at CTIO since 1985.
+ The G filter was specially designed to be a reasonable representation of the passband. defined: by. hypersensitized Illa) emulsions with a two mm thick Schott €1€1385 filter.," The G filter was specially designed to be a reasonable representation of the passband defined by hypersensitized IIIaJ emulsions with a two mm thick Schott GG385 filter,"
+electrous produced by the harder photous.,electrons produced by the harder photons.
+ The temperature at each step through the slab is giveu yw balauciug photoionizatiou heating with the cooling from hydrogen aud helium gas (collisional jonizatious. recombinations. Lines excitation and [ree-free. processes). from metal lines auc from Πο cooling.," The temperature at each step through the slab is given by balancing photoionization heating with the cooling from hydrogen and helium gas (collisional ionizations, recombinations, lines excitation and free-free processes), from metal lines and from $H_2$ cooling."
+ After each iteration the mass averaged temperature T. is computed cousistently aud is used to determiue the gas vertical dispersion aud stratification., After each iteration the mass averaged temperature $T$ is computed consistently and is used to determine the gas vertical dispersion and stratification.
+ The ionization states of metals (C.O.Fe) are computed step by step consistently with the photoionization rates aud charge[n] exchauge 'actious: very lieh ionization states (i.e. potential energies Eo>51 eV) are not considered.," The ionization states of metals (C,O,Fe) are computed step by step consistently with the photoionization rates and charge exchange reactions; very high ionization states (i.e. potential energies $E \gg 54$ eV) are not considered."
+ For uetal abundance we use Z=0.02Z. throughout this paper but we shall discuss brielly its possible variatious., For metal abundance we use $Z=0.02 Z_\odot$ throughout this paper but we shall discuss briefly its possible variations.
+ Πο [ractious. although small. are computec via primordial chemical reactions Palla 1997).," $_2$ fractions, although small, are computed via primordial chemical reactions \citep{cgp97}."
+. From the ionization-recombiuation balance we know that for a small neutral fraction the ratio Napμι depends ou the ratio between the ionization and recombination coefficients and varies inversely with the gas volume deusity., From the ionization-recombination balance we know that for a small neutral fraction the ratio $N_{H\perp}/N_{HI\perp}$ depends on the ratio between the ionization and recombination coefficients and varies inversely with the gas volume density.
+" We define .X as Notice that N is delined at the line of sight value of Ny,=1.6x10!em.? and therefore the corresponding μι value depends on the thickness ratio A/R.", We define $X$ as Notice that $X$ is defined at the line of sight value of $N_{HI}=1.6\times 10^{17} {\hbox{cm}}^{-2}$ and therefore the corresponding $N_{HI\perp}$ value depends on the thickness ratio $h/R$.
+" We derive for jo>1 the following relation where J, is the intensity of the background flux at 912 written in units of LO7 eres 2s ! tsp 1.", We derive for $\eta_0\ge 1$ the following relation where $J_L$ is the intensity of the background flux at 912 written in units of $10^{-22}$ ergs $^{-2}$ $^{-1}$ $^{-1}$ $^{-1}$.
+ For our redshift range (1.75<2< 3.25) we use a constant Jp and a flux spectrum computed for z2.5 by Haardt & Macau (1996. hereatter H-M) for emission by quasars and absorption by intervening clouds.," For our redshift range $1.75<z<3.25$ ) we use a constant $J_L$ and a flux spectrum computed for $z=2.5$ by Haardt $\&$ Madau (1996, hereafter H-M) for emission by quasars and absorption by intervening clouds."
+ A nearly constant ionization rate between z=2 and Lis also required by comparison of results [rom hydrodyuaiical simulations of structure formation with the measured opacity of the Lyiuau-a forest., A nearly constant ionization rate between $z=2$ and 4 is also required by comparison of results from hydrodynamical simulations of structure formation with the measured opacity of the $\alpha$ forest.
+" The H-M flux intensity at 20292.5. Jp£29 5.5. agrees sulficiently well with other estimates (Rauchetal.al.1996) aud we shall use Jj,=5.5 throughout."," The H-M flux intensity at $z\simeq 2.5$, $J_L\simeq 5.5$ , agrees sufficiently well with other estimates \citep{rau97,gia96}
+ and we shall use $J_L =5.5$ throughout."
+ The compression factor ji is unkuown a priori and will be determined by trial and error., The compression factor $\eta_0$ is unknown a priori and will be determined by trial and error.
+ Hence aj) (or equivalently IX ). aud à are parameters to be determiued.," Hence $\eta_0$ (or equivalently $X$ ), and $\alpha$ are parameters to be determined."
+" Iu Figure 3 we show three curves for μι as afunction of Nyy For logs,ην ——1.9.—0.1.0.7. which correspond to X= 2.3. 2.8. and 3.2 respectively."," In Figure 3 we show three curves for $N_{HI\perp}$ as a function of $N_{H\perp}$ for $J_L/\eta_0$ $-1.9,-0.4,0.7$, which correspond to $X = $ 2.3, 2.8, and 3.2 respectively."
+" All curves in Figure 3 are characterized by 3 regious: a first region at high column density where αμ2Nyy. a second regiou where [or a stnallchange in Ny, the HI column density cliauges very rapidly. aud a third region. towards the bottom part of the plot. where the ουδ ος relation is again linear with a slope smaller than unity audwhich in general depends on the variations of 4 with Ny)."," All curves in Figure 3 are characterized by 3 regions: a first region at high column density where $N_{H\perp}\simeq N_{HI\perp}$, a second region where for a smallchange in $N_{H\perp}$ the HI column density changes very rapidly, and a third region, towards the bottom part of the plot, where the $N_{H\perp}$ $N_{HI\perp}$ relation is again linear with a slope smaller than unity andwhich in general depends on the variations of $\eta$ with $N_{H\perp}$ ."
+ For the models in this, For the models in this
+dissipation from such events is Likely to increase (he intrinsic luminosity and radius of these planets.,dissipation from such events is likely to increase the intrinsic luminosity and radius of these planets.
+ Direct imaging and rare transit observations may provide observational verifications that the internal structure of tvpical gas giants may have been modified by GIMs., Direct imaging and rare transit observations may provide observational verifications that the internal structure of typical gas giants may have been modified by GIMs.
+ We shall discuss the GIMs influence on (he mass-radius-Iuminositv of the close-in planets in Paper II., We shall discuss the GIM's influence on the mass-radius-luminosity of the close-in planets in Paper II.
+ We have shown possible mechanisms for the dredge up of heavy elements rom the core and for (he promotion of core growth., We have shown possible mechanisms for the dredge up of heavy elements from the core and for the promotion of core growth.
+ We have not considered here (he long-term sustainability of the stir-up elemental distribution., We have not considered here the long-term sustainability of the stir-up elemental distribution.
+ The follow up evolution of the impacted eas giants radius depends sensitivelv on their cooling efficiency and elemental segregation., The follow up evolution of the impacted gas giant's radius depends sensitively on their cooling efficiency and elemental segregation.
+ Although there are models for contraction and differentiation of initially homogeneous gas eijants (Bodenheimeretal.2003:Baralle2008).. the long-term post GIM evolution still needs to be investigated.," Although there are models for contraction and differentiation of initially homogeneous gas giants \citep{Bodenheimer2003, Baraffe2008}, the long-term post GIM evolution still needs to be investigated."
+ We thank Drs P. Bodenheimer. DP. Garaud. T. Guillot. and. G. Laughlin for useful conversation.," We thank Drs P. Bodenheimer, P. Garaud, T. Guillot, and G. Laughlin for useful conversation."
+ Part of the caleulation is carried ont on the SGI Altix 330 svstem at the Department of Astronomy. Peking University.," Part of the calculation is carried out on the SGI Altix 330 system at the Department of Astronomy, Peking University."
+ This work is supported by NASA NNXOTA- NNNOTAIS88C. NNNOSALAIG. NNNOSAMSAG. and NSECAST-0908807).," This work is supported by NASA NNX07A-L13G, NNX07AI88G, NNX08AL41G, NNX08AM84G, and NSF(AST-0908807)."
+ CDAs involvement was supported by NASA (NNGO5G1496)., CBA's involvement was supported by NASA (NNG05G1496).
+"aul its solution gives where 7 and ¢ are coustauts aud the vertical frequency 4) is defined by The motion in 2 and © is described by the superposition of the circular motion of the guidiug center at dt) at frequency 9. the epicyelic motion represeuted by ""eeceutricity € at frequeucey Ay. aud the forced oscillations of fractional radial amplitudes Cy at frequencies η—ng].","and its solution gives where $i$ and $\zeta$ are constants and the vertical frequency $\nu_0$ is defined by The motion in $R$ and $\phi$ is described by the superposition of the circular motion of the guiding center at $R_0$ at frequency $n_0$ , the epicyclic motion represented by “eccentricity” $e$ at frequency $\kappa_0$ , and the forced oscillations of fractional radial amplitudes $C_k$ at frequencies $k|n_0-n_{pc}|$."
+ The motion in 2 decouples from that in £2 aud © and has only free oscillations at the vertical frequency η., The motion in $z$ decouples from that in $R$ and $\phi$ and has only free oscillations at the vertical frequency $\nu_0$.
+ If there is uo epicyclic or vertical motion (e=7= 0). the orbit is a closed orbit iu the frame rotating at the pattern speed rye. withthe variation of the radius with the azimuthal angle being the stun of terms withfractional amplitude Cy aud period 27/h.," If there is no epicyclic or vertical motion $e = i = 0$ ), the orbit is a closed orbit in the frame rotating at the pattern speed $n_{pc}$, withthe variation of the radius with the azimuthal angle being the sum of terms withfractional amplitude $C_k$ and period $2\pi/k$."
+" If €2$5,Cy. the orbit is approximately a precessing Ixepleriau ellipse with eccentricity €. mean motion iy. auc periapse precession rate If in additiou 7z0. the orbit has an inclination ¢ with respect to the Pluto-Charon orbital plaue aud a uodal precessiou rate Eq. (17))"," If $e \gg \sum_k C_k$, the orbit is approximately a precessing Keplerian ellipse with eccentricity $e$, mean motion $n_0$, and periapse precession rate If in addition $i \neq 0$, the orbit has an inclination $i$ with respect to the Pluto-Charon orbital plane and a nodal precession rate Eq. \ref{n0}) )"
+ shows that the meat motion vp and hence the azimuthal period [πι=2z/1n0 deviate from those of a Ixepleriau orbit., shows that the mean motion $n_0$ and hence the azimuthal period $P_0 = 2\pi/n_0$ deviate from those of a Keplerian orbit.
+ lu Fig., In Fig.
+ 3 the solid line shows Py/—τν1 from Eq. (17))," \ref{fig:period} the solid line shows $P_K/P_0 - 1 =
+n_0/n_K - 1$ from Eq. \ref{n0}) )"
+ as a function of the eeuidine center radius A5 for the best-fit mass ratio of Charon-Pluto. nefny=0.1165. while the dotted lines show the same quantity for fy that are Le (0.0055) from the best-fit value. (," as a function of the guiding center radius $R_0$ for the best-fit mass ratio of Charon-Pluto, $m_c/m_p = 0.1165$, while the dotted lines show the same quantity for $m_c/m_p$ that are $1\,\sigma$ $\pm 0.0055$ ) from the best-fit value. ("
+"The uncertainty in nip dominates that in aj, in the evaluation of Eq. [17]].)",The uncertainty in $m_c/m_p$ dominates that in $a_{pc}$ in the evaluation of Eq. \ref{n0}] ].)
+ Fig., Fig.
+ 1 shows that the azimuthal period is shorter than the Ixeplerian orbital period., \ref{fig:period} shows that the azimuthal period is shorter than the Keplerian orbital period.
+ The orbital parameters of P2 aud P1 in Table 1 were obtained by BCYYS (rom uuperturbe ]xeplerian fits., The orbital parameters of P2 and P1 in Table \ref{table1} were obtained by BGYYS from unperturbed Keplerian fits.
+ As we have shown in this section. the orbits of P2 andP1 are uon-Ixeplerian even if they can be treated as test particles.," As we have shown in this section, the orbits of P2 andP1 are non-Keplerian even if they can be treated as test particles."
+ Nevertheless. the orbital period {5 aud seimiimajor axis « (whicl are iudependent parameters in the fits by BGYYS) in Table 1 should closely resemblethe azimuthaperiod £1 aud guidiug center radius Ry.," Nevertheless, the orbital period $P$ and semimajor axis $a$ (which are independent parameters in the fits by BGYYS) in Table \ref{table1}
+ should closely resemblethe azimuthalperiod $P_0$ and guiding center radius $R_0$ ."
+ With Py=απ]/Gn4in)!anllfpe 35/2 where Dye and dpe are the orbital period aud semimajor axis of Pluto-Charon from Table 1.. we Linc ονο—1=0.0079+0.0038 and. 0.0067+0.0021 for P2 and Pl. respectively.," With $P_K = 2\pi [a^3/G (m_p+m_c)]^{1/2} = P_{pc} (a/a_{pc})^{3/2}$ , where $P_{pc}$ and $a_{pc}$ are the orbital period and semimajor axis of Pluto-Charon from Table \ref{table1}, , we find $P_K/P - 1 = 0.0079
+\pm 0.0038$ and $0.0067 \pm 0.0021$ for P2 and P1, respectively."
+ The values of, The values of
+"As argued in ? and ?7,, the abilityης of the HBI and MTI to reconfigure the magnetic field depends on the level of externally driven turbulence.","As argued in \citet{sharma09} and \citet{ruszkowski10}, the ability of the HBI and MTI to reconfigure the magnetic field depends on the level of externally driven turbulence."
+" In the current case, the turbulence is driven by the structure formation motions."," In the current case, the turbulence is driven by the structure formation motions."
+ Strong turbulence can significantly affect the MTI but some residual radial bias may be present in the field., Strong turbulence can significantly affect the MTI but some residual radial bias may be present in the field.
+" ? showed that even for strong turbulence driving, the MTI operates and significantly contributes to the velocity power spectrum on scales larger than the outer turbulence driving scale Doyz."," \citet{McCourt} showed that even for strong turbulence driving, the MTI operates and significantly contributes to the velocity power spectrum on scales larger than the outer turbulence driving scale $L_{\rm out}$."
+ On scales smaller than [ους the power spectra in the conduction and adiabatic runs are nearly identical., On scales smaller than $L_{\rm out}$ the power spectra in the conduction and adiabatic runs are nearly identical.
+" As these isolated box experiments considered relatively smallouter driving scale (Lout=Mpres/16, where hpres is the pressure scaleheight in a box larger outer driving scales may reduce the Apres/2parameter high),space where the MTI If the Froude parameter is less than unity, then the turbulence does not reset the magnetic field configuration established by the MTI."," As these isolated box experiments considered relatively smallouter driving scale $L_{\rm out}=h_{\rm pres}/16$, where $h_{\rm pres}$ is the pressure scaleheight in a box $h_{\rm pres}/2$ high), larger outer driving scales may reduce the parameter space where the MTI If the Froude parameter is less than unity, then the turbulence does not reset the magnetic field configuration established by the MTI."
+" We define the Froude number here as Fr= co/wpgyL, where o is the gas velocity"," We define the Froude number here as $Fr=\sigma/\omega_{BV}L$ , where $\sigma$ is the gas velocity"
+of this result will be on the possible association of GRBs with global extinctions of biological spices.,of this result will be on the possible association of GRBs with global extinctions of biological spices.
+" These occurred on Earth a factor of 10 times less frequent, about once every 100 Myr."," These occurred on Earth a factor of 10 times less frequent, about once every 100 Myr."
+ These two rates suggest that a typical Galactic GRB pointing to Earth does not cause a major extinction event., These two rates suggest that a typical Galactic GRB pointing to Earth does not cause a major extinction event.
+" We thank Shiho Kobayashi, Ehud Nakar and Elena Rossi for fruitful discussions."," We thank Shiho Kobayashi, Ehud Nakar and Elena Rossi for fruitful discussions."
+" The research was supported by an ERC Advanced Research Grant, by the ISF center for High Energy Astrophysics and by the Israel-France program in Astrophysics erant."," The research was supported by an ERC Advanced Research Grant, by the ISF center for High Energy Astrophysics and by the Israel-France program in Astrophysics grant."
+The rewly redesigned light version of the Space Interferometry Mission. SIALLite. In a long-bascline astrometric interferometer designed to reach sub-nücroarcsecond precision over the course of a five vear nussion.,"The newly redesigned light version of the Space Interferometry Mission, SIM-Lite, is a long-baseline astrometric interferometer designed to reach sub-microarcsecond precision over the course of a five year mission."
+ It has two modes of operation. one for global astrometry with | ceud-ofMission accuracy. aud the other for narrow-anele asroluctry. with a sinele-epoch (1100s visit) accuracybelow 1 citepUinwinds..," It has two modes of operation, one for global astrometry with 4 end-of-mission accuracy, and the other for narrow-angle astrometry, with a single-epoch (1100s visit) accuracybelow 1 \\citep{Unwin08}."
+ The redesigu lowers costs primarily by shortening the baseliue from 91i to 61n and replacing a euide interferometer wi ha telescope star tracker., The redesign lowers costs primarily by shortening the baseline from $9~{\rm m}$ to $6~{\rm m}$ and replacing a guide interferometer with a telescope star tracker.
+ Narrow-angle astrouetry is used to measure the orbital motions of nou-hnuiroux Objects. like neutron stars aud plaucts. ἂν observing the motion of the pareut star.," Narrow-angle astrometry is used to measure the orbital motions of non-luminous objects, like neutron stars and planets, by observing the motion of the parent star."
+ SIN-Lite is desigred to lake narrow angle lueasurenent:4 w alterately switching between a tarect star and anywhere from 3 to 6 reference stars witlin < vl degree radius., SIM-Lite is designed to make narrow angle measurements by alternately switching between a target star and anywhere from 3 to 6 reference stars within a 1 degree radius.
+ The astrometric signal is the m0ion of the arect star relative to the reference star4, The astrometric signal is the motion of the target star relative to the reference stars.
+ To take a concrete example. co idera Calcidate star that is observed 250 times over the cotrse of a 5-vear mission.," To take a concrete example, consider a candidate star that is observed 250 times over the course of a 5-year mission."
+ A peak iu the joint periodogram power distribution (similar to a power spectral density) of the tarec star position constitutes the detection of a planet caididate., A peak in the joint periodogram power distribution (similar to a power spectral density) of the target star position constitutes the detection of a planet candidate.
+ Detection of a planet with a fasC-aar probability requires a signal-to-noise ratio (SNR) of approximately ha., Detection of a planet with a false-alarm probability requires a signal-to-noise ratio (SNR) of approximately 5.8.
+ At 10 pe. the astic1netric signature for an Earth orbiting a oue Soar nuass star at 1 AU is 43pras.. so the SNR τοςmrement calls for a final instrument error of less tran 50 nauo-arcseconcds (uas).," At 10 pc, the astrometric signature for an Earth orbiting a one Solar mass star at 1 AU is 0.3, so the SNR requirement calls for a final instrument error of less than 50 nano-arcseconds (nas)."
+ This corresponds o a siungle-epoch error of ess than 0.52., This corresponds to a single-epoch error of less than 0.82.
+pas.. For 10 nearest ~GO stars. a rolse floor below 35 nas is needed to detec Earths in the habitable zone.," For the nearest $\sim60$ stars, a noise floor below 35 nas is needed to detect Earths in the habitable zone."
+ The key questions. therefore. are whetrer sub-nicroarcsecond siugle-o)xch accuracy is attainable. and. equally iniportant. whether the iustruneut svsteinatic errors do average down well below the sinele-epoch accuracy.," The key questions, therefore, are whether sub-microarcsecond single-epoch accuracy is attainable, and, equally important, whether the instrument systematic errors do average down well below the single-epoch accuracy."
+ This paper reports testbed results relating directly to both the sinele-cpoch ancl eud-of-iuission niuroav anele accuracy., This paper reports testbed results relating directly to both the single-epoch and end-of-mission narrow angle accuracy.
+ The basic eloiueuts of a Alichelson stellar interferometer are shown in Figure 1.., The basic elements of a Michelson stellar interferometer are shown in Figure \ref{fig1}.
+ Starlighto is collected by two spatially separated siderostats., Starlight is collected by two spatially separated siderostats.
+ Each siderostat has au embedded ficlucial tlat marks one cud of the baseline amd also acts as a retroreflector for the mterual aud external ietrologv laser leas., Each siderostat has an embedded fiducial that marks one end of the baseline and also acts as a retroreflector for the internal and external metrology laser beams.
+ Alone each arm. the starlight beam has au anuular footprint while internal metrology occupies the center.," Along each arm, the starlight beam has an annular footprint while internal metrology occupies the center."
+" The quantity of iuterst is the delay ο, which"," The quantity of interest is the delay $x$ , which"
+limits are 21mJy at 100m (86 sources) and 28 mJys at 160um (91 sources).,limits are 21mJy at $\mu$ m (86 sources) and 28 mJys at $\mu$ m (91 sources).
+" We conducted a wide-field spectroscopic survey of this field with Hectospec at MMT in February 2009, confirming cluster membership of 439 galaxies (Hardegree-Ullman et al,"," We conducted a wide-field spectroscopic survey of this field with Hectospec at MMT in February 2009, confirming cluster membership of $439$ galaxies (Hardegree-Ullman et al.,"
+ in prep.), in prep.).
+" In addition, ? obtained redshifts for 630 cluster members with VIMOS/VLT, bringing the total number of unique redshifts for this cluster to 851."," In addition, \citet{Czoske2004}
+ obtained redshifts for 630 cluster members with VIMOS/VLT, bringing the total number of unique redshifts for this cluster to 851."
+" This extensive spectroscopic coverage means that, out of X-ray maps can be used both to separate the AGNs from the starburst galaxies and to identify galaxy groups around clusters, and we have analysed the 0147330201 to characterize the X-ray appearance of A1835."," This extensive spectroscopic coverage means that, out of X-ray maps can be used both to separate the AGNs from the starburst galaxies and to identify galaxy groups around clusters, and we have analysed the 0147330201 to characterize the X-ray appearance of A1835."
+" We decompose the 0.5-2 keV combined image into unresolved and extended emission, following the wavelet technique of ?,, before extracting point source fluxes at the locations of sources."," We decompose the 0.5-2 keV combined image into unresolved and extended emission, following the wavelet technique of \citet{Finoguenov2009}, before extracting point source fluxes at the locations of sources."
+" Additionally, we make use of previously published wide-field data from Spitzer//MIPS 24um and UKIRT/WFCAM J,K bands(?), and Subaru/Suprime-CAM (?).."," Additionally, we make use of previously published wide-field data from /MIPS $\mu$ m and UKIRT/WFCAM J,K \citep{Haines2009}, and Subaru/Suprime-CAM \citep{Okabe2010}."
+" At the redshift of the cluster, z=0.25, the PACS 100 and 160um bands are expected to straddle the peak of the spectral energy distribution for sources with dust temperatures of T 30K. The 1604m/100um color should thus be sensitive to dust temperature differences, with redder objects having cooler dust components."," At the redshift of the cluster, $z=0.25$, the PACS 100 and $\mu$ m bands are expected to straddle the peak of the spectral energy distribution for sources with dust temperatures of T $\sim
+30$ K. The $\mu$ $\mu$ m color should thus be sensitive to dust temperature differences, with redder objects having cooler dust components."
+ We plot observed 160um/100um vs 100um/24um colors for the sources in our cluster in Fig., We plot observed $\mu$ $\mu$ m vs $\mu$ $\mu$ m colors for the sources in our cluster in Fig.
+" 2, and overplot the synthetic colors from commonly used templates from the literature: ? (R09),? (DHO2), ? (L03), ? (CEO1)."," 2, and overplot the synthetic colors from commonly used templates from the literature: \citealt{Rieke2009} (R09),\citealt{Dale2002} (DH02), \citealt{Lagache2003} (L03), \citealt{Chary2001} (CE01)."
+" All the SED templates are poor representations of these galaxies, which appear too red in their 1001m/24um colors. —"," All the SED templates are poor representations of these galaxies, which appear too red in their $\mu$ $\mu$ m colors. –"
+ see ? and ? for a more extended discussion of this 100m excess.," see \citet{Smith2010} and \citet{Rawle2010} for a more extended discussion of this $100\mu{\rm
+ m}$ excess."
+"excess, n order to photonssearc«i for aud characterize any WIA at tlis location.","excess, in order to search for and characterize any WHIM at this location."
+ Our ¢jbservations are from the mosaic made by NAIAENeowte1, Our observations are from the mosaic made by XMM-Newton.
+ dKaastra ct al. (, Kaastra et al. (
+2003) report a simular study of the ceuter of the Coma cluster also using data from he NAIAD imiosaic.,2003) report a similar study of the center of the Coma cluster also using data from the XMM mosaic.
+ Not knowing the relative spatial distributions « he WIIIM and ICAL. it is difficult to predict whether osorving the center or outskirts of a chster will masimize the coutrast of any wari component relative to he donmuuait hot component.," Not knowing the relative spatial distributions of the WHIM and ICM, it is difficult to predict whether observing the center or outskirts of a cluster will maximize the contrast of any warm component relative to the dominant hot component."
+ However both studies do detect a war12 compoucut., However both studies do detect a warm component.
+" We adopt a Coma chster redsift of 0.023. dT,TÜ lau : Land quote crrorbars at the confidence level."," We adopt a Coma cluster redshift of 0.023, $H_{\rm o}=70$ km $^{-1}$ $^{-1}$, and quote errorbars at the confidence level."
+ One degree COLYYCSDOLK sto 1.67 AIpe., One degree corresponds to 1.67 Mpc.
+ The initial results of the NAINENewtou (Jansen et al., The initial results of the XMM-Newton (Jansen et al.
+ 2001) performance verificatio1 observations of the Coma cluster are reported in Bricl et al. (, 2001) performance verification observations of the Coma cluster are reported in Briel et al. (
+2001). Arnaud ot al. (,"2001), Arnaud et al. ("
+2001) and. Neunauu et al. (,2001) and Neumann et al. (
+2001).,2001).
+ In addition to the observations reported there. three additional observations have heen carried out. colpleting the plauned survey of the €ola cluster.," In addition to the observations reported there, three additional observations have been carried out, completing the planned survey of the Coma cluster."
+ We have renained the Coma-bked field frou Bricl et al. (, We have renamed the Coma-bkgd field from Briel et al. (
+2001) Comia-0 in order to avoid confusion with the aditional backeround poiutiugs required to analyze these data.,2001) Coma-0 in order to avoid confusion with the additional background pointings required to analyze these data.
+ Siice. the Coma cluster is a large diffuse source. estimaine the vackeround is non-trivial.," Since the Coma cluster is a large diffuse source, estimating the background is non-trivial."
+ In general we follow he metho of Lumb et al. (, In general we follow the method of Lumb et al. (
+2002).,2002).
+ The observations reported here are perfornie« with the EPIC pu detector (Strüdder ο al., The observations reported here are performed with the EPIC pn detector (Strüdder et al.
+ 2001j uxing its medi filter., 2001) using its medium filter.
+ This confieuration leads to a different detector cotut rate from the cosnüc (CNB) iuxL πι Way N-ray backerouuds in the soft baud. conrpared to the usual thin filter.," This configuration leads to a different detector count rate from the cosmic (CXB) and Milky Way X-ray backgrounds in the soft band, compared to the usual thin filter."
+ The Milkv. Way ckerouinm Lis further comoed of the Local Bubble aud halo COLLvoncuts., The Milky Way background is further composed of the Local Bubble and halo components.
+" We used a LOO ksec observation of the poiu ""OIroe APMOS279|5255 (PIC. Tasinger). pertorned with he suue (Guediunm) filter. to nieasure these backerotud components aud a filter wheel closed observation to measure the internal or particle ackeround."," We used a 100 ksec observation of the point source APM08279+5255 (PI G. Hasinger), performed with the same (medium) filter, to measure these background components and a filter wheel closed observation to measure the internal or particle background."
+ We used the APEC plasina code (Siuith et , We used the APEC plasma code (Smith et al.
+2001) to fit the AGkv Wav backgrouud components ida power-law (DP— 1.15) to fit the CNB., 2001) to fit the Milky Way background components and a power-law $\Gamma=1.45$ ) to fit the CXB.
+ The clement abundance in the APEC code was fixed to solar and teiiperatfure was set to 0.07 keV for the Local Bubble and left as a free parameter for the halo. viclding a 0.22 keV best-fit vahe.," The element abundance in the APEC code was fixed to solar and temperature was set to 0.07 keV for the Local Bubble and left as a free parameter for the halo, yielding a 0.22 keV best-fit value."
+ To ensure the compatibility. of the CNB. we excised from the spectral extraction region the APAI quasar and the N-ray sources identified with Coma ealaxies Gvhich οherwise iucrease the apparent CNB fiux by J).," To ensure the compatibility of the CXB, we excised from the spectral extraction region the APM quasar and the X-ray sources identified with Coma galaxies (which otherwise increase the apparent CXB flux by )."
+ Since t1ος 1s a nudsmiatch iu the Nyy toward the Coma and the APM fields. we use the same spectral shape aud πο1 of he halo and CNB components. but chaiced their Nyy fro1 3.90.3.102 7 (APM) to 9«1017 2 (Coma).," Since there is a mismatch in the $N_H$ toward the Coma and the APM fields, we use the same spectral shape and normalization of the halo and CXB components, but changed their $N_H$ from $3.9\pm0.3\times10^{20}$ $^{-2}$ (APM) to $9\times10^{19}$ $^{-2}$ (Coma)."
+ The Local Bubble is not subject to this absorption., The Local Bubble is not subject to this absorption.
+ We aso checked that the derived results ou he sof Nav chussion from Coma do not depend ol. the posside laree-scae variation in the imtensitv of the Local Buble componcut., We also checked that the derived results on the soft X-ray emission from Coma do not depend on the possible large-scale variation in the intensity of the Local Bubble component.
+ For the spectral aalysis we use the 0.16 keV raiege aud for tje estimate of the detector backgerouud the 15 keV range., For the spectral analysis we use the 0.4–6 keV range and for the estimate of the detector background the 9.5--15 keV range.
+ We do iot use the energies below 0.1 keV. to avoid a complicated treatment of detector backerotnd associated with the eectronic noise (Limb et al.," We do not use the energies below 0.4 keV, to avoid a complicated treatment of detector background associated with the electronic noise (Lumb et al."
+ 2002)., 2002).
+ This prechiCR us frou studving the very soft enission., This precludes us from studying the very soft emission.
+ Although. we could potentially use the energies up to LO keV in the analvsis. we are able to suff&icieutlv. constraiu the hot conr20rent of the Coma cluster already. without using these csnereyv bauds. where removal of the detector background IOC'Onmies critical.," Although, we could potentially use the energies up to 10 keV in the analysis, we are able to sufficiently constrain the hot component of the Coma cluster already without using these energy bands, where removal of the detector background becomes critical."
+ The viguetlug οςnrection is performed taking the source exten aid a receut pu telescope (NRT3) vignetting calibration {Ταmb et al., The vignetting correction is performed taking the source extent and a recent pn telescope (XRT3) vignetting calibration (Lumb et al.
+ 2003) into account. which is mostly import:ut for the absolute flux determination. elven our clicco of the oeucrev range.," 2003) into account, which is mostly important for the absolute flux determination, given our choice of the energy range."
+" Systematic nuicertainty of he fiux deternunation is for both pu aud MOS,", Systematic uncertainty of the flux determination is for both pn and MOS.
+ Iu acditiou to a thorough evaluation of the background. we furtrer reduce the svstematics of our study by using the results of recent calibration efforts of the pu team of MPE oi soft energv response of the pu detector (Iliberl et al.," In addition to a thorough evaluation of the background, we further reduce the systematics of our study by using the results of recent calibration efforts of the pn team of MPE on soft energy response of the pn detector (Haberl et al."
+ 2002. NADTAISAS-5.L release).," 2002, XMMSAS-5.4 release)."
+ Furtherinore. to improve on the detection statistics of the Ward cussion and to reuce the systematic effects of subtraction of the cluster enission. we concentrate on the of the Coma cluster.," Furthermore, to improve on the detection statistics of the warm emission and to reduce the systematic effects of subtraction of the cluster emission, we concentrate on the of the Coma cluster."
+Basie information about the sample galaxies is given in pe,Basic information about the sample galaxies is given in \\ref{T_morph_AG} and \ref{T_morph_NG}.
+istaken fromtheT hird uC [Bright MetR," The morphological type is taken from the Third Reference Catalogue of Bright Galaxies \citep[RC3,][]{RC3}."
+t profial.199] ," If there is no classification in RC3 or it is doubtful, we take the morphological type given in NED; if none of the above types exists, we list the one in \citep[][]{PPP_03} or SIMBAD."
+The Sy sample observations were performed at NAO with the 2-m Ritchey-Chréttien telescope equipped with 1024x Photometrics AT200 CCD camera (CCD chip SITe SIO03AB with a square pixel size of 244m that corresponds to 07309 on the sky) or with 1340x1300 Princeton Instruments VersArray:1300B CCD camera (CCD chip EEV CCD36-40 with a square pixel size of 204m that corresponds to (0/258 on the sky).," The Sy sample observations were performed at NAO with the 2-m Ritchey-Chréttien telescope equipped with $1024\,$$\times$$\,1024$ Photometrics AT200 CCD camera (CCD chip SITe SI003AB with a square pixel size of $24\,\mu\rm m$ that corresponds to $0\farcs309$ on the sky) or with $1340\,$$\times$$\,1300$ Princeton Instruments VersArray:1300B CCD camera (CCD chip EEV CCD36-40 with a square pixel size of $20\,\mu\rm m$ that corresponds to $0\farcs258$ on the sky)."
+ Two galaxies were observed employing a two-channel focal reducer (Jockersetal.2000) and 512«512 Princeton Instruments VersArray:512B CCD camera (CCD chip EEV CCD77-00 with a square pixel size of 24m. which corresponds to 07884 on the sky).," Two galaxies were observed employing a two-channel focal reducer \citep{JCB_00} and $512\,$$\times$$\,512$ Princeton Instruments VersArray:512B CCD camera (CCD chip EEV CCD77-00 with a square pixel size of $24\,\mu\rm m$, which corresponds to $0\farcs884$ on the sky)."
+ Standard Johnson-Cousins BVRclc filters were used., Standard Johnson-Cousins $BVR_{\rm \scriptstyle C}I_{\rm \scriptstyle C}$ filters were used.
+ For a couple of objects we were not able to obtain images of good quality and used archival data., For a couple of objects we were not able to obtain images of good quality and used archival data.
+ The observation log (including archival data) is presented in Table 3.., The observation log (including archival data) is presented in Table \ref{T_obs}.
+ Multiple frames were acquired for all objects of interest., Multiple frames were acquired for all objects of interest.
+ Standard fields were observed two or three times each night at different airmass values not exceeding X 22. The galaxy fields were observed in the same airmass range.," Standard fields were observed two or three times each night at different airmass values not exceeding $X\,$ $\,2$ The galaxy fields were observed in the same airmass range."
+ Zero-exposure frames were taken regularly during the observing runs. and 7c frames of relatively blank night sky regions were taken for fringe frame composition.," Zero-exposure frames were taken regularly during the observing runs, and $I_{\rm \scriptstyle C}$ frames of relatively blank night sky regions were taken for fringe frame composition."
+ Flat field frames were taken in the morning and/or evening twilight., Flat field frames were taken in the morning and/or evening twilight.
+ Both flat fields and {ο blank frames were offset from one exposure to the next. so that stars could be filtered out.," Both flat fields and $I_{\rm \scriptstyle C}$ blank frames were offset from one exposure to the next, so that stars could be filtered out."
+ The primary reduction of the data. as well as most of the surface photometry. was performed by means of packages within the environment.," The primary reduction of the data, as well as most of the surface photometry, was performed by means of packages within the environment."
+ The mean bias level was estimated using the virtual pre-scan bias section., The mean bias level was estimated using the virtual pre-scan bias section.
+ The median of the zero-exposure frames was used to remove the residual bias pattern., The median of the zero-exposure frames was used to remove the residual bias pattern.
+ Dark current correction was not needed., Dark current correction was not needed.
+ The median of the normalized flat field frames in each passband was used for flat fielding., The median of the normalized flat field frames in each passband was used for flat fielding.
+ Median combined blank {ο frames were used for fringe correction: defringing increases the signal-to-noise ratio (S/N) in the outer galaxy regions and allows better estimation of sky background., Median combined blank $I_{\rm \scriptstyle C}$ frames were used for fringe correction; defringing increases the signal-to-noise ratio (S/N) in the outer galaxy regions and allows better estimation of sky background.
+ Cosmic ray events and bad pixels were replaced by a local median value (ricEg;cosMIC command)., Cosmic ray events and bad pixels were replaced by a local median value $\rm \scriptstyle FILTER/COSMIC$ command).
+ All images of a particular object were aligned to match the highest S/N frame. generally the Rc one. and combined to obtain single BVRe/¢ images.," All images of a particular object were aligned to match the highest S/N frame, generally the $R_{\rm \scriptstyle C}$ one, and combined to obtain single $BVR_{\rm \scriptstyle C}I_{\rm \scriptstyle C}$ images."
+ Primary reduction refers both to galaxy and standard field frames., Primary reduction refers both to galaxy and standard field frames.
+ To construct. the point. spread function. (PSF) of the combined frames. we picked up à number of bright. isolated. and unsaturated stars. fitted à 2D Moffat profile (Moffat1969) to them. and weight-averaged the full widths at half maximum (FWHMs) and power-law indices. B. in each passband.," To construct the point spread function (PSF) of the combined frames, we picked up a number of bright, isolated, and unsaturated stars, fitted a 2D Moffat profile \citep{M_69} to them, and weight-averaged the full widths at half maximum (FWHMs) and power-law indices, $\beta$, in each passband."
+ The minimal FWHM (over all passbands) and the corresponding f BiefWsted Cable ic, The minimal FWHM (over all passbands) and the corresponding $\beta$ are listed in Table \ref{T_obs}.
+ dt Chea da Comkuraner (Trujilloetal.2001).. and if 6=1. then Moffat profile Lorentzian.," If $\beta\,$$\rightarrow$$\,\infty$, then Moffat $\,\rightarrow\,$ Gaussian \citep{TAC_01}, and if $\beta=1$, then Moffat $\,\equiv\,$ Lorentzian."
+ We used archival data for the control sample., We used archival data for the control sample.
+ For abou—x half of the galaxies CCD data from the SDSS. the Europea35 Southern Observatory (ESO).NED. and the Isaac Newto2D Group of Telescopes (ING) image archives were used.," For about half of the galaxies CCD data from the SDSS, the European Southern Observatory (ESO),NED, and the Isaac Newton Group of Telescopes (ING) image archives were used."
+ SDSS uses a dedicated 2.5-m telescope and a large-format CCD camera (Gunnetal.1998.2006) at the Apache Por= Observatory in New Mexico to obtain images in five broad bands (ugriz.Fukugitaetal.1996).," SDSS uses a dedicated 2.5-m telescope and a large-format CCD camera \citep{GCR_98,GSM_06} at the Apache Point Observatory in New Mexico to obtain images in five broad bands \citep[$ugriz$,][]{FIG_96}."
+. The imaging data were processed by a photometric pipeline (Luptonetal.2001;Stoughtonetal. 2002).," The imaging data were processed by a photometric pipeline \citep{LGI_01,SLB_02}."
+. The primary reduction of the data from the ING and ESO archives was performed as described above., The primary reduction of the data from the ING and ESO archives was performed as described above.
+ The data taken through NED were reduced by the corresponding authors., The data taken through NED were reduced by the corresponding authors.
+ Digitized Palomar Observatory Sky Survey III. and digitized ESO-Uppsala Survey (Lauberts&Valentijn1989) data were used for the rest of the inactive galaxies.," Digitized Palomar Observatory Sky Survey II, and digitized ESO-Uppsala Survey \citep{LV_89} data were used for the rest of the inactive galaxies."
+" Data sources of the inactive galaxies are listed in refT,,orpliyG.", Data sources of the inactive galaxies are listed in \\ref{T_morph_NG}.
+.Thesuperveningdatareductionde scribedbelowconcernshe ," The supervening data reduction described below concerns both galaxy samples, while calibration is applied only to the Sy galaxies."
+To increase the S/N of the galaxy images we applied the adaptive filter technique (Lorenzetal.1993).. implemented in the Astrophysical Institute of Potsdam (AIP) package.," To increase the S/N of the galaxy images we applied the adaptive filter technique \citep{LRC_93}, implemented in the Astrophysical Institute of Potsdam (AIP) package."
+ The task uses H-transform to calculate the local S/N at each point of the image and determines the size of the impulse response of the filter at this point., The task uses H-transform to calculate the local S/N at each point of the image and determines the size of the impulse response of the filter at this point.
+ Thus. adaptive filter extensively smooths sky background. to a lesser extent galaxy outskirts. and does not treat the highest resolution features.," Thus, adaptive filter extensively smooths sky background, to a lesser extent galaxy outskirts, and does not treat the highest resolution features."
+ Adaptive filter has many advantages over the other most commonly used filters in surface photometry of galaxies (Slavcheva-Mihovaetal.2005): furthermore. it introduces no systematic errors (Capaeciolietal.1988;Tamm&Tenjes 2001).," Adaptive filter has many advantages over the other most commonly used filters in surface photometry of galaxies \citep{SMP_05}; furthermore, it introduces no systematic errors \citep{CHL_88,TT_01}."
+. Aside from intrinsic structures. like dust lanes. star formation regions. etc..," Aside from intrinsic structures, like dust lanes, star formation regions, etc.,"
+ isophotal shape can also be modified in a systematic fashion by contaminating features: contaminating objects (foreground stars or projected/companion galaxies) or other features (diffraction spikes from. bright. stars. scattered light. meteor or satellite trails. ete.).," isophotal shape can also be modified in a systematic fashion by contaminating features: contaminating objects (foreground stars or projected/companion galaxies) or other features (diffraction spikes from bright stars, scattered light, meteor or satellite trails, etc.)."
+ The closer the contaminating feature to the galaxy core. the more carefully we treated it.," The closer the contaminating feature to the galaxy core, the more carefully we treated it."
+We used the following techniques to. clean. out contaminating features: interpolation. PSF subtraction. symmetric replacement. deblending. and annular cleaning. Interpolation.,"We used the following techniques to clean out contaminating features: interpolation, PSF subtraction, symmetric replacement, deblending, and annular cleaning. ."
+ This is the technique used most often., This is the technique used most often.
+ The algorithm m use (within the AIP package) generates masks, The algorithm in use (within the AIP package) generates masks
+wighbors are relatively strougly weighted while the lonec-rauege neighbors have almost no weight.,neighbors are relatively strongly weighted while the longe-range neighbors have almost no weight.
+ So we expect that the so-defined stereoscopic error monotonously decreases with short-rauge weighting towards üeher power indices p21., So we expect that the so-defined stereoscopic error monotonously decreases with short-range weighting towards higher power indices $p>1$.
+ In Fie., In Fig.
+ 10 we show the stereoscopic error calculated with Eqs. (, 10 we show the stereoscopic error calculated with Eqs. (
+"15-16) as a ""uctiou of the power index from p—1 to p=20. for all 1 active regions.","15-16) as a function of the power index from $p=1$ to $p=20$, for all 4 active regions."
+ Indeed. the stereoscopic error uonotonouxlv decreases from p=1 to p=12. because we give proeressivolv less weight to the loug-rauge wisaliguinents.," Indeed, the stereoscopic error monotonously decreases from $p=1$ to $p=12$, because we give progressively less weight to the long-range misalignments."
+ However. frou p=12 towards p=20 the error iuereases again. probably because of the inhomogeneity of the closest neighbors at the shortest range aud the excessive weighting to the very nearest wighbors.," However, from $p=12$ towards $p=20$ the error increases again, probably because of the inhomogeneity of the closest neighbors at the shortest range and the excessive weighting to the very nearest neighbors."
+ However. the plateau iu the range of p225.15 is a good indication that we measure a stable value at the πμ of p 12.," However, the plateau in the range of $p\approx 5-15$ is a good indication that we measure a stable value at the minumum of $p\approx 12$ ."
+" From this we obtain a imisaliguiment contribution of Aasg=9.17 (2007 Apr 30). 5.77 (2007 Mav 9). 11.5° (2007 May. 19). aud 7.2"" (2007 Dec 11). as listed in Table 2."," From this we obtain a misalignment contribution of $\Delta \alpha_{SE}=9.4^\circ$ (2007 Apr 30), $5.7^\circ$ (2007 May 9), $11.5^\circ$ (2007 May 19), and $7.2^\circ$ (2007 Dec 11), as listed in Table 2."
+ A diagram of the various isaliguiment angles is shown in Fie., A diagram of the various misalignment angles is shown in Fig.
+ Ll. which we plot as a function of the iuaxinunun GOES soft N-rav flux during the observing interval.," 11, which we plot as a function of the maximum GOES soft X-ray flux during the observing interval."
+ The GOES flux was dominanted by the cnussion frou the active regions analyzed here. since there was no other comparable active region prescut on the soli disk during the times of the observations.," The GOES flux was dominanted by the emission from the active regions analyzed here, since there was no other comparable active region present on the solar disk during the times of the observations."
+ The case of 2007 May 9 with the lowest GOES shows a single dipolar structure. while the case of 2007 May 19 with the highest COLES fux exhibits 1ultiple dipolar groups.," The case of 2007 May 9 with the lowest GOES shows a single dipolar structure, while the case of 2007 May 19 with the highest GOES flux exhibits multiple dipolar groups."
+ The nüsalisuuieut angles aprss indicate those obtained with the code. Oxprr those with the code (ouly for 2007 Apr 30). ape those optimized with the code. and Adsr indicates the stereosocpic errors.," The misalignment angles $\alpha_{PFSS}$ indicate those obtained with the code, $\alpha_{NLFF}$ those with the code (only for 2007 Apr 30), $\alpha_{PFU}$ those optimized with the code, and $\Delta \alpha_{SE}$ indicates the stereosocpic errors."
+ As a working hvpothnesis we mieht attribute the remaining residuals àvp to the non-potentiality of the active regious. for which we measure the values: ayp=L9 (2007 Apr 30). ayp=5.7 (2007 Mav 9). ayp=85 (2007 Mav 19). aud àyp=7.2 (2007 Dec 11). listed also in Table 2.," As a working hypothesis we might attribute the remaining residuals $\alpha_{NP}$ to the non-potentiality of the active regions, for which we measure the values: $\alpha_{NP}=4.9$ (2007 Apr 30), $\alpha_{NP}=5.7$ (2007 May 9), $\alpha_{NP}=8.8$ (2007 May 19), and $\alpha_{NP}=7.2$ (2007 Dec 11), listed also in Table 2."
+ We consider these values as a new method o quantify the noupotenutialiv of active regions., We consider these values as a new method to quantify the nonpotentialiy of active regions.
+" We find that quiesceut active regious that contain simple dipoles (2007 Apr 30 and Mav 9) have sinall iuisalieunieuts of aypz57.67 with best-fit potential field uodels. while more complex active regions (2007 May. 19 aud Dec 11) have somewhat larger non-poteutial uisaligeniienuts iu the order of z7""9"" compared with best-fit potential field models."," We find that quiescent active regions that contain simple dipoles (2007 Apr 30 and May 9) have small misalignments of $\alpha_{NP} 
+\approx 5^\circ-6^\circ$ with best-fit potential field models, while more complex active regions (2007 May 19 and Dec 11) have somewhat larger non-potential misalignments in the order of $\approx 7^\circ-9^\circ$ compared with best-fit potential field models."
+ While we express thee deeree of noupoteutialitv in terms of a eau niüsaligument anele here. other ucasures are the ratios of the total noupoteutial to the poteutial energy in au active region. which amounts up to EFyp/Ep=1.32 in one flaring active region (Schrijver ct al.," While we express thge degree of nonpotentiality in terms of a mean misalignment angle here, other measures are the ratios of the total nonpotential to the potential energy in an active region, which amounts up to $E_{NP}/E_P \lapprox 1.32$ in one flaring active region (Schrijver et al."
+ 2008)., 2008).
+ Since the four investigated active regions have quite ciffereut misalignnieut aueles. aud also the inferred uoupoteutialitv of the magnetic field varies sguificautlv. we quantity the activity level of the active regious from the soft N-rav. flux measured by GOES.," Since the four investigated active regions have quite different misalignment angles, and also the inferred nonpotentiality of the magnetic field varies significantly, we quantify the activity level of the active regions from the soft X-ray flux measured by GOES."
+ Fie., Fig.
+ 12 shows the GOES light curves for the 1 active regions during the times of stereoscopic triangulation., 12 shows the GOES light curves for the 4 active regions during the times of stereoscopic triangulation.
+ The lowest GOES flux is measured for the AR of 2007 Mav 9. with a level of 10.9 Wan 2 (GOES class AL). while the highest GOES flux is measured for the AR of 2007 May19. which has a flare occurring during tlhe observing period with a GOES flux of 1059 Wan 7 (GOES," The lowest GOES flux is measured for the AR of 2007 May 9, with a level of $10^{-7.6}$ W $^{-2}$ (GOES class A4), while the highest GOES flux is measured for the AR of 2007 May19, which has a flare occurring during the observing period with a GOES flux of $10^{-6.0}$ W $^{-2}$ (GOES"
+and Eq.,and Eq.
+ 17 simplifies to We substitute Eq., \ref{g_inside} simplifies to We substitute Eq.
+ 26 into Eq., \ref{t_vs_gchi_limit} into Eq.
+ 27 to get the limiting Lorentz [actor of the fluid element as |—0: which is greater only bv a numerical [actor thanthe initial Lorentz factor 54 that the [Iuid element received right after being shocked., \ref{g_inside_limit} to get the limiting Lorentz factor of the fluid element as $t\rightarrow 0$: which is greater only by a numerical factor thanthe initial Lorentz factor $\gamma _{0}$ that the fluid element received right after being shocked.
+ To relate the limiting p. n to py. ny. we likewise take Eqs. L8..," To relate the limiting $p$, $n$ to $p_0$, $n_0$ , we likewise take Eqs. \ref{f_inside},"
+ 19. in the limit |gX|>>1 and use Eqs. 26..," \ref{h_inside} in the limit $|g\chi|\gg 1$ and use Eqs. \ref{t_vs_gchi_limit},"
+ 27 with the results to get which again differ only bv numerical factors from their values just after the fluid element is shocked., \ref{g_inside_limit} with the results to get which again differ only by numerical factors from their values just after the fluid element is shocked.
+ ‘This is consistent with the behavior given above by simple physical considerations., This is consistent with the behavior given above by simple physical considerations.
+ For the analogous calculationof (he asymptotic profiles of 5. p. ancl n. we cannot simply apply Eqs. 10.. 11.. 12::," For the analogous calculationof the asymptotic profiles of $\gamma$, $p$, and $n$, we cannot simply apply Eqs. \ref{gammadef}, \ref{pdef}, \ref{ndef}: :"
+ Eqs. 8..," Eqs. \ref{char_powerlaws},"
+ 9 require that 4——2x evervwhere behind (he shock and I. P. and .N diverge as /—0.," \ref{chi} require that $\chi\rightarrow -\infty$ everywhere behind the shock and $\Gamma$, $P$, and $N$ diverge as $t\rightarrow 0$."
+ Instead we take the /—Q limit al a fixed position vr., Instead we take the $t\rightarrow 0$ limit at a fixed position $x$.
+ First we have With Eqs. 10.. ," First we have With Eqs. \ref{gammadef}, ,"
+"27. this gives and Thisis consistent with our «qualitative discussion: the coefficient in(he qualitative relation is a numerical [actor times the constant (—) ""/I?.", \ref{g_inside_limit} this gives and Thisis consistent with our qualitative discussion; the coefficient inthe qualitative relation is a numerical factor times the constant $(-R)^{-m}/\Gamma ^{2}$ .
+ For the p and » profiles.," For the $p$ and $n$ profiles,"
+Radio-loud (RL) narrow-line Seyfert 1 (NLSI) active galactic nuclei (AGNs) have received considerable attention in recent years. since they pose a challenge to current unified schemes (Urry&Padovani1995).,"Radio-loud (RL) narrow-line Seyfert 1 (NLS1) active galactic nuclei (AGNs) have received considerable attention in recent years, since they pose a challenge to current unified schemes \citep{Urry1995}."
+. NLSI AGNs are characterized by an optical spectrum with narrow permitted lines FWHM(Hf) «2000 km/s. a ratio of the equivalent widths of [OILILU5007 to Hf smaller than 3. and a bump caused by Fell (see.e.g..Pogge2000.forà review)..," NLS1 AGNs are characterized by an optical spectrum with narrow permitted lines $\beta$ ) $< 2000$ km/s, a ratio of the equivalent widths of $\lambda$ 5007 to $\beta$ smaller than 3, and a bump caused by FeII \citep[see, e.g.,][for a review]{Pogge2000}."
+ They also exhibit prominent soft X-ray excesses (Bolleretal.1996;Grupe2004).," They also exhibit prominent soft X-ray excesses \citep{Boller1996,Grupe2004}."
+. These properties are indicative of very high (near Eddington) accretion rates and relatively low black hole masses (10*— 10*°M..)., These properties are indicative of very high (near Eddington) accretion rates and relatively low black hole masses $10^6 - 10^8 M_\odot$ ).
+ Only a small percentage (~ 7%) of NLSI are radio-loud., Only a small percentage $\sim 7\%$ ) of NLS1 are radio-loud.
+ In these cases. their flat radio spectra and VLBI variability suggest that several of them could host relativistic jets (Zhouetal.2003:Komossaal.2006:Doiet 2006).," In these cases, their flat radio spectra and VLBI variability suggest that several of them could host relativistic jets \citep{Zhou2003,Komossa2006,Doi2006}."
+. The detection by Fermi/LAT of gamma-rays from a handful of RL-NLSI has provided definitive confirmation of the presence of relativistic jets 1n. these sources (Abdoetal.2009b.d:Foschinietal. 2010:).," The detection by Fermi/LAT of gamma-rays from a handful of RL-NLS1 has provided definitive confirmation of the presence of relativistic jets in these sources \citep{discovery,sample,Foschini2010}."
+. A multiwavelength (MWL) campaign was organized in March-July 2009 to investigate the properties of PMN 70948-0022 with simultaneous observations covering gamma-rays. X-rays. UV. optical. and radio.," A multiwavelength (MWL) campaign was organized in March-July 2009 to investigate the properties of PMN J0948+0022 with simultaneous observations covering gamma-rays, X-rays, UV, optical, and radio."
+This campaign has delivered some interesting. results on the emission processes in this z=—0.585 source (Abdoetal.2009c).,This campaign has delivered some interesting results on the emission processes in this $z=0.585$ source \citep{mwl}.
+. The direct observation of simultaneous variability across the electromagnetic spectrum has confirmed the identification of the gamma-ray source with the lower energy counterpart., The direct observation of simultaneous variability across the electromagnetic spectrum has confirmed the identification of the gamma-ray source with the lower energy counterpart.
+ This coordinated variability is also characteristic of the relativistic jets present in the central regions of blazars., This coordinated variability is also characteristic of the relativistic jets present in the central regions of blazars.
+ Overall. the multiwavelength observational properties of J0948--0022 seem therefore to indicate that the physics of RL-NLSI and blazars are similar. and in particular a relationship between RL-NLSI and the powerful. high aceretion. flat-spectrum radio quasars.," Overall, the multiwavelength observational properties of J0948+0022 seem therefore to indicate that the physics of RL-NLS1 and blazars are similar, and in particular a relationship between RL-NLS1 and the powerful, high accretion, flat-spectrum radio quasars."
+ In this letter. we present additional details of the very long baseline interferometry (VLBI) observations carried out during this campaign.," In this letter, we present additional details of the very long baseline interferometry (VLBI) observations carried out during this campaign."
+ Section 2 describes the procedures of the real-time e-VLBI observations (Szomoru2008).. which were the first application of this technique using a global array of antennas m an astronomical program.," Section 2 describes the procedures of the real-time e-VLBI observations \citep{Szomoru2008}, which were the first application of this technique using a global array of antennas in an astronomical program."
+ Section 3 presents the results. and Sect.," Section 3 presents the results, and Sect."
+ 4 discusses conclusions and prospects., 4 discusses conclusions and prospects.
+ The observations took place on 2009 April 21. May 23. June 10. and July 4 using a subset of the European VLBI Network (EVN) and the Long Baseline Array (LBA) in Australia. and the 34mm Kashima telescope in Japan.," The observations took place on 2009 April 21, May 23, June 10, and July 4 using a subset of the European VLBI Network (EVN) and the Long Baseline Array (LBA) in Australia, and the m Kashima telescope in Japan."
+ For a complete list of participating telescopes and other observational parameters. we refer to Table 1..," For a complete list of participating telescopes and other observational parameters, we refer to Table \ref{t.log}."
+ The initial observations on 21 April were, The initial observations on 21 April were
+onlv relevant for smaller 6. which represents a highly eccentric disk.,"only relevant for smaller $b$, which represents a highly eccentric disk."
+ This example shows how an elliptical wall modilies (he SED. however. only when observations and/or simulations solve the structure of the wall. we will be able (to quantify this effect in the modeling of Colxu Tau/4.," This example shows how an elliptical wall modifies the SED, however, only when observations and/or simulations solve the structure of the wall, we will be able to quantify this effect in the modeling of CoKu Tau/4."
+ Now. we VA changes of the binary svstemeccentricity. e. lor ὁ=θα =0 cos?=0.5 and f=0.28a.," Now, we consider changes of the binary systemeccentricity, $e$ for $\phi=0.5$ $\alpha=0$ , $\cos i=0.5$ and $h=0.28a$."
+" Figure 20. shows the SED of the fiducial model ancl of models with e=0..0.2.0.4. n 0.6 calculated with A, from Artvmowiez&Lubow(1994)."," Figure \ref{fig-Draine-4Myr-exc} shows the SED of the fiducial model and of models with $e = 0., 0.2, 0.4,$ and $0.6$ calculated with $R_{cb}$ from \citet{Artymowicz1}."
+.. As € increases. so does f.," As $e$ increases, so does $R_{cb}$."
+ This figure shows (hat large values of e imply SEDs Irom colder walls. farther away from the stars. which show no LO ja silicate feature.," This figure shows that large values of $e$ imply SEDs from colder walls, farther away from the stars, which show no 10 $\mu m$ silicate feature."
+" The wall surface temperature varies [rom 136.5A [or the smallest. 22... and 80.7A. lor the largest. /?,,,."," The wall surface temperature varies from $136.5\,K$ for the smallest $R_{cb}$, and $80.7\,K$ for the largest $R_{cb}$."
+ In Table 5. we label these models as Series Ecc. where all the parameters are shown.," In Table \ref{table-models} we label these models as Series Ecc, where all the parameters are shown."
+" Figure 21. shows the effects of changing the angle a. lor ὁ=0.5.cos7 =0.5.24, =1.τα and ff=0.28a."," Figure \ref{fig-Draine-4Myr-alfa} shows the effects of changing the angle $\alpha$, for $\phi=0.5$ $\cos i=0.5$ $R_{cb}=1.7a$ and $h=0.28a$."
+ Results are shown for a=0.0.25.0.5. 1.5.1.75.," Results are shown for $\alpha={0,0.25,0.5}$, ${1.5,1.75}$ ."
+ In order to see the clifferences in terms of e. we show results lor e=0 and e=0.2.," In order to see the differences in terms of $e$, we show results for $e=0$ and $e=0.2$."
+ For symmetry arguments (will respect to the ó=0 axis). the configurations for a=0.75 and 1.25 correspond to the same flux.," For symmetry arguments (with respect to the $\phi=0$ axis), the configurations for $\alpha=0.75$ and $1.25$ correspond to the same flux."
+ The same can be said. for the a=0.25 and 1.75 cases.," The same can be said, for the $\alpha=0.25$ and $1.75$ cases."
+ Moreover. the flux for all these cases is similar because the stars. masses are not so different and the binary orbit is circular.," Moreover, the flux for all these cases is similar because the stars' masses are not so different and the binary orbit is circular."
+ The configurations for a=0 and 1 correspond to the same flux. thus. the latter is not shown.," The configurations for $\alpha=0$ and $1$ correspond to the same flux, thus, the latter is not shown."
+ The minimum [lux is given in the a=0.5 model. where the disk is inclined around a line perpendicular to the line that connects both stars. for a=0.5 and a=1.5 the μη band disappears.," The minimum flux is given in the $\alpha=0.5$ model, where the disk is inclined around a line perpendicular to the line that connects both stars, for $\alpha=0.5$ and $\alpha=1.5$ the $10\mu m$ band disappears."
+ The maximum flux occurs in the a=0 model (or a= 1)., The maximum flux occurs in the $\alpha=0$ model (or $\alpha=1$ ).
+" Changing a does nol mocdilv the area of the emitting region. since it only rotates this region in the plane of the sky,"," Changing $\alpha$ does not modify the area of the emitting region, since it only rotates this region in the plane of the sky."
+ Llowever. since the stars are not at the center of the wall. when a changes the distances ol both stars to the visible area of the wall also change. and thus the resulting SED changes.," However, since the stars are not at the center of the wall, when $\alpha$ changes the distances of both stars to the visible area of the wall also change, and thus the resulting SED changes."
+ We can conclude from Figure 21. that the cases are divided in (wo classes., We can conclude from Figure \ref{fig-Draine-4Myr-alfa} that the cases are divided in two classes.
+ The first one does not have a 10401. band and the second one has this feature., The first one does not have a $10\mu m$ band and the second one has this feature.
+ These models. where a changes. for cases e=0 and e=2 are shown in Table 5. as Series a.ecc.," These models, where $\alpha$ changes, for cases $e=0$ and $e=2$ are shown in Table \ref{table-models} as Series $\alpha,ecc$."
+ Figure 22 shows the SED when 7 and fy varies. [or e= 0.6= 0.5.9=0. and H4=1.7.," Figure \ref{fig-Draine-4Myr-i-h} shows the SED when $i$ and $h$ varies, for $e=0$ $\phi=0.5$ $\alpha=0$, and $R_{cb}=1.7a$."
+ The left plot corresponds (of=0.280 and cos;=0.3.0.5.0.7. and the right one is for cos;=0.5 and h=0.14.0.23.0.56.," The left plot corresponds to $h=0.28a$ and $\cos{i}=0.3,0.5,0.7$, and the right one is for $\cos{i}=0.5$ and $h=0.14,0.28,0.56$."
+ The SED depends weakly on 7. but. it has a strong dependence on h. namely. (ie flux increases for larger .," The SED depends weakly on $i$, but, it has a strong dependence on $h$, namely, the flux increases for larger $h$."
+ The reason for the former. is (hat for a perfectly. vertical wall. the flux does not change monotonically with 7.," The reason for the former, is that for a perfectly vertical wall, the flux does not change monotonically with $i$ ."
+ Of course.this effect will change if a realistic shape of the wall is taken.," Of course,this effect will change if a realistic shape of the wall is taken."
+ However. at this moment there are no studies (as far as weknow) which define this shape (see Section 6)). thus. we assume a vertical wall.," However, at this moment there are no studies (as far as weknow) which define this shape (see Section \ref{sec-conclusions}) ), thus, we assume a vertical wall."
+ Also. we are ignoring possible oscillation modes in the wall due to the tidal potential ofthe binary.," Also, we are ignoring possible oscillation modes in the wall \citep{Hayasaki} due to the tidal potential ofthe binary."
+ Η the disk is viewed pole-on there, If the disk is viewed pole-on there
+A wav to describe some global features of our Galaxy is fo perform radio observations that span large areas of the sky.,A way to describe some global features of our Galaxy is to perform radio observations that span large areas of the sky.
+ Below 1.1 CIIz. there are few survevs with large coverage and all-skv. maps such as the £08 MITz inaps of ?T7 aud the 1120 MIIz maps of ??.. and ?..," Below 1.4 GHz, there are few surveys with large coverage and all-sky maps such as the 408 MHz maps of \citet{Haslam1981,Haslam1982} and the 1420 MHz maps of \citet{Reich1982,Reich1986}, , and \citet{Reich2001}."
+ The 15 MITz survey presented in this paper covers of the sky., The 45 MHz survey presented in this paper covers of the sky.
+ This frequency rauge ]xobes mostly the svuchrotron cluission from ligh energy electrons interacting with the large-scale magnetic field of the Galaxy aud the spectral information derived at these frequencies is rolated to the spectral energy distribution of these relativistic clectrous., This frequency range probes mostly the synchrotron emission from high energy electrons interacting with the large-scale magnetic field of the Galaxy and the spectral information derived at these frequencies is related to the spectral energy distribution of these relativistic electrons.
+ Spectral index maps are preseuted in? between the frequencies. LOS. 1120. and 22800. MITz.," Spectral index maps are presented in \citet{Reich2004} between the frequencies 408, 1420, and 22800 MHz."
+ The subject has also received attcution iun its application to the problem of adequately sulracting the foreground radiation from the cosmic nücrowave backeround (7?).., The subject has also received attention in its application to the problem of adequately subtracting the foreground radiation from the cosmic microwave background \citep{Oliveira2008}.
+ huportaut plivsical information can be extracted from the surveys and the spectral iudex maps conceruiug Galactic structure. elobal maeuetic fields aud relativistic ectron distribution.," Important physical information can be extracted from the surveys and the spectral index maps concerning Galactic structure, global magnetic fields and relativistic electron distribution."
+ Iu this paper. we preseut an all-sky PAectral index map by usine our £5 MIIz survey and the (272). all-sky map. together with multi-frequency studies of six. zones that allow us to obtain zero-levol corrections for both maps.," In this paper, we present an all-sky spectral index map by using our 45 MHz survey and the \citep{Haslam1981,Haslam1982} all-sky map, together with multi-frequency studies of six zones that allow us to obtain zero-level corrections for both maps."
+ Iu the Appendix. we estimate the extragalactic backerouncd uou-thermal spectruui based onu a literature compilation.," In the Appendix, we estimate the extragalactic background non-thermal spectrum based on a literature compilation."
+ The £5 MIIz southern aud northern sky maps (77) were combined in the albskw shown in Fig.," The 45 MHz southern and northern sky maps \citep{45south,45north}
+ were combined in the all-sky shown in Fig."
+ 1., 1.
+ Missing data around the north equatorial pole represents ~L%( of the whole sky., Missing data around the north equatorial pole represents $\sim4\%$ of the whole sky.
+ The souhern map data were observed vetween 1982 aud 1991 xvi hau array of 528 E-W dipoles hat produced an effective beam of LG’(a)«2.1°68) ΕΜΠ., The southern map data were observed between 1982 and 1994 with an array of 528 E-W dipoles that produced an effective beam of $4.6^\circ(\alpha)\times2.4^\circ(\delta)$ FWHM.
+ A full descriptioi of the instrument was given Nw? and an experinuenta determination of the beam iu ?.., A full description of the instrument was given by \citet{45ins} and an experimental determination of the beam in \citet{45insex}.
+ The northern data were taken in the periods of 1955- aud 1997-1999 with the MU radar array (??) at an Haneular resolution of 3.6°«23.67.," The northern data were taken in the periods of 1985-1989 and 1997-1999 with the MU radar array \citep{MUradara,MUradarb} at an angular resolution of $3.6^\circ\times3.6^\circ $."
+ Table 1. displays some 6ft the principal characteristics of southern aud northeru mats of the 15 MIIz survey.," Table \ref{tabsurveys}
+ displays some of the principal characteristics of southern and northern parts of the 45 MHz survey."
+ An important characteristic of the all-sky [5 \WIz¢ map is that it consists of ouly two surveys and uses only two rausit iustruiueutswith similar Some relevant features of this map are:, An important characteristic of the all-sky 45 MHz map is that it consists of only two surveys and uses only two transit instrumentswith similar Some relevant features of this map are:
+strength is increased.,strength is increased.
+ The Al/@=5 case can be seen to form a very close but eccentric binary system which goes through a closest approach of ~30 AU., The $M/\Phi=5$ case can be seen to form a very close but eccentric binary system which goes through a closest approach of $\sim 30$ AU.
+ For the binary star formation ease we have also performed simulations which start from larger initial perturbations (4=0.2 and ;1=0.5 percent) and with larger and smaller cloud radii (510! em and 3.107 em)., For the binary star formation case we have also performed simulations which start from larger initial perturbations $A=0.2$ and $A=0.5$ percent) and with larger and smaller cloud radii $5\times 10^{16}$ cm and $3\times 10^{16}$ cm).
+ The runs with larger initial perturbations are progressively less influenced by the magnetic field., The runs with larger initial perturbations are progressively less influenced by the magnetic field.
+ Indeed for the ;4=0.5 percent perturbation all of the calculations up to a mass to flux ratio of ΑΦzc2 formed binaries with separations that are not significantly influenced by the magnetic field strength., Indeed for the $A=0.5$ percent perturbation all of the calculations up to a mass to flux ratio of $M/\Phi \approx 2$ formed binaries with separations that are not significantly influenced by the magnetic field strength.
+ The effect of changing the cloud radius is to change at what stage during the collapse the cloud becomes optically thick (ie., The effect of changing the cloud radius is to change at what stage during the collapse the cloud becomes optically thick (ie.
+ the equation of state changes from isothermal to non-isothermal)., the equation of state changes from isothermal to non-isothermal).
+ Thus dise fragmentation is (relative to the calculations presented here) suppressed for smaller cloud radit and enhanced for larger radii but with similar trends in the influence of the magnetic field., Thus disc fragmentation is (relative to the calculations presented here) suppressed for smaller cloud radii and enhanced for larger radii but with similar trends in the influence of the magnetic field.
+ Results of binary star formation caleulations beginning with a magnetic field oriented perpendicular to the rotation axis (that is. with a field initially in the c. direction) are shown in Figure 9..," Results of binary star formation calculations beginning with a magnetic field oriented perpendicular to the rotation axis (that is, with a field initially in the $x-$ direction) are shown in Figure \ref{fig:binarycoldensBx}."
+ As in Figure 6 some general trends are clear: increasing the magnetic field strength leads to a delayed collapse and increasingly suppresses binary formation., As in Figure \ref{fig:binarycoldens} some general trends are clear: increasing the magnetic field strength leads to a delayed collapse and increasingly suppresses binary formation.
+" In this case. however. the transition from a binary to a single star occurs earlier (that is. a single star is formed at A//=5 in Figure 9 compared to A//@=+ in Figure 6)) and the binary perturbation is increasingly deformed by the magnetic field. which at higher field strengths results in a ""double bar-like collapse” (most evident in the higher tield strength runs in Figure 9.."," In this case, however, the transition from a binary to a single star occurs earlier (that is, a single star is formed at $M/\Phi = 5$ in Figure \ref{fig:binarycoldensBx} compared to $M/\Phi=4$ in Figure \ref{fig:binarycoldens}) ) and the binary perturbation is increasingly deformed by the magnetic field, which at higher field strengths results in a “double bar-like collapse” (most evident in the higher field strength runs in Figure \ref{fig:binarycoldensBx}."
+ As previously. the global trends (delayed collapse and transition to a single star) are the result of the extra support provided to the cloud by magnetic pressure alone.," As previously, the global trends (delayed collapse and transition to a single star) are the result of the extra support provided to the cloud by magnetic pressure alone."
+ This is demonstrated by Figure 10. which shows the results of similar calculations (that. is. with fields initially perpendicular to the rotation axis) but with magnetic tension forces turned off.," This is demonstrated by Figure \ref{fig:binarycoldensBxnotens} which shows the results of similar calculations (that is, with fields initially perpendicular to the rotation axis) but with magnetic tension forces turned off."
+ In this case the transition to a single star occurs for even lower magnetic field strengths (at. AZ/@= 10)., In this case the transition to a single star occurs for even lower magnetic field strengths (at $M/\Phi = 10$ ).
+ This indicates not only that magnetic pressure is providing the dominant role in suppressing fragmentation but also that magnetic tension can act to dilute the effect of magnetic pressure. even binary formation.," This indicates not only that magnetic pressure is providing the dominant role in suppressing fragmentation but also that magnetic tension can act to dilute the effect of magnetic pressure, even binary formation."
+ We note that 22. similarly concluded that magnetic tension forces can act to promote fragmentation. albeit using a simplistic approximation to model the affect of magnetic fields.," We note that \citet{boss00,boss02} similarly concluded that magnetic tension forces can act to promote fragmentation, albeit using a simplistic approximation to model the affect of magnetic fields."
+ Comparing the middle column of Figure IO. to the pressure- calculations with an aligned magnetic field shown in Figure 7 also demonstrates that the effect of magnetic pressure is dependent on the field geometry. acting more like a equivalent thermal pressure when the field is aligned with the rotation axis.," Comparing the middle column of Figure \ref{fig:binarycoldensBxnotens} to the pressure-only calculations with an aligned magnetic field shown in Figure \ref{fig:binarynotens} also demonstrates that the effect of magnetic pressure is dependent on the field geometry, acting more like a equivalent thermal pressure when the field is aligned with the rotation axis."
+ The deformation of the binary perturbation in Figure 9 is not evident in the tension-free calculations (Figure 10). indicating (as one might expect) that this effect is the result of the gus being squeezed by the magnetic field lines.," The deformation of the binary perturbation in Figure \ref{fig:binarycoldensBx} is not evident in the tension-free calculations (Figure \ref{fig:binarycoldensBxnotens}) ), indicating (as one might expect) that this effect is the result of the gas being squeezed by the magnetic field lines."
+ It is this squeezing due to magnetic tension that acts to hold up the rapid transition to a single star with increasing field strength observed in the tension-free runs and thus dilute the effect of magnetic pressure in suppressing fragmentation., It is this squeezing due to magnetic tension that acts to hold up the rapid transition to a single star with increasing field strength observed in the tension-free runs and thus dilute the effect of magnetic pressure in suppressing fragmentation.
+ The magnetic field. being in this case aligned the binary perturbation. effectively acts as a “cushion” between the two stars which prevents their merging.," The magnetic field, being in this case aligned the binary perturbation, effectively acts as a “cushion” between the two stars which prevents their merging."
+ This effect. which we henceforth refer to as “magnetic cushioning”. is graphically illustrated in Figure 11. which shows the magnetic field (arrows in left panel. overlaid on a column density plot) and integrated magnetic pressure (right panel) in the ΑΦ=LO run (corresponding to the second row of Figure 9)) at fp=1.35.," This effect, which we henceforth refer to as “magnetic cushioning”, is graphically illustrated in Figure \ref{fig:magcushion} which shows the magnetic field (arrows in left panel, overlaid on a column density plot) and integrated magnetic pressure (right panel) in the $M/\Phi = 10$ run (corresponding to the second row of Figure \ref{fig:binarycoldensBx}) ) at $t_{\rm ff}=1.35$."
+ The “cushion” formed by the magnetic field between the two stars is clearly evident. and it is this “magnetic cushion” which prevents the binary system from merging to form a single star (and also produces the wonderful symmetry in the spiral arms).," The “cushion” formed by the magnetic field between the two stars is clearly evident, and it is this “magnetic cushion” which prevents the binary system from merging to form a single star (and also produces the wonderful symmetry in the spiral arms)."
+ The results shown in Figure 9 are quantified in Figure 12 which shows the binary separation as a function of time for the magnetic field strengths shown in Figure 9 and may be compared with the corresponding figure (Figure 89) for the runs with the field aligned with the rotation axis., The results shown in Figure \ref{fig:binarycoldensBx} are quantified in Figure \ref{fig:sepBx} which shows the binary separation as a function of time for the magnetic field strengths shown in Figure \ref{fig:binarycoldensBx} and may be compared with the corresponding figure (Figure \ref{fig:sep}) ) for the runs with the field aligned with the rotation axis.
+ As previously. prior to sink particle formation. we compute the separation of two density maxima in opposite hemispheres.," As previously, prior to sink particle formation, we compute the separation of two density maxima in opposite hemispheres."
+ In the stronger field runs CA//>=4 and Al/c>=5) the binary perturbation is strongly deformed by the magnetic field. producing the observed increase in separation observed at (yp1.2.," In the stronger field runs $M/\Phi = 4$ and $M/\Phi=5$ ) the binary perturbation is strongly deformed by the magnetic field, producing the observed increase in separation observed at $t_{\rm ff}\sim 1.2$."
+ The binary separations are in each case smaller than the equivalent runs using a field aligned with the rotation axis. which demonstrates that the effect of the magnetic field of the binary system is stronger in this case.," The binary separations are in each case smaller than the equivalent runs using a field aligned with the rotation axis, which demonstrates that the effect of the magnetic field of the binary system is stronger in this case."
+ The effect of magnetic cushioning is also apparent in the fact that the runs with AZ/=20.10 and 7.5 show a trend of binary separation at closest approach (ip~1.35). in contrast to Figure 8 (although all the separations are smaller than in the aligned-field runs).," The effect of magnetic cushioning is also apparent in the fact that the runs with $M/\Phi = 20, 10$ and $7.5$ show a trend of binary separation at closest approach $t_{\rm ff}\sim 1.35$ ), in contrast to Figure \ref{fig:sep} (although all the separations are smaller than in the aligned-field runs)."
+ We have conducted a study of how magnetic fields affect the collapse of homogenous molecular cloud cores and cores with initial n2=2 density perturbations., We have conducted a study of how magnetic fields affect the collapse of homogenous molecular cloud cores and cores with initial $m=2$ density perturbations.
+ In both cases. the presence of a magnetic field produces a delayed collapsed. with a longer delay for stronger fields.," In both cases, the presence of a magnetic field produces a delayed collapsed, with a longer delay for stronger fields."
+ This affect is easily attributed to the effect of the magnetic pressure on the collapse., This affect is easily attributed to the effect of the magnetic pressure on the collapse.
+ The magnetic field gives extra support to the cloud over the thermal pressure alone: rather than acting like a cloud whose ratio of thermal energy to the magnitude of gravitational energy à.=0.26 tor 0.35 in the axisymmetric models). the effect of the magnetic field is to raise the effective value of à.," The magnetic field gives extra support to the cloud over the thermal pressure alone; rather than acting like a cloud whose ratio of thermal energy to the magnitude of gravitational energy $\alpha=0.26$ (or $0.35$ in the axisymmetric models), the effect of the magnetic field is to raise the effective value of $\alpha$."
+ In the homogenous simulations. we find that a single protostar (sink particle) is formed and surrounded by a dise.," In the homogenous simulations, we find that a single protostar (sink particle) is formed and surrounded by a disc."
+ Stronger magnetic fields lead to a delay in the formation time of the protostar as mentioned above. but they also decrease the rate of accretion onto the disc.," Stronger magnetic fields lead to a delay in the formation time of the protostar as mentioned above, but they also decrease the rate of accretion onto the disc."
+ The dise radius also increases more slowly with time., The disc radius also increases more slowly with time.
+ It is well Known that the rate of infall of mass onto a massive dise is crucial in generating gravitational instabilities (e.g. 222)) and. indeed. we see this effect here.," It is well known that the rate of infall of mass onto a massive disc is crucial in generating gravitational instabilities (e.g. \citealt{bonnell94,whitworth95,hennebelle04}) ) and, indeed, we see this effect here."
+ In the purely hydrodynamical case. the dise surrounding the protostar is gravitationally unstable and exhibits strong spiral density waves soon after the protostar forms (although the instability is not strong enough to force the dise to fragment).," In the purely hydrodynamical case, the disc surrounding the protostar is gravitationally unstable and exhibits strong spiral density waves soon after the protostar forms (although the instability is not strong enough to force the disc to fragment)."
+ With a magnetic field initially aligned with the rotation axis. the slower rate of mass infall onto the dise leads to a weakening of the gravitational instability such that for mass to flux ratios less than A//z10 the spiral features are very weak (Figure 3)).," With a magnetic field initially aligned with the rotation axis, the slower rate of mass infall onto the disc leads to a weakening of the gravitational instability such that for mass to flux ratios less than $M/\Phi \approx 10$ the spiral features are very weak (Figure \ref{fig:axisymcoldens}) )."
+ With a field initially perpendicular to the rotation axis. the gravitational instability is very weak even for A//«.=20 (Figure 55).," With a field initially perpendicular to the rotation axis, the gravitational instability is very weak even for $M/\Phi = 20$ (Figure \ref{fig:axisymcoldens_Bx}) )."
+Now. let us consider particulary solutious for wz0.,"Now, let us consider particular solutions for $\omega\neq0$."
+ Thus. we will use the general equations (17))-(21)).," Thus, we will use the general equations \ref{eins1}) \ref{eins5}) )."
+" with | ifo2Vy>0 and - if ur2W,<0 ", with + if $x+2\Psi_0>0$ and - if $x+2\Psi_0<0$ 
+According to the comments made in this section. we believe that. if significant discrepancies between our results concerning the collisional history of Ceres and Vesta and the observational data obtained from the Dawn Mission were found. they should be linked to a higher degree of collisional evolution in the Main Belt.,"According to the comments made in this section, we believe that, if significant discrepancies between our results concerning the collisional history of Ceres and Vesta and the observational data obtained from the Dawn Mission were found, they should be linked to a higher degree of collisional evolution in the Main Belt."
+ An increase in the collisional activity may be due to If such discrepancies are eventually substantiated. future works should focus on accurately quantifying the degree of collisional evolution suffered by the early massive Main Belt as Well as during the LHB. assuming that the outer planets had initial orbits similar to and substantially different from (e.g. Nice model) their current ones.," An increase in the collisional activity may be due to If such discrepancies are eventually substantiated, future works should focus on accurately quantifying the degree of collisional evolution suffered by the early massive Main Belt as well as during the LHB, assuming that the outer planets had initial orbits similar to and substantially different from (e.g. Nice model) their current ones."
+ From this. the Dawn Mission would be able to give us interesting clues about the initial configuration of the early Solar System and its subsequent dynamical evolution.," From this, the Dawn Mission would be able to give us interesting clues about the initial configuration of the early Solar System and its subsequent dynamical evolution."
+ We have presented a study aimed at analyzing the impactor flux and cratering on Ceres and Vesta produced by Main Belt asteroids., We have presented a study aimed at analyzing the impactor flux and cratering on Ceres and Vesta produced by Main Belt asteroids.
+ To do this. we constructed a statistical code based on Bottke et al. (2005a. 2005b))," To do this, we constructed a statistical code based on Bottke et al. \cite{Bottke2005a, Bottke2005b}) )"
+ that includes catastrophic collisions and noncollisional removal processes. such as the Yarkovsky effect and the orbital resonances.," that includes catastrophic collisions and noncollisional removal processes, such as the Yarkovsky effect and the orbital resonances."
+ Assuming that the dynamical depletion of the early Main Belt was very strong. and owing to that. most Main Belt comminution occurred when the intrinsic collision. probability ¢P;) and the mean impact velocity (V) were comparable to their current values.," Assuming that the dynamical depletion of the early Main Belt was very strong, and owing to that, most Main Belt comminution occurred when the intrinsic collision probability $\langle P_{\text{i}} \rangle$ and the mean impact velocity $\langle V \rangle$ were comparable to their current values."
+ Our main results are the following, Our main results are the following
+inner edge of the disk is at approximately 56 kii when he X-rav flux is at its nmuninmun LEddinetou).,inner edge of the disk is at approximately 56 km when the X-ray flux is at its minimum Eddington).
+" This meaus that. bv asstuing an observer inclination of οὐ, the ueutron star surface is obscured for a disk hickuess of ~60 km."," This means that, by assuming an observer inclination of $60^{\circ}$, the neutron star surface is obscured for a disk thickness of $\sim 60$ km."
+ This thickness is however much arecr than the expected disk scale height which is oulv LOL kan., This thickness is however much larger than the expected disk scale height which is only 0.04 km.
+ Moreover. even considering the plysical size of the inner disk as large as 60 kin. the optical dept[um would drastically reduce to values below 1 after a few au iu height from the disk middle plane.," Moreover, even considering the physical size of the inner disk as large as 60 km, the optical depth would drastically reduce to values below 1 after a few km in height from the disk middle plane."
+ Although the Inner part of the accretion disk could be optically thin. allowing photoclectric absorption aud thus explainiug the ward euerey dependence of the 1 Tz OPO zs amplitude. it cannot explain the very high rus amplitude of the QPO. since the optically thin plasiia cannot completely obscure the neutron star surface.," Although the inner part of the accretion disk could be optically thin, allowing photoelectric absorption and thus explaining the hard energy dependence of the 1 Hz QPO rms amplitude, it cannot explain the very high rms amplitude of the QPO, since the optically thin plasma cannot completely obscure the neutron star surface."
+ The role of interchange instabilities i aciuittius uatter to the maenetosphere is discussed iu detail bv Arous&Lea(1976) and Elsner&Limb(1977)., The role of interchange instabilities in admitting matter to the magnetosphere is discussed in detail by \citet{aro76} and \citet{els77}.
+. The maguectic pressure prevents incoming matter from crossing imaeuctic field lines. but if it is energeticallv avorable (due for example to eravitv) for material o be “inside” the field Hines rather than voutside”. hen iuterchange instabilities can act to move the uaterial inside.," The magnetic pressure prevents incoming matter from crossing magnetic field lines, but if it is energetically favorable (due for example to gravity) for material to be “inside” the field lines rather than “outside”, then interchange instabilities can act to move the material inside."
+ This is often referred to as the Ravleigh-Tavlor mstabilitv. a term that refers to the iustabilitv hat occurs when a more dense fluid overlies a less dense fluid iu a gravitational field.," This is often referred to as the Rayleigh-Taylor instability, a term that refers to the instability that occurs when a more dense fluid overlies a less dense fluid in a gravitational field."
+ Where a plasma is supported against eravity by a maenctic field. this is nore correctly referred to as the Eruskal-Schwiurzschild lustability (Nauskal&Sclavarzschild1951).," Where a plasma is supported against gravity by a magnetic field, this is more correctly referred to as the Kruskal-Schwarzschild instability \citep{kru54}."
+". The couditious uecessary for the ouset of I&ruskal-Schwarzschild instabilities iu the situation where r,,r have since been studied using MIID simulations (Romanova.Wulkarni&Lovelace2008:IulkuniRomanova 2008)."," The conditions necessary for the onset of Kruskal-Schwarzschild instabilities in the situation where $r_m \ne r_c$ have since been studied using MHD simulations \citep{rom08, kul08}."
+. For misaligument angles 0=30° accretion can proceed either stably via fuunel flows. or uustably via interchange iustabilities.," For misalignment angles $\theta \lesssim 30^\circ$ accretion can proceed either stably via funnel flows, or unstably via interchange instabilities."
+ Iu the unstable situation matter accretes via a nunber of “tougues” that peuetrate the maguetopause in the equatorial plane., In the unstable situation matter accretes via a number of “tongues” that penetrate the magnetopause in the equatorial plane.
+ If a certain nunboer of tougues dominate. quasi-periodic oscillatious can emerge in the απο curves.," If a certain number of tongues dominate, quasi-periodic oscillations can emerge in the light curves."
+ Funnel flows can co-exist with accretion bv tougues (ular&Romanova 2008). although their presence should reduce the amplitude of the persistent pulsations bv reducing the azimuthal asvuuuetry (Romanova.I&ulka-rni&Lovelace 2008).," Funnel flows can co-exist with accretion by tongues \citep{kul08}, although their presence should reduce the amplitude of the persistent pulsations by reducing the azimuthal asymmetry \citep{rom08}."
+. This night be consistent with the continued presence of accretion-powered pulsations. with an amplitude that depends ou the phase of the 1 Iz QPO.," This might be consistent with the continued presence of accretion-powered pulsations, with an amplitude that depends on the phase of the 1 Hz QPO."
+ Iu the cases studied by Romanova.νωπά&Lovelace(2008) iunterchiuge instabilitics set im above a critical accretion rate. making them mualikely as a cause for the 1 Tz QPO (lich appears only below a critical rate).," In the cases studied by \citet{rom08} interchange instabilities set in above a critical accretion rate, making them unlikely as a cause for the 1 Hz QPO (which appears only below a critical rate)."
+ Iuterchauge instabilities cannot be completely ruled out. however.," Interchange instabilities cannot be completely ruled out, however."
+ In the standard case studied by Romanova.[E&ulkuni&Lovelace(2008) magnetic pressure and eravity donünate the force equations., In the standard case studied by \citet{rom08} magnetic pressure and gravity dominate the force equations.
+" When or,~ων however. magnetic pressure equals eyavitv."," When $r_m \sim r_c$, however, magnetic pressure equals gravity."
+ At this point other terms start to dominate the force equations and the character of the iutercliauge instability will change, At this point other terms start to dominate the force equations and the character of the interchange instability will change.
+1993).. Baan(1977) showed that sporadic peuctration of the magnetosplere is possible in this regie., \citet{baa77} showed that sporadic penetration of the magnetosphere is possible in this regime.
+ However detailed numerical simulations of the type performed at higher accretion rates have not been done. aud the effect on fuunel flows (and hence the amplitude of the accretion-powered pulsations) is uot known.," However detailed numerical simulations of the type performed at higher accretion rates have not been done, and the effect on funnel flows (and hence the amplitude of the accretion-powered pulsations) is not known."
+" Without further study. periodicitv due to iterchauge stabilities operating in the regime where rj,r. cannot be ruled out as a dmechanisin for the 1 Wz OPO."," Without further study, periodicity due to interchange instabilities operating in the regime where $r_m \sim r_c$ cannot be ruled out as a mechanism for the 1 Hz QPO."
+ As daunatter moves within the «disk (vacially aud agimutthally) it can drag magnetic field lues along with it., As matter moves within the disk (radially and azimuthally) it can drag magnetic field lines along with it.
+ The sheared field lines can temporarily inpede accretion until reconnection establishes a normal flow again., The sheared field lines can temporarily impede accretion until reconnection establishes a normal flow again.
+ The resulting quasi-periodic accretion flow would lead to a correspouding quasi-periodicity in the liehteurve. provided that the accretion funnel and hot spot can respond ou a Ls time scale.," The resulting quasi-periodic accretion flow would lead to a corresponding quasi-periodicity in the lightcurve, provided that the accretion funnel and hot spot can respond on a 1 s time scale."
+ Alagnetie reconnection is oue of the most plausible mechanisius for Type II bursts (cf. ?7?7))., Magnetic reconnection is one of the most plausible mechanisms for Type II bursts (cf. \ref{1Hz:Discussion}) ).
+ Type II bursts are thought to occur in systems where magnetic inhibition causes accreting matter to build up iu a reservoir outside the magnetosphere., Type II bursts are thought to occur in systems where magnetic inhibition causes accreting matter to build up in a reservoir outside the magnetosphere.
+ Once a sufficieut over-deusity of material has accunuulated. mstabilities cause a catastrophic breach of the magnetospheric hamunock. resulting in sudden bursts of accretion (Lewinetal. 1976).," Once a sufficient over-density of material has accumulated, instabilities cause a catastrophic breach of the magnetospheric hammock, resulting in sudden bursts of accretion \citep{lew76}."
+ Dramatic changes in QPO properties iuniediatelv before Type II bursts. with no detectable change iu the spectrum. have been interpreted as inclicating that the QPOs in the persistent cussion are eenerated within the fuel reservoir (Dotanietal.1990).," Dramatic changes in QPO properties immediately before Type II bursts, with no detectable change in the spectrum, have been interpreted as indicating that the QPOs in the persistent emission are generated within the fuel reservoir \citep{dot90}."
+. To what extent this phenomenon is relevant to our 1 IIz OPO is unclear., To what extent this phenomenon is relevant to our 1 Hz QPO is unclear.
+ Tn the following discussion we explore several iiechanisunis that cau produce 1 IIz oscillations., In the following discussion we explore several mechanisms that can produce 1 Hz oscillations.
+ Alv&IKuijpers(1990) used reconnection to explain he QPOs observed in the Rapid Buster. and discussed row the disk would be broken up in blobs’ by magnetic reconnection iustabilities," \citet{aly90} used reconnection to explain the QPOs observed in the Rapid Burster, and discussed how the disk would be broken up in `blobs' by magnetic reconnection instabilities."
+" They predict a frequency a few ines the beat frequency between the rotation rate of the star vy, and the I&epleriau frequency at the immer edge of he disk y.", They predict a frequency a few times the beat frequency between the rotation rate of the star $\nu_s$ and the Keplerian frequency at the inner edge of the disk $\nu_K$.
+" This implies that the frequency of the QPOs should pass through zero at the poiut where rj,1.", This implies that the frequency of the QPOs should pass through zero at the point where $r_m \sim r_c$.
+ An argument in favor of this mechiuisi is that has οσα observed in simulations (Coodson.Wiuglee&&Winelee 1999.. aud Romanovaetal. 20053). Ro," An argument in favor of this mechanism is that has been observed in simulations \citealt{goo97, goo99a, goo99b}, and \citealt{rom05}) )."
+manovaetal(2005) observed a strong quasi-periodicitv associated with outflows of matter when the system was in the propeller τοσο., \citet{rom05} observed a strong quasi-periodicity associated with outflows of matter when the system was in the propeller regime.
+ Unfortunately Romanovaetal.(2005). do not model the ciission 1iechlianisiis. so whether outflows or discrete accretion would seuuiuclv lead to high amplitude QPOs in the N-vay cuission is not clear. Ustvugovaetal.(2006).," Unfortunately \citet{rom05} do not model the emission mechanisms, so whether outflows or discrete accretion would genuinely lead to high amplitude QPOs in the X-ray emission is not clear. \citet{ust06},"
+. using MIID simulations. showed that rapidly rotating accreting stars have strong periodicity of this type linked to strong outflows.," using MHD simulations, showed that rapidly rotating accreting stars have strong periodicity of this type linked to strong outflows."
+ Observationallv. there is evidence for jet formation fou the radio detections in the decay of the 1998 outburst and the peak of the 2002 outburst (Gacusleretal. 1999.. Rupenetal. 20023). although the latter cannot be related with the propeller ouset since it occurs at the maxima fluxes observed for Jls08.," Observationally, there is evidence for jet formation from the radio detections in the decay of the 1998 outburst and the peak of the 2002 outburst \citealt{gae99}, , \citealt{rup02}) ), although the latter cannot be related with the propeller onset since it occurs at the maximum fluxes observed for J1808."
+ The frequency. for this mechauisin. should depend ou accretion rate (as observed. Fie. 7)).," The frequency, for this mechanism, should depend on accretion rate (as observed, Fig. \ref{rms-flux-freq}) )."
+ It should however, It should however
+Alass-loss appears to be an integral part of the star formation process.,Mass-loss appears to be an integral part of the star formation process.
+ Ht is Likely to plav a miajor role in setGne the final mass of the star by reversing the infall and therefore in determination of the IME., It is likely to play a major role in setting the final mass of the star by reversing the infall and therefore in determination of the IMF.
+ Liuminous (7104 L. ) voung stellar objects (YSOs) not only drive the well-known large-scale bipolar molecular outflows (e.g. Lada 1935: Beuther et al., Luminous $>$ $^{4}$ $_{\sun}$ ) young stellar objects (YSOs) not only drive the well-known large-scale bipolar molecular outflows (e.g. Lada 1985; Beuther et al.
+ 2002). but also have small-scale ionized stellar winds.," 2002), but also have small-scale ionized stellar winds."
+" These are weak (5,X few mJv) and compact C 1"")) radio sources (e.g. Tofani et al.", These are weak $S_{\nu}\la$ few mJy) and compact $\la$ ) radio sources (e.g. Tofani et al.
+ 1995)., 1995).
+ They are distinguished from ultra-compact II ll reeions (UCIUIs) observationally by their radio spectral index (hat is usually close to that, They are distinguished from ultra-compact H II regions (UCHIIs) observationally by their radio spectral index that is usually close to that
+baryon) which is deposited at. or near. the photosphere and is rapidly radiated away.,"baryon) which is deposited at, or near, the photosphere and is rapidly radiated away."
+" Such infall energy does not get taken into the star because in the upper atmosphere (where T. 7,) the time it takes the fluid to move inward is much longer (by at least 7./ T) than the time it takes for heat to escape.", Such infall energy does not get taken into the star because in the upper atmosphere (where $T\ll T_c$ ) the time it takes the fluid to move inward is much longer (by at least $T_c/T$ ) than the time it takes for heat to escape.
+ This means infall energy is not important for setting the internal thermal state of the WD: it simply has no influence there., This means infall energy is not important for setting the internal thermal state of the WD; it simply has no influence there.
+ Also. once aceretion has diminished for longer than the time to radiate the infall energy away (such as in DN quiescence). it is no longer relevant to the observed luminosity.," Also, once accretion has diminished for longer than the time to radiate the infall energy away (such as in DN quiescence), it is no longer relevant to the observed luminosity."
+ An additional energy source in the WD interior is slow nuclear burning near the base of the accreted H/He layer., An additional energy source in the WD interior is slow nuclear burning near the base of the accreted H/He layer.
+ This 1s significant when the acereted layer becomes thick. eventually becoming thermally unstable and leading- to a ClassicalM. Nova.," This is significant when the accreted layer becomes thick, eventually becoming thermally unstable and leading to a Classical Nova."
+ The nova energy Is. assumed to be radiated. away in. the explosion.. but we have found that slow burning before that point contributes an amount of. energy to the WD μα.interior: comparable to the energy released by compression.," The nova energy is assumed to be radiated away in the explosion, but we have found that slow burning before that point contributes an amount of energy to the WD interior comparable to the energy released by compression."
+ Our calculations take account of both compressional heating and slow nuclear burning to determine the WD's thermal state., Our calculations take account of both compressional heating and slow nuclear burning to determine the WD's thermal state.
+ We now sketch how compressional heating is included in our stellar model. demonstrate that heating in the envelope dominates that in the core and estimate its magnitude.," We now sketch how compressional heating is included in our stellar model, demonstrate that heating in the envelope dominates that in the core and estimate its magnitude."
+ This is a discussion specific to DNwith low (Mj. and the reader should consult Nomoto (1982) for a complete account.," This is a discussion specific to DNwith low $\timav$, and the reader should consult Nomoto (1982) for a complete account."
+" The simplest estimate of the compressional energy release is the gravitational energy liberated as a fluid element moves down in the WD gravitational field. e=GM/R""."," The simplest estimate of the compressional energy release is the gravitational energy liberated as a fluid element moves down in the WD gravitational field, $g=GM/R^2$."
+" In the non-degenerate outer atmosphere. a fluid element moves a distance of order the scale height. AT /umyg. in the time it takes to replace it by accretion. giving L~(Mogh~(M\kTjim,."," In the non-degenerate outer atmosphere, a fluid element moves a distance of order the scale height, $h=kT/\mu
+m_p g$ , in the time it takes to replace it by accretion, giving $L\sim \timav
+g h\sim \timav k T/\mu m_p$."
+ This exhibits the correct scaling. notably the dependence on j/ which ts a contrasting parameter between the accreted H/He envelope and the C/O core.," This exhibits the correct scaling, notably the dependence on $\mu$ which is a contrasting parameter between the accreted H/He envelope and the C/O core."
+" To calculate the actual heat release we consider the local heat equationds where εν 1s the nuclear burning rate. s is the entropy. and v=—(M)r/Azr pis the slow downward advection speed from accretion,"," To calculate the actual heat release we consider the local heat equation where $\epsilon_N$ is the nuclear burning rate, $s$ is the entropy, and $\vec
+v=-\timav \hat r/4\pi r^2 \rho$ is the slow downward advection speed from accretion."
+ The entropy profile is fixed by the temperature gradient needed to carry the luminosity outward and thus we simultaneously solve equation (1)) with the heat transport equation. using opacities and conductivities from Iglesias Rogers (1996) and Itoh et ((1983) to find the thermal structure of the accreted envelope and the outer edge of the C/O core.," The entropy profile is fixed by the temperature gradient needed to carry the luminosity outward and thus we simultaneously solve equation \ref{eq:heateq}) ) with the heat transport equation, using opacities and conductivities from Iglesias Rogers (1996) and Itoh et (1983) to find the thermal structure of the accreted envelope and the outer edge of the C/O core."
+ For an analytic understanding. we neglect nuclear burning and 9/r and use hydrostatic balance to recast equation (1)) into Entropy decreases inward (1.9. the envelope and core are not convective). so this is an outward L.," For an analytic understanding, we neglect nuclear burning and $\partial/\partial t$ and use hydrostatic balance to recast equation \ref{eq:heateq}) ) into Entropy decreases inward (i.e. the envelope and core are not convective), so this is an outward $L$."
+" In the non-degenerate envelope. where s2KIn(P""7/p)/pini. one More approximation is necessary to obtain an analytic form."," In the non-degenerate envelope, where $s=k\ln(T^{3/2}/\rho)/\mu m_p$, one more approximation is necessary to obtain an analytic form."
+ For an atmosphere in which £L is constant with depth. the envelope satisfies 77?P7.," For an atmosphere in which $L$ is constant with depth, the envelope satisfies $T^{8.5}\propto P^2$."
+ Though L is not constant here. we use this to get an estimate. integrating down to the isothermal core and finding Lz3&7.6M;μμ.," Though $L$ is not constant here, we use this to get an estimate, integrating down to the isothermal core and finding $L\approx 3kT_c \timav/\mu m_p$."
+ Now consider the degenerate C/O core., Now consider the degenerate C/O core.
+ For the (My's and typical M20.6M WD of interest here. the entropy is in the liquid ions at 7.—..10K.," For the $\timav$ 's and typical $M=0.6 M_\odot$ WD of interest here, the entropy is in the liquid ions at $T_c\approx 10^7 \ {\rm K}$."
+ The time takes to transport heat through the interior is ~I0vr<<itM/(MD. so the core is isothermal and any compression is far from Due to uncertainty from the classical novae cycle. we don't know whether the C/O core Is secularly increasing in mass. but if it were. almost all of the work of compression goes into increasing the electron Fermi energy.," The time it takes to transport heat through the interior is $\sim 10^7 {\rm
+yr}\ll M/\timav$, so the core is isothermal and any compression is far from Due to uncertainty from the classical novae cycle, we don't know whether the C/O core is secularly increasing in mass, but if it were, almost all of the work of compression goes into increasing the electron Fermi energy."
+" The integrated heat release in the core would then be L=ΙΣΤ.ΜΗ)fam, (Nomoto 1982) for a 0.6M.. C/O WD. where j;zz1d 18 the ion mean molecular weight."," The integrated heat release in the core would then be $L\approx 15 kT_c\timav/\mu_i
+m_p$ (Nomoto 1982) for a $0.6M_\odot$ C/O WD, where $\mu_i\approx 14$ is the ion mean molecular weight."
+ Due to the mean molecular weight contrast between the accreted envelope and the core. the energy release in the core is about a factor of five smaller than that in the envelope.," Due to the mean molecular weight contrast between the accreted envelope and the core, the energy release in the core is about a factor of five smaller than that in the envelope."
+ Thus. for à given amount of compression of the star. the entropy drop for material the. accreted layer is much. larger than for material: inIne the core.," Thus, for a given amount of compression of the star, the entropy drop for material in the accreted layer is much larger than for material in the core."
+ —=T—4T¥ For this initial study. we dropped the time-dependent term in equation (1)). and presumed that the C/O core mass was constant throughout the classical nova cycle. thus only accounting for the compressional heating and εν in the accreted layer.," For this initial study, we dropped the time-dependent term in equation \ref{eq:heateq}) ), and presumed that the C/O core mass was constant throughout the classical nova cycle, thus only accounting for the compressional heating and $\epsilon_N$ in the accreted layer."
+ This method improves on that of Iben et ((1992) by allowing the accreted envelope mass to change through the 10° year classical nova cycle., This method improves on that of Iben et (1992) by allowing the accreted envelope mass to change through the $10^5$ year classical nova cycle.
+ Early in the cycle. the mass of the accreted layer is small. compressional heating is small. and the WD cools.," Early in the cycle, the mass of the accreted layer is small, compressional heating is small, and the WD cools."
+ in the cycle. the accreted layer becomes thick enough ⋅ compressional heating along with slow hydrogen burning 0 releases a sufficient amount of energy to heat the core.," Later in the cycle, the accreted layer becomes thick enough that compressional heating along with slow hydrogen burning releases a sufficient amount of energy to heat the core."
+ As the WD has a large heat capacity. reaching the equilibrium 7; where the heat exchanged between the envelope and core averages to zero over a single classical nova cycle takes ~10° years.," As the WD has a large heat capacity, reaching the equilibrium $T_c$ where the heat exchanged between the envelope and core averages to zero over a single classical nova cycle takes $\approx 10^8$ years."
+ Since this time is shorter than the time over which (M? changes due to changing orbital period. we construct such equilibrium accretors for a given M and (Mj.," Since this time is shorter than the time over which $\timav$ changes due to changing orbital period, we construct such equilibrium accretors for a given $M$ and $\timav$ ."
+ To do this construction. we first fix 7.‘ at the outer edge of the C/O core at a pressure high enough so that the changing accumulated mass has little direct effect.," To do this construction, we first fix $T_c$ at the outer edge of the C/O core at a pressure high enough so that the changing accumulated mass has little direct effect."
+ With a radiative-zero, With a radiative-zero
+these motions. mav be mathematically modeled through force-free magnetic fields. magnetic fields which satisfv j4.B—0 where j=aB (Priest 1982).,"these motions, may be mathematically modeled through force-free magnetic fields, magnetic fields which satisfy $\bf{j} \times \bf{B} = 0$ where ${\bf j} = \alpha {\bf B}$ \citep{priest82}."
+". lu recent vears. non-linear force-free Ποια, uodeliug has received much attention (Schirijveretal.2006:Moetcealf 2008)."," In recent years, non-linear force-free field modeling has received much attention \citep{sch06,met08}."
+.. A ron-linear orce-free field. is a special class of force-free field. where the scalar fiction a(r) is a function of oositiou. but must be constaut along any Ποια iue.," A non-linear force-free field, is a special class of force-free field, where the scalar function $\alpha(r)$ is a function of position, but must be constant along any field line."
+ Nou-linear force-free fields are of particular interest as they may contain free magnetic energy (Woltjer1958)., Non-linear force-free fields are of particular interest as they may contain free magnetic energy \citep{wolt58}.
+. The coustruction of non-linear orce-free fields (NLFFF) may be broadly split iuto WO STOUps: static and time-dependent models., The construction of non-linear force-free fields (NLFFF) may be broadly split into two groups: static and time-dependent models.
+ Examples of static niodoliug are the “extrapolation” of plotospheric vector magnetic fields iuto the corona (Schriyveretal.2006:Reenier&Priest2009:Jingetal 2010). and direct fitting of NLFFF iodels to observed filaments aud coronal structures (vanBallegooijen2001:Bobractal.jeu 2009).," Examples of static modeling are the “extrapolation"" of photospheric vector magnetic fields into the corona \citep{sch06,reg07,derosa09,wheat09,jing10}, and direct fitting of NLFFF models to observed filaments and coronal structures \citep{bal04,bob08,su09b,sav09}."
+. Static modeling involves constructing either individual NLFFFs or independent sequences of nmaguetic configurations. where there is no direct correlation or evolution between the differcut fields.," Static modeling involves constructing either individual NLFFFs or independent sequences of magnetic configurations, where there is no direct correlation or evolution between the different fields."
+ Tn contrast. time-dependent quasi-static modcling evolves the coronal maenetic ficl through coutiuuous sequences of related non-IBnear force-free fields based on the evolution of a contiuuous nne dependent lower boundary coudition.," In contrast, time-dependent quasi-static modeling evolves the coronal magnetic field through continuous sequences of related non-linear force-free fields based on the evolution of a continuous time dependent lower boundary condition."
+ This boundary condition may be specified through either idealized maguetic field configurations (vanDallegooijeu&vanBallegooijen2005.2006a.2009) or from observations (Mackavetal.2000:Yeates2007. 2008).," This boundary condition may be specified through either idealized magnetic field configurations \citep{bal00,mac03,mac05,mac06a,mac09} or from observations \citep{mac00,yea07,yea08}."
+. Only through using dvuamic models can the build up of free magnetic enerev and maenetic helicity iu the corona be studied. along with the evolution of isolated flux and current systems.," Only through using dynamic models can the build up of free magnetic energy and magnetic helicity in the corona be studied, along with the evolution of isolated flux and current systems."
+ Such a technique has been successfully applied iu the past to model the evolution of the global corona (Mackay.&vanDallegooijeu2006a.b:Yeates&Mackay 2009).. determine the origin and evolution of solar filaments (Yeatesetal.2008:Alackay&vanBallegooijen 2009).. study the formation and litt off of magnetic fiux ropes aud finally the origin of the Suu's Open Maguetie Flux (Alackay&vauDallegooijeu20061:Yeatesetal.201001.," Such a technique has been successfully applied in the past to model the evolution of the global corona \citep{mac06a,mac06b,yea09}, determine the origin and evolution of solar filaments \citep{yea08,mac09}, study the formation and lift off of magnetic flux ropes and finally the origin of the Sun's Open Magnetic Flux \citep{mac06b,yea10b}."
+ Iu this paper we applv iue dependent quasi-static non-linear force-free modeling to model the effect tha the dispersal of an active region has ou the overlving coronal naenetic field., In this paper we apply time dependent quasi-static non-linear force-free modeling to model the effect that the dispersal of an active region has on the overlying coronal magnetic field.
+ A key difference between this paper aud previous studies is the manner in which the time dependent lower boundary coucdition is applied., A key difference between this paper and previous studies is the manner in which the time dependent lower boundary condition is applied.
+ Iu previous studies either idealized magnetic feld distributions or siupliüed surface configurations derived frou observations were prescribed., In previous studies either idealized magnetic field distributions or simplified surface configurations derived from observations were prescribed.
+ Iu the present study. a new technique of prescribing he tine dependent lower boundary condition is presented.," In the present study, a new technique of prescribing the time dependent lower boundary condition is presented."
+ The technique allows the use of observed Lue of sieht compoucut maecuctoerams directly as lower boundary conditions. reproducing as observed. the dispersal of the plotospheric fux of the active reeion.," The technique allows the use of observed line of sight component magnetograms directly as lower boundary conditions, reproducing as observed, the dispersal of the photospheric flux of the active region."
+ From this the effect of the surface motions ou the coronal magnetic field aud the subsequeu enerev input iuto the corona is determined., From this the effect of the surface motions on the coronal magnetic field and the subsequent energy input into the corona is determined.
+ The active region chosen to illustrate this new echnique is NOAA AR S005., The active region chosen to illustrate this new technique is NOAA AR 8005.
+ It is chosen as i was an isolated region and during the period of observations there appears to be no significant Hux enmergeuce and the flux is well balancec hroughout., It is chosen as it was an isolated region and during the period of observations there appears to be no significant flux emergence and the flux is well balanced throughout.
+ The structure of the paper Is as ollows., The structure of the paper is as follows.
+ Iu Section 2 the main properties of the decaying active region are discussed., In Section 2 the main properties of the decaying active region are discussed.
+ In Section 3 he modeling teclinique for both the plotospleric and coronal fields are described., In Section 3 the modeling technique for both the photospheric and coronal fields are described.
+ The results of he simulation are Ooeiven in Section [. while iu Section 5 some simple calculations are compared o the simulations to verify the results.," The results of the simulation are given in Section 4, while in Section 5 some simple calculations are compared to the simulations to verify the results."
+ Finally a discussion aud couchisions are given im Section 6., Finally a discussion and conclusions are given in Section 6.
+ The decaviug active region considered iu this study is NOAA AR so05., The decaying active region considered in this study is NOAA AR 8005.
+ It cmereed ou the far side of the Sun and as it rotated onto the Iib only dispersed magnetic polarities. without sunspots were present.," It emerged on the far side of the Sun and as it rotated onto the limb only dispersed magnetic polarities, without sunspots were present."
+ Tn Figure 1 a full disk SOIIO:MDI (Scherreretal.1995) 96101n lue of sight componcut iiaguetogranm from 18/ Dec 1996 can be seen.," In Figure \ref{fig:fig1} a full disk SOHO:MDI \citep{Scherrer95}
+ $96$ min line of sight component magnetogram from $^{th}$ Dec 1996 can be seen."
+ The active region of interest lies at ceutral meridian just below the equator., The active region of interest lies at central meridian just below the equator.
+ It has a simple magnetic morphology with a single positive and uceative polarity., It has a simple magnetic morphology with a single positive and negative polarity.
+"It has been well established that massive galaxies vpicallv contain massive black holes (DIIS) at their centers. and it is widelv believed that these black holes. while in their so-called ""active galactic nuclei? (AGN) ghases. can have profound influcuce on the galaxics which host them (see for a comprehensive observational review).","It has been well established that massive galaxies typically contain massive black holes (BHs) at their centers, and it is widely believed that these black holes, while in their so-called “active galactic nuclei” (AGN) phases, can have profound influence on the galaxies which host them (see for a comprehensive observational review)."
+ This influence is interred from. among other observations. the correlation vetween the BIT mass aud the bulee mass of the host ealaxy. which was first noticed by aud studied by (1998). and las since been repeatedly confined 2001). as has the somewhat tighter relation between the BIT mass aud the host bulees velocity dispersion. discovered by and (2000).," This influence is inferred from, among other observations, the correlation between the BH mass and the bulge mass of the host galaxy, which was first noticed by and studied by , and has since been repeatedly confirmed , as has the somewhat tighter relation between the BH mass and the host bulge's velocity dispersion, discovered by and ."
+.. The interred BIT mass distribution las also been convincingly luked to the quasar hnuuinositv tuetion 2002)., The inferred BH mass distribution has also been convincingly linked to the quasar luminosity function .
+. Since the work of (1998).. the relatiouship )etween the AGN and host ealaxy has eeuerallv Όσοι characterized as a process of feedback: the galaxy supplies gas to be accreted by the DIT. which cuits some yaction of the accreted mass as mechanical energy to the siroundiues. sole fraction via the broac-line winds. aud sole fraction as the radio jets.," Since the work of , the relationship between the AGN and host galaxy has generally been characterized as a process of feedback: the galaxy supplies gas to be accreted by the BH, which emits some fraction of the accreted mass as mechanical energy to the surroundings, some fraction via the broad-line winds, and some fraction as the radio jets."
+" All these processes can reat the surrounding σας, and since eas which has been reated is less cleuse aud thus accretes more slowly. the ecdhack process is described as ""selfresulatiug (νο, a regative feedback loop)."," All these processes can heat the surrounding gas, and since gas which has been heated is less dense and thus accretes more slowly, the feedback process is described as “self-regulating” (i.e. a negative feedback loop)."
+ Thus there are two modes by which AGN may give cherey to their surroundings: the “mechanical iode. where the ACN inflates local gas bubbles which expand through the interstellar and iuter-galactie ieclia (ISAL/IGAD) 2005): and the velectromaguetic™ inodo. where photons from the ACN accretion region for which the neighboring eas has a low optical depth escape aud directly interact with the rest of the ISAL/TICAL 2005).," Thus there are two modes by which AGN may give energy to their surroundings: the “mechanical” mode, where the AGN inflates local gas bubbles which expand through the inter-stellar and inter-galactic media (ISM/IGM) ; and the “electromagnetic” mode, where photons from the AGN accretion region for which the neighboring gas has a low optical depth escape and directly interact with the rest of the ISM/IGM ."
+. Both of these modes may then interact with the surrounding galactic gas by either (or rather. both) of two mechlanisiis: enerev-based heating) 2003).. or ruonieutunm-based ppressure) 2010).," Both of these modes may then interact with the surrounding galactic gas by either (or rather, both) of two mechanisms: energy-based heating) , or momentum-based pressure) ."
+. have included both miechanisius iu their ougoime work., have included both mechanisms in their ongoing work.
+ Thus there are four conceptual componcuts which feedback models may take iuto account: mechanical-cucrey bbubbles). mechamicalinomenttm wwincls). clectromaguetic-cncrey radiative heating). aud electromaegueticauonieutunm rraciation pressure).," Thus there are four conceptual components which feedback models may take into account: mechanical-energy bubbles), mechanical-momentum winds), electromagnetic-energy radiative heating), and electromagnetic-momentum radiation pressure)."
+ The relative importance of these componcuts has been mich debated among the authors just mentioned., The relative importance of these components has been much debated among the authors just mentioned.
+ A fifth component. the thin radio jets. enüt comparable zannounts of energv aud momentum with the other modes. but these intense beams tend to drill through the ealactic gas. depositing their energv in the IGAL," A fifth component, the thin radio jets, emit comparable amounts of energy and momentum with the other modes, but these intense beams tend to drill through the galactic gas, depositing their energy in the IGM."
+ This makes jets less relevant as a feedback method. uuless they precess. or there is a relative velocity between the ACN and its environment. in which case significaut," This makes jets less relevant as a feedback method, unless they precess, or there is a relative velocity between the AGN and its environment, in which case significant"
+pulsars. we find that the semi-classical model of Donati&Piz-zochero(2004) overestimates the value of the excess angular mometum J. while the quantum approach of Avogadroetal. (2008).. with the SLy4 interaction. roughly agrees with the inferred from observations.,"pulsars, we find that the semi-classical model of \citet{don04} overestimates the value of the excess 	angular mometum $J$, while the quantum approach of \citet{avo08}, with the SLy4 interaction, roughly agrees with the 	value inferred from observations."
+ We studied five heating mechanisms that can be operating in old neutron stars: magnetic field decay. dark matter accretion. crust eracking. vortex creep. and rotochemical heating. and compared them with pulsar observations.," 	We studied five heating mechanisms that can be operating in old neutron stars: magnetic field decay, dark matter 	accretion, crust cracking, vortex creep, and rotochemical heating, and compared them with pulsar observations."
+ We found that magnetic field decay. dark matter accretion. and crust cracking cannot produce detectable heating in. old pulsars.," 	We found that magnetic field decay, dark matter accretion, and crust cracking cannot produce detectable heating in 	old pulsars."
+ Owing to the high yield strain angle2009).. the crust cracking mechanism does not operate in classical pulsars. and. probably only operates in MSPs.," 	Owing to the high yield strain angle, the crust cracking mechanism does not operate in classical pulsars, and 	probably only operates in MSPs."
+ The vortex ereep and rotochemical heating can be important both for classical pulsars and MSPs., The vortex creep and rotochemical heating can be important both for classical pulsars 	and MSPs.
+ In the evolutionary curves with vortex creep. and with the excess angular momentum. J. set to the lowest temperature compatible with the thermal emission. detected. from MSP J0437-4715. the predicted temperature turns out to be very near the observational upper limits for the classical pulsars B1929+10 and B1133+16.," 	In the evolutionary curves with vortex creep, and with the excess angular momentum, $J$, set to the lowest temperature 	compatible with the thermal emission detected from MSP J0437-4715, the predicted temperature turns out to be very near 	the observational upper limits for the classical pulsars B1929+10 and B1133+16."
+ Likewise. rotochemical heating with modified Urea reactions and no Cooper pairing is only 1.7 o below the temperature measured in the MSP JO437-4715.," Likewise, rotochemical heating 	with modified Urca reactions and no Cooper pairing is only 1.7 $\sigma$ below the temperature measured in the 	MSP J0437-4715."
+ However. the temperature prediction of rotochemical heating can be raised 1f a superfluid core is considered in the model.," However, the temperature prediction of rotochemical heating can be raised if a superfluid core is considered 	in the model."
+ The prediction of pinning energies in the inner crust of NSs via the semi-classical model for the vortex-nuclei interaction (Donati&Pizzochero2004) overestimates the temperatures of B0950-08 and B1929+10., 	The prediction of pinning energies in the inner crust of NSs via the semi-classical model for the vortex-nuclei interaction \citep{don04} overestimates the temperatures of B0950+08 and B1929+10.
+ The recent estimations of pinning energies via a quantum approach (Avogadroetal.2008) are consistent with the observations of classical pulsars and MSPs., The recent estimations of pinning energies via a quantum 	approach \citep{avo08} are consistent with the observations of classical pulsars and MSPs.
+ Finally. more stringent constraints on the temperature of some classical pulsars such as BO9SO0+08 could rule out the vortex creep mechanism as the main source of the thermal emission detected in the MSP J0437-4715.," Finally, more stringent constraints 	on the temperature of some classical pulsars such as B0950+08 could rule out the vortex creep mechanism as the main source of 	the thermal emission detected in the MSP J0437-4715."
+The 5-ray pulse prolile observed by CAbdoetal.20000 shows a double-peak structure.,"The $\gamma$ -ray pulse profile observed by \citep{Ab58}
+ shows a double-peak structure."
+ The first peak is offset Grom the radio peak by o~0.08. and he second is at ó~0.57.," The first peak is offset from the radio peak by $\phi\sim0.08$, and the second is at $\phi\sim0.57$ ."
+ The separatiol is No=0.19., The separation is $\Delta\phi=0.49$.
+ The X-ray data fromRATE (Livingstoneetal.2009) are consistent witli the results from (Ixuiperetal.2010)., The X-ray data from \citep{Li09} are consistent with the results from \citep{Ku10}.
+. The spectra slape cali e fitted by a power law such hat most of the emission is non-therma aud the thermal component is coustrained by the upper limit (Ixuiperetal.2010)., The spectral shape can be fitted by a power law such that most of the emission is non-thermal and the thermal component is constrained by the upper limit \citep{Ku10}.
+. The oBELVEC Xav pulse profie shows wo peaks aligued in phase ove |a wide energy range of ~0.5-270 seV.a icd is also very similar to heoL the 7-ray bain.," The observed X-ray pulse profile shows two peaks aligned in phase over a wide energy range of $\sim$ 0.5-270 keV, and is also very similar to that of the $\gamma$ -ray band."
+ We show the inteusity maps for outware Laud ilwaX emission in he uppe“and lower panels of Fie. 6((, We show the intensity maps for outward and inward emission in the upper and lower panels of Fig. \ref{fig:3.1}( (
+C). respectively.,"C), respectively."
+ As seen in the upyer pane. the emission altiuce for the double-peak with Ao=0.19 is ro 0.97-0.98.," As seen in the upper panel, the emission altitude for the double-peak with $\Delta\phi=0.49$ is $r_{ov}\sim0.97$ -0.98."
+ A shift ο [1je refereice ghase 06 is uot necessa‘y in this source., A shift of the reference phase $\delta\phi$ is not necessary in this source.
+ As argued in TCSOs. outward emission cloimilates 1n he ight curve [or a young oulsar with a strong uou-thermal X-ray component. likethe Crab pulsa," As argued in TCS08, outward emission dominates in the light curve for a young pulsar with a strong non-thermal X-ray component, likethe Crab pulsar."
+" PSR J02054-6L19 is the youngest pulsar in our sample (characteristic age 7.2xLO""j y‘) and slows rather stroug uou-thernal racdiatiou in the X-ray band.", PSR J0205+6449 is the youngest pulsar in our sample (characteristic age $\tau_c\sim 5\times 10^3$ yr) and shows rather strong non-thermal radiation in the X-ray band.
+ Thus. it is likely that ouly oward omission contributes to tle observed X-ray. ligit curve in this pulsar.," Thus, it is likely that only outward omission contributes to the observed X-ray light curve in this pulsar."
+ The light curve for 5-rays observed byFern: (Abdoetal.90000) shows an asyiuuetric single peak at o~0.19.," The light curve for $\gamma$ -rays observed by \citep{Ablv}
+ shows an asymmetric single peak at $\phi\sim0.49$."
+ The tail extends ¢owl to ó0.2., The tail extends down to $\phi\sim0.2$.
+ The peak position depeuds slightly ou the euergv rauge above 100 MeV. but tie amount of shift is only ~0.01.," The peak position depends slightly on the energy range above 100 MeV, but the amount of shift is only $\sim0.04$."
+ The X-ray. pulse profile observed by (Abdoetal.2009€). shows a double-peak structure.," The X-ray pulse profile observed by \citep{Ablv}
+ shows a double-peak structure."
+ No peak is seen at the 2-ray peak pliase., No peak is seen at the $\gamma$ -ray peak phase.
+ The separation between first N-ray peak aud tlie 2?-ray peak is Ao~0.32. aud the separation between the secoud X-ray peak ancl the 5-ray peak is Ao~0.1L.," The separation between first X-ray peak and the $\gamma$ -ray peak is $\Delta\phi\sim0.32$, and the separation between the second X-ray peak and the $\gamma$ -ray peak is $\Delta\phi\sim0.14$."
+ The iuteusity imap for outward euiisslou is sliown in the upper panel of Fie. 6((, The intensity map for outward emission is shown in the upper panel of Fig. \ref{fig:3.1}( (
+D).,D).
+ We consider the formation of the >-ray peak auc twὉ X-ray. peaks as being due to outward emission only., We consider the formation of the $\gamma$ -ray peak and two X-ray peaks as being due to outward emission only.
+" Such a solution is possible by choosing an enission altitude of re, 1.01-1.02 wi hoa small phase shilt Jo=0.03.", Such a solution is possible by choosing an emission altitude of $r_{ov}\sim1.01$ -1.02 with a small phase shift $\delta\phi=0.03$.
+ The intensity map for iuwa‘cL emission shows a sudden decrease in the number counts for roy< 1.06., The intensity map for inward emission shows a sudden decrease in the number counts for $r_{ov}<1.06$ .
+ The peak becomes broad and lence the contribution of the iward emission is uot very importaut., The peak becomes broad and hence the contribution of the inward emission is not very important.
+ We show the light cu‘ve of inward aud outward eimissiou Or rue~ 1.01-1.02 in Fig. 6((, We show the light curve of inward and outward emission for $r_{ov}\sim1.01$ -1.02 in Fig. \ref{fig:3.1}( (
+D).,D).
+ The outward emissiou curve is very similar to the observatious in the 5- aud. X-“ay bands., The outward emission curve is very similar to the observations in the $\gamma$ - and X-ray bands.
+ Iu this pulsar. we need to use all the light curves simultaneously. iu order to deteruiue he range of rg.," In this pulsar, we need to use all the light curves simultaneously in order to determine the range of $r_{ov}$."
+" Since emissious with smaller values of ry, are not seen. if 5-ray aid opticalemissjon regions are separated Ni,> 0.10. sunilar to tlie Vela pulsar aud PSR J0659-4-111L we predict tlat an optical pulse profile cannot be observed or will be oulyveryweakly detected."," Since emissions with smaller values of $r_{ov}$ are not seen, if $\gamma$ -ray and opticalemission regions are separated $\Delta r_{ov}>0.10$ , similar to the Vela pulsar and PSR J0659+1414, we predict that an optical pulse profile cannot be observed or will be onlyveryweakly detected."
+ This is cousistent, This is consistent
+The separation between stars and galaxy was done with the colour-colour diagram (U—J) vs. (J—Ks) (see figure 2)).,The separation between stars and galaxy was done with the colour-colour diagram $(U-J)$ vs. $(J-K_s)$ (see figure \ref{fig:star_galaxy}) ).
+" The same colours were also computed from the Pickles(1998) stellar atmosphere models, in order to improve the boundaries between stars and galaxy, especially for the reddest stars which fall out of the main sequence (see figure 2 for full details)."," The same colours were also computed from the \citet{pickles1998} stellar atmosphere models, in order to improve the boundaries between stars and galaxy, especially for the reddest stars which fall out of the main sequence (see figure \ref{fig:star_galaxy} for full details)."
+" Galaxies were selected among the object satisfying the relation: As a cross-check, we run the EAZY code with the Pickles(1998) model stellar atmosphere and checked that the objects identified as galaxies via the two-colour diagram had a x? greater than the x? obtained on the same object with the EAZY galaxy template set."," Galaxies were selected among the object satisfying the relation: As a cross-check, we run the EAZY code with the \citet{pickles1998} model stellar atmosphere and checked that the objects identified as galaxies via the two-colour diagram had a $\chi^2$ greater than the $\chi^2$ obtained on the same object with the EAZY galaxy template set."
+" This criteria was satisfied by all objects previously selected as galaxies with only of the objects selected as stars showing a discordant value for the x”, giving confidence in our method to separate stars from galaxies."," This criteria was satisfied by all objects previously selected as galaxies with only of the objects selected as stars showing a discordant value for the $\chi^2$, giving confidence in our method to separate stars from galaxies."
+ A star-galaxy separation was also available in the original FIRES and FIREWORKS public catalogues., A star-galaxy separation was also available in the original FIRES and FIREWORKS public catalogues.
+" This selection was based on spectroscopy, SED-fitting with stellar templates and visual inspection of the object morphology (Rudnicket 2003)."," This selection was based on spectroscopy, SED-fitting with stellar templates and visual inspection of the object morphology \citep{rudnick2006,rudnick2003}."
+. We verified that those objects selected as stars in the FIRES catalogue were actually falling in the correct region of our color plot., We verified that those objects selected as stars in the FIRES catalogue were actually falling in the correct region of our two-color plot.
+" The availability of IRAC data for all our sample allows us to compute absolute magnitudes in the rest-frame J and H bands with little dependance on the SED templates; in fact, as an extreme case, the rest-frame H band at z=3.5 is shifted to the range 6.7—'7.9um, well bracketed by the channels 3 and 4, centered at 5.8 and 8um. Our final catalogue is composed by total of 14295 galaxies, with redshifts determinations to z=a6.2 and median redshift ZmedC1, distributed over an effective area of 450 arcmin?."," The availability of IRAC data for all our sample allows us to compute absolute magnitudes in the rest-frame $J$ and $H$ bands with little dependance on the SED templates; in fact, as an extreme case, the rest-frame $H$ band at $z=3.5$ is shifted to the range $6.7-7.9\mu$ m, well bracketed by the channels 3 and 4, centered at 5.8 and $\mu$ m. Our final catalogue is composed by a total of 14295 galaxies, with redshifts determinations to $z=6.2$ and median redshift $z_{med}\simeq1$, distributed over an effective area of 450 $^2$."
+" In the redshift range of interest, 1.5«z3.5, there is a total of 3496 objects, of which e6% have spectroscopic redshifts."," In the redshift range of interest, $1.5<z<3.5$, there is a total of 3496 objects, of which $\approx 6\%$ have spectroscopic redshifts."
+" For the measurement of the LF, we adopted three among the most widely used methods, namely the 1/Vinaz (Schmidt1968), the STY maximum likelihood (Sandage,Tammann,&Yahil1979) and the Step-Wise Maximum Likelihood (Efstathiou,Ellis,&Peterson1988).."," For the measurement of the LF, we adopted three among the most widely used methods, namely the $1/V_{max}$ \citep{schmidt1968}, the STY maximum likelihood \citep{styml} and the Step-Wise Maximum Likelihood \citep{swml}."
+ The need to analyze composite samples with different photometric depth was overcome by applying standard techniques available for each chosen method., The need to analyze composite samples with different photometric depth was overcome by applying standard techniques available for each chosen method.
+ Below we summarize them., Below we summarize them.
+" The 1/Vimaxz method, first introduced by Schmidt(1968),, was then generalized by Avni&Bahcall(1980) to allow the simultaneous analysis of composite samples."," The $1/V_{max}$ method, first introduced by \citet{schmidt1968}, was then generalized by \citet{ab1980} to allow the simultaneous analysis of composite samples."
+" This method presents a number of advantages: it is simple to code, it directly estimates the normalization of the LF, and it does not make any assumption on the spatial distribution of galaxies."," This method presents a number of advantages: it is simple to code, it directly estimates the normalization of the LF, and it does not make any assumption on the spatial distribution of galaxies."
+" The draw-back is that it is sensitive to the presence of clustering in the sample, affecting the estimate of the faint end."," The draw-back is that it is sensitive to the presence of clustering in the sample, affecting the estimate of the faint end."
+" In a given redshift interval [z1,z2], the galaxy number density per unit magnitude in the k-th absolute magnitude bin $,(M) is computed as follows: where AM is the width of the magnitude bin and Ngai is the number of galaxies in the redshift interval of interest."," In a given redshift interval $[z1,z2]$, the galaxy number density per unit magnitude in the $k$ -th absolute magnitude bin $\Phi_k(M)$ is computed as follows: where $\Delta M$ is the width of the magnitude bin and $N_{gal}$ is the number of galaxies in the redshift interval of interest."
+" In the analysisproposed by Avni&Bahcall(1980),, Vinaz,i is given by: with n, the number of samples constituting the full catalogue, €); the apparent area in steradians corresponding the the j-th sample, dV/dz the co-moving volume element per steradian and Zmaz,i,j the maximum redshift at which the ;-th galaxy could have been observed within the flux limit of the j-th sample."," In the analysisproposed by \citet{ab1980}, $V_{max,i}$ is given by: with $n_s$ the number of samples constituting the full catalogue, $\Omega_j$ the apparent area in steradians corresponding the the $j$ -th sample, $dV/dz$ the co-moving volume element per steradian and $z_{max,i,j}$ the maximum redshift at which the $i$ -th galaxy could have been observed within the flux limit of the $j$ -th sample."
+" Standard Poisson errors were associated to each $,(M).", Standard Poisson errors were associated to each $\Phi_k(M)$.
+" This method, developed by Efstathiou,Ellis,&Peterson(1988), is based on a maximum likelihood estimate and with a non-parametric form."," This method, developed by \citet{swml}, is based on a maximum likelihood estimate and with a non-parametric form."
+" This method relies on the fact that the number density of galaxies &(M,x) can be factorialized into a function which depends on luminosity only, $(M), and a factor depending on the position f(x) so that ®(M,x)=¢(M)-f(x)."," This method relies on the fact that the number density of galaxies $\Phi(M, {\bf x} )$ can be factorialized into a function which depends on luminosity only, $\phi(M)$, and a factor depending on the position $f({\bf x})$ so that $\Phi(M,{\bf x})=\phi (M)\cdot f({\bf x})$."
+" In particular, this means that the normalization factor of ® is lost during the maximization of the likelihood, and should then be determined in some other way."," In particular, this means that the normalization factor of $\Phi$ is lost during the maximization of the likelihood, and should then be determined in some other way."
+" The LF (M) is expressed as the sum over n steps as: By minimizing the natural logarithm of the likelihood expression, a recursive formula is found, allowing to compute the $i."," The LF $\Phi(M)$ is expressed as the sum over $n$ steps as: By minimizing the natural logarithm of the likelihood expression, a recursive formula is found, allowing to compute the $\phi_i$ ."
+" In order to take into account the non uniform magnitude limits of our data sets, we adopted the modification of the recursive expression as presented in Hilletal. (2010):"," In order to take into account the non uniform magnitude limits of our data sets, we adopted the modification of the recursive expression as presented in \citet{hill2010}: :"
+almost horizontal asviuptote.,almost horizontal asymptote.
+ There are also a nuniber of dwarf. egas-doiinated. and low-surface Irightucss (LSB) galaxies. e.g. DDO 163.," There are also a number of dwarf, gas-dominated, and low-surface brightness (LSB) galaxies, e.g. DDO 168."
+ There is no conspicuous maxinuun. and in some galaxies not even a flat asviuptote to their rotation curve.," There is no conspicuous maximum, and in some galaxies not even a flat asymptote to their rotation curve."
+ It is eenerallv believed. that deviations from the classical dyes is more pronounced in LSBs than iu IISDs. (MeCaush&deBlok1998a.Sanders&Noordermecr2007.Centileetal. 2010)," It is generally believed that deviations from the classical dynamics is more pronounced in LSBs than in HSBs. \cite{mcg98a,san07,gen10}) )"
+) Tweuty-cight members of the sample. of both ISB and LSB types. are located in the Ursa Major cluster of ealaxies. believed to be at the distance of about 15.5 Alpe (Tully&Verheijen 1997)).," Twenty-eight members of the sample, of both HSB and LSB types, are located in the Ursa Major cluster of galaxies, believed to be at the distance of about 15.5 Mpc \cite{tul97}) )."
+ Seven of the galaxies. listed in Table 2.. have central bulees and are treated differeutly. the reason is explained below.," Seven of the galaxies, listed in Table \ref{t3}, have central bulges and are treated differently, the reason is explained below."
+ For a full description of tle saluple the iuterested reader is referred to Saucers aud MeGanel (2002)., For a full description of the sample the interested reader is referred to Sanders and McGaugh (2002).
+ We calculate the rotation speed of a test object circling the galaxy as a function of distance from the galactic ceuter aud the distribution of the detectable matter in the ealaxy., We calculate the rotation speed of a test object circling the galaxy as a function of distance from the galactic center and the distribution of the detectable matter in the galaxy.
+" The procedure we follow is aliiost that of Sanders aud MeGaugh (2002): Iun order to calculate the MOND rotation curves. the first step is determining the Newtoniui accelerationof the detectable matter, gy via the classical Poisson equation."," The procedure we follow is almost that of Sanders and McGaugh (2002): In order to calculate the MOND rotation curves, the first step is determining the Newtonian accelerationof the detectable matter, $g_N$ via the classical Poisson equation."
+ Given the Newtoniau acceleration. the effective acceleration is calculated from the AIOND Eq. (1)).," Given the Newtonian acceleration, the effective acceleration is calculated from the MOND Eq. \ref{gmond}) )."
+ For MOCG. we approximate the ealaxy by a spherically sviuuetric svsten.," For MOG, we approximate the galaxy by a spherically symmetric system."
+ The error conumuitted in this approximation. as described in Binney Tremaine (19857). is on the order of," The error committed in this approximation, as described in Binney Tremaine (1987), is on the order of."
+ Wo asstune a constaut AY./£ ratio throughout the ealaxy. though this is not strictly the case. because of the color eradieut iu spiral galaxies.," We assume a constant $M_*/L$ ratio throughout the galaxy, though this is not strictly the case, because of the color gradient in spiral galaxies."
+ However. in seven bulged spirals we find assigning differeut. M./L to the bulee aud the disk improves the fit to the observed data.," However, in seven bulged spirals we find assigning different $M_*/L$ to the bulge and the disk improves the fit to the observed data."
+ We asstume the II eas is dn co-plauer rotation about the ceuter of the galaxw. an assumption which aay not hold iu galaxies with strong bars (Sauders&MeCGaugh 2002)).," We assume the HI gas is in co-planer rotation about the center of the galaxy, an assumption which may not hold in galaxies with strong bars \cite{san02}) )."
+ Given the observed distribution of the barvonic matter (stellar and gaseous disks. plus a spheroidal bulge. if present}. the effective radial eravitational force. aud subsequeutle the circular speed. is calculated from Eqs. (1))," Given the observed distribution of the baryonic matter (stellar and gaseous disks, plus a spheroidal bulge, if present), the effective radial gravitational force, and subsequently the circular speed, is calculated from Eqs. \ref{gmond}) )"
+ aud. (5)) (6))., and \ref{gmog}) ) - \ref{cdm}) ).
+ Fitting of the caleulated. rotation curves to the observed data points is achieved by adjusting the M./L ratio.. through a least-square U47. defined2 as where σι is the observational uncertainty iu the rotation speeds.," Fitting of the calculated rotation curves to the observed data points is achieved by adjusting the $M_*/L$ ratio, through a least-square $\chi^2$, defined as where $\sigma_i$ is the observational uncertainty in the rotation speeds."
+ The AL./£ ratio of the disk aud of the bulge are our ultimate results., The $M_*/L$ ratio of the disk and of the bulge are our ultimate results.
+ Allmodels trace the observed data with reasonable detail., Allmodels trace the observed data with reasonable detail.
+ The best-fit 4? aud M.E values are listed in Table l., The best-fit $\chi^2$ and $M_*/L$ values are listed in Table \ref{t1}.
+ Figures 1 - 3 show fits of theoretically constructed rotation curves to the observations of 16 galaxies., Figures \ref{f1} - \ref{f3} show fits of theoretically constructed rotation curves to the observations of 46 galaxies.
+ The general trend of USB curves ( steep rise to a maxinmui followed by eradual decline to an almost flat asviiptote]. and of LSB curves ( slow rise often with no asviuptote iu seht) are evident.," The general trend of HSB curves ( steep rise to a maximum followed by gradual decline to an almost flat asymptote), and of LSB curves ( slow rise often with no asymptote in sight) are evident."
+ Seven galaxies have prominent bulec compoucuts., Seven galaxies have prominent bulge components.
+ One expects a bulee with an older population of stars to have a lugher A.E ratio than a disk with a vounger population., One expects a bulge with an older population of stars to have a higher $M_*/L$ ratio than a disk with a younger population.
+ Therefore. to obtain a better ft for these galaxies. we have allowed the model to choose different AL./Ls for the disk and the bulge.," Therefore, to obtain a better fit for these galaxies, we have allowed the model to choose different $M_*/L$ s for the disk and the bulge."
+ The result. shown iu Table 2.. confirms the expectation.," The result, shown in Table \ref{t3}, , confirms the expectation."
+ The minim \?s of Table 2. are detectably lower than those of the corresponding cutries of Table 1 obtained by a single M./L fit., The minimum $\chi^2$ s of Table \ref{t3} are detectably lower than those of the corresponding entries of Table \ref{t1} obtained by a single $M_*/L$ fit.
+ Nevertheless. oue galaxy iu MOND3. NGC 801. and four in MONDI. NGC 801. NGC 5371. UGC 2885. and NGC 5907. predict uuteunably lower AL/£ for the bulee than for the disk.," Nevertheless, one galaxy in MOND3, NGC 801, and four in MOND1, NGC 801, NGC 5371, UGC 2885, and NGC 5907, predict untenably lower $M_*/L$ for the bulge than for the disk."
+ MeGanel (private communication) advises us that this oddity might partially be owing to the low resolution of II data aud/or the sharp rise of c(r) iu the bulge., McGaugh (private communication) advises us that this oddity might partially be owing to the low resolution of HI data and/or the sharp rise of $v(r)$ in the bulge.
+ If ο). ais not quite resolved. one ends to underestimate it. anc the M./L of the bulge with it.," If $v(r)$ is not quite resolved, one tends to underestimate it, and the $M_*/L$ of the bulge with it."
+ lu iuauv of hese galaxies it is indeed πιο obvious if the inner coniponeut is really a “bulec™ iu the classical scuse of a JD component with au rb! profile.," In many of these galaxies it is indeed not obvious if the inner component is really a “bulge"" in the classical sense of a 3D component with an $r^{1/4}$ profile."
+ It js also sak hat ueelecting the fattening of a bulec cads to au error dm its mmass-to-light ratio (Noorcdermecr 2008)., It is also said that neglecting the flattening of a bulge leads to an error in its mass-to-light ratio (Noordermeer 2008).
+ Tn NGC 6916. there is a tinv bulee of the tota ight iu D-band).," In NGC 6946, there is a tiny bulge of the total light in B-band)."
+ It shows up as a small kinematic np iu the immer 1 kpe iu high-resolution Fabry- data (Blais-Ouclletteetal. 2001))., It shows up as a small kinematic bump in the inner 1 kpc in high-resolution Fabry-P\'errot data \cite{bla04}) ).
+. Moreover. the ΤΠ distribution iu NGC 6916 is not sviunetric du the ealactic plane.," Moreover, the HI distribution in NGC 6946 is not symmetric in the galactic plane."
+ It is patchy and seems to deviate from circular orbits (Careginanetal. 1990))., It is patchy and seems to deviate from circular orbits \cite{car90}) ).
+ : ⊺↸∖⊔∖∐⊓⋅↕↸∖↴∖↴∐↕⊺⋜∏⋝↕↸∖⊥∐⋜↧↖⇁↸∖∏∐⋜↧↸⊳↸⊳↸∖↻↑⋜∏⋝↕∙↖↽↕⋜∐⋅∶↴⋁↸∖∖−↴∖↴∙∙⋅↴ ↽ ⋅≽ Suspecting the faihwe of the assuuption of constant ALL. we have followed Barnes etf al. (," Ten entries in Table \ref{t1} have unacceptably large $\chi^2$ s. Suspecting the failure of the assumption of constant $ M_*/L$, we have followed Barnes et al. ("
+2007). aud exanuned a racially varving. A/L—(M./L)y|iv.,"2007), and examined a radially varying, $M_*/L= (M_*/L)_0 + m r$."
+ The best-fit values of the constants (AL./£)y aud i for some of them are displaved in Table 3..," The best-fit values of the constants $
+(M_*/L)_0$ and $m $ for some of them are displayed in Table \ref{variable}. ."
+ The slope. i. Is 1uucli too stall o result in- appreciablyH lower D 47.," The slope, $m$, is much too small to result in appreciably lower $\chi^2$ ."
+ Thed reason forJ the failure should lie elsewhere., The reason for the failure should lie elsewhere.
+ For example. the assumptionof cold unobservable molecular gas in the galactic disk (TiretE Combes a2009) leadsto better fits⋅ with∙ lower D X7.," For example, the assumptionof cold unobservable molecular gas in the galactic disk (Tiret Combes 2009) leadsto better fits with lower $\chi^2$ ."
+ Wow realistic are the inferred AL./£ ratios?, How realistic are the inferred $M_*/L$ ratios?
+ Stellar population svuthesis (SPS) models predict a linear, Stellar population synthesis (SPS) models predict a linear
+frequency of 161.7 Hz (vertical dotted lines in Fig.,frequency of 161.7 $\mu$ Hz (vertical dotted lines in Fig.
+ 5 (top))., \ref{psdfull} (top)).
+ For each of them a daily modulation can be observed at + 111.57 wHz and twice this frequency (see the inset of Fig., For each of them a daily modulation can be observed at $\pm$ 11.57 $\mu$ Hz and twice this frequency (see the inset of Fig.
+ (top))., \ref{psdfull} (top)).
+" In particular, the region where we expect to find the acoustic-mode power excess (centered around 900 4Hz) is completely dominated by this peak structure."," In particular, the region where we expect to find the acoustic-mode power excess (centered around 900 $\mu$ Hz) is completely dominated by this peak structure."
+" Thus, we need to interpolate the data in the best way to reduce these effects as much as possible as it was already done in HD 175726 (?).. "," Thus, we need to interpolate the data in the best way to reduce these effects as much as possible as it was already done in HD 175726 \citep{2009A&A...506...33M}. ."
+For that star the algorithm proposed by (?) was used giving excellent results., For that star the algorithm proposed by \citep{Holds89} was used giving excellent results.
+" In our case, we adopted a different strategy based on the ""inpainting"" technique using a Multi-Scale Discrete CosineTransform (see Appendix A for a more detailed explanation on the algorithm)."," In our case, we adopted a different strategy based on the “inpainting” technique using a Multi-Scale Discrete CosineTransform (see Appendix A for a more detailed explanation on the algorithm)."
+ This algorithm has been tested and the results have been verified (seeforfurtherdetails?).., This algorithm has been tested and the results have been verified \citep[see for further details][]{2010arXiv1003.5178S}.
+" Moreover, different groups have also used other interpolation techniques (linear, based on wavelets, etc) leading to results for the global parameters of the modes (such as the mean large separation) that agree within the error bars."," Moreover, different groups have also used other interpolation techniques (linear, based on wavelets, etc) leading to results for the global parameters of the modes (such as the mean large separation) that agree within the error bars."
+ An independent measurement of the surface rotation can be obtained through the signature left by the active regions crossing the visible stellar disk., An independent measurement of the surface rotation can be obtained through the signature left by the active regions crossing the visible stellar disk.
+ It provides a first estimation of the stellar rotation., It provides a first estimation of the stellar rotation.
+" Moreover, combined with the measured vsini=19.0+1.5 (see Sect."," Moreover, combined with the measured v $\sin i = 19.0\,\pm\,1.5$ (see Sect."
+" 2) and assuming a given radius of the star, we can also estimate the inclination axis of the star, which can be very helpful to better constrain the fitting techniques to characterise the acoustic modes (???).."," 2) and assuming a given radius of the star, we can also estimate the inclination axis of the star, which can be very helpful to better constrain the fitting techniques to characterise the acoustic modes \citep{2003ApJ...589.1009G,2006MNRAS.369.1281B,2008A&A...486..867B}."
+" The light curve of HD 170987 shows a periodic modulation of about 4 days (see Fig. 4)),"," The light curve of HD 170987 shows a periodic modulation of about 4 days (see Fig. \ref{fig:lcurve}) ),"
+ which could be due to the combination of the stellar magnetic activity (star spots) and the rotation of the star., which could be due to the combination of the stellar magnetic activity (star spots) and the rotation of the star.
+ It is interesting to analyse this periodicity in a more detailed way by studying the low-frequency range of its power spectrum., It is interesting to analyse this periodicity in a more detailed way by studying the low-frequency range of its power spectrum.
+ This is shown in Fig., This is shown in Fig.
+ 6 where a clear peak appears at 2.70 uHz., \ref{figt} where a clear peak appears at 2.70 $\mu$ Hz.
+" This frequency corresponds to a period of 4.3 days, which can be associated to the rotation period of the surface of the star."," This frequency corresponds to a period of 4.3 days, which can be associated to the rotation period of the surface of the star."
+ At lower frequency —around 0.32 uHz- another structure appears which could be related to low-frequency noise., At lower frequency –around 0.32 $\mu$ Hz– another structure appears which could be related to low-frequency noise.
+" Indeed, considering that the radius is estimated to ~ 2 Ro and that the observed v sini is of ~kms,, the second peak is very unlikely to be produced by the rotation of the stellar surface."," Indeed, considering that the radius is estimated to $\sim$ 2 $R_\odot$ and that the observed v $\sin i$ is of $\sim$, the second peak is very unlikely to be produced by the rotation of the stellar surface."
+ Spot modelling of the light curve was performed according to the method exposed in ?.., Spot modelling of the light curve was performed according to the method exposed in \citet{2009A&A...506..245M}.
+" Since all the stars analyzed in that paper, with a spectral type close to HD 170987, show evidence of differential rotation, we used the fitting method D (Table 2 of ? ?)) with a fixed value of the differential rotation profile."," Since all the stars analyzed in that paper, with a spectral type close to HD 170987, show evidence of differential rotation, we used the fitting method D (Table 2 of \citeauthor{2009A&A...506..245M} \citeyear{2009A&A...506..245M}) ) with a fixed value of the differential rotation profile."
+ This method has proven to be efficient for the determination of the rotation period and of the spot lifetime., This method has proven to be efficient for the determination of the rotation period and of the spot lifetime.
+ Measurements of the star inclination and differential rotation are much less precise., Measurements of the star inclination and differential rotation are much less precise.
+ The best fit model is presented in Fig., The best fit model is presented in Fig.
+ 7 for a 60-day long sample., \ref{taches} for a 60-day long sample.
+" Like similar type of stars, the modelling is not able toreproduce the sharpest features of the light curve."," Like similar type of stars, the modelling is not able toreproduce the sharpest features of the light curve."
+ Fitting them, Fitting them
+order effect that varies with bolometer temperature and it is important that this is understood as correcting for it is critical in both the astrometric and photometric accuracy achievable in SPIRE maps.,order effect that varies with bolometer temperature and it is important that this is understood as correcting for it is critical in both the astrometric and photometric accuracy achievable in SPIRE maps.
+ The first assumption has been that the bolometer thermal response is represented by an ’RC’ type response characterised by a single time constant., The first assumption has been that the bolometer thermal response is represented by an 'RC' type response characterised by a single time constant.
+ Investigation into the presence of any slower bolometer time constants on scan map and spectrometer time domain data has so far revealed no evidence for a slow response to optical power and only the correction for a first order fast time response is currently being performed by the pipeline processing., Investigation into the presence of any slower bolometer time constants on scan map and spectrometer time domain data has so far revealed no evidence for a slow response to optical power and only the correction for a first order fast time response is currently being performed by the pipeline processing.
+ High energy particles hitting the bolometers cause signal glitches as the energy is absorbed and the temperature of the bolometer increases., High energy particles hitting the bolometers cause signal glitches as the energy is absorbed and the temperature of the bolometer increases.
+ The response to these glitches gives a direct measurement of the bolometer and electronics time constants., The response to these glitches gives a direct measurement of the bolometer and electronics time constants.
+ To measure the photometer bolometer time-constant the data from a large number of photometer glitches have been collected and compared to the theoretical impulse response modelled as a delta-function convolved with the photometer electronics and bolometer transfer functions., To measure the photometer bolometer time-constant the data from a large number of photometer glitches have been collected and compared to the theoretical impulse response modelled as a delta-function convolved with the photometer electronics and bolometer transfer functions.
+" Applying this process to sets of individual (single event) glitches, reveals that the photometer glitch response for most photometer bolometers can be best described by a two-component transfer function, with a fast time constant τι ο 6 ms and a second slow component with a time constant of το 7105 ms and amplitude about 6.6 (Fig."," Applying this process to sets of individual (single event) glitches, reveals that the photometer glitch response for most photometer bolometers can be best described by a two-component transfer function, with a fast time constant $\tau_1$ $\approx$ 6 ms and a second slow component with a time constant of $\tau_2$ $\approx$ 105 ms and amplitude about 6.6 (Fig."
+ 1)., 1).
+" A second type of common, concurrent glitches are also seen which are most likely due to interaction between particles and the bolometer array substrate."," A second type of common, concurrent glitches are also seen which are most likely due to interaction between particles and the bolometer array substrate."
+ Applying the same process to these reveals that the common glitch response is best described by a first-order-only bolometer transfer function with a time constant of about 6 ms., Applying the same process to these reveals that the common glitch response is best described by a first-order-only bolometer transfer function with a time constant of about 6 ms.
+" This disparity in the system response between single-event and common glitches, and how this relates to the absence of any slow response to optical power, is not yet understood and is currently being investigated."," This disparity in the system response between single-event and common glitches, and how this relates to the absence of any slow response to optical power, is not yet understood and is currently being investigated."
+" The frequency response correction of both the electronics and the bolometers amounts to no more than of the signal for standard scan map data and given our knowledge of the electronics and bolometers any uncertainties in this correction can, again, be safely ignored in the overall calibration uncertainty."," The frequency response correction of both the electronics and the bolometers amounts to no more than of the signal for standard scan map data and given our knowledge of the electronics and bolometers any uncertainties in this correction can, again, be safely ignored in the overall calibration uncertainty."
+" The conversion from voltage to flux density in the photometer is based on the assumption that the small signal responsivity can be related to the bolometer operating point voltage nearly linearly i.e. that we can write the inverse responsivity in the form where S is the flux density, K1,Ko,K3 are constants and V is the measured voltage across the bolometer."," The conversion from voltage to flux density in the photometer is based on the assumption that the small signal responsivity can be related to the bolometer operating point voltage nearly linearly i.e. that we can write the inverse responsivity in the form where S is the flux density, $K_1,K_2,K_3$ are constants and $V$ is the measured voltage across the bolometer."
+" K is in units of Jy/V, Κο is in Jy and Ka is in V. Integrating this equation gives where now Vo is another[Vo arbitraryKy] constant that we chose to represent the bolometer voltage measured under a standard observing condition on a dark region of the sky."," $K_1$ is in units of Jy/V, $K_2$ is in Jy and $K_3$ is in V. Integrating this equation gives where now $V_0$ is another arbitrary constant that we chose to represent the bolometer voltage measured under a standard observing condition on a dark region of the sky."
+ The various parameters have been determined from measurements in-flight using PCAL to measure the relative responsivity while viewing sources of different brightnesses in order to change the bolometer voltage., The various parameters have been determined from measurements in-flight using PCAL to measure the relative responsivity while viewing sources of different brightnesses in order to change the bolometer voltage.
+ The calibration curves were measured for each detector in two steps., The calibration curves were measured for each detector in two steps.
+" First, the change in voltage AVpgcaAr observed for PCAL flashes were measured when the telescope was pointed at regions of the sky with different flux densities ranging from dark sky to the peak of Sgr A (for the nominal voltage settings) or Sgr B2 (for the lower gain bright source mode settings)."," First, the change in voltage $\Delta V_{PCAL}$ observed for PCAL flashes were measured when the telescope was pointed at regions of the sky with different flux densities ranging from dark sky to the peak of Sgr A (for the nominal voltage settings) or Sgr B2 (for the lower gain bright source mode settings)."
+ The plot of l/AVpcaAr versus V gives the unscaled version of Eq. 1;;, The plot of $1/\Delta V_{PCAL}$ versus $V$ gives the unscaled version of Eq. \ref{equation1};
+ an example is shown in Fig., an example is shown in Fig.
+ 2 where the telescope was pointed at the dark sky and different locations in the region of SgrA to vary the flux onto the detectors., \ref{figure_calibration} where the telescope was pointed at the dark sky and different locations in the region of SgrA to vary the flux onto the detectors.
+ The calibration curve itself is then scaled using data from observations in which Neptune (or a source calibrated using Neptune) is scanned across each detector in each array., The calibration curve itself is then scaled using data from observations in which Neptune (or a source calibrated using Neptune) is scanned across each detector in each array.
+ The feedhorns and bolometers are not inherently sensitive to polarisation and any instrinsic polarisation in the calibration targets will have negligible effect on the overall photometric calibration., The feedhorns and bolometers are not inherently sensitive to polarisation and any instrinsic polarisation in the calibration targets will have negligible effect on the overall photometric calibration.
+ The initial absolute calibration scale was set using the asteroid Ceres and a model of its submillimetre flux density based on the Standard Thermal Model (Lebofsky et, The initial absolute calibration scale was set using the asteroid Ceres and a model of its submillimetre flux density based on the Standard Thermal Model (Lebofsky et
+determined.,determined.
+ In this caleulation we set 7;=LO’ K. The only parameter sensitive to 7; is may in the leptonic scenario., In this calculation we set $T_e=10^7$ K. The only parameter sensitive to $T_e$ is $n_{\rm ISM}$ in the leptonic scenario.
+ The 2o upper limit of rg is 0.007 * for T;=10* IX. The dotted line in the right panel of Fig., The $2\sigma$ upper limit of $n_{\rm ISM}$ is $0.007$ $^{-3}$ for $T_e=10^7$ K. The dotted line in the right panel of Fig.
+ 1. indicates the thermal emission for these parameters., \ref{fig:lepton} indicates the thermal emission for these parameters.
+ since the ISM density will be more essential lor the discussion of the hybrid ancl hadronic models. we will discuss the 7; dependence of these results in detail in the subsection 2.4.," Since the ISM density will be more essential for the discussion of the hybrid and hadronic models, we will discuss the $T_e$ dependence of these results in detail in the subsection 2.4."
+ For the 2-D confidence regions of the parameters. we onlv show combinations with relatively large correlation.," For the 2-D confidence regions of the parameters, we only show combinations with relatively large correlation."
+" There are very weak correlations among a,. £L. W, and D."," There are very weak correlations among $\alpha_e$, $E_c^e$, $W_e$ and $B$."
+" The weak correlation between a, and V, is mostly due to the [acts that electrons near the high- cutoff are well-constrained by observations and low-energy electrons contribute the most to WW; for a,>2.", The weak correlation between $\alpha_e$ and $W_e$ is mostly due to the facts that electrons near the high-energy cutoff are well-constrained by observations and low-energy electrons contribute the most to $W_e$ for $\alpha_e>2$.
+" The correlation between 0, and ZZ. is caused by the well observed spectral shape in hard XN-ravs aud TeV 5-ravs.", The correlation between $\delta_e$ and $E_c^e$ is caused by the well observed spectral shape in hard X-rays and TeV $\gamma$ -rays.
+ The combination of the A-ray and οταν data helps to determine the model parameters., The combination of the X-ray and $\gamma$ -ray data helps to determine the model parameters.
+ In the following we will see Chat X-ray data alone lead to poorly constrained aud highly correlated parameters., In the following we will see that X-ray data alone lead to poorly constrained and highly correlated parameters.
+ This is an example to show the importance of global fit to the multi-wavelengthi data., This is an example to show the importance of global fit to the multi-wavelength data.
+ The parameters and 47 values of the best fit model are compiled in Tables 1 and 2.. respectively.," The parameters and $\chi^2$ values of the best fit model are compiled in Tables \ref{table:best} and \ref{table:chi2}, respectively."
+ Since nis is not well constrained. ils 20 upper limit instead of the best fit value is listed in Table 1..," Since $n_{\rm ISM}$ is not well constrained, its $2\sigma$ upper limit instead of the best fit value is listed in Table \ref{table:best}."
+ The best-fit parameters are consistent with previous studies (Aharonianal.2008:Tanakaοἱ2008) and can be readily accommodated with typical SNRs (Ellisonetal.2010)..," The best-fit parameters are consistent with previous studies \citep{2006A&A...449..223A,2008ApJ...683L.163L,2008ApJ...685..988T} and can be readily accommodated with typical SNRs \citep{2010ApJ...712..287E}."
+ The overall V? of the fit is relatively large with the reduced 4? ~466.9/232=2.01., The overall $\chi^2$ of the fit is relatively large with the reduced $\chi^2$ $\sim 466.9/232=2.01$.
+ Such a high value of 4? shows that svstematic errors dominate., Such a high value of $\chi^2$ shows that systematic errors dominate.
+ This is nol surprising given the relatively complex structure of the SNR. uncertainties related to the particle acceleration process. and our rather simple prescription of the emission model.," This is not surprising given the relatively complex structure of the SNR, uncertainties related to the particle acceleration process, and our rather simple prescription of the emission model."
+ The, The
+From this limit we derive for the apparent velocity This is in accord with the observation that detected EGRET sources. Le. primarily MeV — GeV peaked emitters. often show rapid superluminal motion. whereas the only TeV peaked source with good VLBI coverage. namely Mkn 421. displayed subluminal motion m the time range of 1994 to 1997 (Piner et al. 1999)).,"From this limit we derive for the apparent velocity This is in accord with the observation that detected EGRET sources, i.e. primarily MeV – GeV peaked emitters, often show rapid superluminal motion, whereas the only TeV peaked source with good VLBI coverage, namely Mkn 421, displayed subluminal motion in the time range of 1994 to 1997 (Piner et al. \cite{pin99}) )."
+ All spectra exhibit a strong annihilation line at an energy of D-511 keV. The copiously produced positrons cool down to thermal energies before annihilating. and a relatively narrow line is produced at a given time.," All spectra exhibit a strong annihilation line at an energy of $D\cdot 511\,$ keV. The copiously produced positrons cool down to thermal energies before annihilating, and a relatively narrow line is produced at a given time."
+ However. since the Doppler factor may vary during the integration time of current detectors (typically three weeks (1)," However, since the Doppler factor may vary during the integration time of current detectors (typically three weeks (!)"
+ in case of COMPTEL). the line may look like a broad hump in the data.," in case of COMPTEL), the line may look like a broad hump in the data."
+ The spectra shown so far illustrate the effect of the blast wave deceleration and the thus modulated injection rate., The spectra shown so far illustrate the effect of the blast wave deceleration and the thus modulated injection rate.
+ To this we oppose in Fig., To this we oppose in Fig.
+ 6 and Fig., \ref{spece} and Fig.
+ 7 the case of the blast wave interacting with an isolated gas cloud.," \ref{specg}
+ the case of the blast wave interacting with an isolated gas cloud."
+ The light curves of high energy sources often display phases of activity lasting for months or years. on which rapid fluctuations are superimposed (e.g. Quinn et al. 1999)).," The light curves of high energy sources often display phases of activity lasting for months or years, on which rapid fluctuations are superimposed (e.g. Quinn et al. \cite{qui99}) )."
+ We propose that the secular variability is related to the existence or non-existence of a relativistic blast wave in the sources. and thus to the availability of free energy in the system.," We propose that the secular variability is related to the existence or non-existence of a relativistic blast wave in the sources, and thus to the availability of free energy in the system."
+ Strong emission on the other hand requires that the available kinetic energy of the blast wave is converted into radiation., Strong emission on the other hand requires that the available kinetic energy of the blast wave is converted into radiation.
+ Depending on the distribution of the interstellar matter the kinetic energy of the blast wave may be tapped only every now and then., Depending on the distribution of the interstellar matter the kinetic energy of the blast wave may be tapped only every now and then.
+ The observed fast variability. of high-energy 5-ray sources would thus be caused by density inhomogeneities in the interstellar medium of the sources., The observed fast variability of high-energy $\gamma$ -ray sources would thus be caused by density inhomogeneities in the interstellar medium of the sources.
+ The situation displayed in Fig., The situation displayed in Fig.
+ 6 would then correspond to one of the rapid outbursts in the >-ray light curves of AGN. which would generally follow the gas density variations in the volume traversed by the blast wave.," \ref{spece} would then correspond to one of the rapid outbursts in the $\gamma$ -ray light curves of AGN, which would generally follow the gas density variations in the volume traversed by the blast wave."
+ The large collection area of currently operating Ceerenkov telescopes permits studies of spectral variability on very short timescales., The large collection area of currently operating Čeerenkov telescopes permits studies of spectral variability on very short timescales.
+ The observational situation is. disturbing. however. since data taken with different telescopes can be seemingly inconsistent.," The observational situation is disturbing, however, since data taken with different telescopes can be seemingly inconsistent."
+ As an example. both the CAT team and the HEGRA group performed observation of Mkn 501 in the time of March 1997 to October 1997.," As an example, both the CAT team and the HEGRA group performed observation of Mkn 501 in the time of March 1997 to October 1997."
+ The CAT data show a statistically significant correlation between the spectral hardness and flux in the energy range between 0.33 TeV and 13 TeV. The HEGRA data on the other hand give no evidence in favor of a correlation in the energy range between | TeV and 10 TeV. Both behaviours could be reproduced with our model., The CAT data show a statistically significant correlation between the spectral hardness and flux in the energy range between 0.33 TeV and 13 TeV. The HEGRA data on the other hand give no evidence in favor of a correlation in the energy range between 1 TeV and 10 TeV. Both behaviours could be reproduced with our model.
+ In Fig., In Fig.
+ 8 we show the spectral variability at TeV energies only., \ref{tevspec} we show the spectral variability at TeV energies only.
+ This plot is essentially a blow-up of Fig. 6.., This plot is essentially a blow-up of Fig. \ref{spece}.
+ It is obvious that the rise time of an outburst is much smaller than the fall time., It is obvious that the rise time of an outburst is much smaller than the fall time.
+ Also in the declining phase the source undergoes strong spectral evolution. a consequence of the proton energy loss timescale," Also in the declining phase the source undergoes strong spectral evolution, a consequence of the proton energy loss timescale"
+for inagiug iu ultra deep fields. this may of course be very difficult to achieve with JAVST itself.,"for imaging in ultra deep fields, this may of course be very difficult to achieve with JWST itself."
+ IHToscever. such follow-up observations may for mstance be feasible with the upcoming 12 m European Extremely Lareescopot?.," However, such follow-up observations may for instance be feasible with the upcoming 42 m European Extremely Large."
+", Another caveat associated with weak or abscut [OTTASO07 cinissiou as a criterion for pop III galaxies is that pop II/I galaxies with very low ionization parameters may also display very low |OIII| fluxes.", Another caveat associated with weak or absent $\lambda$ 5007 emission as a criterion for pop III galaxies is that pop II/I galaxies with very low ionization parameters may also display very low [OIII] fluxes.
+ In the framework of a single spherical nebulae surrounding the star-forming core of the galaxy. this could correspond to the situation where stellar feedback has increased the inner radius of the cloud to large radii," In the framework of a single spherical nebulae surrounding the star-forming core of the galaxy, this could correspond to the situation where stellar feedback has increased the inner radius of the cloud to large radii."
+ Towever. such objects may no longer qualify as type A objects but may rather belong to type B (sce Fig. 3)).," However, such objects may no longer qualify as type A objects, but may rather belong to type B (see Fig. \ref{schematic}) ),"
+ with a drastically lower nebular fluxes and potentially distinct spectral characteristics as a result., with a drastically lower nebular fluxes and potentially distinct spectral characteristics as a result.
+ A detailed investigation of the properties of type D galaxies would be required to settle this issue., A detailed investigation of the properties of type B galaxies would be required to settle this issue.
+ One max. perhaps. argue that narrowband surveys are more likely to be successful in isolating pop III ealaxies than any broadband survey.," One may, perhaps, argue that narrowband surveys are more likely to be successful in isolating pop III galaxies than any broadband survey."
+" Aside from ""elassieal? narrowband features of pop IIT objects. likeo very high Ίνα or |ÉHeHl|AT1G10 equivalent widths. one iav for instance look for the Balmer jump at 0.36 jun (c.g.Schacrer2002.2003:Inoue 2011b)."," Aside from “classical” narrowband features of pop III objects, like very high $\alpha$ or $\lambda1640$ equivalent widths, one may for instance look for the Balmer jump at $0.36\ \mu$ m \citep[e.g.][]{Schaerer a,Schaerer b,Inoue b}."
+. All objects with spectra dominated by nebular cussion should cisplayv a sudden juup in the coutimuun level at the short-wavelength side compared to the long wavelength side of this limit., All objects with spectra dominated by nebular emission should display a sudden jump in the continuum level at the short-wavelength side compared to the long wavelength side of this limit.
+ Since metal-enriched stars are less eficieut in jonizine hydrogen aud have redder stell contiuua. hey wil relatively speaking be less dominated o» nebula enmüssiou (see Fie. 2)).," Since metal-enriched stars are less efficient in ionizing hydrogen and have redder stellar continua, they will – relatively speaking – be less dominated by nebular emission (see Fig. \ref{fneb}) )."
+ Therefore. oue may conjecture that the objects with the most pronounced Bahuer discoutinuities are likely to be pop III galaxies. aud that suitably placed niurowbaud filters would be able to pick up on this.," Therefore, one may conjecture that the objects with the most pronounced Balmer discontinuities are likely to be pop III galaxies, and that suitably placed narrowband filters would be able to pick up on this."
+ However. our models mdicate hat the difference between the Bahucr jumps between x» II aud pop III galaxies are so small that dt will be next to impossible to tell them apart using lis feature.," However, our models indicate that the difference between the Balmer jumps between pop II and pop III galaxies are so small that it will be next to impossible to tell them apart using this feature."
+ While the Balmer jump iav still be used for ideutifvius extremely young. nebular-douinuated objects at low inctallicities (albeit not necessarily just yop III objects). the redshift range over which this eature can be efficiently probed with a given setup of JIWST narrowband filters is very small (A:~ 0.1).," While the Balmer jump may still be used for identifying extremely young, nebular-dominated objects at low metallicities (albeit not necessarily just pop III objects), the redshift range over which this feature can be efficiently probed with a given setup of JWST narrowband filters is very small $\Delta z \sim 0.1$ )."
+ This makes the nuuber of potential targets iu the correspouding volume so tim. that it would seem iore time-cficicut to use broadbaud criteria to pre-select pop III ealaxy caudidates and then to eo directly for follow-up spectroscopy.," This makes the number of potential targets in the corresponding volume so tiny, that it would seem more time-efficient to use broadband criteria to pre-select pop III galaxy candidates and then to go directly for follow-up spectroscopy."
+" Dronuuetal.(2001) argue that very massive pop III stars are sclfsimular. in the sense that a fixed mass of eas converted into either 100AY... 300AL. or 500AL, stars would produce ucarly identical SEDs."," \citet{Bromm et al. b} argue that very massive pop III stars are self-similar, in the sense that a fixed mass of gas converted into either $100\ M_\odot$, $300\ M_\odot$ or $500\ M_\odot$ stars would produce nearly identical SEDs."
+" Indeed. the lifetimes. ZAMS temperatures and the time-averaged. lnass-hormalized Lyinan coutiuuuni fluxes of AMz100- AZ, stars differ by no more than LO230 (Schaerer2002)."," Indeed, the lifetimes, ZAMS temperatures and the time-averaged, mass-normalized Lyman continuum fluxes of $M\approx 100$ -500 $M_\odot$ stars differ by no more than 10–30 \citep{Schaerer a}."
+. ILowever. the approximation of selfsmnuibuitv deteriorates once pop III stars at AL<100AL. are considered. aud does not hold for the Well ionizing spectrin and the corresponding recombination lines. or for stars that have evolved off the main sequence. as discussed. im detail by Schaerer(2002).," However, the approximation of self-similarity deteriorates once pop III stars at $M<100\ M_\odot$ are considered, and does not hold for the HeII ionizing spectrum and the corresponding recombination lines, or for stars that have evolved off the main sequence, as discussed in detail by \citet{Schaerer a}. ."
+ As secu in Fie. L.," As seen in Fig. \ref{Mmin},"
+ the mass-to-lght ratios produced by our pop IIL.1 (—10037..) aud pop IIL2 (~10AZ.) IMEs differ bv a factor of z2 25., the mass-to-light ratios produced by our pop III.1 $\sim 100\ M_\odot$ ) and pop III.2 $\sim 10\ M_\odot$ ) IMFs differ by a factor of $\approx 2$ –3.
+ This is mostly due to the inclusion of low-mass stars which contribute substantially to the mass and not to the hnunuinosity., This is mostly due to the inclusion of low-mass stars which contribute substantially to the mass and not to the luminosity.
+ Hence. the colours produced by the different. pop III IMEs remain verv similar (Figs. 6.. 7..8," Hence, the colours produced by the different pop III IMFs remain very similar (Figs. \ref{typeA_Inoue_crit}, \ref{typeA_MIRIcrit},"
+ aud 11))., \ref{typeA_MIRIcolcol} and \ref{typeC_colcrit}) ).
+ For a fixed star formation history. galaxies with pop IIT.2 IAIFs can continue to sline for longer than those with pop IIL.1 IMIFs. but their colours are otherwise identical.," For a fixed star formation history, galaxies with pop III.2 IMFs can continue to shine for longer than those with pop III.1 IMFs, but their colours are otherwise identical."
+ Pop IIT ealaxies with Kroupa(2001) TMIFs display slightly less extreme colours. but this is lavecly caused by a shift in the time scales op IIL2 ealaxies will display the sale colours. just at a slehtly higher age.," Pop III galaxies with \citet{Kroupa} IMFs display slightly less extreme colours, but this is largely caused by a shift in the time scales – pop III.2 galaxies will display the same colours, just at a slightly higher age."
+ It is oulv in the case of a purely stellar SED (our case €: Fie. 113) , It is only in the case of a purely stellar SED (our case C; Fig. \ref{typeC_colcrit}) )
+that Pop IIL1 aud pop IIL2 ect separated from the Iroupa(2001) IME. even at very. low ages., that Pop III.1 and pop III.2 get separated from the \citet{Kroupa} IMF even at very low ages.
+ Hence. the detection of objects with extreme colours would support the idea of a pop IIT TIF with more hieli-1uass stars thu predicted by the Iroupa(2001) IME. but the prospects of diseutaugliue a pop HELI IME from a pop IIEL.2 IME using integrated colours aloue secur bleak.," Hence, the detection of objects with extreme colours would support the idea of a pop III IMF with more high-mass stars than predicted by the \citet{Kroupa} IMF, but the prospects of disentangling a pop III.1 IMF from a pop III.2 IMF using integrated colours alone seem bleak."
+ Yeedrasil is a new spectral svuthesis model which includes pop IL IP aud III stars. dark stars. nebular cluission and dust extinction.," Yggdrasil is a new spectral synthesis model which includes pop I, II and III stars, dark stars, nebular emission and dust extinction."
+ Using this model. we derive the stellar population masses of the faimtest calaxics detectable by JAVST through broadband imaging in ultra deep fields.," Using this model, we derive the stellar population masses of the faintest galaxies detectable by JWST through broadband imaging in ultra deep fields."
+ Asstuine 100 Lh exposures per filter. we find that JWST may be able to detect pop III galaxies with stellar masses as low as ~LO’AJ. at i=1η. whereas the corresponding huit for pop II/I ealaxics is ~LO°AL...," Assuming 100 h exposures per filter, we find that JWST may be able to detect pop III galaxies with stellar masses as low as $\sim 10^5\ M_\odot$ at $z=10$, whereas the corresponding limit for pop II/I galaxies is $\sim 10^6 M_\odot$."
+ We also arene that the broadband fixes of voung (age =10 Myr) pop TIT galaxies are likely to be significantly affected by nebular euissiou. unless the fraction of ionizing radiation escapiug directly iuto the intergalactic medimnu is extremely high (fii2 0.95).," We also argue that the broadband fluxes of young (age $\lesssim 10$ Myr) pop III galaxies are likely to be significantly affected by nebular emission, unless the fraction of ionizing radiation escaping directly into the intergalactic medium is extremely high $f_\mathrm{esc}\gtrsim 0.95$ )."
+ We also discuss different strategies for selecting Pop III ealaxy candidates based on their JWST broadbaud colours. both in the limiting cases of Lavine cominant (low fice) aud neeligible (very high fic} nebular cluission.," We also discuss different strategies for selecting Pop III galaxy candidates based on their JWST broadband colours, both in the limiting cases of having dominant (low $f_\mathrm{esc}$ ) and negligible (very high $f_\mathrm{esc}$ ) nebular emission."
+ Iu the former case. we particularly highlieht the possibility of adding damaging in two MIBI filters (E560 and FTFOW) to the NTRCain filter sets usually discussed in the context of IWST ultra deep fields. since this should allow for a clean selection of pop III galaxies at zz7 8.," In the former case, we particularly highlight the possibility of adding imaging in two MIRI filters (F560W and F770W) to the NIRCam filter sets usually discussed in the context of JWST ultra deep fields, since this should allow for a clean selection of pop III galaxies at $z\approx 7$ –8."
+ Selecting targets for spectroscopic follow-up using colour criteria based solely on NIRCam filters would be more economical in terms of exposure time. but also increases the risk of interlopers.," Selecting targets for spectroscopic follow-up using colour criteria based solely on NIRCam filters would be more economical in terms of exposure time, but also increases the risk of interlopers."
+ Iu the case ofpop III galaxies dominated by direct star light (very. high fic). we argue that iiaeiue m the NIRCam F200W and ELLINV filters should allow pop III candidates to be selected throughout," In the case ofpop III galaxies dominated by direct star light (very high $f_\mathrm{esc}$ ), we argue that imaging in the NIRCam F200W and F444W filters should allow pop III candidates to be selected throughout"
+work (lrastevetal.2008) we have reported predictions on the gravitational masses. racii. maximal rotational (Ixepler) frequencies. and thermal properties of (rapidly) rotating neutron stars.,"work \citep{KLW2} we have reported predictions on the gravitational masses, radii, maximal rotational (Kepler) frequencies, and thermal properties of (rapidly) rotating neutron stars."
+ In (his work. using the nuclear constrained EOSs we calculate the momenta of inertia for both spherically-svinmetric (static) and (rapidly) rotating neutron stars using well established formalisms in the literature.," In this work, using the nuclear constrained EOSs we calculate the momenta of inertia for both spherically-symmetric (static) and (rapidly) rotating neutron stars using well established formalisms in the literature."
+ Such studies are important and timely as thev are related to the astrophysical observations in the near future., Such studies are important and timely as they are related to the astrophysical observations in the near future.
+ In. particular. (he moment of inerüa of pulsar οἱ in the extremely relativistic neutron star binary P5R. J0737-3039 (ήσανetal.2003) may be determined in a lew vears through detailed measurements of the periastron advance (Bejgerοἱal.2005).," In particular, the moment of inertia of pulsar $A$ in the extremely relativistic neutron star binary PSR J0737-3039 \citep{Burgay2003} may be determined in a few years through detailed measurements of the periastron advance \citep{BBH2005}."
+. In this section. we first outline the theoretical tools one uses to extract. information about the EOS of neutron-rich nuclear matter from heavy-ion collisions.," In this section, we first outline the theoretical tools one uses to extract information about the EOS of neutron-rich nuclear matter from heavy-ion collisions."
+ We put the special emphasis on exploring the clensily-clepencence of the symmetry energv as the study on ihe EOS of svimmetric nuclear matter with heavy-ion reactions is better known to the astrophysical community. and it has been extensively reviewed. see e.g.. Rels.," We put the special emphasis on exploring the density-dependence of the symmetry energy as the study on the EOS of symmetric nuclear matter with heavy-ion reactions is better known to the astrophysical community and it has been extensively reviewed, see e.g., Refs."
+ (Danielewicz.Lynch2002:Steineretal.2005;Lattimer&Prakash2007) lor recent reviews.," \citep{Danielewicz:2002pu,Steiner:2004fi,Lattimer:2007} for recent reviews."
+ We will then stummarize the latest constraints on the density dependence of (he svametry energv extracted [rom studying isospin diffusion and isoscaling in heavy-ion reactions al intermediate energies., We will then summarize the latest constraints on the density dependence of the symmetry energy extracted from studying isospin diffusion and isoscaling in heavy-ion reactions at intermediate energies.
+ Finally. we address (he question of what kind of isospin-asviinelry. especially [or dense matter. can be reached in heavy-ion reactions.," Finally, we address the question of what kind of isospin-asymmetry, especially for dense matter, can be reached in heavy-ion reactions."
+ lleavy-ion reactions are a unique means {ο create in terrestrial laboratories dense nuclear malter similar to those found in the core of neutron stus., Heavy-ion reactions are a unique means to create in terrestrial laboratories dense nuclear matter similar to those found in the core of neutron stars.
+ Depending on the beam energy. impact parameter and the reaction svstem. various hadrons and/or partons may be created during the reaction.," Depending on the beam energy, impact parameter and the reaction system, various hadrons and/or partons may be created during the reaction."
+ To extract information about the EOS of dense matter [rom reactions requires careful modelling of the reaction dynamics and selection of sensitive observables., To extract information about the EOS of dense matter from heavy-ion reactions requires careful modelling of the reaction dynamics and selection of sensitive observables.
+ Among the available tools. computer simulations based on the Boltzmann-Uehline-Uhlenbeck (BUU) transport theory have been very. useful. see. e.g.. Refs.," Among the available tools, computer simulations based on the Boltzmann-Uehling-Uhlenbeck ) transport theory have been very useful, see, e.g., Refs."
+ (BertschDanielewiez.Lacey.&Lynch2002). for reviews.," \citep{Bertsch,Danielewicz:2002pu} for reviews."
+ The evolution of the phase space distribution Iunction. f;(r.p./) of nucleon / is governed by both the mean field potential Ü and the collision integral {ο via the equation Normally. effects of the collision integral £445; via both elastic and inelastic channels," The evolution of the phase space distribution function $f_i(\vec{r},\vec{p},t)$ of nucleon $i$ is governed by both the mean field potential $U$ and the collision integral $I_{collision}$ via the equation Normally, effects of the collision integral $I_{collision}$ via both elastic and inelastic channels"
+There are two possible sites for the r-process here in those cases where } is low - in the jet aud in the disk wind.,There are two possible sites for the $r$ -process here in those cases where $Y_e$ is low - in the jet and in the disk wind.
+ The jet has the merit of originating in the vicinity of the black hole where the neutron excess is likely to be greatest., The jet has the merit of originating in the vicinity of the black hole where the neutron excess is likely to be greatest.
+ If magnetic fields drive the heating and initial acceleration of the jet. Chen this neutron excess will likely be preserved.," If magnetic fields drive the heating and initial acceleration of the jet, then this neutron excess will likely be preserved."
+ If neutrinos drive the outflow (AlacFaclven&Wooslev1999).. Y. will be reset to some extent by the weak interactions.," If neutrinos drive the outflow \citep{mac99}, $Y_e$ will be reset to some extent by the weak interactions."
+ llowever. if the details of the outflow above the black hole are similar to spherically svaumnietric neutrino driven ultra-relativistie winds. then only v7 annihialion will be important and the flow will remain neutron rich (Pruet.Fuller.&Cardall2001).," However, if the details of the outflow above the black hole are similar to spherically symmetric neutrino driven ultra-relativistic winds, then only $\nu \bar \nu$ annihilation will be important and the flow will remain neutron rich \citep{pru01}."
+. It is clear that the rapid expansion of the jet will be avorable to freezing out. with a large abundance of Iree neutrons and (his is conducive to the r-process (Llolfman. 1997)., It is clear that the rapid expansion of the jet will be favorable to freezing out with a large abundance of free neutrons and this is conducive to the $r$ -process \citep{hof97}.
+. However. a pure nucleonic Jet. coming Iron the adiabatic expansion of a fireball with energy loading 7=Eiycsgu/pc:100 is too muc 10f a good thing when it comes to entropy. and rapid expansion.," However, a pure nucleonic jet coming from the adiabatic expansion of a fireball with energy loading $\eta = E_{\rm internal}/\rho c^2 \gtaprx 100$ is too much of a good thing when it comes to entropy and rapid expansion."
+ The entropy was given in eq., The entropy was given in eq.
+ 12. (Qian&Woosley1996) The internal energve. (ere/enm)oj0 is and hence. fora=200 and AT~ 2 MeV The tine scale for the expansion. |fc. is less Chan a millisecond.," \ref{entropy} \citep{qia96}
+ The internal energy (erg/gm) is and hence, for $\eta = 200$ and $kT \sim$ 2 MeV The time scale for the expansion, $r/c$, is less than a millisecond."
+ The hieh entropy implies a low densitv at the Cemiperature when » and p can start to recombine., The high entropy implies a low density at the temperature when $n$ and $p$ can start to recombine.
+ Coupled with the rapid expaisjon. one mieht conclude that the jet will remain pure nucleons and light nuclei and (hat an r-process is impossible.," Coupled with the rapid expansion, one might conclude that the jet will remain pure nucleons and light nuclei and that an $r$ -process is impossible."
+ This conclusion would be erroneous on several counts., This conclusion would be erroneous on several counts.
+ First. the jet may not always have such high energy loading.," First, the jet may not always have such high energy loading."
+ There nav be cases where no gamnma-ray burst is produced, There may be cases where no gamma-ray burst is produced
+Diffuse LIutrared Background Experiment (DIRBE) iu comparison with the COLe distributio1 and the second a sticky particles model based on star counts from the Two Micron All Sky Survey (2MASS)). but neither of these models account for more than about of the 3-kpe arum maser population (a third of the level of associion of either the expauding ring or t1e ellipse).,"Diffuse Infrared Background Experiment (DIRBE) in comparison with the CO distribution and the second a sticky particles model based on star counts from the Two Micron All Sky Survey (2MASS)), but neither of these models account for more than about of the 3–kpc arm maser population (a third of the level of association of either the expanding ring or the ellipse)."
+ lore complex asvnunetric models which incorporate several arm-like features emanating from the bar [are no better: the models of ?. (a two climeusional quasi-steady state model with wear aud Lar liner Lindblad radius arms aud immer ultra-haruouie arms). ? (a quasi-equilibriui1 flow model based on COBE/DIRBE near infrared luminosities) and ? (a three d[umuensioual N-body inodel based on COBE/DIRBE near infrared luminosities incorporating the | arm aud connecting) arms). each associate no more than of the 3-kpC arn iasers.," More complex asymmetric models which incorporate several arm-like features emanating from the bar fare no better: the models of \citet{mulder86} (a two dimensional quasi-steady state model with near and far inner Lindblad radius arms and inner ultra-harmonic arms), \citet{englmaier99} (a quasi-equilibrium flow model based on COBE/DIRBE near infrared luminosities) and \citet{fux99} (a three dimensional N-body model based on COBE/DIRBE near infrared luminosities incorporating the $^{-1}$ arm and connecting arms), each associate no more than of the 3–kpc arm masers."
+ The details of all the models discussed above are sumanarlsed in Table 2. and their loci are shown in Figures 10. aud 11.., The details of all the models discussed above are summarised in Table \ref{modeltable} and their loci are shown in Figures \ref{modelcomparisonpart1} and \ref{modelcomparisonpart2}.
+ The models presentec| are as published. incorporating the authorW. Chosen best parameters (e.g. patter speeds. bar orientations. solar distauces) aud also iuclucdiug a F error margin in the loci of the structures.," The models presented are as published, incorporating the authors' chosen best parameters (e.g. pattern speeds, bar orientations, solar distances) and also including a $^{-1}$ error margin in the loci of the structures."
+ A direct comparison between the types of mioclels is complicated by the simple fact that the fur arm models cover more of theLe domai[un increasing their likelihood of associating sources.," A direct comparison between the types of models is complicated by the simple fact that the four arm models cover more of the domain, increasing their likelihood of associating sources."
+ However. within tve elliptical/lateral two arm models. that of ? provides the best fit to the sources.," However, within the elliptical/lateral two arm models, that of \citet{rodriguez08} provides the best fit to the sources."
+ This model has a large corotation 'aclius ol 5-7 kpe. as a result of a slow bar (where that bar is mace up of both a bulge aud ouger component).," This model has a large corotation radius of $-$ 7 kpc, as a result of a slow bar (where that bar is made up of both a bulge and longer component)."
+ The inclusion of a long bar component. clearly influential ou hieli-niiss star [ortjalion (Section ??)). may be responsible for this higher association.," The inclusion of a long bar component, clearly influential on high-mass star formation (Section \ref{barimp}) ), may be responsible for this higher association."
+ Within the iore c«xnplex imull-arln uiodels. the resonance features of 2? mprovide the higlest level of associatiou.," Within the more complex multi-arm models, the resonance features of \citet{mulder86} provide the highest level of association."
+ The ack of siguiicant 6.7-GHz methanol maser emission towards the . alcd connecting” arms in these 11oclels suggests these may be primarily gas features. uot tuπο igh-mass star [orration.," The lack of significant 6.7–GHz methanol maser emission towards the $^{-1}$ and `connecting' arms in these models suggests these may be primarily gas features, not undergoing high-mass star formation."
+ This Davy be due to the broad range of velocities. which botl features exhibit. incdicatine[n] highly tur»ulent material.," This may be due to the broad range of velocities, which both features exhibit, indicating highly turbulent material."
+" We similarly do uot see 6.7-GHz metlaned maser ejssion towards te broad raige of velocities of material towards the Galactic Centre (he metha1iol velocity range is ~200 1 compared with the CO and spanning ~500kkinss 1) and O elulsslOli Is see towards Balas ""Clump 2° (2).. a feature with a high range of velocitjes. but it js seen towards tle snaller velocity range of Bania's ""Clunp 1. (5).."," We similarly do not see 6.7–GHz methanol maser emission towards the broad range of velocities of material towards the Galactic Centre (the methanol velocity range is $\sim$ $^{-1}$ compared with the CO and spanning $\sim$ $^{-1}$ ) and no emission is seen towards Bania's `Clump 2' \citep{bania77}, a feature with a high range of velocities, but it is seen towards the smaller velocity range of Bania's `Clump 1' \citep{green10mmb2}."
+ In addition to the near aud far 3-kpc arm population we [iu the maserLe clistril demonstrates a pronuinent feature near —22* longitude which we i-—lerpret as the approximate tangent of the 3-kpe arms., In addition to the near and far 3–kpc arm population we find the maser distribution demonstrates a prominent feature near $-$ $^{\circ}$ longitude which we interpret as the approximate tangent of the 3–kpc arms.
+ Given tlie association of 6.7-GHz inetli10] imnasers with the paralle sectious of the 3-kpe arms and the taugent. we can expect the entire structure to exhibit met taser emissiou aud form au approximately oval structure in the ¢omain.," Given the association of 6.7-GHz methanol masers with the parallel sections of the 3–kpc arms and the tangent, we can expect the entire structure to exhibit methanol maser emission and form an approximately oval structure in the domain."
+ When this structure is combined with the spiral arms aud other features. essentially all the imaser emission is acc for.," When this structure is combined with the spiral arms and other features, essentially all the maser emission is accounted for."
+ This oval structure can eitlier relate physically to au elliptical rit£e ora circular ring with, This oval structure can either relate physically to an elliptical ring or a circular ring with radial
+The origin of high energy particles seen radiating m astroplvsical svnchrotron sources is süll a major unresolved problem of high-energv. and. plasma astrophysics.,The origin of high energy particles seen radiating in astrophysical synchrotron sources is still a major unresolved problem of high-energy and plasma astrophysics.
+ These sources include supernova remnants (SNRs) and extragalactic radio sources by jets etc. (e.g..Novamia1995;Meisenheimeretal.1989:CarilliandBarthel1996:Junor 1999).," These sources include supernova remnants (SNRs) and extragalactic radio sources by jets etc. \citep[e.g.,][]{Koy95,Mei89,Car96,Jun99}."
+. The shock acceleration has been discussed as one of important. processes producing the particles. and there are many theoretical ancl observational efforts on the hish-energy particle acceleration/heating so far.," The shock acceleration has been discussed as one of important processes producing the high-energy particles, and there are many theoretical and observational efforts on the high-energy particle acceleration/heating so far."
+ The dilbusive/Fermi acceleration process particularly met, The diffusive/Fermi acceleration process particularly met
+map (Fig. 14)),map (Fig. \ref{fig:2201}) )
+ shows a jet in the same position angle. /.e..e—40°.," shows a jet in the same position angle, $i.e., \sim -40\degr$."
+ Polarization was detected onlv in (he jet. with the inferred 5 field being perpendicular to the jet direction.," Polarization was detected only in the jet, with the inferred $B$ field being perpendicular to the jet direction."
+ The data for the LBLs. which belong primarily to the 1-Jy sample (2).. have been obtained from ?..," The data for the LBLs, which belong primarily to the 1-Jy sample \citep{KuehrSchmidt90}, have been obtained from \citet{Gabuzda00}."
+ Data lor four additional LBLs belonging to our ZZE0-1-RGD sample have been included in (he analvsis., Data for four additional LBLs belonging to our +RGB sample have been included in the analysis.
+ OF these. 1147+245 also belongs to the 1-Jv sample.," Of these, 1147+245 also belongs to the 1-Jy sample."
+ The data for the HWBLs comes primarily [rom our ZZEO-L-RGD sample., The data for the HBLs comes primarily from our +RGB sample.
+ The redshilt distribution Lor the LBL and HBL samples is presented in Fig. 15.., The redshift distribution for the LBL and HBL samples is presented in Fig. \ref{fig:redshift}.
+ The distributions of the core fractional polarization. the inner-jet (ry10 pe from the core) fractional polarization. the difference between the core polarization angle and the jet direction |.—0| and the difference between the inner jet polarization angle and the jet direction |y;—8| for LBLs and HDLs are presented in Figs. 17.. 18.. 19..," The distributions of the core fractional polarization, the inner-jet $r < 10$ pc from the core) fractional polarization, the difference between the core polarization angle and the jet direction $|\chi_c-\theta|$ and the difference between the inner jet polarization angle and the jet direction $|\chi_j-\theta|$ for LBLs and HBLs are presented in Figs. \ref{fig:m_c}, \ref{fig:chi_c}, \ref{fig:m_j},"
+ and 20.. respectively.," and \ref{fig:chi_j}, respectively."
+ Table G tabulates the values for (hese pauzuneters., Table \ref{tab:versus} tabulates the values for these parameters.
+ Note that multiple values lor a single source are included in the histograms if they. [all in different bins. but not if thev [all in the same bin.," Note that multiple values for a single source are included in the histograms if they fall in different bins, but not if they fall in the same bin."
+ The Sv-1 galaxy 2201044 was excluded from the analvsis., The Sy-1 galaxy 2201+044 was excluded from the analysis.
+ Table 8. summarises (he results of the comparison between LBLs and HIDLs., Table \ref{tab:KS} summarises the results of the comparison between LBLs and HBLs.
+ It has the following columus: Cl. (, It has the following columns: Col. (
+1) the property being compared. (2) (he corresponding (wo-siclecl ]xohnnogorovSmürnov (IX.5) statistic. which specifies here (he maximun deviation between,"1) the property being compared, (2) the corresponding two-sided Kolmogorov–Smirnov (K–S) statistic, which specifies here the maximum deviation between"
+o test this conclusion. but also to male iieasuremoents of the equation of state »uanneter e—p/p of this component.,"to test this conclusion, but also to make measurements of the equation of state parameter $w\equiv p/\rho$ of this component."
+ In principle the proposed SZE survey cun do just thist.. because the equation of state of the dark energv detcuines low its energy deusity evolves poxFfSATA whore Rois the scale factor.," In principle the proposed SZE survey can do just \cite{haiman00}, because the equation of state of the dark energy detemines how its energy density evolves $\rho_{Q}\propto 
+R^{-3(1+w)}$, where $R$ is the scale factor."
+ This evolution affects the expansion history of the universe. which is coupled o the volunc-redslift relation aud the erowth rate of deusity perturbations).," This evolution affects the expansion history of the universe, which is coupled to the volume-redshift relation and the growth rate of density \cite{peebles80}."
+ Fig 25 coutains a plot of the O3; aud c dependence of the vields of the SZE survey in the case that the mass limit is taken to be a constant AL..., Fig \ref{fig:paramdepend} contains a plot of the $\Omega_{M}$ and $w$ dependence of the yields of the SZE survey in the case that the mass limit is taken to be a constant $M_{lim}=2\times10^{14}h_{50}^{-1}M_{\odot}$ .
+ Note (lower left) that increasing c from -1 (the A case) to -0.2 ( with fixed Qa;= 0.3) decreases the number of cluster expected at itermeciate redshifts. but iucreases the nunuboer expected at high redshift.," Note (lower left) that increasing $w$ from -1 (the $\Lambda$ case) to -0.2 ( with fixed $\Omega_{M}=0.3$ ) decreases the number of cluster expected at intermediate redshifts, but increases the number expected at high redshift."
+ This is explained, This is explained
+both 2 ccontinuum emission and lline emission. while the 2.18 images (Lemaireetal.1996) were not corrected for instrumental scattered light.,"both 2 continuum emission and line emission, while the 2.18 images \citep{LFG96} were not corrected for instrumental scattered light."
+ We believe this has a minimal effect on our results. since the eemission Is so tightly confined to the filaments.," We believe this has a minimal effect on our results, since the emission is so tightly confined to the filaments."
+ Values of gereater than 4.8 mag are seen over an extended region outside the ffilaments., Values of greater than 4.8 mag are seen over an extended region outside the filaments.
+ Another issue is the ~3.3 ccontinuum emission that underlies the 3.294 HEF emission., Another issue is the $\sim$ 3.3 continuum emission that underlies the 3.29 IEF emission.
+ Without an {0 image of7023.. we do not know how much of the increase in un Figures 7 and 8 is due to a redder underlying continuum emission.," Without an $L'$ image of, we do not know how much of the increase in in Figures \ref{f7} and \ref{f8} is due to a redder underlying continuum emission."
+ In particular. we cannot derive a 3.29 IIEF feature-to-continuum ratio. which is the usual metric of the strength of the 3.29 HEF emission. Sellgrenetal.(1996).," In particular, we cannot derive a 3.29 IEF feature-to-continuum ratio, which is the usual metric of the strength of the 3.29 IEF emission. \citet{SWA96},"
+. however. presented measurements of the narrowband 3.29 ssurface brightness. in their Table 4 and Figure 5. which are not corrected for the underlying ~3.3 ccontinuum emission.," however, presented measurements of the narrowband 3.29 surface brightness, in their Table 4 and Figure 5, which are not corrected for the underlying $\sim$ 3.3 continuum emission."
+ These are the results with which we compare the data shown in Figure 8.., These are the results with which we compare the data shown in Figure \ref{f8}.
+ The positions measured in the reflection nebula sample of Sellgrenetal.(1996) were chosen by first searching for the brightest regions of K emission., The positions measured in the reflection nebula sample of \citet{SWA96} were chosen by first searching for the brightest regions of $K$ emission.
+ The K filter was chosen for initial mapping of each reflection nebula because the high thermal background at 3.29 mmakes K-band data vastly easier to obtain than 3.29 nnarrowband data., The $K$ filter was chosen for initial mapping of each reflection nebula because the high thermal background at 3.29 makes $K$ -band data vastly easier to obtain than 3.29 narrowband data.
+ Thus. the positions where Sellgrenetal.(1996) measured K—[3.29 μπα] are biased toward being near the 2 ccontinuum emission peak.," Thus, the positions where \citet{SWA96} measured $K-$ [3.29 ] are biased toward being near the 2 continuum emission peak."
+ Figure 7. shows that the 2 ccontinuum emission peak in hhas the smallest value of tthat we were able to measure., Figure \ref{f7} shows that the 2 continuum emission peak in has the smallest value of that we were able to measure.
+ It is intriguing that the highest 3.29 IIEF feature-to-continuum ratio measured by Sellgrenetal.(1996) is «9. corresponding to &—|3.29 jm]] = 4.5 mag.," It is intriguing that the highest 3.29 IEF feature-to-continuum ratio measured by \citet{SWA96}
+ is $\sim$ 9, corresponding to $K-$ [3.29 ] = 4.5 mag."
+ Since we measure == 5.5 mag in the ffilaments of7023.. we can speculate that the 3.29 IIEF feature-to-continuum ratio could be as high as ~23 if the underlying continuum color is unchanged.," Since we measure = 5.5 mag in the filaments of, we can speculate that the 3.29 IEF feature-to-continuum ratio could be as high as $\sim$ 23 if the underlying continuum color is unchanged."
+ This can only be confirmed by future observations that provide an estimate of the ~3.3 ccontinuum emission underlying the 3.29 HEF emission., This can only be confirmed by future observations that provide an estimate of the $\sim$ 3.3 continuum emission underlying the 3.29 IEF emission.
+ We note. however. that Figure 6 of Sellgrenetal.(1996) shows no correlation between K—L’ and 3.29 IIEF feature-to-continuum ratio.," We note, however, that Figure 6 of \citet{SWA96} shows no correlation between $K-L'$ and 3.29 IEF feature-to-continuum ratio."
+ Furthermore. support for a high feature-to-continuum ratio in the 3.29 ffilaments in iis provided by the spectra of Joblinetal.(1996).. who find 3.29 ffeature-to-continuum ratios of up to 20. 34. and 37 in the reflection nebula2023.. theBar.. and the reflection nebula surrounding NGC 1333/8VS-3 3)). respectively.," Furthermore, support for a high feature-to-continuum ratio in the 3.29 filaments in is provided by the spectra of \citet{Joblin96}, who find 3.29 feature-to-continuum ratios of up to 20, 34, and 37 in the reflection nebula, the, and the reflection nebula surrounding NGC 1333/SVS-3 ), respectively."
+ Figures 6.. 7.. and 8 may hold out the promise of making significant progress 1n understanding the emission mechanisms. and/or the relationship between the 2 ccontinuum emitters and the 3.29 IIEF carriers.," Figures \ref{f6}, \ref{f7}, and \ref{f8} may hold out the promise of making significant progress in understanding the emission mechanisms, and/or the relationship between the 2 continuum emitters and the 3.29 IEF carriers."
+ The difference in the spatial distributions of 2 ccontinuum emission and 3.29 HEF emission in7023.. with the 2 ccontinuum emission strongest in neutral gas and the 3.29 IIEF emission strongest in dense molecular filaments. is one essential clue.," The difference in the spatial distributions of 2 continuum emission and 3.29 IEF emission in, with the 2 continuum emission strongest in neutral gas and the 3.29 IEF emission strongest in dense molecular filaments, is one essential clue."
+ Even more intriguing. the smooth increase of the ccolor with increasing projected distance from the illuminating star of ppoints towards the stellar flux playing a crucial role in determining the relative strengths of the 2 ccontinuum emission and the 3.294 HEF emission.," Even more intriguing, the smooth increase of the color with increasing projected distance from the illuminating star of points towards the stellar flux playing a crucial role in determining the relative strengths of the 2 continuum emission and the 3.29 IEF emission."
+ , 
+Hippareos data in By. Vy and Hp bands (covering the time nterval 1990-1993) are used to adjust the period P and the epoch Το so as to increase the time span.,"Hipparcos data in $B_{\rm T}$ , $V_{\rm T}$ and $H_{\rm P}$ bands (covering the time interval $-$ 1993) are used to adjust the period $P$ and the epoch $T_{\rm 0}$ so as to increase the time span."
+ We find no evidence for a change in the orbital period and a formal upper limit for dP/dt of 0.0025 sec γι]., We find no evidence for a change in the orbital period and a formal upper limit for dP/dt of 0.0025 sec $^{-1}$.
+ We compute the solution using mode 2 of the Wilson-Devinney code. which ts appropriate for detached binaries with o constraints on the potentials.," We compute the solution using mode 2 of the Wilson-Devinney code, which is appropriate for detached binaries with no constraints on the potentials."
+ As the potentials of both stars are permitted to vary. the stellar sizes are constrained by the eclipses alone.," As the potentials of both stars are permitted to vary, the stellar sizes are constrained by the eclipses alone."
+ The relative luminosity of the secondary star is calculated from the other parameters (12 being coupled with the temperatures)., The relative luminosity of the secondary star is calculated from the other parameters $L_{2}$ being coupled with the temperatures).
+ We use a square root limb darkening law with the coefficients interpolated from van Hamme (1993) tables for the appropriate temperatures and surface gravities., We use a square root limb darkening law with the coefficients interpolated from van Hamme (1993) tables for the appropriate temperatures and surface gravities.
+ In doing this we assume a solar metallicity and this ts later and independently confirmed by the atmospheric analysis., In doing this we assume a solar metallicity and this is later and independently confirmed by the atmospheric analysis.
+ The bolometric albedo A and gravity brightening coefficients f are fixed to unity values appropriate for radiative atmospheres., The bolometric albedo $A$ and gravity brightening coefficients $\beta$ are fixed to unity values appropriate for radiative atmospheres.
+ V570 Per light and radial velocity curves do not show a clear signature of an eccentric orbit., V570 Per light and radial velocity curves do not show a clear signature of an eccentric orbit.
+ The circularity of the orbit is confirmed by initial modeling runs during which eccentricity is allowed to vary and nevertheless remains consistent with zero., The circularity of the orbit is confirmed by initial modeling runs during which eccentricity is allowed to vary and nevertheless remains consistent with zero.
+ After a few of such trials e Is set to zero., After a few of such trials $e$ is set to zero.
+ In the explicit standard use of the Wilson-Devinney code. the light curves are treated as flux values normalized to unity out of eclipses. separately for each photometric band.," In the explicit standard use of the Wilson-Devinney code, the light curves are treated as flux values normalized to unity out of eclipses, separately for each photometric band."
+ This causes a loss of information on the color temperature of the stars and therefore the necessity to assume a temperature for the primary (Τι)., This causes a loss of information on the color temperature of the stars and therefore the necessity to assume a temperature for the primary $T_{\rm 1}$ ).
+ However. the limb darkening coefficients do depend on the absolute temperature.," However, the limb darkening coefficients do depend on the absolute temperature."
+" The goodness of the choice for T, can be assessed by deriving a series of solutions where different values for T, are adopted and by selecting the one giving the smallest residuals.", The goodness of the choice for $T_{\rm 1}$ can be assessed by deriving a series of solutions where different values for $T_{\rm 1}$ are adopted and by selecting the one giving the smallest residuals.
+" This method confirms our choice for T, within its sensitivity of ~ 100 K. The results of the simultaneous photometric and spectroscopic solution and corresponding physical quantities are. given in Table 2. while the model radial velocity and light curves are over-plotted to the observational data in Figure 1."," This method confirms our choice for $T_{\rm 1}$ within its sensitivity of $\sim$ 100 K. The results of the simultaneous photometric and spectroscopic solution and corresponding physical quantities are given in Table 2, while the model radial velocity and light curves are over-plotted to the observational data in Figure 1."
+ The formal accuracy of the solution is on the masses and ~ on the radii for both components., The formal accuracy of the solution is on the masses and $\sim$ on the radii for both components.
+ The system parameters which are mainly dependent on the radial velocity curves (1.6. a. Mj. M». q. Vy) are very well constrained.," The system parameters which are mainly dependent on the radial velocity curves (i.e. $a$, $M_{\rm 1}$, $M_{\rm 2}$, $q$, $V_{\rm \gamma}$ ) are very well constrained."
+ The radii. which are dependent on the eclipse fitting. even if determined with a formal high accuracy. anyway suffer from the low inclination of the orbit and the consequent shallowness of the eclipses.," The radii, which are dependent on the eclipse fitting, even if determined with a formal high accuracy, anyway suffer from the low inclination of the orbit and the consequent shallowness of the eclipses."
+" No departure from spherical symmetry is found (R4 /R,,,, ~ 0.99 for both stars) and the system is detached (Αα ~ 0.17).", No departure from spherical symmetry is found $R_{pole}$ $R_{point}$ $\sim$ 0.99 for both stars) and the system is detached $R_{\rm 1}$ $a$ $\sim$ 0.17).
+" The orbital solution gives a difference in the temperatures of the two components of 290 K with à formal uncertainty of 15 K. However neither 7, nor 7» are known with this accuracy.", The orbital solution gives a difference in the temperatures of the two components of 290 K with a formal uncertainty of 15 K. However neither $T_{\rm 1}$ nor $T_{\rm 2}$ are known with this accuracy.
+ We remember that 7) 1s assumed to be 68505100 K. To check the accuracy of the derived orbital. solution. we compare the corresponding distance with the Hipparcos parallax.," We remember that $T_{\rm 1}$ is assumed to be $\pm$ 100 K. To check the accuracy of the derived orbital solution, we compare the corresponding distance with the Hipparcos parallax."
+ In caleulating the distance. we assume the usual relation RZAv/Eg 23.1.," In calculating the distance, we assume the usual relation $R$ $A_V$ $E_{B-V}$ =3.1."
+ Our determination of the distance. d212323 pe. is in very good agreement with the Hipparcos parallax 117155 pe. Lo limits).," Our determination of the distance, $\pm$ 3 pc, is in very good agreement with the Hipparcos parallax $^{104}_{132}$ pc, $\sigma$ limits)."
+ The spectral analysis of V570 Per allows to obtain an independent determination of the atmospheric parameters of both components: Tay. loge. Vie. sin and. most Important. the metallicity [Fe/H] and microturbulent velocity. which cannot be determined or guessed from the orbital analysis.," The spectral analysis of V570 Per allows to obtain an independent determination of the atmospheric parameters of both components: $T_{\rm eff}$, $\log g$, $V_{\rm rot}$ $\sin i$ and, most important, the metallicity [Fe/H] and microturbulent velocity, which cannot be determined or guessed from the orbital analysis."
+ The analysis is performed via y fitting procedure. using the large gric of synthetic Kuruez spectra of Munari et al. (," The analysis is performed via $\chi^2$ fitting procedure, using the large grid of synthetic Kurucz spectra of Munari et al. ("
+2004. in submission).,"2004, in submission)."
+ The spectra cover the 5500 wwavelength range with the same resolving power of the observed spectra (Rp = 200000)., The spectra cover the 500 wavelength range with the same resolving power of the observed spectra $R_{\rm P}~$ = 000).
+ The spectra are calculated with the revised solar abundances by Grevesse Sauval (1998) and the new opacity distribution functions (ODFs) of Castelli Kuruez (2004). throughout the whole HR diagram for 12 different rotational velocities. Τομ: ranging from 3500 to 5500 K. loge from 0.0 to 5.0 and |Fe/H] from -2.5 to «0.5.," The spectra are calculated with the revised solar abundances by Grevesse Sauval (1998) and the new opacity distribution functions (ODFs) of Castelli Kurucz (2004), throughout the whole HR diagram for 12 different rotational velocities, $T_{\rm eff}$ ranging from 3500 to 500 K, $\log g$ from 0.0 to 5.0 and [Fe/H] from $-$ 2.5 to +0.5."
+ The grid includes also a complete set of spectra calculated for a-enhanced chemical composition ({a/Fe]=+0.4) and for different values of micro-turbulent velocity (1. 2. 4 km κος).," The grid includes also a complete set of spectra calculated for $\alpha$ -enhanced chemical composition $\alpha$ /Fe]=+0.4) and for different values of micro-turbulent velocity (1, 2, 4 km $^{-1}$ )."
+ As discussed by Zwitter et al. (, As discussed by Zwitter et al. (
+2004). the grid steps are small enough to allow a safe interpolation between adjacent spectra.,"2004), the grid steps are small enough to allow a safe interpolation between adjacent spectra."
+ The y fitting procedure is performed on two high S/N spectra of V570 Per taken at different phases. near 0.34 and 0.66. when the difference in radial velocity between the components i$ nearmaximum. so that the lines of the two components are well separated.," The $\chi^2$ fitting procedure is performed on two high S/N spectra of V570 Per taken at different phases, near 0.34 and 0.66, when the difference in radial velocity between the components is nearmaximum, so that the lines of the two components are well separated."
+ Using spectra obtained at two different phases allows to account for the blending of different lines from the two components., Using spectra obtained at two different phases allows to account for the blending of different lines from the two components.
+ We concentrate our analysis on the same wavelength range 4890-5590 uused to derive the radial velocities., We concentrate our analysis on the same wavelength range 4890–5590 used to derive the radial velocities.
+ To remain independent, To remain independent
+(2008))).,).
+ LHIadronic models explain the huieh energv emission as an outcome of the svnchrotron proton emission ancl proton-photon interactions with the svnehrotron photons. (svnchrotron proton. blazar model) (Mannheim(1998):Mückeetal.(2003):Botteher.Reimer.&Marscher (2009111.," Hadronic models explain the high energy emission as an outcome of the synchrotron proton emission and proton-photon interactions with the synchrotron photons (synchrotron proton blazar model) \cite{mannheim98,mucke03,boettcher09}) )."
+ 3€2279 (2= 0.536) is à well studied source at various energv bands with several simultaneous multi wavelength. campaigns(Maraschictal.(1994):Hartmanοἱal.(1996):Woehrleet(1998):Botteher (2007))).," 279 $z=0.536$ ) is a well studied source at various energy bands with several simultaneous multi wavelength \cite{maraschi94,hartman96,wehrle98,boettcher07}) )."
+ lt was the first blazar observed at. gamma ray. energies by the satellite based experiment. (lartmanctal. (1992))) and the first’ ESO to be detected: at (ιοίΓον (very high energv (VIIE)) gamma rav energies by ground based atmospheric: Cherenkov experiment Nbort (2008)))., It was the first blazar observed at gamma ray energies by the satellite based experiment \cite{hartman92}) ) and the first FSRQ to be detected at GeV-TeV (very high energy (VHE)) gamma ray energies by ground based atmospheric Cherenkov experiment \cite{albert08}) ).
+ The thax of το is also known to be strongly variable at racio. infrared. optical. UV. X-ray and eamma-rav energies (Alakinoctal.(1989). and references therein).," The flux of 279 is also known to be strongly variable at radio, infrared, optical, UV, X-ray and gamma-ray energies \cite{makino89} and references therein)."
+ During 2006. 2279 undergone a dramatic [lare and was observed by WEBT (Whole Earth blazar Telescope) campaign (Bottcherctal. (2007))) at radio. near infrared and optical frequencies ancl simultancoush monitored. by(RANT) at X-ray energies (Chatterjeeetal. (2008))) and by telescope at VILE eammia ravs (Albertetal. (2008))).," During 2006, 279 undergone a dramatic flare and was observed by WEBT (Whole Earth blazar Telescope) campaign \cite{boettcher07}) ) at radio, near infrared and optical frequencies and simultaneously monitored by at X-ray energies \cite{chatterjee08}) ) and by telescope at VHE gamma rays \cite{albert08}) )."
+ Vhe SED of 2279 is modelled. as. synchrotron and inverse Compton emission. using these simultaneous observations at various energies., The SED of 279 is modelled as synchrotron and inverse Compton emission using these simultaneous observations at various energies.
+ The explanation of the Compton hump as an outcome of SSC process due to he cooling of a relativistic power-law electron distribution is not verv successful (Bottcher.Reimer.&Alarscher (2009)))., The explanation of the Compton hump as an outcome of SSC process due to the cooling of a relativistic power-law electron distribution is not very successful \cite{boettcher09}) ).
+ This interpretation had a serious drawbacks in explaining the high energy. ganuna ray. (7 E6/0V) detection rom 2279 since it require very high radiation. energy density compared. to magnetic energy. densitv(Maraschi.CGhisellini.&Celotti (1992)))., This interpretation had a serious drawbacks in explaining the high energy gamma ray $> 1GeV$ ) detection from 279 since it require very high radiation energy density compared to magnetic energy \cite{maraschi92}) ).
+ It. also. require unusually ow magnetic field to explain the VILE Botteher.temer.&Alarscher (2009)))., It also require unusually low magnetic field to explain the VHE \cite{boettcher09}) ).
+ However in case of BL Lac objects. the models considering the SSC process are. able ο. reproduce ganuna rav emission successfully (Stecker.Dhattacharvva.Sahavanathan.&Bhatt (2005))).," However in case of BL Lac objects, the models considering the SSC process are able to reproduce gamma ray emission successfully \cite{stecker96,coppi99,bhattacharyya05}) )."
+ Other alternative explanation for the Compton hump of 2279 is he EC mechanism with the target photons external to the jet., Other alternative explanation for the Compton hump of 279 is the EC mechanism with the target photons external to the jet.
+ Target photons from the accretion disk will be strongly deboosted due to high Lorentz factor of the jet (L10 25) and may not play an important role (however see Boettcher.Alause.&Sehlickeiser(1997) for an alternative)., Target photons from the accretion disk will be strongly deboosted due to high Lorentz factor of the jet $\Gamma\sim 10-25$ ) and may not play an important role (however see \cite{boettcher97} for an alternative).
+ The photons from BLE clouds cannot be boosted to VILI gamma ravs since the scattering of Lyman alpha line emission (the dominant emission from BLE clouds (Erancisetal.(1991) ))) will result. in Wlein-Nishina regime (CGhisellini&Tavecchio (2000):: section 83.2)) predicting à steep spectrum contrary to the observed hard. spectrum., The photons from BLR clouds cannot be boosted to VHE gamma rays since the scattering of Lyman alpha line emission (the dominant emission from BLR clouds \cite{francis91}) )) will result in Klein-Nishina regime \cite{ghisellini09}; section \ref{sec:ec}) ) predicting a steep spectrum contrary to the observed hard spectrum.
+ The third option. infrared photons from the dusty torus. can be a plausible candidate for the target. photons in EC process since the scattering of these photons to VILE gamma ray. energies. will still be in Thomson regime (CGhisellini&Tavecchio(2009):: section 83.2)).," The third option, infrared photons from the dusty torus, can be a plausible candidate for the target photons in EC process since the scattering of these photons to VHE gamma ray energies will still be in Thomson regime \cite{ghisellini09}; section \ref{sec:ec}) )."
+ The importance of Compton scattering of infrared photons from. dusty torus was first. studied. in detail by Blazejowskietal.(2000) to explain high energy gamma ravs detected by£CRET., The importance of Compton scattering of infrared photons from dusty torus was first studied in detail by \cite{blazejowski00} to explain high energy gamma rays detected by.
+ Vhough EC process is widelv accepted to explain the ganima ray emission of 2279. Lindfors.Valtaoja.Türler(2005) argues in favour of SSC mechanism based on observed radio-ganima rav. correlation ancl the quaclratic dependence of the svnchrotron and inverse Compton peak fluxes during the Iare.," Though EC process is widely accepted to explain the gamma ray emission of 279, \cite{lindfors05} argues in favour of SSC mechanism based on observed radio-gamma ray correlation and the quadratic dependence of the synchrotron and inverse Compton peak fluxes during the flare."
+ In the present paper we use the simultaneous observation of 2279 at radio. near infrared. optical. X-ray and very hieh cnerey (VILE) eamma ray energies 3ottcher.etal.(2007):Chatterjee(2008):Albertetal.(2008))) to deduce the physical parameters of the source.," In the present paper we use the simultaneous observation of 279 at radio, near infrared, optical, X-ray and very high energy (VHE) gamma ray energies \cite{boettcher07,chatterjee08,albert08}) ) to deduce the physical parameters of the source."
+ We consider svnchrotron. svnchrotron self Compton (SSC) and external Compton (EC) emission from a power-aw distribution of electrons.," We consider synchrotron, synchrotron self Compton (SSC) and external Compton (EC) emission from a power-law distribution of electrons."
+ We assume a magnetic field which is in equipartition with the particle energy., We assume a magnetic field which is in equipartition with the particle energy.
+ In the next section we outline the model and in section &3. we show the dausible target. photons for the inverse Compton process o explain the observed Bux at dillerent. energy. bands., In the next section we outline the model and in section \ref{sec:ic} we show the plausible target photons for the inverse Compton process to explain the observed flux at different energy bands.
+ La section Ed. we present the estimated. physical parameters of 2279 and discuss implications of the present model and he results., In section \ref{sec:result} we present the estimated physical parameters of 279 and discuss implications of the present model and the results.
+" A cosmology with (3,,=0.3. O4=0.7 and Ποτοkms+Alpe Djs used in this work."," A cosmology with $\Omega_m = 0.3$, $\Omega_\Lambda = 0.7$ and $H_0 = 70\;km\;s^{-1}\;Mpc^{-1}$ is used in this work."
+ We assume the emission region to be a spherical blob moving down the jet at relativistic speed with Lorentz factor E., We assume the emission region to be a spherical blob moving down the jet at relativistic speed with Lorentz factor $\Gamma$.
+ Since the jets of blazars are aligned. towards the line of sight. of the observer. we assume the Doppler factor 6=LV.," Since the jets of blazars are aligned towards the line of sight of the observer, we assume the Doppler factor $\delta \approx \Gamma$ ."
+ The radiation from emission region is due to svnchrotron and inverse Compton cooling of relativistic electrons. described by a broken power-law distribution (quantities with prime are measured at the rest frame of the emission region) where 5/ is the Lorentz factor of the electron. p ane q are the power-law indices before and after the break corresponding to the Lorentz factor 5; and A is the particle normalisation.," The radiation from emission region is due to synchrotron and inverse Compton cooling of relativistic electrons described by a broken power-law distribution (quantities with prime are measured at the rest frame of the emission region) where $\gamma^\prime$ is the Lorentz factor of the electron, $p$ and $q$ are the power-law indices before and after the break corresponding to the Lorentz factor $\gamma_b^\prime$ and $K$ is the particle normalisation."
+ In. the blob frame protons are assumed. to be cold. and contribute only to the inertia of the jet., In the blob frame protons are assumed to be cold and contribute only to the inertia of the jet.
+ The emission region is permeatecl with taneglecl magnetic field D., The emission region is permeated with tangled magnetic field $B$.
+" M we assume the equipartition condition between. the magnetic field energy. densitv(g) and the particle energy density. then the equipartition magnetic field D, will satisfy where m is the electron rest mass and e is the velocity of licht."," If we assume the equipartition condition between the magnetic field energy $U_B$ ) and the particle energy density, then the equipartition magnetic field $B_{eq}$ will satisfy where $m$ is the electron rest mass and $c$ is the velocity of light."
+ We neglect thecontribution of protons in equation (2)) due to the assumption they are cold and do not contribute to the particle energy. density., We neglect thecontribution of protons in equation \ref{eq:equipart}) ) due to the assumption they are cold and do not contribute to the particle energy density.
+ “Phe size of the emission region A can be approximated [rom the observed variability timescale £4; as where z is the redshift of the source., The size of the emission region $R^\prime$ can be approximated from the observed variability timescale $t_{var}$ as where $z$ is the redshift of the source.
+curve closely.,curve closely.
+ The quasar ος 279 is one of the most active blazars known., The quasar 3C 279 is one of the most active blazars known.
+ The ejection of 5-rav emitGng plasmoids from the nucleus of 3C 279 might be a frequent event., The ejection of $\gamma$ -ray emitting plasmoids from the nucleus of 3C 279 might be a frequent event.
+ We do therefore expect (to observe (he quasi-exponential decay phase following such ejection events onlv in a temporary quiescent phase. in which it is not overwhelmed by subsequent ejection events.," We do therefore expect to observe the quasi-exponential decay phase following such ejection events only in a temporary quiescent phase, in which it is not overwhelmed by subsequent ejection events."
+ This might be the reason why these quasi-exponential decavs in (he light. curves of this blazar are not more frequently observed., This might be the reason why these quasi-exponential decays in the light curves of this blazar are not more frequently observed.
+ In. addition to one or (wo occurrences cluring the 2006 WEBT campaign. several more quasi-exponentially decaving light curve segments can also be identified in the 2007 WEBT campaign data (Larionovetal.2008).. indicating (hat these are rare. but not unique events.," In addition to one or two occurrences during the 2006 WEBT campaign, several more quasi-exponentially decaying light curve segments can also be identified in the 2007 WEBT campaign data \citep{larionov08}, indicating that these are rare, but not unique events."
+ We thank the anonvamous referee for a very constructive report., We thank the anonymous referee for a very constructive report.
+ We thank Matthew G. Baring for stimulating discussions., We thank Matthew G. Baring for stimulating discussions.
+ This work was supported by NASA through. Guest Observer Grant NNNOSADGTG., This work was supported by NASA through XMM-Newton Guest Observer Grant NNX08AD67G.
+first X-ray detection of this source was with the κ{ΟΠ observatory (hence the catalogue name. 2E 2206.6|4517: llarris et aL.,"first X-ray detection of this source was with the $\textit{Einstein}$ observatory (hence the catalogue name, 2E 2206.6+4517; Harris et al.,"
+ 1993)., 1993).
+ Coaclelition of six IPC observations spanning 26.6 ksec vield a te detection with 115 net counts., Coaddition of six IPC observations spanning 26.6 ksec yield a $\sigma$ detection with 115 net counts.
+ The source is undetected in à 1.5 ksec LIAE observation., The source is undetected in a 1.5 ksec HRI observation.
+ No laring activity is evident during these observations., No flaring activity is evident during these observations.
+ We analvzed thirty nine observations from the ROSAT sublic data archive (16 PSPC: 23 HB) which ποιάος this SOULCC position within the field of view., We analyzed thirty nine observations from the ROSAT public data archive (16 PSPC; 23 HRI) which included this source position within the field of view.
+ The total observing ime was 250 hours., The total observing time was $\approx$ 50 hours.
+ The observations span two Lares with almost complete light curves. the tail end of a third. Dare. and show variability in the low-level N-rav. emission.," The observations span two flares with almost complete light curves, the tail end of a third flare, and show variability in the low-level X-ray emission."
+ The ]are detected with the PSPC has a longdecay time. possibly oxoviding evidence for significant continual heating during he flare decay.," The flare detected with the PSPC has a long decay time, possibly providing evidence for significant continual heating during the flare decay."
+ Ehus. a further in-depth study of [ares and quiescent emission. from 26 2206.6|4517. could. provide a useful test of current models.," Thus, a further in-depth study of flares and quiescent emission from 2E 2206.6+4517 could provide a useful test of current models."
+" detect the Xeray souree 2E 2206.6 14517in 16 "" ROSAT Position Sensitive rh 03/07/90 NOTA.UT and 23. observations with rh 03/07/90 NUTS. the instrument bandpass of 0.1 - 2.4 rh 03/07/90 SOT5.66 continuous X-ray. observations (217 τα 03/07/90 NUTO.NG with the PSPC between 18-22 June 1990 ra 0l uu in-orbit. calibration period). 30-31 rh 05/07/90 SOTT.O 29 May 1993 and 02 June."," We detect the X-ray source 2E 2206.6+4517 in 16 observations with the ROSAT Position Sensitive Proportional Counter (PSPC) and 23 observations with the HRI, within the instrument bandpass of 0.1 - 2.4 keV. Extensive and continuous X-ray observations $\approx$ 17 hrs) were made with the PSPC between 18-22 June 1990 (during the ROSAT in-orbit calibration period), 30-31 December 1991, and 29 May 1993 and 02 June."
+ The LIRR rhllo270 05/07 £90 AR. Lac began thirteen. days rh LIST6 10/06/92 of the PSPC observations., The HRI calibration observations of AR Lac began thirteen days after the completion of the PSPC observations.
+ Fourteen r 1202061 10/12 2-8 July 1990. include the source .," Fourteen HRI pointings between 2-8 July 1990, include the source 2E 2206.6+4517."
+" ee , additional IRI observations are ΕΟΟ Ίο 03/06 archive over the period of June 1992 ch202160 06/06 ου for a total observing time of 33 rh202161 09/07/96 50273.20 hours.", Eight additional HRI observations are available in the archive over the period of June 1992 through November 1996 for a total observing time of 33 hours.
+ WeC mures measuredCondcou rates using the ΗΛΙΟΙΟ data analysis software and Corrected. for vignetting. due to ⋅⋅⋠ ↿↓↕⋖⋅↓⋜⊔⋅⋏∙≟∢⊾↓⋅⋜⋯⋏∙≟∢⊾∪⇂∪∐−⋜∟∖↓⊳∖⋜⋯⋏∙≟↓⋖⊾⊳∖↿∖−≻∖∕↙⇥∖∕ 39)., We measured count rates using the IRAF/PROS data analysis software and corrected for vignetting due to the large range of off-axis angles $<$ $\theta$ $<$ ).
+ The observations were subdivided into multiple time bins to achieve à. higher temporal. resolution. while preserving à minimal 2e detection for cach bin., The observations were subdivided into multiple time bins to achieve a higher temporal resolution while preserving a minimal $\sigma$ detection for each bin.
+ An annular region was centered on the source to correct for the background. count rate for most cases., An annular region was centered on the source to correct for the background count rate for most cases.
+ For detections near the cdge of the field. a nearby circular background region was chosen.," For detections near the edge of the field, a nearby circular background region was chosen."
+ A log of the ROSAT observations is given in Table. 1. which includes the exposure time and oll-axis angle., A log of the ROSAT observations is given in Table \ref{tab_pspc} which includes the exposure time and off-axis angle.
+ Spectral fitting of the PSPC cata was done with the NSPEC software. package., Spectral fitting of the PSPC data was done with the XSPEC software package.
+ Source ancl background. counts were extractec using the NSELECT task from the FOOLS package. and the ancillary response files were construcκl with PCARE to account for oll-axis vignetting.," Source and background counts were extracted using the XSELECT task from the FTOOLS package, and the ancillary response files were constructed with PCARF to account for off-axis vignetting."
+ We ignored the lowest 11 spectral energy bins to avold scatterecl solar EUV contamination., We ignored the lowest 11 spectral energy bins to avoid scattered solar EUV contamination.
+ The highest 56 spectral channels were also omitted. due to poor statistics. a result. of a marked decrease in the effective area of the instrument at these energies.," The highest 56 spectral channels were also omitted due to poor statistics, a result of a marked decrease in the effective area of the instrument at these energies."
+ The energy bin distribution oversamples the intrinsic spectral resolution of the PSPC. so we grouped the jns by a factor of 5 to improve the statistics.," The energy bin distribution oversamples the intrinsic spectral resolution of the PSPC, so we grouped the bins by a factor of 5 to improve the statistics."
+ Alultiple optical spectra were taken by Perry. Berlind with the Tillinghast 607 telescope and FAST spectrograph (Fabricant ct al..," Multiple optical spectra were taken by Perry Berlind with the Tillinghast 60"" telescope and FAST spectrograph (Fabricant et al.,"
+ 1998) at the Fred. Lawrence. Whipple Observatory on Mount Hopkins., 1998) at the Fred Lawrence Whipple Observatory on Mount Hopkins.
+ X slit width of.. a 300 ines/nun grating. and the Loral CCD with 15/2. pixels oovided a resolution of z5 X.," A slit width of, a 300 lines/mm grating, and the Loral CCD with $\mu$ m pixels provided a resolution of $\approx 5$ $\rm{\AA}$."
+ X 3.5 minute exposure was aken on UT 16 May. 1998 ancl (wo ten minute exposures were acquired on UT 30 May 1998 anc UT 24 June 1998., A 3.5 minute exposure was taken on UT 16 May 1998 and two ten minute exposures were acquired on UT 30 May 1998 and UT 24 June 1998.
+ An observation of Feige 34 was used for extinction correction ancl [ux calibration., An observation of Feige 34 was used for extinction correction and flux calibration.
+ Standard bias subtraction. Iat-Fielding. the extraction of one-dimensional spectra. ancl wavelength calibration were performed in LAL.," Standard bias subtraction, flat-fielding, the extraction of one-dimensional spectra, and wavelength calibration were performed in IRAF."
+"emission region through the causality argument. D.«cAr~ em ~13r,.","emission region through the causality argument, $D<c\Delta t\sim 2.4\times
+10^{14} $ cm $\sim 13 r_g$."
+ The highly ionized absorber is observed to be outflowing at v;16.500 km s! during the average 2009 Long Look observation.," The highly ionized absorber is observed to be outflowing at $\upsilon_{out}\sim 16,500$ km $^{-1}$ during the average 2009 Long Look observation."
+ In the time-resolved spectral analysis of the same exposure. the outflown£& absorber is found to be variable over very short time scales (Fig. 7))," In the time-resolved spectral analysis of the same exposure, the outflowing absorber is found to be variable over very short time scales (Fig. \ref{FIG12}) )."
+ With the much lower S/N of each temporal slice. it was not possible to assess whether the absober varied in ionization state. column density. or blueshift.," With the much lower S/N of each temporal slice, it was not possible to assess whether the absober varied in ionization state, column density, or blueshift."
+ However. the very short time scale variability suggests that the absorber is very compact and very close to the X-ray source and that we are possibly observing rapid mass ejections from the inner regions of the accretion disk. e.g.. the base of an aceretion disk wind.," However, the very short time scale variability suggests that the absorber is very compact and very close to the X-ray source and that we are possibly observing rapid mass ejections from the inner regions of the accretion disk, e.g., the base of an accretion disk wind."
+ We note in particular how the absorption features do develop in à very short time scale. from the sixth to the seventh interval. together with the strong continuum flare.," We note in particular how the absorption features do develop in a very short time scale, from the sixth to the seventh interval, together with the strong continuum flare."
+" By making some assumptions. we can try to estimate the mass outflow rate M, associated to the highly ionized absorber."," By making some assumptions, we can try to estimate the mass outflow rate $\dot{M}_{out}$ associated to the highly ionized absorber."
+ In general. the mass outflow rate can be written às where A(7) Is a geometrical factor that accounts for how the flow diverges. and pC) and v0) are the density and velocity of the flow.," In general, the mass outflow rate can be written as where $A(r)$ is a geometrical factor that accounts for how the flow diverges, and $\rho(r)$ and $\upsilon_{out}(r)$ are the density and velocity of the flow."
+ By assuming a spherically symmetric. isotropic. steady flow with a constant velocity. one can write where my is the hydrogen atom mass. r the absorber distance from the central SMBH. 7s the absorber number density. Cj=dQ/4z is the solid angle occupied by the flow as seen by the central point source. and Fy the volume filling factor.," By assuming a spherically symmetric, isotropic, steady flow with a constant velocity, one can write where $m_{\rm H}$ is the hydrogen atom mass, $r$ the absorber distance from the central SMBH, $n$ the absorber number density, $C_f=d\Omega/4\pi$ is the solid angle occupied by the flow as seen by the central point source, and $F_{\rm V}$ the volume filling factor."
+ If we use appropriate values for PG 1126-041. we cal rearrangethe expression and write where ay)=1/10!! em™. ro=r/10ry.," If we use appropriate values for PG 1126-041, we can rearrangethe expression and write where $n_{11}=n/10^{11}$ $^{-3}$, $r_{10}=r/10r_g$."
+" In the expression above. the geometrical covering factor 0«C,<1 and volume filling factor Fy<| are unknown. and depend on the wind geometry and duty cycle."," In the expression above, the geometrical covering factor $0<C_f<1$ and volume filling factor $F_{\rm{V}}\leq 1$ are unknown, and depend on the wind geometry and duty cycle."
+ The density and the radial distice of the flow are also unknown., The density and the radial distance of the flow are also unknown.
+ Assuming that the outflowir£ absorber is a thin spherical shell of gas 1 photoronization. equilibrium. one can use the definition. of the ionizatior parameter €=Li/nr- to replace the unknown density and distance with the ratio £L;/£.," Assuming that the outflowing absorber is a thin spherical shell of gas in photoionization equilibrium, one can use the definition of the ionization parameter $\xi\equiv
+L_{ion}/n r^2$ to replace the unknown density and distance with the ratio $L_{ion}/\xi$."
+ However. in addition to the strongg assumptions made regard the spherical geometry of the outflow. there are uncertainties the ronizing continuum luminosity. and especially in the actual value of the ionization. parameter that may affect our mass outflow rate estimates.," However, in addition to the strong assumptions made regarding the spherical geometry of the outflow, there are uncertainties in the ionizing continuum luminosity, and especially in the actual value of the ionization parameter that may affect our mass outflow rate estimates."
+ In our spectral analysis results we quoted only the statistical errors on £: the systematic errors related to the input parameters used in the phototorization code in our case) can be much larger. especially for highly ionized absorbers where only a few spectral transitions are available (e.g.. Fe XXV. Fe XXVI.," In our spectral analysis results we quoted only the statistical errors on $\xi$: the systematic errors related to the input parameters used in the photoionization code in our case) can be much larger, especially for highly ionized absorbers where only a few spectral transitions are available (e.g., Fe XXV, Fe XXVI)."
+ The major dependence of the best-fit € value is related to the slope Γι adopted for the ionizing power-law continuum., The major dependence of the best-fit $\xi$ value is related to the slope $\Gamma_{ion}$ adopted for the ionizing power-law continuum.
+" This is because €e 71,/ng. where n, is the ionizing photons number density. my the photoionized gas"," This is because $\xi\propto n_{\gamma}/n_{\rm{H}}$ , where $n_{\gamma}$ is the ionizing photons number density, $n_{\rm{H}}$ the photoionized gas"
+between quadrants.,between quadrants.
+" Comparing the observed counts-in-cells distribution between the GQED and NBD, based on the least squares distance alone, fy(N) for 2D projected cells tend to follow the GQED while fy(N) for 3D spherical cells tend to follow the NBD."," Comparing the observed counts-in-cells distribution between the GQED and NBD, based on the least squares distance alone, $f_V(N)$ for 2D projected cells tend to follow the GQED while $f_V(N)$ for 3D spherical cells tend to follow the NBD."
+" However, in most cases, the GQED and NBD both fall within the measured quadrant to quadrant variations, and we note that the observed fy(NV) for individual quadrants within a sample may be closer to the GQED or NBD."," However, in most cases, the GQED and NBD both fall within the measured quadrant to quadrant variations, and we note that the observed $f_V(N)$ for individual quadrants within a sample may be closer to the GQED or NBD."
+ This shows that the SDSS catalog is unable to distinguish between the GQED and NBD because the difference between the GQED and NBD is smaller than the quadrant to quadrant variations., This shows that the SDSS catalog is unable to distinguish between the GQED and NBD because the difference between the GQED and NBD is smaller than the quadrant to quadrant variations.
+" Since the NBD has been shown to be unphysical (Saslaw; 1996),, on physical grounds we use the GQED for further analysis."," Since the NBD has been shown to be unphysical \citep{1996ApJ...460...16S}, on physical grounds we use the GQED for further analysis."
+" Comparing the values of N and b, between redshift ranges, we note that the values of N and correspondingly b are lower in the low redshift range than in the high redshift range for the same magnitude cutoff (samples 1b and 2b)."," Comparing the values of $\overline{N}$ and $b$, between redshift ranges, we note that the values of $\overline{N}$ and correspondingly $b$ are lower in the low redshift range than in the high redshift range for the same magnitude cutoff (samples 1b and 2b)."
+" This is despite the presence of the SDSS ""great wall"" in the low redshift range.", This is despite the presence of the SDSS “great wall” in the low redshift range.
+" On a closer look, we note that the “great wall” spans"," On a closer look, we note that the “great wall” spans"
+GRBs have been exposed (o be diverse both in their observed. prompt emission ancl allerglow emission.,GRBs have been exposed to be diverse both in their observed prompt emission and afterglow emission.
+ As a result. the physical interpretation has remained controversial.," As a result, the physical interpretation has remained controversial."
+ In this work. based on the X-ray data (Evansetal.2010) of 150 well-monitored GRBs with measured redshifts. we find. in the rest [rame. some underlving global features concealed in (he observed complex X-rav. light curves and the spectral evolution. which could clarify the diversity. issue.," In this work, based on the X-ray data \citep{evans10} of 150 well-monitored GRBs with measured redshifts, we find, in the rest frame, some underlying global features concealed in the observed complex X-ray light curves and the spectral evolution, which could clarify the diversity issue."
+ Generally speaking. (wo distinct aud consequent emission inereclients unequivocally exist in (he X-ray light curves as explicitly revealed by (he spectral differences shown in Figure 3.. which can be simply recognized as (he prompt emission wilh apparent spectral softening and the afterglow emission without apparent spectral evolution. respectively.," Generally speaking, two distinct and consequent emission ingredients unequivocally exist in the X-ray light curves as explicitly revealed by the spectral differences shown in Figure \ref{fig:restframelc2}, which can be simply recognized as the prompt emission with apparent spectral softening and the afterglow emission without apparent spectral evolution, respectively."
+ The most straightforward global features are the early plateau (ox/%! with ay0 on average) with a mild spectral softening (with a photon index increase AT~0.5 on average) and the late single power-law decline (xπι wilh a;—1.4 on average) without spectral variation (wilh a constant photon index P2 on average) up to /~I0* s. which are the previously known prompt enission and allerelow emission. respectively.," The most straightforward global features are the early plateau $\propto t^{\alpha_1}$ with $\alpha_1\sim 0$ on average) with a mild spectral softening (with a photon index increase $\Delta\Gamma\sim 0.5$ on average) and the late single power-law decline $\propto t^{\alpha_4}$ with $\alpha_4 \sim -1.4$ on average) without spectral variation (with a constant photon index $\Gamma\sim 2$ on average) up to $t\sim 10^{7}$ s, which are the previously known prompt emission and afterglow emission, respectively."
+" The transition rom the prompt emission to the alterglow emission undergoes two distinct. phases: a steep decline (x/? wilh à»~—2.6 on average)wilh a severe spectral softening (AT~1.0 on average) ancl a shallow decline (x/* with a,~—0.9 on average) wilh a mild spectral hardening to agree wilh the invariant spectra in the afterglow emission.", The transition from the prompt emission to the afterglow emission undergoes two distinct phases: a steep decline $\propto t^{\alpha_2}$ with $\alpha_2 \sim -2.6$ on average)with a severe spectral softening $\Delta\Gamma\sim 1.0$ on average) and a shallow decline $\propto t^{\alpha_3}$ with $\alpha_3 \sim -0.9$ on average) with a mild spectral hardening to agree with the invariant spectra in the afterglow emission.
+ The temporal feature is consistent with the canonical behavior previously proposed. (Zhangal.2006:Nouseket 2006).," The temporal feature is consistent with the canonical behavior previously proposed \citep{zhang06,nousek06}."
+. The spectral feature supports the speculation that the steep decline is connected with the prompt emission. while the shallow decline is more relevant to the afterglow emission.," The spectral feature supports the speculation that the steep decline is connected with the prompt emission, while the shallow decline is more relevant to the afterglow emission."
+ On the other hand. the erratic A-ray flares reportedly existing in half of the bursts are not self-explanatory in our analvsis.," On the other hand, the erratic X-ray flares reportedly existing in half of the bursts are not self-explanatory in our analysis."
+" Since (hey exhibit significant spectral softening and their peak luminosity follows a steep power-law decline Ly,x1?* statistically (Chincarinietal.2010).. thev should be connected with the steep declines. which also explains (heir analogy (o the prompt eniission."," Since they exhibit significant spectral softening and their peak luminosity follows a steep power-law decline $L_{\rm pk}\propto t^{-2.7}$ statistically \citep{chincarini10}, they should be connected with the steep declines, which also explains their analogy to the prompt emission."
+ Furthermore. as proposed earlier the spectral feature could be a crucial signature for diseriminating the prompt emission from the alterelow emission.," Furthermore, as proposed earlier the spectral feature could be a crucial signature for discriminating the prompt emission from the afterglow emission."
+ For (is reason. the emission detected by theFermi satellite which normally decays as a single power law with an index a~—1.4 and shows no significant spectral evolution could be more convincingly interpreted as (he emission of the external forward shock (e.g..Zou 2010)..," For this reason, the high-energy emission detected by the satellite which normally decays as a single power law with an index $\alpha \sim -1.4$ and shows no significant spectral evolution could be more convincingly interpreted as the emission of the external forward shock \citep[e.g.,][]{zou09,kumar09,gao09,ghisellini10}. ."
+ Interestingly. despite being composed of complex spikes. the superposition of the prompt," Interestingly, despite being composed of complex spikes, the superposition of the prompt"
+Bedke (1991).,Bedke (1994).
+ Although there are discrepancies between tle adopted classifications and those of de Vaucoucurs et al. (, Although there are discrepancies between the adopted classifications and those of de Vaucouleurs et al. (
+1991) for individual galaxies. particularly Virgo Cluster galaxies. the main conclusiois of this oper are not altered if these types are substituted (Ixoopuiaun Ixeunev L998).,"1991) for individual galaxies, particularly Virgo Cluster galaxies, the main conclusions of this paper are not altered if these types are substituted (Koopmann Kenney 1998b)."
+ Figure d presens the relationship othe central concentration parameter aud pe in the two environments., Figure 1 presents the relationship of the central concentration parameter and $\frac{F_{H\alpha}}{F_{R24}}$ in the two environments.
+ [ubde types are indicaed with different svinbols. separated by whole class. so that Sa stands for Sa-Sab eaaxies. Sc or Se-Sed galaxies. ete;," Hubble types are indicated with different symbols, separated by `whole' class, so that Sa stands for Sa-Sab galaxies, Sc for Sc-Scd galaxies, etc."
+ Lines euclose Sa aud Se galaxies., Lines enclose Sa and Sc galaxies.
+ Several peculiar eaaxies. for which the IIubble classification is obviously 1radequate (described further at the eud ¢X this section). are Xotted with separate sviubols.," Several peculiar galaxies, for which the Hubble classification is obviously inadequate (described further at the end of this section), are plotted with separate symbols."
+" For the isolates siuuple. the IIubbe type is well correlated with both C30 and E,Foi"," For the isolated sample, the Hubble type is well correlated with both C30 and $\frac{F_{H\alpha}}{F_{R24}}$."
+" Manu-Whitnuey tess show that differewes dn he distributions of C30 a1 pu"" for different IIubble types are significant at the level.", Mann-Whitney tests show that differences in the distributions of C30 and $\frac{F_{H\alpha}}{F_{R24}}$ for different Hubble types are significant at the level.
+" Furthermore. there is a eood correlation sotaveen, C30 and i (r=-0.7| Spearnan p—-0.79. significaice level 5 x  "," Furthermore, there is a good correlation between C30 and $\frac{F_{H\alpha}}{F_{R24}}$ (r=-0.74 Spearman $\rho$ =-0.79, significance level 5 x $^{-7}$ )."
+These results are consistent with those of Hashimoto et al. (, These results are consistent with those of Hashimoto et al. (
+908). who have shown that central concentration aud star formation activity are well correlate for their sample of 16377 field galaxies.,"1998), who have shown that central concentration and star formation activity are well correlated for their sample of 16377 field galaxies."
+ The xesenuce of a good correlation between these simple paramecrs shows that classification bv IIubble type is 1062ningful for most isolated σαlaxies., The presence of a good correlation between these simple parameters shows that classification by Hubble type is meaningful for most isolated galaxies.
+ Tn the Virgo Cluster. there is a mich weaser correlation between the two quantitative ouwaneters (r—-0.15. Spearman p=-0.17. significance 3.x? 1).," In the Virgo Cluster, there is a much weaker correlation between the two quantitative parameters (r=-0.45, Spearman $\rho$ =-0.47, significance 3 x $^{-4}$ )."
+ The weaker correlation is maar cause by a wider range in star formation activity as a function of concentration. with oth reduced and cnhanced levels of activity compared to the isolated sample.," The weaker correlation is primarily caused by a wider range in star formation activity as a function of concentration, with both reduced and enhanced levels of activity compared to the isolated sample."
+ While Virgo Sb and Se galaxies span simular rauges m C30 as isolated Sb aud Sc galaxies. the distribution for Virgo galaxies classified. as Sa is shifted to lower vaues than isolated ealaxies classified as Sa.," While Virgo Sb and Sc galaxies span similar ranges in C30 as isolated Sb and Sc galaxies, the distribution for Virgo galaxies classified as Sa is shifted to lower values than isolated galaxies classified as Sa."
+ The statistical difference in the distributions im the two environnients is at ιο level., The statistical difference in the distributions in the two environments is at the level.
+ This is caused by two effects: of the Virgo Sa sample have C30 iore tvpica of isolated Sb-Sc. and the distribution does not exteud to the hieh coicentrations seen iu isolated Sa galaxies.," This is caused by two effects: of the Virgo Sa sample have C30 more typical of isolated Sb-Sc, and the distribution does not extend to the high concentrations seen in isolated Sa galaxies."
+ Why have galaxies with concentrations more tvpical of isolated later-tvpe galaxies beeu classified as Sa in the Virgo Cluster?, Why have galaxies with concentrations more typical of isolated later-type galaxies been classified as Sa in the Virgo Cluster?
+ Virgo galaxies with looe DuFray fux less han -2.2 tend to be classified as Sa. regardless of their central conceitrations!," Virgo galaxies with log $\frac{F_{H\alpha}}{F_{R24}}$ flux less than -2.2 tend to be classified as Sa, regardless of their central concentrations!"
+ In fact. while Sa iux Sh galaxies in Vireo have statistically similar distributions of C30. they have siguificautlv different level) distributions in star formation activity.," In fact, while Sa and Sb galaxies in Virgo have statistically similar distributions of C30, they have significantly different level) distributions in star formation activity."
+" Πορσος, the classification of Virgo yalasrics as Sa ds based more on the weakness of star formation than on the underlying stellar light distribution."," Hence, the classification of Virgo galaxies as Sa is based more on the weakness of star formation than on the underlying stellar light distribution."
+ The morphology of these ealaxies is therefore not well described by the Iubble classification. since the stellar mass distribution aud star formation activity are not well correlated.," The morphology of these galaxies is therefore not well described by the Hubble classification, since the stellar mass distribution and star formation activity are not well correlated."
+ The completeness of our Vireo Sa suuple is ~ to Br < LEX hus we estimate that οδοί of briehtoO (D. < 13) VireooO oOgalaxies classified as Sa are acnally siuall to intermediate couceutration systems with reduced star formation., The completeness of our Virgo Sa sample is $\sim$ to $_T$ $<$ 13; thus we estimate that $\pm$ of bright (B $<$ 13) Virgo galaxies classified as Sa are actually small to intermediate concentration systems with reduced star formation.
+ The absence of high couceutration Sa galaxies is ikelv to be at least partially due to the incompleteness of our sample., The absence of high concentration Sa galaxies is likely to be at least partially due to the incompleteness of our sample.
+ Measurements of concentration for some of the Sa galaxies which are not in our sample are listed in other studies (Okara e al 19sLl: Doselli ct al., Measurements of concentration for some of the Sa galaxies which are not in our sample are listed in other studies (Okamura et al 1984; Boselli et al.
+ 1997): these uunubers jugeest that there are some higher conccutration galaxies classified as Sa in Vireo. but also additional low concentration galaxies classified as Sa.," 1997); these numbers suggest that there are some higher concentration galaxies classified as Sa in Virgo, but also additional low concentration galaxies classified as Sa."
+ Given the small number of Virgo SO ealaxies In our sample complete to D. « 13). direct conclusions about the distributiou of C30 for these galaxies with respect to isolated galaxies cannot be drawn.," Given the small number of Virgo S0 galaxies in our sample complete to B $<$ 13), direct conclusions about the distribution of C30 for these galaxies with respect to isolated galaxies cannot be drawn."
+ Tow does the star formation activity iu the Virgo Cluster compare to isolated galaxies as a fuuction of C30?, How does the star formation activity in the Virgo Cluster compare to isolated galaxies as a function of C30?
+ Iu the median. ReeHDI of Vireo Cluster spirals has Όσοι reduced by factors of," In the median, $\frac{F_{H\alpha}}{F_{R24}}$ of Virgo Cluster spirals has been reduced by factors of"
+"has adopted a source term, g(rgai) (the stellar mass loss rate per unit volume) of q(rgai)ος for Τραι€[2"",10""], for n=0, 2 and 3.","has adopted a source term, $q(r_{gal})$ (the stellar mass loss rate per unit volume) of $q(r_{gal}) \propto r_{gal}^{-\eta}$ for $r_{gal} \in [2^{\prime \prime}, 10^{\prime \prime}]$, for $\eta = 0$, 2 and 3."
+ He finds that the modelΤρι with 7=3 closely fits observed surface brightness profile., He finds that the model with $\eta = 3$ closely fits observed surface brightness profile.
+" From the definition of q(rgat), dM/dlnrgaixTAad(rgat)ox "," From the definition of $q(r_{gal})$, $d\dot{M}/dlnr_{gal} \propto r_{gal}^3q(r_{gal}) \propto r_{gal}^{3-\eta}$."
+It is evident from Fig., It is evident from Fig.
+" 9 that in therange of Τραι that reaQuataert considers (about 0.08 to 0.39pc), dM/dinrga is a power law."," \ref{fig:dmdot_dlnr_v_r} that in therange of $r_{gal}$ that \citet{quataert:2004} considers (about 0.08 to 0.39pc), $d\dot M/dlnr_{gal}$ is a power law."
+" Further, depending on which particular set of mass function parameters that we consider, η&2.6—2.8 for stellar collisions."," Further, depending on which particular set of mass function parameters that we consider, $\eta \cong 2.6 - 2.8$ for stellar collisions."
+ These values of 7 are very close to the model most consistent with the surface brightness profile., These values of $\eta$ are very close to the model most consistent with the surface brightness profile.
+" Using the the 2-10keV luminosity as measured by (Baganoffetal.2003),, we estimate the total mass loss rate at a radius of rgar~1.5"" (0.06pc)."," Using the the 2-10keV luminosity as measured by \citep{baganoff:2003}, we estimate the total mass loss rate at a radius of $r_{gal} \sim 1.5''$ (0.06pc)."
+" We use the word ""total"" to indicate the mass loss rate integrated over Galactic radius.", We use the word “total” to indicate the mass loss rate integrated over Galactic radius.
+" By using this total mass loss rate as an upper limit, we will be able to constrain the PDMF in the Galactic center by precluding any PDMFs with total mass loss rates greater than this value."," By using this total mass loss rate as an upper limit, we will be able to constrain the PDMF in the Galactic center by precluding any PDMFs with total mass loss rates greater than this value."
+" We will do this by integrating our calculated mass loss rate profiles (e.g., Figs."," We will do this by integrating our calculated mass loss rate profiles (e.g., Figs."
+ and 10)) over Τραι., \ref{fig:dmdot_dlnr_v_r} and \ref{fig:dir_indir}) ) over $r_{gal}$.
+" Unbound material at a radius rgo; has a dynamical timescale of where the characteristic velocity at radius gai, Uchar(Ygal), is taken as the velocity dispersion as given in 3.."," Unbound material at a radius $r_{gal}$ has a dynamical timescale of where the characteristic velocity at radius $r_{gal}$, $v_{char}(r_{gal})$, is taken as the velocity dispersion as given in \ref{sec:coll_rates}. ."
+ The electron density at radius Τραι may therefore be estimated by: where m; is the proton mass., The electron density at radius $r_{gal}$ may therefore be estimated by: where $m_p$ is the proton mass.
+The meelanis by which jets are lauuchecl was discussed in DOS: the following sumaumry 'elterates the salient poluts of the complex cyuamics by which unbound outflows emerge from he accretion low.,The mechanism by which jets are launched was discussed in D05; the following summary reiterates the salient points of the complex dynamics by which unbound outflows emerge from the accretion flow.
+ Iu the region of the accretion [low uea‘the mareinally stable orbit (the iuuer OLUS 'eelon i he terminology of DHIS). both pressure gradients aud the Lorentz force act to lift uaterial away f‘ol the equatorial plane.," In the region of the accretion flow near the marginally stable orbit (the inner torus region in the terminology of DHK), both pressure gradients and the Lorentz force act to lift material away from the equatorial plane."
+ Soue of this maerial is lauuched magueto-centrifugally li a 1naunner 'ereininisceut of the scenario of Blaudford. Payne (1982). geierating the coronal wind: some of tlis material. which has too 1MLC augular nonmentiunm to penerate the centrifugal oarrier. also becomes pzut of the massive ft11el-wall jet.," Some of this material is launched magneto-centrifugally in a manner reminiscent of the scenario of Blandford Payne (1982), generating the coronal wind; some of this material, which has too much angular momentum to penetrate the centrifugal barrier, also becomes part of the massive funnel-wall jet."
+ Tlere is also evideuce liat he low-augular noiuenutunm fuunel outlow originates cleeper il Ιe acc‘ello1 flow., There is also evidence that the low-angular momentum funnel outflow originates deeper in the accretion flow.
+ Sonie of this inae‘ial is produced in a eraviiohydromagetic (GHM) iperacioἩ in tlie ergosyhere (Puiy Coonii. 1990: Punsly. 2005). aud possibly i a process similar ο tha proposed by BlaneMord Zuajeς (1977) where conditious in the eremphere appro:ich the [orce-[ree limit.," Some of this material is produced in a gravitohydromagnetic (GHM) interaction in the ergosphere (Punsly Coroniti, 1990; Punsly, 2005), and possibly in a process similar to that proposed by Blandford Znajek (1977) where conditions in the ergosphere approach the force-free limit."
+ The maerial iu t1e [tunnel outflow. is accelerated by a relaively strong. »redoiinauly radial Lorentz force: gas pressure gradients iu the funnel do not conjbute siguilicautly.," The material in the funnel outflow is accelerated by a relatively strong, predominantly radial Lorentz force; gas pressure gradients in the funnel do not contribute significantly."
+ The ‘eis also a signilicaut amount of entrainment at the interface between the unuel wall aid corona., There is also a significant amount of entrainment at the interface between the funnel wall and corona.
+" 'Turuing now to the question of Lorentz factor. W. in the Πας] outflow. F""usgure | shows Wü0) for the Evy. Ry. and Spy models at the end of the simulations. as well as lime-averagec values. (M(r.t); for each model."," Turning now to the question of Lorentz factor, $W$, in the funnel outflow, Figure \ref{f1} shows $W(r,\theta)$ for the ${\rm E}_{\rm vf}$, ${\rm R}_{\rm vf}$, and ${\rm S}_{\rm vf}$ models at the end of the simulations, as well as time-averaged values, $\langle W(r,\theta) \rangle_t$ for each model."
+ Elevated values of the Lorentz [actor are foiud in compact. 101. evactated knots that ascend the Γιο] racially: the combination of low cdesity and higl eniperawe in the knots vields high specific enthalpy. h=14€cP/p..," Elevated values of the Lorentz factor are found in compact, hot, evacuated knots that ascend the funnel radially; the combination of low density and high temperature in the knots yields high specific enthalpy, $h=1+\epsilon+P/\rho$."
+" The Iighest. values o"" Lorentz actor reach the maxiuum allowed by the code (Mina,=00. aud test runs show that mucl ügher vaues could be reached with a higher ceiling [see Appendix]). aud these a'e only. found at arge raci. suggesting the presence of au extensive region where the knots are gradwally accelerate o higher Lorentz factors."," The highest values of Lorentz factor reach the maximum allowed by the code $W_{\rm max} = 50$, and test runs show that much higher values could be reached with a higher ceiling [see Appendix]), and these are only found at large radii, suggesting the presence of an extensive region where the knots are gradually accelerated to higher Lorentz factors."
+ The time-averaged plots coutain three interesting features: i) maxima at arge raclii. ij) asyiumetry. aid ili) spiu-depen«ent collimation.," The time-averaged plots contain three interesting features: i) maxima at large radii, ii) asymmetry, and iii) spin-dependent collimation."
+ i) Tine-averaging snoothes out the joassage of the knots aud reiuforces the notion of au exlensive acceleration Zoje: the Lorentz [actor increases smoothly [rom the base of the jet. near μις. OUlward to r/Azz 100—0600. where tle Inaxiia are found.," i) Time-averaging smoothes out the passage of the knots and reinforces the notion of an extensive acceleration zone; the Lorentz factor increases smoothly from the base of the jet, near $r_{\rm ms}$, outward to $r/M\approx$ 400—600, where the maxima are found."
+ ii) The asyuunetr vi1 the north aid sout1 components of the jets. which is seen both in the instantaneous aud tiuje-averaged. plots. cde;»euds on the details of injection at the base of the jet and on the time iuterval ‘or data collection (the dominance of one or the other funnel will reverse," ii) The asymmetry in the north and south components of the jets, which is seen both in the instantaneous and time-averaged plots, depends on the details of injection at the base of the jet and on the time interval for data collection (the dominance of one or the other funnel will reverse"
+"the line of sight, the distribution of magnetic perturbations in y-coordinate may be different from that we find for the Sun, if the magnetic features are confined to the equatorial zone.","the line of sight, the distribution of magnetic perturbations in $y$ -coordinate may be different from that we find for the Sun, if the magnetic features are confined to the equatorial zone."
+" The features would normally be seen closer to the side of the limb opposite to the visible pole, producing a skewed distribution of Ay depending on whether the spots or bright areas are the dominating contributors."," The features would normally be seen closer to the side of the limb opposite to the visible pole, producing a skewed distribution of $\Delta y$ depending on whether the spots or bright areas are the dominating contributors."
+" However, giant near-polar spots appear to be common on very active, fast-rotating stars."," However, giant near-polar spots appear to be common on very active, fast-rotating stars."
+" The research described in this paper was carried out at the Jet Propulsion Laboratory, California Institute of Technology, under a contract with the National Aeronautics and Space Administration."," The research described in this paper was carried out at the Jet Propulsion Laboratory, California Institute of Technology, under a contract with the National Aeronautics and Space Administration."
+" Observations at the 150-foot tower telescope on Mt. Wilson have been supported over the years by grants from the National Aeronautics and Space Administration, the National Science Foundation (NSF) and Office of Naval Research."," Observations at the 150-foot tower telescope on Mt. Wilson have been supported over the years by grants from the National Aeronautics and Space Administration, the National Science Foundation (NSF) and Office of Naval Research."
+ Research is currently supported by the NSF through grant AGS-0958779., Research is currently supported by the NSF through grant AGS-0958779.
+"Although most hieh-luninosity, extragalactic radio sources belonging to Fanaroff-Rilev class II (FRIT) have two reasonably svuuimetric lobes of radio ciissiol ou opposite sides of the parent optical object. the radio jets in these sources are overwhelminely onc-sided (cf.","Although most high-luminosity, extragalactic radio sources belonging to Fanaroff-Riley class II (FRII) have two reasonably symmetric lobes of radio emission on opposite sides of the parent optical object, the radio jets in these sources are overwhelmingly one-sided (cf."
+ Bridle Perley 198D., Bridle Perley 1984).
+ This is particularly true for quasars where radio jets are detected much more frequently than iu radio galaxies., This is particularly true for quasars where radio jets are detected much more frequently than in radio galaxies.
+ For example. iu a study of 12 extended 3CR quasars with the Verv Laree Array (VLA). Bridle et al (," For example, in a study of 12 extended 3CR quasars with the Very Large Array (VLA), Bridle et al. ("
+1991) made deep tuages with typical ruis. noises of 204/-Jv/bewun.,1994) made deep images with typical rms noises of $\mu$ Jy/beam.
+ They detected jets in all 12 sources and candidate counter-jets in 7., They detected jets in all 12 sources and candidate counter-jets in 7.
+ However. there is no unaunubieuous counter-jet in auv of these sources.," However, there is no unambiguous counter-jet in any of these sources."
+ In a simular study of à suuple of 13 radio galaxies bv Ferniui et al. (, In a similar study of a sample of 13 radio galaxies by Fernini et al. (
+1993. 1997). only two have radio jets aud neither of these show evidence of anv counter jet.,"1993, 1997), only two have radio jets and neither of these show evidence of any counter jet."
+ Although radio jets are not frequently detected in FRI radio galaxies. two of the good exaiuples of counter-jets seen iu FRII radio sources are associated with relatively nearby racio ealaxies.," Although radio jets are not frequently detected in FRII radio galaxies, two of the good examples of counter-jets seen in FRII radio sources are associated with relatively nearby radio galaxies."
+ One of these is he well-studied classical double-lobed source. Cyvguus. A. which is at a τουμα! of 0.056.," One of these is the well-studied classical double-lobed source, Cygnus A, which is at a redshift of 0.056."
+ A jet towards the north-west fom he uucleus aud a faint counter-jet towards the south-east can be seen on both parsec and kiloparsec scales aloug a position anele. PA=2a L+2” (Perley. Dreher Cowan 1981: Linfield 1985: Corili et al.," A jet towards the north-west from the nucleus and a faint counter-jet towards the south-east can be seen on both parsec and kiloparsec scales along a position angle, $=$ $\pm$ $^\circ$ (Perley, Dreher Cowan 1984; Linfield 1985; Carilli et al."
+ 1996: Ivichbamn. Alef Witzel 1996: Sorathia et al.," 1996; Krichbaum, Alef Witzel 1996; Sorathia et al."
+ 1996)., 1996).
+ Another FRIT radio galaxy where a well-collamated jet aux couuter-je have been detected is 30353. which is also at a relatively stall redshift of 0.03 (Swain. Bridle αι 1995).," Another FRII radio galaxy where a well-collimated jet and counter-jet have been detected is 3C353, which is also at a relatively small redshift of 0.03 (Swain, Bridle Baum 1998)."
+ Radio jets are believed to )o signatures of the beams carving energv from the nucleus to the outer lobes. and there have been a muuber of sugeestious to explain he apparent asvuuuectry of he jets despite the outer obes being reasonably svnunetric.," Radio jets are believed to be signatures of the beams carrying energy from the nucleus to the outer lobes, and there have been a number of suggestions to explain the apparent asymmetry of the jets despite the outer lobes being reasonably symmetric."
+ These iuchide the yossibilitics that (1) the jets are intrinsically sviuauictric mit appear one-sided or highly asviuuuetre because the xiehtuess of the receding oue is diminished due to bulk relativistic motion while the approaching jet is boosted for he same reason (Rees 1978: Blaudford Kounigl 1979: Scheuer Readhead 1979): (3) the active bea on the couuter-jet side has a low svuchrotron chussivity and Gi} the enerev supply ids iutriusicallv asviuuetrce. but supplies energy. alternately on opposite sides of the nucleus (cs," These include the possibilities that (i) the jets are intrinsically symmetric but appear one-sided or highly asymmetric because the brightness of the receding one is diminished due to bulk relativistic motion while the approaching jet is boosted for the same reason (Rees 1978; Blandford Könnigl 1979; Scheuer Readhead 1979); (ii) the active beam on the counter-jet side has a low synchrotron emissivity and (iii) the energy supply is intrinsically asymmetric, but supplies energy alternately on opposite sides of the nucleus (cf."
+ Saikia Wiita 1982: Ruduick Edgar 1981)., Saikia Wiita 1982; Rudnick Edgar 1984).
+ There appears to be a consensus of opinion that the observed asvnunuetry or onc-s;dedness of the jets is due to bulk relativistic motion., There appears to be a consensus of opinion that the observed asymmetry or one-sidedness of the jets is due to bulk relativistic motion.
+ Although early arguineuts relied ou the strong correlation of jet detection with core streugth (cf., Although early arguments relied on the strong correlation of jet detection with core strength (cf.
+ Saikia 1981) aud the parsec- and kiloparsec-scale jets beiug ou the same side (cf., Saikia 1984) and the parsec- and kiloparsec-scale jets being on the same side (cf.
+ Bridle Perley 1981). the strongest evidence comes from the correlation of depolarization asvuunetry with jet sidecuess (Laing 1988: Gairineton et al.," Bridle Perley 1984), the strongest evidence comes from the correlation of depolarization asymmetry with jet sidedness (Laing 1988; Garrington et al."
+ 1988: Carrineton. Conwav Loahv 1991).," 1988; Garrington, Conway Leahy 1991)."
+ This would also be consistent with jets )olug seen more frequently in quasars and also beige very asviuuetric and one-sided in these. since quasars are believed to be inclined at smaller angles to the liue of sieht than radio ealaxies (cf.," This would also be consistent with jets being seen more frequently in quasars and also being very asymmetric and one-sided in these, since quasars are believed to be inclined at smaller angles to the line of sight than radio galaxies (cf."
+ Darthel 1989)., Barthel 1989).
+ Iu this paper. usiugg observations made with miullaresec resolution. we report the detection of two-sided jets iu," In this paper, using observations made with milliarcsec resolution, we report the detection of two-sided jets in"
+"on March 9, 2009.","on March 9, 2009."
+" The R300V grating was used, centered at 540 nm and effectively covering from 430 to 810 nm at a resolving power R~1500."," The R300V grating was used, centered at 540 nm and effectively covering from 430 to 810 nm at a resolving power $\sim$ 1500."
+" The slit width was 1"", while the seeing was 1"".8."," The slit width was $''$, while the seeing was $''$ .8."
+ HD 192281 was chosen as the standard star to account for flux and sensitivity calibration., HD 192281 was chosen as the standard star to account for flux and sensitivity calibration.
+" A low-resolution (resolving power R ranging from 290 and 570), 40 min spectrum of M1-75 was taken with ACAM on the WHT on June 11, 2009, with the 400 lines mm! transmission VPH (Volume Phase Holographic) disperser, covering the wavelength range between 350 and 950 nm."," A low-resolution (resolving power R ranging from 290 and 570), 40 min spectrum of M1-75 was taken with ACAM on the WHT on June 11, 2009, with the 400 lines $^{-1}$ transmission VPH (Volume Phase Holographic) disperser, covering the wavelength range between 350 and 950 nm."
+" The 1"" wide slit was positioned at P.A.—0? in order to get the light from the two central star candidates (see section 4).", The $''$ wide slit was positioned at $^\circ$ in order to get the light from the two central star candidates (see section 4).
+" The seeing was 2"".8, and the standard star was HD 338808."," The seeing was $''$ .8, and the standard star was HD 338808."
+" The spectra were de-biased, flat-fielded, distortion-corrected and wavelength calibrated (from copper-argon and copper-neon arc lamps) using standard routines."," The spectra were de-biased, flat-fielded, distortion-corrected and wavelength calibrated (from copper-argon and copper-neon arc lamps) using standard routines."
+" After extraction of the selected nebular and central star features from the orthogonal 2-D spectra, the 1-D spectra were telluric and sensitivity corrected using the spectrum of the spectrophotometric standard star."," After extraction of the selected nebular and central star features from the orthogonal 2-D spectra, the 1-D spectra were telluric and sensitivity corrected using the spectrum of the spectrophotometric standard star."
+ The GHaFAS [Nm] 658.3 nm integrated image of M 1-75 is shown in Fig. 2.., The $\alpha$ FAS ] 658.3 nm integrated image of M 1–75 is shown in Fig. \ref{F2}.
+" While no central star (CSPN) is visible in this image, the nebula clearly shows two pairs of lobes with different orientations."," While no central star (CSPN) is visible in this image, the nebula clearly shows two pairs of lobes with different orientations."
+" They are nested in a central, brighter rim resembling a horseshoe."," They are nested in a central, brighter rim resembling a horseshoe."
+" Both systems of lobes appear distorted and fragmented, and their faint outer edges are difficult to track near the poles."," Both systems of lobes appear distorted and fragmented, and their faint outer edges are difficult to track near the poles."
+ The lobes of M 1-75 were the subject of a spatio-kinematical model by ?.., The lobes of M 1–75 were the subject of a spatio-kinematical model by \citet{dobrincic08}.
+" From two slit spectra, approximately along each pair of lobes, a [Nu] image from ?,, and simple assumptions such as ballistic expansion, they determined the larger and smaller lobes to lie at inclinations of 87° and 65°, respectively, and to be likely co-eval, with kinematical ages of 2700 and 2400 years per kpc of distance to the nebula, respectively."," From two slit spectra, approximately along each pair of lobes, a ] image from \citet{manchado96a}, and simple assumptions such as ballistic expansion, they determined the larger and smaller lobes to lie at inclinations of $^{\circ}$ and $^{\circ}$, respectively, and to be likely co-eval, with kinematical ages of 2700 and 2400 years per kpc of distance to the nebula, respectively."
+" Fabry-Perot interferometry (and GHaFAS, in particular) represents a significant step forward in spatio-kinematical modelling of planetary nebulae."," Fabry-Perot interferometry (and $\alpha$ FAS, in particular) represents a significant step forward in spatio-kinematical modelling of planetary nebulae."
+" Not only it allows for a resolution in wavelength comparable to high-resolution echelle spectrographs, but the series of ""Doppler-map"" images it produces span the whole nebula, instead of being limited by a narrow slit whose orientation has to be decided based on previous images."," Not only it allows for a resolution in wavelength comparable to high-resolution echelle spectrographs, but the series of “Doppler-map” images it produces span the whole nebula, instead of being limited by a narrow slit whose orientation has to be decided based on previous images."
+" From the spatio-kinematical point of view, a single GHaFAS data cube encompasses everything that is needed (i.e. information of the emission both in the plane of the sky and along the line of sight), its quality being only limited by seeing."," From the spatio-kinematical point of view, a single $\alpha$ FAS data cube encompasses everything that is needed (i.e. information of the emission both in the plane of the sky and along the line of sight), its quality being only limited by seeing."
+" In particular, it is noteworthy to remark the faint, high velocity emission regions offset from the axis of the larger lobes (see Fig. 3))"," In particular, it is noteworthy to remark the faint, high velocity emission regions offset from the axis of the larger lobes (see Fig. \ref{F3}) ),"
+ near the polar caps (especially in the southwest region)., near the polar caps (especially in the southwest region).
+" Those certainly would have remained unnoticed, had we been constrained by a narrow slit oriented along the lobes’ “expected” axis."," Those certainly would have remained unnoticed, had we been constrained by a narrow slit oriented along the lobes' “expected” axis."
+" The data cube, with Doppler-shift images spread across 48 channels, allowed us to build a spatio-kinematical model of both systems of lobes (see ? for a detailed description of the method)."," The data cube, with Doppler-shift images spread across 48 channels, allowed us to build a spatio-kinematical model of both systems of lobes (see \citealt{santander04b} for a detailed description of the method)."
+ Our first approach for each pair of lobes consisted of fitting a Solf-Ulrich (?) surface to the data., Our first approach for each pair of lobes consisted of fitting a Solf-Ulrich \citep{solf85} surface to the data.
+" This analytical model is described, in spherical coordinates, by: where r is the angular distance to the centre of the nebula (i.e. the adopted central star, see section 4), tD~! the"," This analytical model is described, in spherical coordinates, by: where $r$ is the angular distance to the centre of the nebula (i.e. the adopted central star, see section 4), $tD^{-1}$ the"
+ , 
+nass would still be =3 Mj.,mass would still be $\stackrel{<}{_\sim}3$ $M_{\rm Jup}$.
+ Our conclusions regarding the nature of this object could be tested by future observations al higher spectral resolution (~1000). with an instrument such as NIRSPEC (AleLeanοἱal.1998).. particularly if the wavelength range were extended to include the J-band.," Our conclusions regarding the nature of this object could be tested by future observations at higher spectral resolution $\sim1000$ ), with an instrument such as NIRSPEC \citep{mcl98}, particularly if the wavelength range were extended to include the $J$ -band."
+ Measurement of kev [αν ratios of H3O anc CII; features (see (2003))) would then facilitate more accurate spectral tvping. and the gravity estimate could be refined using the IX I doublet lines near 1.25 jan. Comparison of the resulting Thy and logg with the COND isochrones might then lead (o a better constraint on (he age which. in the present analvsis. has been assumed to be 1. Myr based on sky location in the p Oph cloud.," Measurement of key flux ratios of $_2$ O and $_4$ features (see \citet{mcl03}) ) would then facilitate more accurate spectral typing, and the gravity estimate could be refined using the K I doublet lines near 1.25 $\mu$ m. Comparison of the resulting $T_{\rm eff}$ and $\log g$ with the COND isochrones might then lead to a better constraint on the age which, in the present analysis, has been assumed to be 1 Myr based on sky location in the $\rho$ Oph cloud."
+ Our detection of this object suggests that the initial mass function for the p Oph star-Iorming region extends well into the planetarv-11nass regime., Our detection of this object suggests that the initial mass function for the $\rho$ Oph star-forming region extends well into the planetary-mass regime.
+ Object #44450 is. however. only one of about a dozen objects within our 1?x9’ region of study whose near-infrared SEDs are similarly indicative of planuetarv-mnass brown dwarls. and we plan to obtain near infrared spectra of these objects in the near future.," Object 4450 is, however, only one of about a dozen objects within our $1^\circ\times9'$ region of study whose near-infrared SEDs are similarly indicative of planetary-mass brown dwarfs, and we plan to obtain near infrared spectra of these objects in the near future."
+ We thereby hope to accumulate some statistics on the low end of the brown dwarl mass function to further our goal of constraining the formation models., We thereby hope to accumulate some statistics on the low end of the brown dwarf mass function to further our goal of constraining the formation models.
+ In this regard. we also look [orwarcl (ο the expected wealth of information from the Wide-field Infrared Survey. Explorer (WISE) mission. which will provide an inventory of low-mass cold brown dwarls in the immediate solar neighborhood.," In this regard, we also look forward to the expected wealth of information from the Wide-field Infrared Survey Explorer (WISE) mission, which will provide an inventory of low-mass cold brown dwarfs in the immediate solar neighborhood."
+ We thank the referee (A. Mueneh) for helpful comments., We thank the referee (A. Muench) for helpful comments.
+ The data presented herein were obtained at the W.M. heck Observatory. which is operated as a scientific partnership among (he California Institute of Technology. the University of California and the National Aeronautics and Space Administration.," The data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration."
+ The Observatory was made possible by the generous financial support of the W.M. Weck Foundation., The Observatory was made possible by the generous financial support of the W.M. Keck Foundation.
+ The authors wish to recognize and, The authors wish to recognize and
+"We also have first-integral forms of the proton-IHuid Euler equation(2):: and the cüllerence-Euler(6):: = AL | Cus where C5,1 and Ci, are the integration8 constants for the 1proton and. dillerence-Euler equations.i respectively.","We also have first-integral forms of the proton-fluid Euler equation: = and the difference-Euler: = M + C_d, where $C_\rmp$ and $C_d$ are the integration constants for the proton and difference-Euler equations, respectively."
+1 Finally. for the equation of state we choose an energy functional € given by k=(11) a Ovo-fluicl analogue of à polvtropie model (Prix.Comer&Andersson2002:Passamonti.HaskellAnclersson2009).," Finally, for the equation of state we choose an energy functional $\clE$ given by =, a two-fluid analogue of a polytropic model \citep{prix_ca,passa_sf}."
+. This could also be written in terms of the polvtropic indices Ny=Γον1). x=(n.pj.," This could also be written in terms of the polytropic indices $N_\rmx=1/(\gamma_\rmx-1)$, $\rmx=\{\rmn,\rmp\}$."
+ Note that can be generalised to include crosseterms. corresponding to svmnmetry energy. and entrainment.," Note that can be generalised to include cross-terms, corresponding to symmetry energy and entrainment."
+ The chemical potentials are then defined. from the energy functional byου... where the index x represents one particle species (neutrons or protons) ancl the v index represents the other.," The chemical potentials are then defined from the energy functional by, where the index x represents one particle species (neutrons or protons) and the y index represents the other."
+ Phe other equations of the svstem depend on the details of the magnetic-field configuration: we discuss these next., The other equations of the system depend on the details of the magnetic-field configuration; we discuss these next.
+ The simplest case is when the charged component of the stellar matter is governed by normal MID., The simplest case is when the charged component of the stellar matter is governed by normal MHD.
+ The magnetic force Fi is then the familiar Lorentz force. given hy bDB..(13) where D is the magnetic field andJ the current.," The magnetic force $\fmag$ is then the familiar Lorentz force, given by , where $\bB$ is the magnetic field and $\bj$ the current."
+ Most studies of NS equilibria have used the normal-MITD. equations and have additionally adopted a single-Luicl model. where the charged component is the entire [uid this ansatz has been used by innumerable authors.," Most studies of NS equilibria have used the normal-MHD equations and have additionally adopted a single-fluid model, where the charged component is the entire fluid — this ansatz has been used by innumerable authors."
+" Itecent examples include Tomimura&IZriguchi(2005):IxiuchiYoshida(2008):(2008):Lander&Jones(2009). and. Ciolfi.Ferrari&Gualtieri(2010).. but there are decades of earlier work on single-Iuid models see. οσοι, Ferraro(1954). and Roxbureh(1966)."," Recent examples include \citet{tomi_eri,kiuchi,haskell,1f_eqm} and \citet{ciolfi}, but there are decades of earlier work on single-fluid models — see, e.g., \citet{ferraro} and \citet{roxburgh}."
+.. An obvious way to progress is to work with a two-Llluic model. but still in normal MIID: the protons are a normal Luicl and he neutrons superlluid.," An obvious way to progress is to work with a two-fluid model, but still in normal MHD: the protons are a normal fluid and the neutrons superfluid."
+ This is not just a stepping stone towards the superlluid/superconducting model of it already allows us to study the role of stratification in magnetised stars., This is not just a stepping stone towards the superfluid/superconducting model of \citet{baym_pp} — it already allows us to study the role of stratification in magnetised stars.
+" Furthermore. the interior magnetic ield strength in magnetars may. well exceed the second critical field 44,5~1015 G at whieh superconductivity is destroved: in his case magnetar matter could indeed be predominantly normal protons and superlluid neutrons 2011).."," Furthermore, the interior magnetic field strength in magnetars may well exceed the second critical field $H_{c2}\sim 10^{16}$ G at which superconductivity is destroyed; in this case magnetar matter could indeed be predominantly normal protons and superfluid neutrons \citep{GAL}."
+ Since only the proton fluid. is magnetised. the derivations for all equations with magnetic terms proceed in the same manner as the single-I[uid case. with factors of p replaced by. pi: otherwise they are identical.," Since only the proton fluid is magnetised, the derivations for all equations with magnetic terms proceed in the same manner as the single-fluid case, with factors of $\rho$ replaced by $\rho_\rmp$; otherwise they are identical."
+ For this reason we do not give the details here. but the essentially identical single-Iuid derivation may be found in various papers (see. e.g. (2009))).," For this reason we do not give the details here, but the essentially identical single-fluid derivation may be found in various papers (see, e.g., \citet{1f_eqm}) )."
+ “Phe non-magnetic equations of the svstem are those discussed in the previous subsection., The non-magnetic equations of the system are those discussed in the previous subsection.
+ In the mixed-field (or purelypoloidal field) case we have a Poisson equation for the ó-component of the magnetic vector, In the mixed-field (or purelypoloidal field) case we have a Poisson equation for the $\phi$ -component of the magnetic vector
+instantaneously.,instantaneously.
+ The trend with Z and Uww is mainly caused by changes in the H2 and total hydrogen density distributions., The trend with $Z$ and $\UMW$ is mainly caused by changes in the $\H2$ and total hydrogen density distributions.
+" On the other hand, if n, is large (>10+ and the scatter is generated by the time averaging, then cm?)the scatter changes only weakly with Z and Uww (~0.09 dex versus ~0.13 dex)."," On the other hand, if $\nc$ is large $\gtrsim{}10^4$ $^{-3}$ ) and the scatter is generated by the time averaging, then the scatter changes only weakly with $Z$ and $\UMW$ $\sim{}0.09$ dex versus $\sim{}0.13$ dex)."
+" We note that in order to observe a significant change in the relation, metallicities <0.3Z and interstellar UV Uy,fields Uyw>10 are required."," We note that in order to observe a significant change in the relation, metallicities $\leq{}0.3\,Z_\odot$ and interstellar UV fields $\UMW\geq{}10$ are required."
+ Star-forming galaxies at high redshifts should therefore be the natural candidates to test our predictions., Star-forming galaxies at high redshifts should therefore be the natural candidates to test our predictions.
+ The authors thank A. Leroy and F. Bigiel for stimulating discussions and the anonymous referee for comments that helped improve the manuscript., The authors thank A. Leroy and F. Bigiel for stimulating discussions and the anonymous referee for comments that helped improve the manuscript.
+" The authors are grateful to the Aspen Center for Physics and to the participants and organizers of the workshop ""Star Formation in Galaxies: From Recipes to Real Physics""."," The authors are grateful to the Aspen Center for Physics and to the participants and organizers of the workshop ""Star Formation in Galaxies: From Recipes to Real Physics""."
+" N. Y. G and A. V. K. were supported by the NSF grants AST-0507596 and AST-0708154, and by the Kavli Institute for Cosmological Physics at the University of Chicago through the NSF grant PHY-0551142 and an endowment from the Kavli Foundation."," N. Y. G and A. V. K. were supported by the NSF grants AST-0507596 and AST-0708154, and by the Kavli Institute for Cosmological Physics at the University of Chicago through the NSF grant PHY-0551142 and an endowment from the Kavli Foundation."
+" The simulations used in this work have been performed on the Joint Fermilab - KICP Supercomputing Cluster, supported by grants from Fermilab, Kavli Institute for Cosmological Physics, and the University of Chicago."," The simulations used in this work have been performed on the Joint Fermilab - KICP Supercomputing Cluster, supported by grants from Fermilab, Kavli Institute for Cosmological Physics, and the University of Chicago."
+ This work made extensive use of the NASA Astrophysics Data System and arXiv.org preprint server., This work made extensive use of the NASA Astrophysics Data System and arXiv.org preprint server.
+ In order to test for possible resolution effects we have rerun one of our simulations (Z=1 and Uyw= 1) at two times better (32 pc) and also two times worse (125 pc) spatial resolution., In order to test for possible resolution effects we have rerun one of our simulations $Z=1$ and $\UMW=1$ ) at two times better (32 pc) and also two times worse (125 pc) spatial resolution.
+ We show in Fig., We show in Fig.
+" A6 the scale dependence of the scatter in the relation for the three different resolutions (32, 65, 125 "," \ref{fig:resTest} the scale dependence of the scatter in the relation for the three different resolutions (32, 65, 125 pc)."
+"As in Fig. 4,,"," As in Fig. \ref{fig:scaleScatter},"
+" we present both the scatter (1) due to SFRUy, time averaging (using a high nc) and (2) due to the non-linearpc). scaling between SFRs and H» density (using instantaneous SFR).", we present both the scatter (1) due to SFR time averaging (using a high $\nc$ ) and (2) due to the non-linear scaling between SFRs and $\H2$ density (using instantaneous SFR).
+ We find that the amount of scatter as a function of scale is similar in each of the three simulations and that there is no apparent significant systematic trend with resolution., We find that the amount of scatter as a function of scale is similar in each of the three simulations and that there is no apparent significant systematic trend with resolution.
+" As discussed in section 3.2,, the scatter in the relation decreases rather slowly (ος 179?) with averaging scale 1, much slower than expected from a V9? scaling, unless the gas is arranged in a 1D configuration."," As discussed in section \ref{sect:scale}, the scatter in the relation decreases rather slowly $\propto{}l^{-0.5}$ ) with averaging scale $l$, much slower than expected from a $V^{-0.5}$ scaling, unless the gas is arranged in a 1D configuration."
+" The scatter σι in the relation can be formally written as where p,=+[dxp,(x) is the spatial average of the star formation rate density and the brackets (...) denote the", The scatter $\sigma_l$ in the relation can be formally written as where $\overline{\dot{\rho}_*}=\frac{1}{V}\int{}dx\dot{\rho}_*(x)$ is the spatial average of the star formation rate density and the brackets $\langle \ldots \rangle$ denote the
+parütion functions currently do this.,partition functions currently do this.
+ Failing to estimate partition functions and strength [functions consistently can lead to inaccurate predictions of final equilibrium parameters., Failing to estimate partition functions and strength functions consistently can lead to inaccurate predictions of final ${\it equilibrium}$ parameters.
+ In equilibrium. the hard won rates no longer matter and only the partition functions govern (he final quantities of interest (17. abundances. etc).," In equilibrium, the hard won rates no longer matter and only the partition functions govern the final quantities of interest $Y_e$, abundances, etc.)."
+ llere we acldress the reliability of our caleulatecl rates., Here we address the reliability of our calculated rates.
+ We have shown in section 3 that with a simple prescription some gross features of the strength distribution. in particular the total strength and centroid of the distribution. may be estimated.," We have shown in section 3 that with a simple prescription some gross features of the strength distribution, in particular the total strength and centroid of the distribution, may be estimated."
+ However. shell model aud experimental results twpically show a rich structure in the strength distribution. with this structive varving markedly [rom nucleus to nucleus (e.g. Canrieretal.(1999))).," However, shell model and experimental results typically show a rich structure in the strength distribution, with this structure varying markedly from nucleus to nucleus (e.g. \cite{lan1}) )."
+ Do the rates depend sensilively on the finer features of the strength distribution?, Do the rates depend sensitively on the finer features of the strength distribution?
+ Or. alternatively. can a computationally simple method give a good estimate of the weak interaction rates?," Or, alternatively, can a computationally simple method give a good estimate of the weak interaction rates?"
+ These «questions can be addressed by comparing our relatively simply derived rates wilh rates based on more detailed large dimension shell model calculations., These questions can be addressed by comparing our relatively simply derived rates with rates based on more detailed large dimension shell model calculations.
+ In figures 7 and 8 we show the log of the ratio of our calculated beta decay rates to those calculated in LangankeanclMartinez-Pinedo(2000) αἱ several temperature/densitv. points relevant [οι stellar collapse. and [οι all nuclei in the range A=G0-65 for which LAIP caleulated rates.," In figures \ref{comp1} and \ref{comp2} we show the log of the ratio of our calculated beta decay rates to those calculated in \cite{lan2} at several temperature/density points relevant for stellar collapse, and for all nuclei in the range A=60-65 for which LMP calculated rates."
+ The horizontal axis in this figure is the log of the beta decay rate ealeulated by LMD., The horizontal axis in this figure is the log of the beta decay rate calculated by LMP.
+ We have presented (he comparison in this wav because nuclei with verv small rates will nol be so important in determining the evolution., We have presented the comparison in this way because nuclei with very small rates will not be so important in determining the evolution.
+ The comparison is generally remarkably favorable. will (vpical results dillering by less than a factor of two.," The comparison is generally remarkably favorable, with typical results differing by less than a factor of two."
+ Figures 9..10. and 11. are the same as the previous three figures. but now electron capture rates are being compared.," Figures \ref{comp4}, \ref{comp5} and \ref{comp6} are the same as the previous three figures, but now electron capture rates are being compared."
+ Again.5 the comparison is favorable. the rates as caleulated in the (wo schemes being5 within a [actor of three or so.," Again, the comparison is favorable, the rates as calculated in the two schemes being within a factor of three or so."
+ This is surprising given the potential sensitivity of electron capture rates lo the placement and the width of the Gamow-Teller distribution., This is surprising given the potential sensitivity of electron capture rates to the placement and the width of the Gamow-Teller distribution.
+ The fact that à simple prescription and rough estimates of the strength do a reasonably eood job in getting (he rates relevant lor stellar collapse is not to sav (hat all rates are accuratelv determined at all temperatures ancl densiües with a simple model., The fact that a simple prescription and rough estimates of the strength do a reasonably good job in getting the rates relevant for stellar collapse is not to say that all rates are accurately determined at all temperatures and densities with a simple model.
+ Some of our calculated rates deviate sienilicantlv [rom more detailed calculations., Some of our calculated rates deviate significantly from more detailed calculations.
+ This tvpically occurs lor o.(9)) decay when thermal population of the states in a collective Gamow-Teller resonance resulis in an exponential sensitivity. or for lepton capture when (he maximum lepton energies are just on the edee of being able to reach (he resonance in (he daughter.," This typically occurs for $\beta^-(\beta^+)$ decay when thermal population of the states in a collective Gamow-Teller resonance results in an exponential sensitivity, or for lepton capture when the maximum lepton energies are just on the edge of being able to reach the resonance in the daughter."
+ EnergyStercoscopicSystem...Wintou," \citep[see][for a review]{Gaensler06}. \citep[\textit{High Energy Stereoscopic System}, \citep{HESSscanII,Gallant07,Funk07a}."
+ s-ravs +-ravs. E>0.1 clectr," $\gamma$ $\gamma$ $E > 0.1$ \citep[e.g.,][]{deJager08}. \citep[bulk Lorentz factor $\Gamma_W \sim
+ 10^6$ in the Crab Nebula;][]{kc84a,kc84b},"
+"ous can produce pliotous in different cucreyoO. ranges:c» while een anion s photous with MA several Aiev. Hivor« omipton scatteringenergies of the wubicut photon field can produce hieh energy photons. up to tens of TeV, electrons allerespousibl fae theQDPYNA PWN«CGaeusler ceonussjon (here forafter 5-rav clectrous) are likely cexperiment.survey energetic than those generating the X-ray one (A-ray electrons). their syuchrotron radiation ling at infrared. optical. or ultraviolet frequencies."," electrons can produce photons in different energy ranges: while synchrotron radiation yields photons with energies up to several MeV, inverse Compton scattering of the ambient photon field can produce high energy photons, up to tens of TeV. The electrons responsible for the $\gamma$ -ray emission (here after $\gamma$ -ray electrons) are likely less energetic than those generating the X-ray one (X-ray electrons), their synchrotron radiation lying at infrared, optical, or ultraviolet frequencies."
+ For typical nebular magnetic field intensities (D-—1 100 pC). svuchrotron photous with euergv ~1 keV are produced bv electrous with Loreutz factor ~0.3 3. 4107.," For typical nebular magnetic field intensities $B \sim 1$ –100 $\mu$ G), synchrotron photons with energy $\sim$ 1 keV are produced by electrons with Lorentz factor $\sim$ 0.3–3 $\times 10^8$."
+" The Cosmic Dackgronud Radiation, the lieht. and the provide the target dust-scattoredphotons for inverse Compton starlightscattering. with Ypical plitou euereies around 10"" eV. 10 eV. and 1 eV. respectively,"," The Cosmic Background Radiation, the dust-scattered light, and the starlight provide the target photons for inverse Compton scattering, with typical photon energies around $10^{-3}$ eV, $10^{-2}$ eV, and 1 eV, respectively."
+" Iu the Phonon photons D PM n m 1. baeec ""memefera Wt novutz ""tor M um"," In the Thomson regime, photons with energy $\sim$ 1 TeV are produced by electrons with Lorentz factor $\sim$ 0.1–3 $\times
+10^7$."
+" ue then catterent energies, the coodug tue ot. he X-ravfo electrons ds μαμα than the oue of the Tiene t"," Due to their different energies, the cooling time of the X-ray electrons is smaller than the one of the $\gamma$ -ray electrons."
+" Sensthe fraces desMESnt MUNlustory of the nepua,."" DomWACLCAS the rm cunission traces a longer history, possibly up to the ul birth."," Therefore, the X-ray emission traces the recent history of the nebula, whereas the $\gamma$ -ray emission traces a longer history, possibly up to the pulsar birth."
+ different ΠοΠιο of the Οπου». ogether with the Theiuteraction with the iubieut mediuu. can lead to the significant projected angular separation. sometimes nieasured between the peaks of the 2- aud ray briehtuess profiles (e.g.Cs.00.7.Aharonianetal.2006d).," The different lifetime of the electrons, together with the interaction with the ambient medium, can lead to the significant projected angular separation, sometimes measured between the peaks of the $\gamma$ - and X-ray brightness profiles \citep[e.g. G18.0--0.7,][]{HESSJ1825b}."
+. Since the sourceο of injected --....electrons. the 5lua euerev ..loss rate dubbed spin-down bIuuinuosity. Cülsur- as time goes by. we expect a differcut evolution l the 5- and N-vay hiuinosities. following the particle are an theο... lur spi ONlowls," Since the source of injected electrons, the pulsar rotational energy loss rate dubbed spin-down luminosity, decreases as time goes by, we expect a different evolution of the $\gamma$ - and X-ray luminosities, following the particle aging and the pulsar spin-down."
+titutoE. this paper we address first the correlations betweenBassi P P, In this paper we address first the correlations between the
+titutoE. this paper we address first the correlations betweenBassi P PW, In this paper we address first the correlations between the
+titutoE. this paper we address first the correlations betweenBassi P PWN, In this paper we address first the correlations between the
+quantitatively: the shell width/leieth varies as the 1/3 power of the (small)ratio 2c;ος of trausvese/longitudinal 11ΟmMeLtum imparted to the shell by the jet. aud also decreases as the —2/3 power of Within a thousand. years. ile width/leugth ratio of the bow shock drops below the 1:10 rati typicaly associated with bipolar uolecular flows.,"quantitatively: the shell width/length varies as the $1/3$ power of the (small)ratio $\beta c_s/v_s$ of transverse/longitudinal momentum imparted to the shell by the jet, and also decreases as the $-2/3$ power of Within a thousand years, the width/length ratio of the bow shock drops below the 1:10 ratio typically associated with bipolar molecular flows."
+ This progressive elougation is strikinely illustrat in Fietre 1)). which portrays ile outflow shape at uiltigle scales: these correspond to successi factors of five increase in the age of tle jet.," This progressive elongation is strikingly illustrated in Figure \ref{fig:multscale}) ), which portrays the outflow shape at multiple scales; these correspond to successive factors of five increase in the age of the jet."
+ Ow couclusion tlal a singe leading bow shock is inable to produce large-scale molecUlal oulllows Is consistent with the finimes of previous autlors [rom oth ποιαμάνας (Masson1993) and numerical Sldies (see references in Leeetal.(2000b)))., Our conclusion that a single leading bow shock is unable to produce large-scale molecular outflows is consistent with the findings of previous authors from both semianalytic \citep{Masson1993} and numerical studies (see references in \cite{Lee2000b}) ).
+" Advocates of ""|χιye jet” models argue tha internal bow slocks or jet walerine can hep widen the bow-shock shells.", Advocates of “pure jet” models argue that internal bow shocks or jet wandering can help widen the bow-shock shells.
+ Because iuterual bow shocks prodice much less trausvese momentui flux than leading bow sli)cks (ων a facor cορ where A vds the velocity variatiou in the jet beam that leads to the internal shock) .lowevel. we do nol expec their contribution Oo provide a major boost to the trausvelse shell expansion. and this is indeec what our uuie‘ica simulatioiis show (Leeetal.(2000b):: see also Stone&Norma1(19932.1 3))).," Because internal bow shocks produce much less transverse momentum flux than leading bow shocks (by a factor $\sim \Delta v/v_j$, where $\Delta v$ is the velocity variation in the jet beam that leads to the internal shock), however, we do not expect their contribution to provide a major boost to the transverse shell expansion, and this is indeed what our numerical simulations show \cite{Lee2000b}; see also \cite{Stone1993a, Stone1993b}) )."
+ Jet precession. if st[liciently rapid to produce a smooth shell. is indistiigulshiable f|'om an intrinsically wide-angle wixl distribution.," Jet precession, if sufficiently rapid to produce a smooth shell, is indistinguishable from an intrinsically wide-angle wind distribution."
+ Thus. we contend that the relatively weak collijatiou of ‘Classical” bipolar ot(flows cannot be reconciled with the idea that an optical jet represents the whole angular extent of the primary wiud from a protostar: we 1iler that a least wide-augle wind iaist be present.," Thus, we contend that the relatively weak collimation of “classical” bipolar outflows cannot be reconciled with the idea that an optical jet represents the whole angular extent of the primary wind from a protostar; we infer that at least wide-angle wind must be present."
+" Athough jet bow shocks ¢annot explain everyhiug. they can explain molecular shell featlres associated with ""spim seen in posiion-velocity diag""ms of high. spatial resolution CO maps (Leeeal.2000a)."," Although jet bow shocks cannot explain everything, they can explain molecular shell features associated with “spurs” seen in position-velocity diagrams of high spatial resolution CO maps \citep{Lee2000a}."
+. The analysis of tlus paper [ocuses on leadiug bow shocks. because the initial conditlous are very well defiued.," The analysis of this paper focuses on leading bow shocks, because the initial conditions are very well defined."
+ Internal bow shocks car. however. be analyzed in a similar fashion.," Internal bow shocks can, however, be analyzed in a similar fashion."
+" Inapwe jet model. the cocoo1 gas enveloping the jet eam (but still interior to the leadiug JOW shock) will be at low density /ο. compared to the ambent ISM (it chiely cousists of shockec jel eas expauded in volume). aix will have significant fo""ward motion v (va. where vag is the velocity of the uext shock cowustream)."," In a pure jet model, the cocoon gas enveloping the jet beam (but still interior to the leading bow shock) will be at low density $\rho_e$ compared to the ambient ISM (it chiefly consists of shocked jet gas expanded in volume), and will have significant forward motion $v_e$ $\sim v_{s0}$, where $v_{s0}$ is the velocity of the next shock downstream)."
+ In a moclel wiere the jet is a dense core within a wiud. again the surrouiuxking envelope gas would |ave low density {N asstuuption) aud lewe Orwalc [velocity (e< ο).," In a model where the jet is a dense core within a wide-angle wind, again the surrounding envelope gas would have low density (by assumption) and large forward velocity $v_e\le v_j$ )."
+ TIe internal shock speed :Sp=Qa+v9)/2 is the mean value of I ipsitre:un (09) aud dowusream (04) jet speeds., The internal shock speed $v_{si}=(v_1+v_2)/2$ is the mean value of the upstream $v_2$ ) and downstream $v_1$ ) jet speeds.
+ Equatious (1))-(15)) wou call carry through. vvih he replacement of e;—vc;te. and p—pe.," Equations \ref{eq:jetM}) \ref{eq:uR}) ) would all carry through, with the replacement of $v_s\rightarrow v_{si} - v_e $, and $\rho\rightarrow \rho_e$."
+ The input mass aud trausverse ruonientum fluxes ake the same formi as vefore. except with p—p; aud e;—(es0)/2= Ac/2.," The input mass and transverse momentum fluxes take the same form as before, except with $\rho\rightarrow \rho_j$ and $v_s\rightarrow (v_2-v_1)/2=\Delta
+v/2$ ."
+" The shell shape still ha su xA"" belavior at large distance (with appropriately imocdilied coefficient). aud the »ositiou-velocity relation still approches the form «pyx264 fagain with modified coelficieut|."," The shell shape still has a $|z|\propto R^3$ behavior at large distance (with appropriately modified coefficient), and the position-velocity relation still approches the form $v_R\propto
+|z|^{-2/3}$ (again with modified coefficient)."
+square of the distance.,square of the distance.
+" The HESS-like sources are distributed in a thin disk, in a volume that is larger than the visibility limit."," The HESS-like sources are distributed in a thin disk, in a volume that is larger than the visibility limit."
+ The y-ray flux due to the HESS source population above 6 per cent of the Crab flux in the region —30?«I|30? and —2?«I2? is where $5=57.8x10?cm?s! is the maximum flux detected by HESS from a source.," The $\gamma$ -ray flux due to the HESS source population above 6 per cent of the Crab flux in the region $-{30}^{o}<l<{30}^{o}$ and $-{2}^{o}<l<{2}^{o}$ is where $S_2= 57.8 \times {10}^{-12} \, {{cm^2}} \, {s^{-1}}$ is the maximum flux detected by HESS from a source."
+" From the PSR and SNR distributions plotted in Fig.(2)) there are 84 SNRs and 5324 PRSs in the region —30?«|30° and —2?«b<2° , whereas in the Milagro region (40°<|100° and —5°<b« 5?) there are 36 SNRs and 1358 PSRs."," From the PSR and SNR distributions plotted in \ref{fig2}) ) there are 84 SNRs and 5324 PRSs in the region $-{30}^{o}<l<{30}^{o}$ and $-2^o<b<2^o$ , whereas in the Milagro region ${40}^{o}<l<{100}^{o}$ and $-5^o<b<5^o$ ) there are 36 SNRs and 1358 PSRs."
+" Assuming that supernova remnants and pulsars are the radio counterparts of high energy gamma ray sources, SNRs and PWNe, the Milagro region has 26 percent of the sources which are in the HESS region."," Assuming that supernova remnants and pulsars are the radio counterparts of high energy gamma ray sources, SNRs and PWNe, the Milagro region has 26 percent of the sources which are in the HESS region."
+" Also, the Milagro Galactic diffuse emission is measured for a threshold energy of about 3.5 TeV, whereas the flux from unresolved sources calculated in Eq.(8)) refers to the HESS threshold energy of 200 GeV. In order to estimate the contributionof unresolved"," Also, the Milagro Galactic diffuse emission is measured for a threshold energy of about 3.5 TeV, whereas the flux from unresolved sources calculated in \ref{eqnfluxtotal}) ) refers to the HESS threshold energy of 200 GeV. In order to estimate the contributionof unresolved"
+Although greatly outnumbered by spiral galaxies. elliptical galaxies contain ~20(Fukugita.Hogan. 1998) and dominate the high mass end of the galaxy mass function (e.g..Nakamura 2003).,"Although greatly outnumbered by spiral galaxies, elliptical galaxies contain $\sim 20$, 1998) and dominate the high mass end of the galaxy mass function (e.g., 2003)."
+ Their status at the top of the food chain is a strong motivation for their study às it provides insight in the processes governing the build-up of galaxies over cosmic time. the nature of the most luminous galaxies at z-3 and beyond. the formation of the first stars. and the origins of supermassive black holes.," Their status at the top of the food chain is a strong motivation for their study as it provides insight in the processes governing the build-up of galaxies over cosmic time, the nature of the most luminous galaxies at $z\sim 3$ and beyond, the formation of the first stars, and the origins of supermassive black holes."
+ Despite vigorous research efforts it is still not known how and when elliptical galaxies were formed., Despite vigorous research efforts it is still not known how and when elliptical galaxies were formed.
+ Separating the time of the formation of their from the time of theirassembly. a partial answer is that most of the stars in the most massive ellipticals were formed at high redshift.," Separating the time of the formation of their from the time of their, a partial answer is that most of the stars in the most massive ellipticals were formed at high redshift."
+" This result is supported by a large number of studies of nearby and distant galaxies. and applies to ellipticals in the general field as well as those in rich clusters (e.g..Bower.Lucey. 1992: 1997: 1998. 2003: 19983. 19050: 1999, 2002: 2001a: 2005)."," This result is supported by a large number of studies of nearby and distant galaxies, and applies to ellipticals in the general field as well as those in rich clusters (e.g., 1992; 1997; 1998, 2003; 1998a, 1998b; 1999, 2002; 2001a; 2005)."
+" Even when accounting for possible selection effects due to morphological evolution (""progenitor bias""; 2001). itis very difficult to fit the observed evolution of massive ellipticals with formation redshifts substantially lower than z2."," Even when accounting for possible selection effects due to morphological evolution (“progenitor bias”; 2001), it is very difficult to fit the observed evolution of massive ellipticals with formation redshifts substantially lower than $z\sim 2$."
+ Although there is general agreement that the mean age is high. there are strong indications that not all stars in all elliptical galaxies formed at high redshift: Trager et ((2000) and others (e.g.. 2005) find evidence for a “sprinkling” of recent star formation in otherwise old ellipticals. and recent studies of the evolution of the Fundamental Plane relation 1987) find evidence that low mass early- galaxies have lower luminosity-weighted ages than high mass ones (Treu et 22005: van der Wel et 22005).," Although there is general agreement that the mean age is high, there are strong indications that not all stars in all elliptical galaxies formed at high redshift: Trager et (2000) and others (e.g., 2005) find evidence for a “sprinkling” of recent star formation in otherwise old ellipticals, and recent studies of the evolution of the Fundamental Plane relation 1987) find evidence that low mass early-type galaxies have lower luminosity-weighted ages than high mass ones (Treu et 2005; van der Wel et 2005)."
+" Much less is known about the assembly of elliptical galaxies. that is. when and how the galaxies were ""put together""."," Much less is known about the assembly of elliptical galaxies, that is, when and how the galaxies were “put together”."
+" In the literature. a distinction is often made between models in which elliptical galaxies form from the collapse of primordial gas clouds (“monolithic collapse"": e.g..Eggen.Lynden-Bell. 1962: 1999). or in mergers of smaller galaxies (“hierarchically”: e.g. 1972:White 1991)."," In the literature, a distinction is often made between models in which elliptical galaxies form from the collapse of primordial gas clouds (“monolithic collapse”; e.g., 1962; 1999), or in mergers of smaller galaxies (“hierarchically”; e.g. 1972; 1991)."
+ Within the framework of a merger model the question of assembly can be specified as. first. the typical age of the last major merger: and second. the nature of the progenitor galaxies.," Within the framework of a merger model the question of assembly can be specified as, first, the typical age of the last major merger; and second, the nature of the progenitor galaxies."
+ In many individual cases we can readily answer both questions: abundant evidence leaves little doubt that some ellipticals formed recently through the merger of spiral galaxies. or will do so in the near future (e.g.. 1972:Schweizer 1982; 1996).," In many individual cases we can readily answer both questions: abundant evidence leaves little doubt that some ellipticals formed recently through the merger of spiral galaxies, or will do so in the near future (e.g., 1972; 1982; 1996)."
+ However. it is not clear whether such mergers are exceptional or responsible for the formation of the majority of ellipticals.," However, it is not clear whether such mergers are exceptional or responsible for the formation of the majority of ellipticals."
+ Among the observational argumentsfor recent mergers are the detection of “fine-structure” (e.g.. shells and ripples) in a large number of nearby ellipticals 1983: Schweizeret 11990): the presence of kinematically-decoupled cores 1988: 1988): small-scale dust seen in the majority of ellipticals (e.g.. 1995); (some) studies of the redshift evolution of galaxy morphologies. and of close pairs (e.g.. 2000: 2002: 2003): and strong evolution of the mass density of red galaxies to zoI. inferred from the COMBO-17 survey 2004).," Among the observational arguments recent mergers are the detection of “fine-structure” (e.g., shells and ripples) in a large number of nearby ellipticals 1983; Schweizeret 1990); the presence of kinematically-decoupled cores 1988; 1988); small-scale dust seen in the majority of ellipticals (e.g., 1995); (some) studies of the redshift evolution of galaxy morphologies, and of close pairs (e.g., 2000; 2002; 2003); and strong evolution of the mass density of red galaxies to $z\sim 1$, inferred from the COMBO-17 survey 2004)."
+The distribution of galaxy luminosities. Le. the galaxy lummosity function (LF). and its environmental dependence have often been used to provide strong constraints on theories of galaxy evolution ???).,"The distribution of galaxy luminosities, i.e. the galaxy luminosity function (LF), and its environmental dependence have often been used to provide strong constraints on theories of galaxy evolution ."
+. Galaxy environment can be important in shaping several galaxy properties. such as colors. morphologies. and star-formation rates reviews).," Galaxy environment can be important in shaping several galaxy properties, such as colors, morphologies, and star-formation rates ."
+. Since galaxy SERs are strictly related to their total infrared (IR) emission (?).. powerful constraints on how galaxies evolve in relation to their environment are expected to be obtained from the analysis of the galaxy IR ΤΕ».," Since galaxy SFRs are strictly related to their total infrared (IR) emission , powerful constraints on how galaxies evolve in relation to their environment are expected to be obtained from the analysis of the galaxy IR LFs."
+ Following early studies of the IR. properties of cluster galaxies with/RAS and therein).. the (2).. theAKART satellite (?).. and now the (?).. have only recently allowed a precise derivation of galaxy IR LFs at various redshifts and in various environments in a precise way.," Following early studies of the IR properties of cluster galaxies with and , the , the satellite , and now the , have only recently allowed a precise derivation of galaxy IR LFs at various redshifts and in various environments in a precise way."
+ Most determinations of galaxy IR LFs in cluster and supercluster environments have so far been based on data., Most determinations of galaxy IR LFs in cluster and supercluster environments have so far been based on data.
+ analyzed the IR LFs of the rich nearby clusters Coma and A3266., analyzed the IR LFs of the rich nearby clusters Coma and A3266.
+ According to their analysis. the bright end of the IR LF has a universal form for local rich clusters. and cluster and field IR LFs have similar values of their characteristic luminosities. Li.," According to their analysis, the bright end of the IR LF has a universal form for local rich clusters, and cluster and field IR LFs have similar values of their characteristic luminosities, $\rm{L}_{\rm IR}^{\star}$."
+ Nearby rich clusters have a lower star-forming galaxy fraction than field galaxies. although this fraction increases with cluster-centric distance.," Nearby rich clusters have a lower star-forming galaxy fraction than field galaxies, although this fraction increases with cluster-centric distance."
+ These results were confirmed by?.. who analyzed a larger sample of clusters in the redshift range0.," These results were confirmed by, who analyzed a larger sample of clusters in the redshift range."
+9.. They noted the similarity in the shape of the cluster and field IR LFs. and confirmed that there is an increase in the fraction of luminous IR galaxies (LIRGs!)) with cluster-centric distance. out to 1.5 virialradii.. where it is still below the field value.," They noted the similarity in the shape of the cluster and field IR LFs, and confirmed that there is an increase in the fraction of luminous IR galaxies ) with cluster-centric distance, out to 1.5 virial, where it is still below the field value."
+ also noted the absence of LIRGs from the central regions of a sample of 32 X-ray selected clusters., also noted the absence of LIRGs from the central regions of a sample of 32 X-ray selected clusters.
+ RecentHerschel observations provide evidence of a lack of IR galaxies not only at the bright end but also the faint end of the IR LF of the nearby Virgo cluster relative to the field(?)., Recent observations provide evidence of a lack of IR galaxies not only at the bright end but also the faint end of the IR LF of the nearby Virgo cluster relative to the field.
+. The environmental dependence of the fraction of high-SFR galaxies may not be a simple function of cluster-centric distance., The environmental dependence of the fraction of high-SFR galaxies may not be a simple function of cluster-centric distance.
+ detected a large-scale in the IR. connecting a rich and a poor cluster at z-0.2.," detected a large-scale in the IR, connecting a rich and a poor cluster at $\rm{z} \sim 0.2$."
+ They observed that the fraction of high-SFR galaxies is largest in the filament. i.e. larger than in the cluster core. but also larger than in other. lower density. regions of the supercluster.," They observed that the fraction of high-SFR galaxies is largest in the filament, i.e. larger than in the cluster core, but also larger than in other, lower density, regions of the supercluster."
+ The filament detected by was the first to be found via IR. observations. (with Spitzery; observations have recently revealed other large-scale structure filaments traced by IR-emitting galaxies(??).. but the analysis of their galaxy populations is still ongoing.," The filament detected by was the first to be found via IR observations (with ); observations have recently revealed other large-scale structure filaments traced by IR-emitting galaxies, but the analysis of their galaxy populations is still ongoing."
+ observed that the medium- and low-density regions of another (more distant. z~ 0.8) supercluster host comparable fractions of star-forming galaxies. while red mid-IR emitters are preferentially located in medium-density environments. such as galaxy filaments.," observed that the medium- and low-density regions of another (more distant, $\rm{z} \sim 0.8$ ) supercluster host comparable fractions of star-forming galaxies, while red mid-IR emitters are preferentially located in medium-density environments, such as galaxy filaments."
+ Both and argued that star-formation ts triggered in galaxies in the infall regions around clusters., Both and argued that star-formation is triggered in galaxies in the infall regions around clusters.
+ came to the same conclusion after analyzing the IR galaxy population in à z=0.165 supercluster., came to the same conclusion after analyzing the IR galaxy population in a $\rm{z}=0.165$ supercluster.
+" They also found that while the IR galaxies prefer to live in medium-density environments. their SFRs are not particularly high for their stellar masses )). Le. they have normal specific SFRs (SSFR= SFR/M,)."," They also found that while the IR galaxies prefer to live in medium-density environments, their SFRs are not particularly high for their stellar masses ), i.e. they have normal specific SFRs $\mbox{sSFR} \equiv \rm{SFR}/\rm{M}_{\star}$ )."
+"al 22 Gllz (1.3 em) if the electron clensity is 4xLO? Ὁ, ",at 22 GHz (1.3 cm) if the electron density is $4\times 10^5$ $^{-3}$.
+This means that the formulae lor oplically thin eniission (e.g.Mezger&Henderson100) are not useful for the II] regions in our sample. aud probably not for small. high-density. HIT regions in general.," This means that the formulae for optically thin emission \citep[e.g.][]{MezgerHenderson1967} are not useful for the HII regions in our sample, and probably not for small, high-density HII regions in general."
+ If the gas is optically Chick to Iree-free emission. and the electron density is estimated from the radio continuum by incorrectly. assuming optically thin emission. then the derived electron densitv is just (hat density. required to make the optical depth unity.," If the gas is optically thick to free-free emission, and the electron density is estimated from the radio continuum by incorrectly assuming optically thin emission, then the derived electron density is just that density required to make the optical depth unity."
+ This can be understood by a simple example., This can be understood by a simple example.
+" The radio Iree-Iree emission through a slab of unilorm gas density and temperature is. The optically thin approximation amounts to replacing |l—exp(—7,)] by 7, so that. whereas. in the optically thick limit. Thus if the HI] region is optically Chick (equation 7)) but 7, is derived assuming equation 6.. then (he derived value of (he optical dept is 1. and the electron density derived from equation 4 is the density required to make z,=1."," The radio free-free emission through a slab of uniform gas density and temperature is, The optically thin approximation amounts to replacing $[1-\exp(-\tau_\nu)]$ by $\tau_\nu$ so that, whereas, in the optically thick limit, Thus if the HII region is optically thick (equation \ref{eq:thickEmission}) ) but $\tau_\nu$ is derived assuming equation \ref{eq:thinEmission}, then the derived value of the optical depth is 1, and the electron density derived from equation \ref{eq:opticalDepth}
+ is the density required to make $\tau_\nu=1$."
+ Thus at 22 Gllz (1.3 em). the electron density derived for an IIII region with a size of 0.01 pe is never be greater than about 10 C if the optically thin formulae are used. — reearclless of the true density.," Thus at 22 GHz (1.3 cm), the electron density derived for an HII region with a size of 0.01 pc is never be greater than about $10^5$ $^{-3}$ if the optically thin formulae are used – regardless of the true density."
+ This uiderestimate also allects quantities derived [rom (he electron density such as the number of ionizing photons required to balance recombinations which in (urn is often used (o derive the spectral tvpes ol the exciting stars., This underestimate also affects quantities derived from the electron density such as the number of ionizing photons required to balance recombinations which in turn is often used to derive the spectral types of the exciting stars.
+ Formulas for spherical WI reeions that have a density gradient such Chat the central part of the II] region is optically thick are derived in Keto(2003)., Formulas for spherical HII regions that have a density gradient such that the central part of the HII region is optically thick are derived in \citet{Keto2003}.
+. Llowever. these formulae require a specilication of the size of the optically thick region.," However, these formulae require a specification of the size of the optically thick region."
+ As shown in table 4. the line widths decrease with increasing [frequency in all cases.," As shown in table \ref{CMrrl}, the line widths decrease with increasing frequency in all cases."
+ This suggests the increasing importance of pressure broadening al lower frequencies., This suggests the increasing importance of pressure broadening at lower frequencies.
+ For sources with more than one centimeter RRL observed. the electron densities increase with (he increasing frequency of the centimeter line used in (he analvsis.," For sources with more than one centimeter RRL observed, the electron densities increase with the increasing frequency of the centimeter line used in the analysis."
+ The one exception is the, The one exception is the
+Compton scattering in the framework of uuparticle physics; aud we fiud that wuparticle physics plays a negligible role in the cosmic rav phoon aud ο propagation.,"Compton scattering in the framework of unparticle physics, and we find that unparticle physics plays a negligible role in the cosmic ray photon and $e^\pm$ propagation."
+ Let us close with several couunuenuts., Let us close with several comments.
+ The work is supported in part by the Science Foundation of Shandong University at Weihai uuder Craut No., The work is supported in part by the Science Foundation of Shandong University at Weihai under Grant No.
+ 0000507300020., 0000507300020.
+ (Seubach1999)., \citep{sembach99}.
+.. (/=98727. b—10738) (Wakker, $l=98\fdg27$ $b=40\fdg38$ \citep{wakwoer91}.
+ Wakkeretal.(1999) D—10 6«109NE.(D/10kpey? 0.080.19(D/10kpe)AL.wr|. al.2000) ," \citet{wakker99} ${\rm D}\sim10$ $6\times10^{6}\;{\rm
+M}_{\sun}\;({\rm D}/10\;{\rm kpc})^{2}$ $0.08-0.19\; ({\rm
+D}/10\;{\rm kpc})\;{\rm M}_{\sun}\;{\rm yr}^{-1}$ \citep{wakker99, gibson00} \citep{vanwoerden99}."
+"μὴ=172kins+ Vise=133lans5| j=(L1+Q8)ς Dj(19.2£(8)ς1075.ci3, ", $\vlsr=-172\;\kms$ $\vlsr=-133\;\kms$ $=(4.1\pm0.8)\times10^{18}\;\cmm{-2}$ $=(19.2\pm0.8)\times10^{18}\;\cmm{-2}$ $-$ $-$ 
+Ineasurelment is recovered by these three 2lem absorption components.,measurement is recovered by these three 21cm absorption components.
+ This is the first detection oL warm phase gas in absorption in an extragalactic system. aud the highest redshift detection of warLnu neutral dhase gas known.," This is the first detection of warm phase gas in absorption in an extragalactic system, and the highest redshift detection of warm neutral phase gas known."
+ It is also the first time limits on temperature have been made iua redshifted DLA system which show a two temperature distribution of the neutral gas. comparable to that found in our own Galaxy.," It is also the first time limits on temperature have been made in a redshifted DLA system which show a two temperature distribution of the neutral gas, comparable to that found in our own Galaxy."
+Fey=0.01 in which only NCL had a MDBII (i.e... £pjj=0 for NC2).,"$F_{\rm BH}=0.01$ in which only NC1 had a MBH (i.e., $F_{\rm BH}=0$ for NC2)."
+ 22 shows that ihe merger model with a AIBIT only for NCI has a significantly higher central density Chan that in the model with MDIIs in both NCT and NC., 2 shows that the merger model with a MBH only for NC1 has a significantly higher central density than that in the model with MBHs in both NC1 and NC2.
+ This strongly suggests that dvnamical heating [rom binary MDIIs sinking into the merger remnant of the (wo NC's is important in reducing the inner densitv of the remnant., This strongly suggests that dynamical heating from binary MBHs sinking into the merger remnant of the two NCs is important in reducing the inner density of the remnant.
+ We also observe that the merger model with a MBI only for NCT has a smaller inner stellar density in comparison wilh the model with no AIDIIL. revealing that a single MDII can also heat up the remnant {ο some extent as it sinks io the center.," We also observe that the merger model with a MBH only for NC1 has a smaller inner stellar density in comparison with the model with no MBH, revealing that a single MBH can also heat up the remnant to some extent as it sinks to the center."
+ 33 illustrates that the final stellar profile of the three sequential major merger events with £pi=0.025 has a significantly lower central density than when £j=0 (i.e.. no MDII: compare Figure 1 by Makino Ebisuzaki 1996).," 3 illustrates that the final stellar profile of the three sequential major merger events with $F_{\rm BH}=0.025$ has a significantly lower central density than when $F_{\rm BH}=0$ (i.e., no MBH; compare Figure 1 by Makino Ebisuzaki 1996)."
+" 33 also reveals that the inner stellar densities of NC* merger remnants in the sequential model with Fi,=0.025 can become proeressivelv lower as NC (ancl possibly inspiraling globular cluster) merging proceeds: the central stellar density in the final merger remnant (1.e.. after three sequential merger events) is a factor of ~ 16 smaller than in the original NC for this sequence."," 3 also reveals that the inner stellar densities of NC merger remnants in the sequential model with $F_{\rm BH}=0.025$ can become progressively lower as NC (and possibly inspiraling globular cluster) merging proceeds: the central stellar density in the final merger remnant (i.e., after three sequential merger events) is a factor of $\sim$ $16$ smaller than in the original NC for this sequence."
+ 44 shows that pxc ad r=005νο is lower in the sequential merger models with larger Fpj owing to stronger dynamical heating., 4 shows that ${\rho}_{\rm NC}$ at $r=0.05R_{\rm NC}$ is lower in the sequential merger models with larger $F_{\rm BH}$ owing to stronger dynamical heating.
+ We do however note that the model with Fey=0.01 does not show a significant decrease of its stellar densities al r=0.2Rxc owing io (he much less effective dvnamical effects of the MBIIs on the NC's., We do however note that the model with $F_{\rm BH}=0.01$ does not show a significant decrease of its stellar densities at $r=0.2R_{\rm NC}$ owing to the much less effective dynamical effects of the MBHs on the NCs.
+ This is in accord with the binary. MDII scouring of the core regions in massive elliptical galaxies. in which the loss cone tvpically dominates only the inner lew percent of the stellar distribution (e.g.. Trujillo et al.," This is in accord with the binary MBH scouring of the core regions in massive elliptical galaxies, in which the loss cone typically dominates only the inner few percent of the stellar distribution (e.g., Trujillo et al."
+ 2004)., 2004).
+ These simulations also suggest that larger. more massive NC's will be more cdifIicult to observe as distinct NCs as their stellar densities may become comparable or less than that of their host galaxy.," These simulations also suggest that larger, more massive NCs will be more difficult to observe as distinct NCs as their stellar densities may become comparable or less than that of their host galaxy."
+ Furthermore. 44 shows that pxce al r=ονο (corresponding to the scale-radii of the original NC's) is lower in (lie sequential merger models with larger Pj. though the dependence is weaker than that of νο at r=0.05Rye on Fey.," Furthermore, 4 shows that ${\rho}_{\rm NC}$ at $r=0.2R_{\rm NC}$ (corresponding to the scale-radii of the original NCs) is lower in the sequential merger models with larger $F_{\rm
+ BH}$, though the dependence is weaker than that of ${\rho}_{\rm NC}$ at $r=0.05R_{\rm NC}$ on $F_{\rm BH}$."
+ This result implies (hat more massive NC merger remnants wilh lower central densities are more susceptible to tidal destruction by the external gravitational fields of their host galaxies owing to their lower mean stellar densities., This result implies that more massive NC merger remnants with lower central densities are more susceptible to tidal destruction by the external gravitational fields of their host galaxies owing to their lower mean stellar densities.
+ 50 [ar we have focused on the internal density. profiles of NC's and have not discussed theirprojected ((wo-climensional. 2D) racial density profiles. which can be directly compared wilh recent observational studies for (3) the origin of the apparent MBII-NC: connection (e.g... Cotte οἱ al.," So far we have focused on the internal density profiles of NCs and have not discussed their (two-dimensional, 2D) radial density profiles, which can be directly compared with recent observational studies for (i) the origin of the apparent MBH-NC connection (e.g., Côtté et al."
+" 2006) and (ii) the observed Fiji,—Mii relation (Graham Spitler 2009).", 2006) and (ii) the observed $F_{\rm BH}-M_{\rm sph}$ relation (Graham Spitler 2009).
+view. because the wavelength carrics information about longitudinal stratification only (Equation (21))). while the amplitude is influenced by both longitudinal aud transverse density profiles (Equation (38))).,"view, because the wavelength carries information about longitudinal stratification only (Equation \ref{eq:wavelength}) )), while the amplitude is influenced by both longitudinal and transverse density profiles (Equation \ref{eq:amptot}) ))."
+ Figure 3. displavs the kink wave amplitude computed ποια. Equation (51)) for different combinations of paralcters., Figure \ref{fig:ampli} displays the kink wave amplitude computed from Equation \ref{eq:amptotres}) ) for different combinations of parameters.
+ In Figure 3((a) we determine the effect of PAR when the remainingquantities are kept constant. while Fieure 23((b) aud (c) show the effect of w aud Pia!Pape respectively.," In Figure \ref{fig:ampli}( (a) we determine the effect of $l/R$ when the remainingquantities are kept constant, while Figure \ref{fig:ampli}( (b) and (c) show the effect of $\omega$ and $\rho_{\rm i 0} / \rho_{\rm apex}$, respectively."
+ Both //P? aud i have a similar effect on the amplitude as both of them control the eficiency of the resonant damping idu frontof the increase of the amplitude due to stratification., Both $l/R$ and $\omega$ have a similar effect on the amplitude as both of them control the efficiency of the resonant damping in frontof the increase of the amplitude due to stratification.
+ When both 7/R and zw get larger. the effect of resouaut daauping becomes strougor.," When both $l/R$ and $\omega$ get larger, the effect of resonant damping becomes stronger."
+ Ou the contrary. pio/Papx controls the role of longitudinal stratification through A (Equation (11))).," On the contrary, $\rho_{\rm i 0} / \rho_{\rm apex}$ controls the role of longitudinal stratification through $\Lambda$ (Equation \ref{eq:deflamstrat}) ))."
+ As pin/Papex grows CX decreases). the increase in amplitude due to stratification becomes more muportaunt.," As $\rho_{\rm i 0} / \rho_{\rm apex}$ grows $\Lambda$ decreases), the increase in amplitude due to stratification becomes more important."
+ For the set of parameters used iu Figure 3((b). the approximate critical frequency for a constaut amplitude according to Equation (52)) is earlsthy8100.," For the set of parameters used in Figure \ref{fig:ampli}( (b), the approximate critical frequency for a constant amplitude according to Equation \ref{eq:relationamp0}) ) is $\omega_{\rm crit} L / {\vk}_0 \approx 100$."
+ We see that. as expected. for the critical frequency. amplitude is coustaut for small 2. but then imereases slieltly with lavecr :.," We see that, as expected, for the critical frequency, amplitude is constant for small $z$, but then increases slightly with larger $z$."
+ We stress again that. for fixed frequency. the amplitude is determined by the combined effect of longitudiual stratification aud transverse inhomogeneity. while the wavelength is only affected by longitudinal stratification.," We stress again that, for fixed frequency, the amplitude is determined by the combined effect of longitudinal stratification and transverse inhomogeneity, while the wavelength is only affected by longitudinal stratification."
+ Uere the aim is to show how the analytical results derived in the previous sections for the cigcumeode problem are related to the time-dependent problem., Here the aim is to show how the analytical results derived in the previous sections for the eigenmode problem are related to the time-dependent problem.
+ Iu addition. we go bevoud the TT aud TB approximations used in the derivation of the analytical expressions. allowing us further insight.," In addition, we go beyond the TT and TB approximations used in the derivation of the analytical expressions, allowing us further insight."
+ To study the problem of huear propagating waves. Equations (1)) and (2)) are nunmericallv solved in cylindrical coordinates.," To study the problem of linear propagating waves, Equations \ref{eq:b1}) ) and \ref{eq:b2}) ) are numerically solved in cylindrical coordinates."
+ Since we are only interested iu the kink mode we assuue an aziuuthal dependence of the form exp(fig)with m=|. and the equations are integrated iu the radial and longitudinal coordinates (7 and + du our case. respectively).," Since we are only interested in the kink mode we assume an azimuthal dependence of the form $\exp \left( i m \varphi\right)$with $m=1$, and the equations are integrated in the radial and longitudinal coordinates $r$ and $z$ in our case, respectively)."
+ A driver at the footpoiut of the loop ἐν= 0) with a single frequency ος is introduced to excite propagating waves along the tube.," A driver at the footpoint of the loop $z=0$ ), with a single frequency $\omega$, is introduced to excite propagating waves along the tube."
+ The spatial for of the driver in the radial direction is chosen. to siuplifv thines. to be the eigenimode of a loop with no resonant laver.," The spatial form of the driver in the radial direction is chosen, to simplify things, to be the eigenmode of a loop with no resonant layer."
+ This enables us to mainly excite the resonantly damped modes. while the direct excitation of Alfvénn modes is very weak.," This enables us to mainly excite the resonantly damped modes, while the direct excitation of Alfvénn modes is very weak."
+ Thus. the driver at +=0 is nuplemented through the periodic variation iu time aud he radial dependence of the kink mode cigentunctions.," Thus, the driver at $z=0$ is implemented through the periodic variation in time and the radial dependence of the kink mode eigenfunctions."
+ The smmlatious are donewith the code MoLMIID (see iu a uniformi exid of 1000«100 points in the r-: plane. and in the range O<r/R10 aud «cHR 100.," The simulations are donewith the code MoLMHD \citep[see][for
+further details about the numerical method]{bonaetal10} in a uniform grid of $1000\times100$ points in the $r$ $z$ plane, and in the range $0<r/R<10$ and $0<z/R<100$ ."
+ For couvenicuce. m the code tine is expressed in units of τν= Res. with ου the internal Alfvóun velocity at := 0. so that the dimensionless Trequencev Iswr7a.," For convenience, in the code time is expressed in units of $\ta = R / v_{\rm A i 0}$ , with $v_{\rm A i 0}$ the internal Alfvénn velocity at $z=0$ , so that the dimensionless frequency is $\omega \ta$ ."
+ Note that iu the previous sections. the," Note that in the previous sections, the"
+A number of studies looking at the influence of e-printing on citation counts across disciplines (e.g..Lawrence2001) and in Astronomy and Physics in particular (e.g..Schwarz&Kennicutt2004;Metcalfe2005:Hennekenetal.2006) found that papers freely available online. particularly through the arXiv e-printserver!.. are cited more often than those not.,"A number of studies looking at the influence of e-printing on citation counts across disciplines \citep[e.g.,][]{2001Nature.411.521L} and in Astronomy and Physics in particular \citep[e.g.,][]{2004BAAS...36.1654S,2005BAAS...37..555M,2006JEPub...9....2H}
+ found that papers freely available online, particularly through the arXiv e-print, are cited more often than those not."
+ This difference was studied in more detail by Kurtzetal. (2005)... who proposed three (non-exclusive) effects potentially responsible for higher citations rates of journal articles also published as astro-ph e-prints.," This difference was studied in more detail by \citet{2005IPM....41.1395K}, who proposed three (non-exclusive) effects potentially responsible for higher citations rates of journal articles also published as astro-ph e-prints."
+ These are defined as (seeKurtzetal.2005): Kurtzetal.(2005) found that open access to older issues of astronomical journals through the Astrophysics Data System (ADS.Kurtzetal.2000) did not lead to increased citation counts. plausibly denying a significant effect of the OA postulate.," These are defined as \citep[see][]{2005IPM....41.1395K}: \citet{2005IPM....41.1395K} found that open access to older issues of astronomical journals through the Astrophysics Data System \citep[ADS,][]{2000A&AS..143...41K} did not lead to increased citation counts, plausibly denying a significant effect of the OA postulate."
+ The EA postulate was supported by a significant increase in citation rate for recent papers., The EA postulate was supported by a significant increase in citation rate for recent papers.
+ Concerning the SB postulate. Kurtzetal.(2005) note that significantly fewer papers not posted to astro-ph are among top 200 cited articles in the year 2003 Astrophysical Journal article than expected from the combined ΟΑΓΕΑ effect alone.," Concerning the SB postulate, \citet{2005IPM....41.1395K} note that significantly fewer papers not posted to astro-ph are among top 200 cited articles in the year 2003 Astrophysical Journal article than expected from the combined OA+EA effect alone."
+ It is thus established that EA and SB play a significant role in gathering citations., It is thus established that EA and SB play a significant role in gathering citations.
+ We observed that on many days the arXiv astro-ph listing is headed by one or more articles submitted à few seconds after the passing of the deadline for submissions for the next mailing., We observed that on many days the arXiv astro-ph listing is headed by one or more articles submitted a few seconds after the passing of the deadline for submissions for the next astro-ph mailing.
+ One example — among many others — is the astro- listing of Sep 24. 2007 headed by three articles received within 20ss after this list was started.," One example – among many others – is the astro-ph listing of Sep 24, 2007 headed by three articles received within s after this list was started."
+ Considering that the number of new articles posted to astro-ph on a typical day is about 35. this temporal clustering close to the submission deadline is conspicuous and suggests a peculiar form of self-promotion.," Considering that the number of new articles posted to astro-ph on a typical day is about $35$, this temporal clustering close to the submission deadline is conspicuous and suggests a peculiar form of self-promotion."
+ Apparently. a subset of authors expects higher visibility and thus also higher citation rates for articles listed at or close to the top of an astro-ph mailing.," Apparently, a subset of authors expects higher visibility and thus also higher citation rates for articles listed at or close to the top of an astro-ph mailing."
+" In the following we study the dependence of citation counts on articles"" positions in the astro-ph listings.", In the following we study the dependence of citation counts on articles' positions in the astro-ph listings.
+ In Sect., In Sect.
+ 2. we describe how we reconstructed the daily astro-ph mailings for a span of 2.5 years and count citations to e-prints.," \ref{sec:data}
+ we describe how we reconstructed the daily astro-ph mailings for a span of 2.5 years and count citations to e-prints."
+ We analyze the dependence of citations rates on astro-ph listing position in Sect., We analyze the dependence of citations rates on astro-ph listing position in Sect.
+ 3. and discuss three possible explanations for the effect we find in our dataset., \ref{sec:analysis} and discuss three possible explanations for the effect we find in our dataset.
+ We present our conclusions in Sect. 4.., We present our conclusions in Sect. \ref{sec:disc-concl}.
+ Using the SPIRES High-Energy Physics Literature we reconstructed the daily astro-ph mailings 1n the period from July 2002 to December 2005., Using the SPIRES High-Energy Physics Literature we reconstructed the daily astro-ph mailings in the period from July 2002 to December 2005.
+ The starting date of our analysis is fixed by the date from which on astro-ph is automatically added to SPIRES HEP as soon as it Is posted., The starting date of our analysis is fixed by the date from which on astro-ph is automatically added to SPIRES HEP as soon as it is posted.
+ Restricting SPIRES HEP queries by the date a record was added allows us to request astro-ph listings for single days., Restricting SPIRES HEP queries by the date a record was added allows us to request astro-ph listings for single days.
+ Older astro-ph articles were added to SPIRES HEP later in 2002., Older astro-ph articles were added to SPIRES HEP later in 2002.
+ They were easily recognized by their arXiv identifier and discarded from our dataset., They were easily recognized by their arXiv identifier and discarded from our dataset.
+ Remaining articles added out of order and occasional small gaps were also found based on, Remaining articles added out of order and occasional small gaps were also found based on
+"be as large as9%,, with a reverse maximum-minimum setup (—6% in summer and —15% in winter) than Figure 4..","be as large as, with a reverse maximum-minimum setup $-6\%$ in summer and $-15\%$ in winter) than Figure \ref{fig4}."
+" In such a case, it is possible that the GFS model forecast is closer to actual situation in summer, rather than in winter as suggested in the figure."," In such a case, it is possible that the GFS model forecast is closer to actual situation in summer, rather than in winter as suggested in the figure."
+" We also attempt to identify any annual variation (Figure 5)), but no significant feature can be noted, possibly due to insufficient temporal coverage."," We also attempt to identify any annual variation (Figure \ref{fig5}) ), but no significant feature can be noted, possibly due to insufficient temporal coverage."
+" As described before, we generated three additional models and will use them to compare with the ISCCP data for July 2006."," As described before, we generated three additional models and will use them to compare with the ISCCP data for July 2006."
+" The result is shown in Figure 6 and, in our opinion,"," The result is shown in Figure \ref{fig6} and, in our opinion,"
+for h(Umm!) derived in In Appenclix (12)) and (E)) we show that the polarisation pseudo power spectra are rotationally invariant.,"for $h(\ell,\ell',m,m')$ derived in In Appendix \ref{app:pseudopol}) ) and \ref{app:rotinv}) ) we show that the polarisation pseudo power spectra are rotationally invariant."
+ For this reason we will in the rest of the paper put the centre of the Gabor window on the north pole., For this reason we will in the rest of the paper put the centre of the Gabor window on the north pole.
+ This makes the calculations easier while keeping the ecnerality of the In this section we will study the Gabor kernel for à Ciaussian and a top-hat Gabor window., This makes the calculations easier while keeping the generality of the In this section we will study the Gabor kernel for a Gaussian and a top-hat Gabor window.
+" The Gaussian is defined as where we will use 5 and 15 degrees EWIIM (corresponding to 0=2.12 and σξ 6.38"") and a cut-oll angle θε:=30.", The Gaussian is defined as where we will use $5$ and $15$ degrees FWHM (corresponding to $\sigma=2.12^\circ$ and $\sigma=6.38^\circ$ ) and a cut-off angle $\theta_C=3\sigma$.
+ We will also be comparing with a top-hat window covering the same area on the sky (66 is the same)., We will also be comparing with a top-hat window covering the same area on the sky $\theta_C$ is the same).
+ The top-hat window is defined as Studying equation (5)) and (6)) one sees that the ££ and D modes are mixing when only a portion of the sky is observed., The top-hat window is defined as Studying equation \ref{eq:pcle}) ) and \ref{eq:pclb}) ) one sees that the $E$ and $B$ modes are mixing when only a portion of the sky is observed.
+ The kernel AS.) is the kernel which takes full sky CF modes to the pseudo coelIicients CF and similarly for CP.," The kernel $K_2(\ell,\ell')$ is the kernel which takes full sky $C^E_\ell$ modes to the pseudo coefficients $\tilde
+C^E_\ell$ and similarly for $C^B_\ell$."
+ The kernel A.ος} is the one which causes the mixing.," The kernel $K_{-2}(\ell,\ell')$ is the one which causes the mixing."
+ In figure (1)) and (2)) we have plotted the kernels AS(6.0) together with Woof.) fora 5 and 15 degree ENLEM Gaussian Gabor window with 67:=3o.," In figure \ref{fig:k2g5}) ) and \ref{fig:k2g15}) ) we have plotted the kernels $K_2(\ell,\ell')$ together with $K_{-2}(\ell,\ell')$ for a $5$ and $15$ degree FWHM Gaussian Gabor window with $\theta_C=3\sigma$."
+ One can see that the diagonal of the AS(f kernel is about an order of magnitude larger than the diagonal of the mixing kernel ἐν2(6(. (7).," One can see that the diagonal of the $K_2(\ell,\ell')$ kernel is about an order of magnitude larger than the diagonal of the mixing kernel $K_{-2}(\ell,\ell')$ ."
+ This means that CF would dominate CF and CF would dominate C? provided that the two spectra CF and CE. were of the same order of magnitudo., This means that $C^E_\ell$ would dominate $\tilde C^E_\ell$ and $C^B_\ell$ would dominate $\tilde C^B_\ell$ provided that the two spectra $C^E_\ell$ and $C^B_\ell$ were of the same order of magnitude.
+ Llowever as discussed before. the CF are expected to be considerably smaller than CF in most cosmological models.," However as discussed before, the $C^B_\ell$ are expected to be considerably smaller than $C^E_\ell$ in most cosmological models."
+ For the E mode this is not a problem as CF will then hardly be allected., For the $E$ mode this is not a problem as $C^E_\ell$ will then hardly be affected.
+ Phe problem is the 2B mode which in this case will be dominated by the £ The separation of £and 2 modes of polarisation on the cut sky was already discussed in, The problem is the $B$ mode which in this case will be dominated by the $E$ The separation of $E$and $B$ modes of polarisation on the cut sky was already discussed in
+the Mask Preparation Software. see Garillietal.(1999))) from pixel to mau with respect to the quadrant optical axis using the CCD to mask transformation The purpose of the entire MMU system is to perform the following operations: Iu all operatious the MALU will only. kuow masks by their individual six-digit code. aud will iorimally handle entire mask sets.,"the Mask Preparation Software, see \citet{Garilli}) ) from pixel to mm with respect to the quadrant optical axis using the CCD to mask transformation The purpose of the entire MMU system is to perform the following operations: In all operations the MMU will only know masks by their individual six-digit code, and will normally handle entire mask sets."
+ A mask set is defined as a set of four masks. with the same ast 2 digits in the id. used for the same observation in the f£. VIMOS or NIRMOS quacranuts.," A mask set is defined as a set of four masks, with the same last 5 digits in the id, used for the same observation in the 4 VIMOS or NIRMOS quadrants."
+ All operations (except contingency recovery and maintenance operations managed. autonomously yy the MALL operator) are executed after receipt from MCS of a Job order in the Dorm of a file rausferred into a staging area of the MHC. aud terminate upou placing a Termination report iu he same staging area. to be retrieved by MCS that las also access to the tables describing the current Content of the SC (SCT) ancl of the ICs (ICT).," All operations (except contingency recovery and maintenance operations managed autonomously by the MMU operator) are executed after receipt from MCS of a Job order in the form of a file transferred into a staging area of the MHCU, and terminate upon placing a Termination report in the same staging area, to be retrieved by MCS that has also access to the tables describing the current content of the SC (SCT) and of the ICs (ICT)."
+" The SCT. for each quadrant. is a list of 100 ileus containiug the codes of the masks currently stored or ""empty. with no information about he slot where the mask are located."," The SCT, for each quadrant, is a list of 100 items containing the codes of the masks currently stored or ""empty"", with no information about the slot where the mask are located."
+ The ICT. for each Instrument. Cabinet. is a lst of 15 ens which combine the slot number with the code of the contained mask.," The ICT, for each Instrument Cabinet, is a list of 15 items which combine the slot number with the code of the contained mask."
+ The information about the slit positions aucl sizes is contained in the Machine Slit Files (MSF) transmitted together with a MMJ., The information about the slit positions and sizes is contained in the Machine Slit Files (MSF) transmitted together with a MMJ.
+ The full manufacturing cvele is a [step operation (Fig., The full manufacturing cycle is a 4–step operation (Fig.
+ 9)., 9).
+"so-called mate. which is au No.ON square. sviunietnc matrix whose (k.{) αιιν is giveu bv The superscript denotes thetranspose,","so-called ”, which is an $N \times
+N$ square, symmetric matrix whose $(k,j)$ -entry is given by The superscript ${}^*$ ” denotes the."
+ Matrix Fis unitary. ic. with J the ideutitv matrix.," Matrix $\Fb$ is unitary, i.e., with $\Ib$ the identity matrix."
+ Another useful property is that with By means of Eq. (9)), Another useful property is that with By means of Eq. \ref{eq:FF}) )
+ it can be shown that where Fy=RLE| aud Fy=Z[F] and where [| aud πι are the veal and the imaginary parts of a complex quantity., it can be shown that where $\Fb_{\Rmatc} \equiv \Rmatc[\Fb]$ and $\Fb_{\Imatc} \equiv \Imatc[\Fb]$ and where $\Rmatc[.]$ and $\Imatc[.]$ are the real and the imaginary parts of a complex quantity.
+ Iudeed. and A isa real matrix. heuce Eq. (11))," Indeed, and $\Hb$ isa real matrix, hence Eq. \ref{eq:orth}) )"
+ has to hold., has to hold.
+ From Eq. (3)), From Eq. \ref{eq:DFT}) )
+ if is also possible to see that ‘ry is a real quautity., it is also possible to see that $\xh_0$ is a real quantity.
+" In the case that Nis an even προ, the same holds for a,)4 where Ny=N/2| aud represeuts the siuallest integer greater than +."," In the case that $N$ is an even number, the same holds for $\xh_{N_{\dag}+1}$ where $N_{\dag} = \lceil N/2 \rceil$ and $\lceil z \rceil$ represents the smallest integer greater than $z$."
+ Finally. dealing with 2. ouly half of this array ean be considered. since £yà=p. ko1.2.....NyOy; where 6= Lif?=j aud zero otherwise.," Finally, dealing with $\xhb$, only half of this array can be considered, since $\xh_{N-k} = \xh^*_k$, $k=1, 2, \ldots,
+N_{\dag}-\delta_{2 N_{\dag}}^N$, where $ \delta_i^j = 1$ if $i=j$ and zero otherwise."
+" With his notation. the periodogram of a is defined as where fy is the Ath ον of the array a[0904seery,al? aud |.| denotes the nox."," With this notation, the periodogram of $\xb$ is defined as where $\xhu_k$ is the $k$ th entry of the array $\xhbu^T = [\xh_0, \xh_1, \ldots, \xh_{N_{\dag}-1}]$ and $| . |$ denotes the norm."
+" By iicaus of a coluun array obtained by the colui coucatenation of £4=Ria) aud a,=τα]. Eq. (13))"," By means of a column array obtained by the column concatenation of $\xhbu_{\Rmatc} \equiv \Rmatc[\xhbu]$ and $\xhbu_{\Imatc} \equiv \Imatc[\xhbu]$ , Eq. \ref{eq:period2}) )"
+ can be rewritten in the form Iu Sect., can be rewritten in the form In Sect.
+ 5 it is shown that this periodogram is equivalent to that obtainable by mucaus the least-squaresfit of model (1)) with M=N.¢; =j aud fy=iN.," \ref{sec:fit} it is shown that this periodogram is equivalent to that obtainable by means the least-squaresfit of model \ref{eq:model}) ) with $M=N$, $t_j = j$ and $f_k = k/N$ ."
+" Au important point to stress is that. if a is the realization of a (not necessarily Cassia) radon process. then cach t, is given by the sun of [IN random variables."," An important point to stress is that, if $\xb$ is the realization of a (not necessarily Gaussian) random process, then each $\xhu_k$ is given by the sum of $N$ random variables."
+ This is because of the linearity of theoperator F., This is because of the linearity of the $\Fb$.
+ Thanks to the theorein. therefore. the cutrics of @ can be expected to be Gaussian random quantities.," Thanks to the , therefore, the entries of $\xhb$ can be expected to be Gaussian random quantities."
+" As a consequence.the eutries of 2 cam also be expected to be Gaussian random quantities with covariance matrix Cz—[zzδι| evel by Here. E[.] denotes theoperator, is the covariance matrix of a. aud FE the matrix obtained by the first rows of the F."," As a consequence,the entries of $\zhb$ can also be expected to be Gaussian random quantities with covariance matrix $\Cb_{\zhb} = {\rm E}[\zhb \zhb^T]$ given by Here, ${\rm E}[.]$ denotes the, $\Cb_{\xb} = {\rm E}[\xb \xb^T]$ is the covariance matrix of $\xb$, and $\Fbu$ the matrix obtained by the first$N_{\dag}$ rows of the $\Fb$."
+ From Eqs. (16)), From Eqs. \ref{eq:Ce}) )
+" and (11)). it is easy to deduce that. if a is the realization of a stand white-noise process. ναι. C,=I. then Cz isa diagonal matrix with (Cz)44= 1. (C;)sv,=0 and (Cz)ha=O.5."," and \ref{eq:orth})), it is easy to deduce that, if $\xb$ is the realization of a standard white-noise process, i.e., $\Cb_{\xb} = \Ib$, then $\Cb_{\zhb}$ is a diagonal matrix with $(\Cb_{\zhb})_{11}=1$ , $(\Cb_{\zhb})_{N_{\dag} N_{\dag}}=0$ and $(\Cb_{\zhb})_{kk}=0.5$."
+ In other words. the cutrics of 2 are mutualle uncorrelated.," In other words, the entries of $\zhb$ are mutually uncorrelated."
+" In tuu. this micas that W[r,| is uncorrelated with Z[r,]"," In turn, this means that $\Rmatc[\xhu_k]$ is uncorrelated with $\Imatc[\xhu_k]$."
+" If x is a colored (uot necessarily stationary) noise process. C, 1 no a diagonal matrix. then this holds also for Cz."," If $\xb$ is a colored (not necessarily stationary) noise process, $\Cb_{\xb}$ is not a diagonal matrix, then this holds also for $\Cb_{\zhb}$."
+" Uowever. from a it is possible to obtain an array gy containing mutually uncorrelated cutrics by imieaus of the traustormation The matrix C,7 cau be computed via with U the orthogonal matrix whose cols contain the cigenvectors of Cy and X a diagonal matrix containing⋅⋅ the corresponding: eigeuvalues: (Àj;No "," However, from $\xb$ it is possible to obtain an array $\yb$ containing mutually uncorrelated entries by means of the transformation The matrix $\Cb_{\xb}^{-1/2}$ can be computed via with $\Ub$ the orthogonal matrix whose columns contain the eigenvectors of $\Cb_{\yb}$ and $\Sigmab$ a diagonal matrix containing the corresponding eigenvalues $\{ \lambda_l \}_{l=0}^{N - 1}$."
+"This decomposition is particularly simple aud computationally CHicient if Cy. is achreulent matrix since d$ can be dciagonalized according to where ""diag|q| denotes a diagonal matrix whose diagonal eutiies are given by the array q a cis the firstcolumn of Cy.", This decomposition is particularly simple and computationally efficient if $\Cb_{\xb}$ is a matrix since it can be diagonalized according to where ${\rm diag}[\qb]$ ” denotes a diagonal matrix whose diagonal entries are given by the array $\qb$ and $\cb$ is the firstcolumn of $\Cb_{\xb}$ .
+ Becauseof this. the decomposition (185)) can be directly computed with U=F aud X= diagle|: hence. y=FS ΔΕ ," Becauseof this, the decomposition \ref{eq:eig})) can be directly computed with $\Ub = \Fb$ and $\Sigmab = {\rm diag}[\cb]$ ; hence, $\yb = \Fb^* \Sigmab^{-1/2} \Fb$ ."
+This meaus that the whitening operation can be performed im the harmonic domain through the following procedure: a) computation of z. tthe DET of ad: b) computation. of eq[rj=κ~ALjXl/2 ER and ο)," This means that the whitening operation can be performed in the harmonic domain through the following procedure: a) computation of $\xhb$ , the DFT of $\xb$ ; b) computation of $\{ \yh_k \} = \{\xh_k / \lambda^{1/2}_k \}$ ; and c)"
+Iu deriving these expressions. we have made use of the identities,"In deriving these expressions, we have made use of the identities"
+ (Cole&Lacey1996: 5=(logp)fd(logr). isLin 53 (Diubinski&Carlhere 5.," \citep{Cole96,Moore99,Bullock2001,vdB02,Navarro04}; $\gamma
+\equiv d(\log{\rho})/d(\log{r})$ $\gamma \sim -1$ $\gamma \sim -3$ \citep{Dubinski91,NFW,Navarro97,Huss99,MacMillan06}."
+ (Dalcanton&οσα2001:Merrittetal.," $\gamma$ \citep{Dalcanton01,Merritt05}."
+2005).. Tavlor&Navarro(2001) ln(2005) pfa?. à1.9 for over two aud a half decades in radius in spite of a varving density profile (see also Austinetal.(2005) for an analytic exploration of this behavior)., \citet{Taylor01} \citet{Dehnen05} $\rho/\sigma^3$ $\alpha \sim 1.9$ for over two and a half decades in radius in spite of a varying density profile (see also \citet{Austin05} for an analytic exploration of this behavior).
+ There is also evidence for a universality in the relation between the slope of the deusitv profile and the velocity anisotropy., There is also evidence for a universality in the relation between the slope of the density profile and the velocity anisotropy.
+" The anisotropy parameter. ο) is defined as where a, is the velocity dispersion in the taugeutial direction aud oa, is the velocity dispersion iu the radial direction."," The anisotropy parameter, $\beta$ is defined as where $\sigma_{\phi}$ is the velocity dispersion in the tangential direction and $\sigma_r$ is the velocity dispersion in the radial direction."
+ 9 describes the degree of anisotropy of the velocity dispersion in a spherical halo. which iu siuulations ranges four isotropic in the core (2= 0) to radially anisotropic at the outskirts (0= 1)(Cole& 2001).," $\beta$ describes the degree of anisotropy of the velocity dispersion in a spherical halo, which in simulations ranges from isotropic in the core $\beta=0$ ) to radially anisotropic at the outskirts $\beta=1$ \citep{Cole96,Carlberg97,Fukushige01}."
+. Hansen&Moore(2006) fluc a direct correlation between J aud the density profile slope for halos with a wide variety of initial concditious (i.e. isolated collapse. morecrs).," \citet{Hansen06}
+ find a direct correlation between $\beta$ and the density profile slope for halos with a wide variety of initial conditions (i.e. isolated collapse, mergers)."
+ Hausenetal.(2006) find another universality. this time in the velocity distribution function. for a siuilu variety of initial couclitions.," \citet{Hanseneveryone06} find another universality, this time in the velocity distribution function, for a similar variety of initial conditions."
+ The radial orbit instability (ROI) may be the cause of the apparent uuiversalitv of these dark matter halo properties., The radial orbit instability (ROI) may be the cause of the apparent universality of these dark matter halo properties.
+ The ROI occurs im anisotropic spherical systenis composed of particles with predominautly radial orbits., The ROI occurs in anisotropic spherical systems composed of particles with predominantly radial orbits.
+ The instability arises when particles iu precessiug clongated loop orbits experience a torque due to a slight asvuuuectry., The instability arises when particles in precessing elongated loop orbits experience a torque due to a slight asymmetry.
+ This torque causes them to lose some aneular momentum aud move towards the ceuter of thesystem., This torque causes them to lose some angular momentum and move towards the center of thesystem.
+Particles withsmall enough angular momenta become trapped iu box orbits. alguing themsclves witli the erowing bar aud reimforciug the initial perturbation.,"Particles withsmall enough angular momenta become trapped in box orbits, aligning themselves with the growing bar and reinforcing the initial perturbation."
+The threshold / is constant across the field of each plate. but varies from plate to plate according to the Duetuation of the sky density for incividual plates (Macdox et al.,"The threshold $t$ is constant across the field of each plate, but varies from plate to plate according to the fluctuation of the sky density for individual plates (Maddox et al."
+ 1990a.b).," 1990a,b)."
+" When working with APAI cata it is convenient to define an ""isophotal APAL magnitude. mi. of an image as This APAL parameter is related to the actual isophotal magnitude. 6,,. which is given by where Z ds the magnitude zero-point of the photographic plate."," When working with APM data it is convenient to define an `isophotal APM magnitude', $m_{\rm iso}$, of an image as This APM parameter is related to the actual isophotal magnitude, $b_{\rm iso}$, which is given by where $Z$ is the magnitude zero-point of the photographic plate."
+ Also. the isophotal size of the image is given by din. and idus ave the three main APAL parameters that we use to measure the SB profile of galaxies.," Also, the isophotal size of the image is given by $I_{\rm iso}$, $t$ and $A_{\rm iso}$ are the three main APM parameters that we use to measure the SB profile of galaxies."
+ In former work on APM ealaxy survey. Alacidox et al. (," In former work on APM galaxy survey, Maddox et al. ("
+1990b) used these parameters together with the assumption of Gaussian galaxy profiles to caleulate a total APAL magnitude myo. which was then calibrated to b; using CCD magnitudes.,1990b) used these parameters together with the assumption of Gaussian galaxy profiles to calculate a total APM magnitude $m_{\rm tot}$ which was then calibrated to $b_{J}$ using CCD magnitudes.
+ For faint ealaxies near the magnitude limit of the ADM survey this is a reasonable approximation. because the galaxy images are small and the profile shape is dominated by the secing disc.," For faint galaxies near the magnitude limit of the APM survey this is a reasonable approximation, because the galaxy images are small and the profile shape is dominated by the seeing disc."
+ jut for brighter galaxies an exponential profile is a better approximation (see Section 5))., But for brighter galaxies an exponential profile is a better approximation (see Section \ref{sec_discussion}) ).
+ In the following subsections. we mainly use these three fundamental APAL parameters to investigate the SB of galaxies.," In the following subsections, we mainly use these three fundamental APM parameters to investigate the SB of galaxies."
+ Additionallv. there are two other APAL parameters that ave related to the profile of galaxies.," Additionally, there are two other APM parameters that are related to the profile of galaxies."
+ Both of themare density weighted. parameters., Both of themare density weighted parameters.
+ One is στι defined as where. {ωνis the density at. location . ancl y relative to the image centroid.," One is $\sigma^{2}$, defined as where, $D_{\rm obj}(x,y)$is the density at location $x$ and $y$ relative to the image centroid."
+" Theoreticallv. 07. together with la, can constrain the SD. profile of the image. and judge whether it is best fit by a Gaussian. or exponential. Or even anm rLl profile."," Theoretically, $\sigma^{2}$, together with $A_{\rm iso}$ can constrain the SB profile of the image, and judge whether it is best fit by a Gaussian, or exponential, or even an $r^{1/4}$ profile."
+⋅ ⇁⋅Unfortunately. in. practice. this oiis not possible. because a7E changes by only about 104. between these three profiles. and this is the same order as the observational error in 07 (see Section 5)).," Unfortunately, in practice this is not possible, because $\sigma^{2}$ changes by only about $10\%$ between these three profiles, and this is the same order as the observational error in $\sigma^{2}$ (see Section \ref{sec_discussion}) )."
+ Additionally. 62 can be seriously allected by other factors. such as the ellipticity. and substructure of the image. and any other asvnunetry of the profile⋅ which. will. tend to increase. 2στ. rp.," Additionally, $\sigma^{2}$ can be seriously affected by other factors, such as the ellipticity and substructure of the image, and any other asymmetry of the profile which will tend to increase $\sigma^{2}$."
+"This. means that the measurement of o7 is not as robust as cdi Or fi,"" and so we do not use o72 in the measurement of SB in this paper."," This means that the measurement of $\sigma^{2}$ is not as robust as $A_{\rm iso}$ or $I_{\rm iso}$, and so we do not use $\sigma^{2}$ in the measurement of SB in this paper."
+ In Section 5.2.. we consider the distribution of a> [or bright galaxies which have a large image and. hence a smaller error in 07. and show that an exponential prolile is actually a reasonable approximation to the SB profile of most galaxies.," In Section \ref{sec_profiles}, we consider the distribution of $\sigma^{2}$ for bright galaxies which have a large image and hence a smaller error in $\sigma^{2}$, and show that an exponential profile is actually a reasonable approximation to the SB profile of most galaxies."
+ Another useful parameter is the ellipticity of the image. which is derived from censity-weightecl second moments.," Another useful parameter is the ellipticity of the image, which is derived from density-weighted second moments."
+ This is can be easily converted to the apparent axial ratio of he APA image (Macdox et al., This is can be easily converted to the apparent axial ratio of the APM image (Maddox et al.
+ 1990a: Lambas et al., 1990a; Lambas et al.
+ 1992). On the assumption that the areal profiles of an .APM image have similar shapes. then e is independent. of other profile parameters. and. will not allect the measurement of SB.," 1992), On the assumption that the areal profiles of an APM image have similar shapes, then $e$ is independent of other profile parameters, and will not affect the measurement of SB."
+ However. e will alleet the estimation of 6 as discussed in Section 5.2..," However, $e$ will affect the estimation of $\sigma^{2}$ as discussed in Section \ref{sec_profiles}."
+ If the image of a galaxy is represented by à circularly-symmetric pure exponential surface brightness profile. The central peak value po. and the characteristic size ro fully describe the image.," If the image of a galaxy is represented by a circularly-symmetric pure exponential surface brightness profile, The central peak value $p_{0}$, and the characteristic size $r_{0}$ fully describe the image."
+ In this case. ro is the exponential scale length.," In this case, $r_{0}$ is the exponential scale length."
+ The total luminosity of the image. fy... is simply the integral of por).," The total luminosity of the image, $I_{\rm tot}$, is simply the integral of $p(r)$."
+ “Phe integration can be divided into three parts which contribute to £i. where dow. fis. and Zuaaq represent the fux lost. due to saturation. the observed flux and the unobserved. Lux bevond the isophotal size of the galaxy respectively.," The integration can be divided into three parts which contribute to $I_{\rm tot}$ , where $I_{\rm sat}$ , $I_{\rm iso}$ and $I_{\rm field}$ represent the flux lost due to saturation, the observed flux and the unobserved flux beyond the isophotal size of the galaxy respectively."
+ οσο can be written as: and, These can be written as: and
+and the ionized eas in the Bar likely occurs in the regions localized to the cluster sources.,and the ionized gas in the Bar likely occurs in the regions localized to the cluster sources.
+ Iu addition. the well-known mincavitv (marked with the erav circle in Figures la aud. £0). a regiou depressed of radio coutiumuu. is located SW of Ser A* (Yusef-Zadeli.Morris.&Ekers1989).," In addition, the well-known minicavity (marked with the gray circle in Figures 4a and 4b), a region depressed of radio continuum, is located SW of Sgr A* \citep{yusef89}."
+. This region is likely created by a fast wind outflow from the sources iu the vicinity sweeping up the eas in the orbiting ionized streams. as suggested by the morphology of a bright [FeITI| line i- the infrared (Lutz.Krabbe.&Cenzel1993).," This region is likely created by a fast wind outflow from the sources in the vicinity sweeping up the gas in the orbiting ionized streams, as suggested by the morphology of a bright [FeIII] line in the infrared \citep{lutz93}."
+ However. the iuiuicavitv is a few arescc displaced from the hieh-kinetic-teniperature region. which in fact is best seen as a depression associated with W300 line ciission.," However, the minicavity is a few arcsec displaced from the high-kinetic-temperature region, which in fact is best seen as a depression associated with $\alpha$ line emission."
+ The uiduicavity cau also be explained by Melia.Coker.&Yuset-Zadehli(1996) as a downstream. focuced flow due to the accretion by the SMDBII at Ser A*.," The minicavity can also be explained by \cite{melia96}
+ as a downstream, focuced flow due to the accretion by the SMBH at Sgr A*."
+ Alternatively. the hieh 4uperature iu the Dar could result from a collision between two ionized orbiting flows. as suggested iu the ανασα] model of Zhaoetal.(2009).," Alternatively, the high temperature in the Bar could result from a collision between two ionized orbiting flows, as suggested in the dynamical model of \cite{zhao09}."
+.. Α computation using their Weplerian orbital model (see Section 3.2) shows that the specific kinetic energy reaches a nmnaxiuun (see Figure 1d) in the Bar region where high kinetic temperatures (see Figure 19) are observed., A computation using their Keplerian orbital model (see Section 3.2) shows that the specific kinetic energy reaches a maximum (see Figure 4d) in the Bar region where high kinetic temperatures (see Figure 4c) are observed.
+ The distribution of specific kinetic euergv (Figure td) shows a substantial increase from 2«105 ere bat a regiou near IRS 5 in the Northern Aru to 12«10! ereoOo & Fat TRS 33 in the Dar., The distribution of specific kinetic energy (Figure 4d) shows a substantial increase from $2\times10^{14}$ erg $^{-1}$ at a region near IRS 5 in the Northern Arm to $12\times10^{14}$ erg $^{-1}$ at IRS 33 in the Bar.
+ However. we note that a high kinetic energy of the orbiting flow docs not have to correspond to a hieh kinetic temperature: the ionized flows can release their kinetic enerey to shocks if the ionized flows collide with cach other.," However, we note that a high kinetic energy of the orbiting flow does not have to correspond to a high kinetic temperature; the ionized flows can release their kinetic energy to shocks if the ionized flows collide with each other."
+ If they collide with 16 strong stellar winds from the clusters of stars in the Bar area. then the eas could be heated by 1ο shocks to higher temperatures.," If they collide with the strong stellar winds from the clusters of stars in the Bar area, then the gas could be heated by the shocks to higher temperatures."
+ In addition. re model also shows that the ionized flows of ιο Northern aud Eastern Arms do collide at ie region with high kinetic temperatures (see Figure 21d of Zhaoetal.(2009) and the 3D structure of the κρίνα arius in Figure 5).," In addition, the model also shows that the ionized flows of the Northern and Eastern Arms do collide at the region with high kinetic temperatures (see Figure 21d of \cite{zhao09} and the 3D structure of the minispiral arms in Figure 5)."
+ The observed high kinetic temperatures in the Dar suggest that the ionized eas could be further heated by the shocks generated from the collision., The observed high kinetic temperatures in the Bar suggest that the ionized gas could be further heated by the shocks generated from the collision.
+ liucnmiaties of the ionized eas iu both the Northern Ari and Western Arc have been explained as Ixepleriau circular motion of a circumunclear ring (Schwarz.Bregman.&vanCorkom1050:&Duschl2000;Liszt—2003).," Kinematics of the ionized gas in both the Northern Arm and Western Arc have been explained as Keplerian circular motion of a circumnuclear ring \citep{schw89,lacy91,voll00,lisz03}."
+. Based on the VLA) observations of W920. Roberts&Coss(1993) found that the ionized riug consists of three separated kincmatic components. the Western Arc. Northern Anu. aud Dar.," Based on the VLA observations of $\alpha$, \cite{rob93} found that the ionized ring consists of three separated kinematic components, the Western Arc, Northern Arm, and Bar."
+ The radial velocities of the Western Arc can be fitted with a Iseplerian model with circular motion., The radial velocities of the Western Arc can be fitted with a Keplerian model with circular motion.
+ Using the observations of the Drackett-5 line. Pauuiuxl.Maillard.&Morris(2001) found that the kinematics of the ionized eas iu the Northern Ax can be explained as a bundle of elliptical orits," Using the observations of the $\gamma$ line, \cite{paum04} found that the kinematics of the ionized gas in the Northern Arm can be explained as a bundle of elliptical orbits."
+ Iu addition. the Eastern Avi may be interrupted by a high-velocity component “tip”.," In addition, the Eastern Arm may be interrupted by a high-velocity component “tip”."
+" Also. the Bar component extends to the IRS6/31 region which is affected by two other kinematic features. the “Western Bridec™ and the ""Dar Overlay.”"," Also, the Bar component extends to the IRS6/34 region which is affected by two other kinematic features, the “Western Bridge” and the “Bar Overlay.”"
+ The VLA observations of the W920 liue sugeest that he loci of the three ninispiral arms can be fitted with three families of elliptical orbits: the orbits of doth the Northern and Eastern Avms have veh eccentricities while the Western Arc is in rearly circular motion (Zhaoetal.2009)., The VLA observations of the $\alpha$ line suggest that the loci of the three minispiral arms can be fitted with three families of elliptical orbits; the orbits of both the Northern and Eastern Arms have high eccentricities while the Western Arc is in nearly circular motion \citep{zhao09}.
+. The VLA measurements of the proper motions and he radial velocity resolved the degeneracies iu the inclination angles. sugecsting that the two ionized flows from the Northern aud Eastern Arms collide in the Dar regio-," The VLA measurements of the proper motions and the radial velocity resolved the degeneracies in the inclination angles, suggesting that the two ionized flows from the Northern and Eastern Arms collide in the Bar region."
+ On the other hand. Miziceetal.(2007) suggested that the nonu-Ixepleriau motions might be nuportaut iu the kinciatics of the filameuts observed ucar Ser À*.," On the other hand, \cite{muzic07} suggested that the non-Keplerian motions might be important in the kinematics of the filaments observed near Sgr A*."
+ Tn addition. Voller&Duschl(2000) noticed that a substantial aunount of velocity dispersion along with additional iusvards radial motion is required iu order to fif the observed. velocities of the ionized flows with their Iseplerian circulaz-orbit model.," In addition, \cite{voll00} noticed that a substantial amount of velocity dispersion along with additional inwards radial motion is required in order to fit the observed velocities of the ionized flows with their Keplerian circular-orbit model."
+" Using the Keplerian model described in Zhaoetal.(2009) along with the orbital parameters derived from fittine the VLA [92 data and proper motion measurements, we recomputed the 3D locations of the three ionized aris."," Using the Keplerian model described in \cite{zhao09} along with the orbital parameters derived from fitting the VLA $\alpha$ data and proper motion measurements, we recomputed the 3D locations of the three ionized arms."
+ Figure 5 shows a 3D visualization of the three ionized flows in the Northern. Eastern. and Western Arxius (Are). which illustrates that the Northern. Arm," Figure 5 shows a 3D visualization of the three ionized flows in the Northern, Eastern, and Western Arms (Arc), which illustrates that the Northern Arm"
+Decause HectoNLADP is à survey of red galaxies. the large-scale features of the map are more sharply defined (han they would be if we included the less strongly clustered blue objects (Davis Geller 1976: Zehavi et al.,"Because HectoMAP is a survey of red galaxies, the large-scale features of the map are more sharply defined than they would be if we included the less strongly clustered blue objects (Davis Geller 1976; Zehavi et al."
+ 2011)., 2011).
+ The much larger BOSS survev (White et al., The much larger BOSS survey (White et al.
+ 2011) eeared toward exploitation of the baryon acoustic peak. select even redder objects essentially eliminating objects with 2<0.45.," 2011) geared toward exploitation of the baryon acoustic peak, select even redder objects essentially eliminating objects with $z \lesssim 0.45$."
+ The BOSS survey is optimized to take advantage of a sparsely sampled large volume at greater redshilt than the range we sample with HectoMAD., The BOSS survey is optimized to take advantage of a sparsely sampled large volume at greater redshift than the range we sample with HectoMAP.
+ In contrast. the dense sampling of HectoMAP. enables a focus on clusters οἱ galaxies and (heir relationship with laree-sale structure.," In contrast, the dense sampling of HectoMAP enables a focus on clusters of galaxies and their relationship with large-sale structure."
+ Based on our earlier SIIELS survey (Geller et al., Based on our earlier SHELS survey (Geller et al.
+ 2010: IXurtz et al., 2010; Kurtz et al.
+ 2011). HectoMADP should contain ~50 clusters with rest-frame line-of-sight velocity dispersions 2GOO km ," 2011), HectoMAP should contain $\sim 50$ clusters with rest-frame line-of-sight velocity dispersions $\gtrsim 600$ km $^{-1}$."
+The survey should. vield a robust determination of the cluster mass function in the range Q.2<z«0.55., The survey should yield a robust determination of the cluster mass function in the range $0.2 < z < 0.55$.
+ Svstems of galaxies are evident throughout this range in Figure 4: it is. ol course. harder to see them by eve at greater redshift because the line-ol-sight. velocity dispersion is a small fraction of the mean redshift.," Systems of galaxies are evident throughout this range in Figure \ref{fig:hectomap}; it is, of course, harder to see them by eye at greater redshift because the line-of-sight velocity dispersion is a small fraction of the mean redshift."
+ The set of svstems extracted. from the full survey. promise. among many other applications. a strong test of cluster catalogs based on identification of the red sequence in photometric data.," The set of systems extracted from the full survey promise, among many other applications, a strong test of cluster catalogs based on identification of the red sequence in photometric data."
+ As an example. Figure 5 shows à small test of the GAIBCG catalog (lao et al.," As an example, Figure \ref{fig:gmbcg} shows a small test of the GMBCG catalog (Hao et al."
+ 2010)., 2010).
+ In this catalog. clusters are selected from the SDSS photometric catalog with an algorithm based on identilving the red sequence and the brightest cluster member.," In this catalog, clusters are selected from the SDSS photometric catalog with an algorithm based on identifying the red sequence and the brightest cluster member."
+" The figure shows a sample of the highest confidence GAIBCG clusters (with at least 15 members) that lie within the most complete HectoMAP region right ascension range 13.3:33""of*""Mδρυ9""15.33.oh", The figure shows a sample of the highest confidence GMBCG clusters (with at least 15 members) that lie within the most complete HectoMAP region right ascension range $13.33^h < \alpha_{2000} < 15.33^h$.
+ The vellow bar is centered at the mean photometric redshift flor each GAIBCG cluster in the range of this portion of HectoNLADP and the bar extends for zz/(1+2)~0.01., The yellow bar is centered at the mean photometric redshift for each GMBCG cluster in the range of this portion of HectoMAP and the bar extends for $\pm \Delta{z}/(1+z) \sim 0.01$.
+ The black points are galaxies in (he HectoATAP survey: these red galaxies are the ones that populate a typical cluster red sequence., The black points are galaxies in the HectoMAP survey; these red galaxies are the ones that populate a typical cluster red sequence.
+ The correspondence between the redshift survey and this cluster catalog is poor., The correspondence between the redshift survey and this cluster catalog is poor.
+ Many ol the GMDCG clusters do not exist at all or they are superpositions of cuts through the cosmic web., Many of the GMBCG clusters do not exist at all or they are superpositions of cuts through the cosmic web.
+ Some of the svstems which appear as clean “fingers” in HectoMAP. ave missed bv the GAIBCG scheme., Some of the systems which appear as clean “fingers” in HectoMAP are missed by the GMBCG scheme.
+ The discrepancies underscore the importance of spectroscopy [or understanding structure in the universe and its evolution., The discrepancies underscore the importance of spectroscopy for understanding structure in the universe and its evolution.
+ As we complete HeetoMAP we will extend this test: it is reminiscent of the davs of testing the Abell (Abell. Corwin Olowin 1989) catalog where many of the least rich svstenis are also superpositions.," As we complete HectoMAP we will extend this test; it is reminiscent of the days of testing the Abell (Abell, Corwin Olowin 1989) catalog where many of the least rich systems are also superpositions."
+ Perhaps the most striking features of the HIectoMAP clisplay are the structures that appear to cross (the survey roughly perpendicular (ο the line-of-sight at several recshilts., Perhaps the most striking features of the HectoMAP display are the structures that appear to cross the survey roughly perpendicular to the line-of-sight at several redshifts.
+ Because (he survey is incomplete. we cannot apply an objective measure of these structures (and features in other orientations) at this stage.," Because the survey is incomplete, we cannot apply an objective measure of these structures (and features in other orientations) at this stage."
+ Characteristics of the LectoNIAP walls, Characteristics of the HectoMAP walls
+This workshop was the first mecting devoted exclusively to “Cosmic Flows” since the conference held at the Institut d'Astroplivsique cde Paris exactly 6 νους carlicr.,This workshop was the first meeting devoted exclusively to “Cosmic Flows” since the conference held at the Institut d'Astrophysique de Paris exactly 6 years earlier.
+ The lnajor issucs addressed in Victoria were sunmiarized by Strauss in his introductory review: One of the long-taudiug goals of observational cosmolosv has been to identify a für sample of the Universe., The major issues addressed in Victoria were summarized by Strauss in his introductory review: One of the long-standing goals of observational cosmology has been to identify a fair sample of the Universe.
+ Within the context of cosmic flows. this amounts to identifving the depth at which the local Universe is at rest with respect to the CAIB. and thus contains all of the luajor attractors responsible for the 600 uunotiou of the Local Croup.," Within the context of cosmic flows, this amounts to identifying the depth at which the local Universe is at rest with respect to the CMB, and thus contains all of the major attractors responsible for the 600 motion of the Local Group."
+ In practice. this can be done by selecting a sample of galaxies for which distances are known and determimineg their uct bulk peculiar velocity im the CMD frame.," In practice, this can be done by selecting a sample of galaxies for which distances are known and determining their net bulk peculiar velocity in the CMB frame."
+ Comine into this workshop. the situation to a depth of <0hEATpe was fanly clear: existing peculiar velocity catalogs indicated that the CALB-frame bulk motion of this volume is ~3004+1001ns1.," Coming into this workshop, the situation to a depth of $\lesssim 50
+\hmpc$ was fairly clear: existing peculiar velocity catalogs indicated that the CMB-frame bulk motion of this volume is $\sim 300\pm100
+\kms$."
+" A tidal analysis indicates that amost of this imotion is due to unasses bevond x50A1NIpe,"," A tidal analysis indicates that most of this motion is due to masses beyond $\lesssim
+50 \hmpc$."
+ However. neselv-presented observational results ou these scales favor the lower cud of this range. aud thus more local sources.," However, newly-presented observational results on these scales favor the lower end of this range, and thus more local sources."
+ Toury Dressler obtained precise distances to nearby (zz30 INIpc)) ealaxies with the Surface Brightucss Fluctuation. (SBF) technique., Tonry Dressler obtained precise distances to nearby $\lesssim 30$ ) galaxies with the Surface Brightness Fluctuation (SBF) technique.
+ They modeled the flow feld with Virgo aud a Creat Attractor (GÀ) and found a GÀ iufall velocity of 2894137 aat the position of the LC. (, They modeled the flow field with Virgo and a Great Attractor (GA) and found a GA infall velocity of $289\pm137$ at the position of the LG. (
+This is about half of that originally claimed by the 7 Samurai.),This is about half of that originally claimed by the 7 Samurai.)
+ Furthermore. Toury Dressler claim. ouly a simall (<200kimis 13 vesidual motion arising frou sources bevoud thei volune.," Furthermore, Tonry Dressler claim only a small $\lesssim 200 \kms$ ) residual motion arising from sources beyond their volume."
+" The ENEAR project. a Fundamental Plane (FP)-based ""sou-of 7-Sammral survey. fiuds somewhat intermediate results: a bulk flow of 313553 aat a depth of ~505+Mpe (Weener)."," The ENEAR project, a Fundamental Plane (FP)-based `son-of 7-Samurai' survey, finds somewhat intermediate results: a bulk flow of $343\pm53$ at a depth of $\sim 50 \hmpc$ (Wegner)."
+ The “Shellfow project (Courteau) using the Tull-Fisher (TF) distance indicator. found that galaxies in an all-sky shell ceutered ou 55 hhave a low bulk motion of 80+150huis|.," The `Shellflow' project (Courteau) using the Tully-Fisher (TF) distance indicator, found that galaxies in an all-sky shell centered on $55$ have a low bulk motion of $80 \pm 150\kms$."
+ Taken together these new results sugeest that most of the mass responsible for the motion of the LG les within 60IMpe. aud eau presumably be identified with the CA.," Taken together these new results suggest that most of the mass responsible for the motion of the LG lies within 60, and can presumably be identified with the GA."
+ Woudt showed results from galaxy redshift survevs in the Zone of Avoidance which indicate zu excess of galaxies in the GA xeeion centered on Abell 3627 atV, Woudt showed results from galaxy redshift surveys in the Zone of Avoidance which indicate an excess of galaxies in the GA region centered on Abell 3627 at.
+ A3627 appears to have the richness of Coma. aud its position coincides with the peak of the CA as inferred frou POTENT recoustructious.," A3627 appears to have the richness of Coma, and its position coincides with the peak of the GA as inferred from POTENT reconstructions."
+ The observational situation on larger (>60h!Mpe) scales remains complicated., The observational situation on larger $>60 \hmpc$ ) scales remains complicated.
+ At the 793 Paris iectine. the controversial results of Lauer Posti (LP) were dissected.," At the '93 Paris meeting, the controversial results of Lauer Postman (LP) were dissected."
+ Iu. Victoria. a umber of recent results probing simular scales were presented: these were based mainly on sparse samples of ealaxy clusters.," In Victoria, a number of recent results probing similar scales were presented; these were based mainly on sparse samples of galaxy clusters."
+ The FP-based SMAC sample (Smith) aud the TF-based LPLOk sample (Willick) both find a large bulk flow. 600+200kinsE (but in a direction different from LP). whereas the TF-based SCT|ID (Caovanelli Dale) sample vields a very low bulk motion (~100+ 5)," The FP-based SMAC sample (Smith) and the TF-based LP10k sample (Willick) both find a large bulk flow, $\sim 600\pm200 \kms$ (but in a direction different from LP), whereas the TF-based SCI+II (Giovanelli Dale) sample yields a very low bulk motion $\sim 100 \pm 100$ )."
+ The nearby SNIa (Riess) give au intermediate result ~270]au«|. in good directional agreemeut with SATAC/LPLOk.," The nearby SNIa (Riess) give an intermediate result $\sim 270 \kms$, in good directional agreement with SMAC/LP10k."
+ The EFAR survev (Colless) finds low peculiar motions. but is restricted to two regions in the sky which do uot sample well the directions of the SATAC/LPLOk flows.," The EFAR survey (Colless) finds low peculiar motions, but is restricted to two regions in the sky which do not sample well the directions of the SMAC/LP10k flows."
+ One nuelt wonder if these different eroups nieasure the same peculiar velocities for common clusters., One might wonder if these different groups measure the same peculiar velocities for common clusters.
+ The answer is a qualified yes: formally — giveu the large errors — the 4? is acceptable. but there may be a trend for some surveys have systematically lower peculiar velocitics than others (Lucey).," The answer is a qualified yes: formally – given the large errors – the $\chi^2$ is acceptable, but there may be a trend for some surveys have systematically lower peculiar velocities than others (Lucey)."
+ It is important to note that what is quoted is the bulk flow of a eiven sparse sample and not that of a volue., It is important to note that what is quoted is the bulk flow of a given sparse sample and not that of a volume.
+ I argued that eiven the sampling aud the presence of immternal flows which act like au extra noise term all of these surveys (with the exception of LP) are consistent with each other and with a mean larec-scale bulls flow of ~300£001s.1., I argued that given the sampling and the presence of internal flows — which act like an extra noise term — all of these surveys (with the exception of LP) are consistent with each other and with a mean large-scale bulk flow of $\sim 300-400 \kms$.
+ This is somewhat larger than the bulk motions obtained on smaller scales. but the difference max not be statistically siguificaut.," This is somewhat larger than the bulk motions obtained on smaller scales, but the difference may not be statistically significant."
+ Ultimately. larger and denser samples will be needed to deteruiue the contribution from the largest scales.," Ultimately, larger and denser samples will be needed to determine the contribution from the largest scales."
+In the simulation. the background magnetic field B is the flux rope magnetic field defined as Equations (1)) to (4)).,"In the simulation, the background magnetic field $\vec B$ is the flux rope magnetic field defined as Equations \ref{B_vec}) ) to \ref{rho_def}) )."
+ We consider a parallel-tvpe magnetic flux rope model with the tvpical values of Ry=0.1 AU and By=20 nT near the Earth., We consider a parallel-type magnetic flux rope model with the typical values of $R_0=0.1$ AU and $B_0=20$ nT near the Earth.
+ For a time dependence of [Iux rope radius aud magnetic field. we use the model by Shimazu Vandas (2002). which can very elleclively reproduce observations.," For a time dependence of flux rope radius and magnetic field, we use the model by Shimazu Vandas (2002), which can very effectively reproduce observations."
+" In this model. when the fIux rope is located near the Sun. for instance. at 0.1. AU from the Sun. A, and By will be 0.01 AU and 2.000 nT. Cosmic ray particles are traced by the Duneman-Doris Method."," In this model, when the flux rope is located near the Sun, for instance, at 0.1 AU from the Sun, $R_0$ and $B_0$ will be 0.01 AU and 2,000 nT, Cosmic ray particles are traced by the Buneman-Boris Method."
+ The number of particles AJ is 10.000.," The number of particles $M$ is 10,000."
+ All these particles are initially located at the flux rope edge with randomly distributed penetration angles., All these particles are initially located at the flux rope edge with randomly distributed penetration angles.
+ When a particle escapes [rom the edge. a new particle is sel up immediately at the edge and we trace its orbit.," When a particle escapes from the edge, a new particle is set up immediately at the edge and we trace its orbit."
+ A different set of the random munbers for magnetic field irregularities ave selected [or each Figure & depicts the simulated particles on the r—4? plane for a case where (he flux rope is located al 0.1 AU from the Sun aud the cosmicrav rigiditv is GO GV., A different set of the random numbers for magnetic field irregularities are selected for each Figure \ref{cospos_60GV010} depicts the simulated particles on the $r-\varphi$ plane for a case where the flux rope is located at 0.1 AU from the Sun and the cosmicray rigidity is 60 GV.
+ In this case. the parameter fj is 0.0669.," In this case, the parameter $f_0$ is 0.0669."
+ Solid and dashed circles in the figure show the flux rope edge ancl the boundary of the forbidden region shown in Figure 7.. respectively.," Solid and dashed circles in the figure show the flux rope edge and the boundary of the forbidden region shown in Figure \ref{penet}, respectively."
+ The grav ny. points are cosnic rav particles., The gray tiny points are cosmic ray particles.
+" Panels a). b). c). and d) are snapshots for the time /= 2.5. 250. 750. and 2.500 i,' where uy is the proton evclotron frequency («y=0.479 ! in this study)."," Panels a), b), c), and d) are snapshots for the time $t=$ 2.5, 250, 750, and 2,500 $\omega_p^{-1}$, where $\omega_p$ is the proton cyclotron frequency $\omega_p=0.479$ $^{-1}$ in this study)."
+ Panel a) shows that almost all cosmic rav particles are located. outside of the forbidden region and the effect of magnetic field irregularities is small., Panel a) shows that almost all cosmic ray particles are located outside of the forbidden region and the effect of magnetic field irregularities is small.
+ ILowever. after approximately nine minutes later. in Panel b). the effect of magnetic field irregularities clearly appears. ancl a large number of particles penetrate into the forbidden region in the [lux rope.," However, after approximately nine minutes later, in Panel b), the effect of magnetic field irregularities clearly appears, and a large number of particles penetrate into the forbidden region in the flux rope."
+ In Panel c). some particles reach the center of the flux rope.," In Panel c), some particles reach the center of the flux rope."
+ Ii Panel d). cosmic ταν particle distribution is nearly uniform.," In Panel d), cosmic ray particle distribution is nearly uniform."
+ Penetration into the forbidden region is due not to drift motion but to the diffusion process caused by particle scattering and magnetic field line random walk (Jokipii Parker Figure 9 shows a time evolution of a mean radial distance of GO GV cosmic ray. particles in the (vpical flix rope. whose parameters are Ry=0.1 AU and By=20 nT near the Earth.," Penetration into the forbidden region is due not to drift motion but to the diffusion process caused by particle scattering and magnetic field line random walk (Jokipii Parker Figure \ref{rm_sim} shows a time evolution of a mean radial distance of 60 GV cosmic ray particles in the typical flux rope, whose parameters are $R_0=0.1$ AU and $B_0=20$ nT near the Earth."
+ The mean radial distance at time / is caletlated as where [ds simulation time step number until time / and pj;1j is a distance from the [ας rope axis of (he 7th cosmic rav. particle at the simulation (me step j., The mean radial distance at time $t$ is calculated as where $k$ is simulation time step number until time $t$ and $\rho_{ij}$ is a distance from the flux rope axis of the $i$ th cosmic ray particle at the simulation time step $j$.
+ In this definition. the entire simulation time step up to time / is used to caleulate (p).," In this definition, the entire simulation time step up to time $t$ is used to calculate $\left<\rho\right>$ ."
+ This definition is based on a continuous Cosmic rav penetration from the outside of the flux rope., This definition is based on a continuous cosmic ray penetration from the outside of the flux rope.
+ The thin dotted, The thin dotted
+"Compton parameter of electrons with 77, is given by where e;%i&p-2τι ityabKNCX_Me» Sy. dnd⋅ cxumm(p-2y(p-D if bxLEN—op.","Compton parameter of electrons with $\gamma_\ob'$ is given by where $c_7\approx{3\over8}(p-2)/(p-1)$ if $\nukn_\ob<\nu_c\ll\nu_m$ , and $c_8\approx(p-2)/(p-1)$ if $\nu_c\ll\nukn_\ob<\nu_m$."
+" The case of i4,«pEN is neglected provided the condition of eq. (19)).", The case of $\nu_m<\nukn_\ob$ is neglected provided the condition of eq. \ref{eq:cond1}) ).
+ Thus with helps of eqs. CLE. (13)). (16))," Thus with helps of eqs. \ref{eq:num}) ), \ref{eq:nuc}) ), \ref{eq:oob}) )"
+ and (25)). we now can calculate Ya.," and \ref{eq:fastcoolYm}) ), we now can calculate $Y_\ob$ ."
+" In the case of high eg value. ep=107. with the other parameters being Fs,=f1ng64I. we get. see Appendix. Here the redshift effect will add a term (1+2)* on the r.h.s.. We have also calculated the ορ value in a while parameter space of eg=1072—107 and n=107—em? for both cases of CE.r)=(IOerg.1098) and CIOerg.1075). with the results presented in Table I.."," In the case of high $\eps_B$ value, $\eps_{B}=10^{-2}$, with the other parameters being $E_{54}=t_3=n_0=\eps_{e,-1}=1$, we get, see Appendix, Here the redshift effect will add a term $(1+z)^{-5/8}$ on the r.h.s.. We have also calculated the $Y_\ob$ value in a while parameter space of $\eps_B=10^{-5}-10^{-2}$ and $n=10^{-2}-10^2\rm cm^{-3}$ for both cases of $(E,t)=(10^{54}{\rm erg},10^3{\rm s})$ and $(10^{53}{\rm erg},10^2{\rm s})$, with the results presented in Table \ref{tab}."
+ For LAT energy range. we should take ο 2100 MeV in eq.(7)). which. then. gives a lower limit to the upstream field. Note redshift effect is added hereinafter. and Y; has been derived in refsec:KN..," For LAT energy range, we should take $h\nu_\ob=$ 100 MeV in \ref{eq:iclimitenergy}) ), which, then, gives a lower limit to the upstream field, Note redshift effect is added hereinafter, and $Y_\ob$ has been derived in \\ref{sec:KN}."
+" Plugging the Ya, values calculated in the cases of (Ε.Ε)=CIO""erg.107) and (0erg.1073). into eq.28)). it is easy to find that the constraint by the latter case is less stringent. Le.. the value of lower limit to B, is smaller."," Plugging the $Y_\ob$ values calculated in the cases of $(E,t)=(10^{54}{\rm erg},10^3{\rm s})$ and $(10^{53}{\rm erg},10^2{\rm s})$ into \ref{eq:Buic}) ), it is easy to find that the constraint by the latter case is less stringent, i.e., the value of lower limit to $B_u$ is smaller."
+ Moreover. in the former case. we have The result of more accurate calculation in eg=107 case. Yon(s=0)1.2 (see Appendix). has been used.," Moreover, in the former case, we have The result of more accurate calculation in $\eps_B=10^{-5}$ case, $Y_\ob(z=0)=1.2$ (see Appendix), has been used."
+" We see that the relatively relaxed lower limit to B, appears when taking the lowest allowed value. ej= 107, but the limit is in the order of mG. and insensitive to the ej value."," We see that the relatively relaxed lower limit to $B_u$ appears when taking the lowest allowed value, $\eps_B=10^{-5}$ , but the limit is in the order of mG, and insensitive to the $\eps_B$ value."
+ One may argue that the ej>10? assumption is motivated by X-ray afterglow observations on day scale. and may not hold at early time. ~10°s. that is addressed in this paper.," One may argue that the $\eps_B\ga10^{-5}$ assumption is motivated by X-ray afterglow observations on day scale, and may not hold at early time, $\sim10^3$ s, that is addressed in this paper."
+" Indeed. if relaxing the eg>107 assumption. the lower limit to B, does not change much. and only varies slowly with ep as; ~ερ16"," Indeed, if relaxing the $\eps_B>10^{-5}$ assumption, the lower limit to $B_u$ does not change much, and only varies slowly with $\eps_B$ as $\sim\eps_{B}^{1/6}$."
+" Furthermore. with /744,=100 MeV. eq.(10)) gives an upper limit to the upstream field. It can be seen that taking (Ε.Ε)=(IO""erg.1075) gives more stringent constraint than CE.£)2CIO""erg.1075)."," Furthermore, with $h\nu_\ob=100$ MeV, \ref{eq:synlimitenergy}) ) gives an upper limit to the upstream field, It can be seen that taking $(E,t)=(10^{54}{\rm erg},10^3{\rm s})$ gives more stringent constraint than $(E,t)=(10^{53}{\rm
+erg},10^2{\rm s})$."
+ As we expect ερS107. then we have the relatively relaxed constraint It is interesting to note that the constraints by IC and synchrotron/jitter cooling together limit the upstream field in a closed range that spans about two orders of magnitude.," As we expect $\eps_B\la10^{-2}$, then we have the relatively relaxed constraint It is interesting to note that the constraints by IC and synchrotron/jitter cooling together limit the upstream field in a closed range that spans about two orders of magnitude."
+ We have shown that 10°s-seale. >100 MeV GRB emission recently revealed by Fermi-LAT provides stringent constraints on the upstream magnetic field.," We have shown that $10^3$ s-scale, $>100$ MeV GRB emission recently revealed by Fermi-LAT provides stringent constraints on the upstream magnetic field."
+ A lower limit to. the magnetic field is obtained by requiring the acceleration time of electrons producing >100 MeV synchrotron photons to be shorter than their energy-loss time due to IC scattering the afterglow photons., A lower limit to the magnetic field is obtained by requiring the acceleration time of electrons producing $>100$ MeV synchrotron photons to be shorter than their energy-loss time due to IC scattering the afterglow photons.
+ By scanning the possible afterglow parameter space. the lower limit for the magnetic. field is given in eq. (29)).," By scanning the possible afterglow parameter space, the lower limit for the magnetic field is given in eq. \ref{eq:Blower}) )."
+ Interestingly. an upper limit to the magnetic field is also obtained by requiring the acceleration time to be shorter than the energy-loss time due to synchrotron/jitter radiation upstream.," Interestingly, an upper limit to the magnetic field is also obtained by requiring the acceleration time to be shorter than the energy-loss time due to synchrotron/jitter radiation upstream."
+ Given a maximum equipartition value of magnetic field downstream. the upper limit for the field is given in eq. (32)).," Given a maximum equipartition value of magnetic field downstream, the upper limit for the field is given in eq. \ref{eq:Bupper}) )."
+" Combining both lower and upper limits. the upstream magnetic field is limited in a closed range with two orders of magnitude uncertainty. 10?10275, mG. There should be another lower limit to B, by requiring that the acceleration timeof 100-MeV emitting electrons is shorter than the dynamical time of the afterglow shock."," Combining both lower and upper limits, the upstream magnetic field is limited in a closed range with two orders of magnitude uncertainty, $10^0n_0^{9/8}{\rm mG}\la B_u\la10^2n_0^{3/8}$ mG. There should be another lower limit to $B_u$ by requiring that the acceleration timeof 100-MeV emitting electrons is shorter than the dynamical time of the afterglow shock."
+" As LW06. we neglect this ""age limit” because the constraint is much less stringent than the one by IC cooling."," As LW06, we neglect this ""age limit"" because the constraint is much less stringent than the one by IC cooling."
+" Even assuming the ~10 GeV photons from the Fermi-LAT GRBs is produced by external shock synchrotron radiation. the constraint to B, by this argument. B,70.1 mG (Piran&Nakar2010).. is still much less stringent than the cooling limit here."," Even assuming the $\sim10$ GeV photons from the Fermi-LAT GRBs is produced by external shock synchrotron radiation, the constraint to $B_u$ by this argument, $B_u>0.1$ mG \citep{pn10}, is still much less stringent than the cooling limit here."
+ The lower limit to upstream field by Fermi-LAT observations is larger than the previous constraint using X-ray afterglow observations (LW06) by orders of magnitude., The lower limit to upstream field by Fermi-LAT observations is larger than the previous constraint using X-ray afterglow observations (LW06) by orders of magnitude.
+ It can be seen thatthis high amplitude field 1snot likely to be provided by a magnetized wind from the GRB progenitor (see discussion in LW06). therefore theonly reasonable origin of this field is due to magnetic field amplification upstream. most likely by the streaming of high energy shock accelerated particles (e.g..Couchetal.2008:Medvedev& 2010)..," It can be seen thatthis high amplitude field isnot likely to be provided by a magnetized wind from the GRB progenitor (see discussion in LW06), therefore theonly reasonable origin of this field is due to magnetic field amplification upstream, most likely by the streaming of high energy shock accelerated particles \citep[e.g.,][]{Couch08,Medvedev09,Pelletier09,Lemoine10,Niemiec10}. ."
+ However. compared to the common few g/G-scale interstellar medium field. this," However, compared to the common few $\mu$ G-scale interstellar medium field, this"
+the dispersion of virial mass as function of stellar mass to be relatively high (see Figure 2). and also. that the dispersion of the velocity dispersion with respect to virial mass. presents a different behaviour from that presented by the dispersion of the velocity dispersion with respect to stellar mass (see Figures 3 and 4).,"the dispersion of virial mass as function of stellar mass to be relatively high (see Figure 2), and also, that the dispersion of the velocity dispersion with respect to virial mass, presents a different behaviour from that presented by the dispersion of the velocity dispersion with respect to stellar mass (see Figures 3 and 4)."
+ In what follows we shall analyse each one of these possibilities: Following what we have expressed above. cases 3 and 4 would play an important role in the behaviour of the intrinsic dispersion of these masses with respect to the velocity dispersion.," In what follows we shall analyse each one of these possibilities: Following what we have expressed above, cases 3 and 4 would play an important role in the behaviour of the intrinsic dispersion of these masses with respect to the velocity dispersion."
+ Neither one of these cases may be discarded a priori. we only know that the final result of the action that one or the other may cause is different.," Neither one of these cases may be discarded a priori, we only know that the final result of the action that one or the other may cause is different."
+ Which value. then. is the correct one for the mass?," Which value, then, is the correct one for the mass?"
+ or which one of equations 4 or 5 is the one that produces the appropriate results?, or which one of equations 4 or 5 is the one that produces the appropriate results?
+ The answer to this question is complicated and. at present. we do not have enough elements to elucidate as to what is the correct answer.," The answer to this question is complicated and, at present, we do not have enough elements to elucidate as to what is the correct answer."
+ As a consequence of this. we shall use both masses in order to investigate the universality of the FJR and that of the rest of the structural relations.," As a consequence of this, we shall use both masses in order to investigate the universality of the FJR and that of the rest of the structural relations."
+ An interesting point that comes up when comparing virial and stellar mass (see | above) is that the slope of the fit is approximately equal to 1., An interesting point that comes up when comparing virial and stellar mass (see 1 above) is that the slope of the fit is approximately equal to 1.
+ This result is different from those of recent papers (Cappellarietal.2006:; Robertsonetal.2006:: Koopmans&Treu2010)) in which they find that the slope ts up to larger. meaning that within the virial radius there ts up to a of the total mass in the form of dark matter.," This result is different from those of recent papers \cite{cap06}; ; \cite{rob06}; \cite{koo10}) ) in which they find that the slope is up to larger, meaning that within the virial radius there is up to a of the total mass in the form of dark matter."
+ What we find in this paper contradicts this result and shows that the amount of dark matter within the virial radius is negligible., What we find in this paper contradicts this result and shows that the amount of dark matter within the virial radius is negligible.
+ We consider that. given its relevance. this result should be dealt with in a profound manner which we will do in a forthcoming paper. since this topic is beyond the scope of the present paper.," We consider that, given its relevance, this result should be dealt with in a profound manner which we will do in a forthcoming paper, since this topic is beyond the scope of the present paper."
+ An interesting property of the SDSS sample used in this work is that it contains galaxies in the redshift range 0.01 0.35. so that. we can study the behaviour of the structural properties as function. of redshift.," An interesting property of the SDSS sample used in this work is that it contains galaxies in the redshift range $0.01 < z < 0.35$ , so that, we can study the behaviour of the structural properties as function of redshift."
+ In this case. and in order to avoid the geometrical effect. we can study the behaviour of the intrinsic dispersion of the FJR in two ways: the first one consists in considering samples of galaxies of approximately equal magnitude but with different redshift. and the second in considering samples of galaxies with approximately equal mass but with different redshift.," In this case, and in order to avoid the geometrical effect, we can study the behaviour of the intrinsic dispersion of the FJR in two ways; the first one consists in considering samples of galaxies of approximately equal magnitude but with different redshift, and the second in considering samples of galaxies with approximately equal mass but with different redshift."
+ Figure 8 shows the behaviour of the intrinsic dispersion of the FJR as function of the redshift., Figure 8 shows the behaviour of the intrinsic dispersion of the FJR as function of the redshift.
+ Each symbol and colour represent à sample of galaxies of approximately constant magnitude., Each symbol and colour represent a sample of galaxies of approximately constant magnitude.
+ In some cases the errors and the narrow redshifts ranges where the samples are distributed make it difficult to observe any trend., In some cases the errors and the narrow redshifts ranges where the samples are distributed make it difficult to observe any trend.
+" However. for M,~—22.0 and M,~-23.0 we have data of relatively good quality so we are able to see a smooth trend (see Figure 9)."," However, for $M_{g} \sim -22.0$ and $M_{g} \sim -23.0$ we have data of relatively good quality so we are able to see a smooth trend (see Figure 9)."
+ The run test confirm an underlying trend with a confidence level of approximately905c., The run test confirm an underlying trend with a confidence level of approximately.
+". In order to characterise the behaviour of the intrinsic dispersion as a function of the redshift. in Figure 9 we have fitted the points of the sample with M,~—22 to a straight line whose equation Is: Similarly the straight line equation obtained from. the sample with M,~—23 is: For the mass case. in Figure 10 we can see the behaviour of the intrinsic dispersion of the FJR as function of the redshift."," In order to characterise the behaviour of the intrinsic dispersion as a function of the redshift, in Figure 9 we have fitted the points of the sample with $M_{g} \sim -22$ to a straight line whose equation is: Similarly the straight line equation obtained from the sample with $M_{g} \sim -23$ is: For the mass case, in Figure 10 we can see the behaviour of the intrinsic dispersion of the FJR as function of the redshift."
+ In this graph each symbol and colour represent a sample of galaxies with approximately constant virial mass., In this graph each symbol and colour represent a sample of galaxies with approximately constant virial mass.
+ In this case it also happens that the errors and the narrow redshifts ranges where some samples are distributed make it difficult to observe any trend., In this case it also happens that the errors and the narrow redshifts ranges where some samples are distributed make it difficult to observe any trend.
+ However. for the samples with logCMagia/M&)~11.3 and log(Myiria/M&)~11.8. the data are of sufficiently good quality so that here we cansee a clear," However, for the samples with ${\bf M_{virial}}/{\bf M_{\odot}}) \sim 11.3$ and ${\bf M_{virial}}/{\bf M_{\odot}}) \sim 11.8$, the data are of sufficiently good quality so that here we cansee a clear"
+strong lines (AAL12 49022. 5016. 5876) appear blended with Ie or Si lines belonging1. to the companions. making them less useful for the analysis of the line profile variability.,"strong lines $\lambda\lambda$ 4121, 4922, 5016, 5876) appear blended with Fe or Si lines belonging to the companions, making them less useful for the analysis of the line profile variability."
+ For star D the spectral line profiles of and appear stronglv variable. as presented in Figs.22 and 3.," For star B the spectral line profiles of and appear strongly variable, as presented in 2 and 3."
+ Usually. magnesium does not present large. cquivalent width variations in magnetic peculiar stars1997).," Usually, magnesium does not present large equivalent width variations in magnetic peculiar stars."
+.. However. line profile. variability of the magnesium doublet a ΑΕΙ. has been reported. by(1999)... whe» published. the first Mg map or an Ap star (CUVir).," However, line profile variability of the magnesium doublet at $\lambda$ 4481 has been reported by, who published the first Mg map for an Ap star Vir)."
+ Ehey found that. even. without showing important. equivalen width variations. AJ4SI exhibits notorious profile variations. similar o those observed. in VVcl for the components A and D. The profile shape sugecss that Ale is more abundant in the region lacing the com»»ion star.," They found that, even without showing important equivalent width variations, $\lambda$ 4481 exhibits notorious profile variations, similar to those observed in Vel for the components A and B. The profile shape suggests that Mg is more abundant in the region facing the companion star."
+ A more complete study. of line. profiles. that would allow to determine the surface distribution. of chenvical elements. will require a larger number of spectra obt:üned at different: phases.," A more complete study of line profiles, that would allow to determine the surface distribution of chemical elements, will require a larger number of spectra obtained at different phases."
+" We can assume that the| stellar rotation in the short-period eccentric binary AB is svnchronized with the angular velocity at. periastron. so that the rotational period dlillers from the orbital period by a factor 1.163. (Po,=1.5846 d. Poa=1.38622 d)."," We can assume that the stellar rotation in the short-period eccentric binary AB is synchronized with the angular velocity at periastron, so that the rotational period differs from the orbital period by a factor 1.163 $P_{orb}=1.5846$ d, $P_{rot}=1.3622$ d)."
+ This fact implies that the stellar surface visible at a given orbital phase varies with time. completing one evele in 9.7 days. which corresponds to about 6 orbital Cycles.," This fact implies that the stellar surface visible at a given orbital phase varies with time, completing one cycle in 9.7 days, which corresponds to about 6 orbital cycles."
+ A multi-epoch study of Vel would be of great interest to prove whether the location of stellar spots is fixed on the star surface or is related to the position of the stellar companion., A multi-epoch study of Vel would be of great interest to prove whether the location of stellar spots is fixed on the star surface or is related to the position of the stellar companion.
+ In our spectra of the C and D components the most interesting. and rather unexpected fact is the presence of spectral lines tvpical of Lle\In stars.," In our spectra of the C and D components the most interesting, and rather unexpected fact is the presence of spectral lines typical of HgMn stars."
+ The line A3984 is present in both stars € and D. while strong and lines are present in the spectrum of component D. Although comparatively weak. some lines have been also identified in both components.," The line $\lambda$ 3984 is present in both stars C and D, while strong and lines are present in the spectrum of component D. Although comparatively weak, some lines have been also identified in both components."
+ In the reconstructed spectrum of star D we identified lines at 4046 and 4514 and more than two dozens lines., In the reconstructed spectrum of star D we identified lines at 4046 and 4514 and more than two dozens lines.
+ No noticeable lines ofextscli..extscii..extscii.. extiscii.. extscii.. were detected in stars € or D. Phe measured central wavelength of line at 3984. and for stars C ancl D. respectively. indicates that heavy Lg isotopes are predominant in these components.," No noticeable lines of, were detected in stars C or D. The measured central wavelength of line at 3984, and for stars C and D, respectively, indicates that heavy Hg isotopes are predominant in these components."
+ We will return to this point in the next section., We will return to this point in the next section.
+ The weakness of spectral lines for the two less massive stars makes impossible to study. line profile variations in their spectra., The weakness of spectral lines for the two less massive stars makes impossible to study line profile variations in their spectra.
+ The abundances were determined with the synthetic spectrum. method. using ALLASO model atmospheres and the SYNTPLILE code1993)., The abundances were determined with the synthetic spectrum method using ATLAS9 model atmospheres and the SYNTHE code.
+. Stellar temperatures of companions C and D were estimated from the correlation of excitation potentials of Fe lines with abundances2008)., Stellar temperatures of companions C and D were estimated from the correlation of excitation potentials of Fe lines with abundances.
+. Alicroturbulent velocities were derived from the correlation of equivalent widths of Fe lines with abundances., Microturbulent velocities were derived from the correlation of equivalent widths of Fe lines with abundances.
+ In case of stars A and. D. however. it was not. possible to apply the same procedure since Le lines are weaker and they appear broader because of the higher rotation.," In case of stars A and B, however, it was not possible to apply the same procedure since Fe lines are weaker and they appear broader because of the higher rotation."
+ Estimatecl temperatures for these stars were derived. from the observed masses and radii: interpolating in the grid. of Geneva theoretical stellar models for solar composition.," Estimated temperatures for these stars were derived from the observed masses and radii, interpolating in the grid of Geneva theoretical stellar models for solar composition."
+ 44 shows the position of stars A and D in the mass-ratio diagram., 4 shows the position of stars A and B in the mass-ratio diagram.
+ Both stars are locate close the ZAAIS., Both stars are located close the ZAMS.
+" As is shown in this figure. the masses ancl radii determined from light ancl RY curves corresponc to temperatures of 139004500Why and 132004450 for components A ancl Bo respectively,"," As is shown in this figure, the masses and radii determined from light and RV curves correspond to temperatures of $\pm$ K and $\pm$ K for components A and B, respectively."
+ Surface. gravity KNwere calculated: directly from masses and radii., Surface gravity were calculated directly from masses and radii.
+ Phe adopte atmospheric pareumeters are listed in 233, The adopted atmospheric parameters are listed in 3.
+ The results of the abundance analysis are presented in 44., The results of the abundance analysis are presented in 4.
+ For comparison we list in the last column the solar, For comparison we list in the last column the solar
+behind an altered distribution function that created the conditions for a new instability.,behind an altered distribution function that created the conditions for a new instability.
+" described in more detail that each spiral pattern seems to be a strongly unstable true mode of the disc in which it grows; it lasts for only a few rotations, changing the disc properties as it decays so as to provoke new instabilities of the altered disc having different pattern speeds."," \cite{Sell00}
+ described in more detail that each spiral pattern seems to be a strongly unstable true mode of the disc in which it grows; it lasts for only a few rotations, changing the disc properties as it decays so as to provoke new instabilities of the altered disc having different pattern speeds."
+ The discovery of resonance scattering among the stars of the solar neighbourhood suggests that this mechanism may occur in real galaxies., The discovery of resonance scattering among the stars of the solar neighbourhood \citep{Sell10b} suggests that this mechanism may occur in real galaxies.
+ I am actively engaged in addressing the many details of this picture that remain unclear., I am actively engaged in addressing the many details of this picture that remain unclear.
+ Simple models of disc galaxies support many global linear instabilities (e.g.," Simple models of disc galaxies support many global linear instabilities \citep[\eg][]{Kaln78,Jala07}."
+" The bar-forming mode is generally the fastest ∙∙growing, but has almost no spirality."," The bar-forming mode is generally the fastest growing, but has almost no spirality."
+" These studies are therefore important to understand stability, but do not appear promising for spiral generation."," These studies are therefore important to understand stability, but do not appear promising for spiral generation."
+" 'The *density wave"" theory for spiral modes, described in detail by (1996). invokes a more specific galaxy model with a cool outer disc and hot inner disc."," The “density wave” theory for spiral modes, described in detail by \cite{BL96}, invokes a more specific galaxy model with a cool outer disc and hot inner disc."
+" The local stability parameter, Q=or/oR,crit (Toomred[1964).,, is postulated to be Q=1 in the outer disc and to rise steeply to Q>2 near the centre."," The local stability parameter, $Q = \sigma_R
+/ \sigma_{R,\rm crit}$ \citep{Toom64}, is postulated to be $Q \ga 1$ in the outer disc and to rise steeply to $Q > 2$ near the centre."
+" Under these specific conditions, Bertin Lin find slowly evolving spiral modes that grow by their mechanism."," Under these specific conditions, Bertin Lin find slowly evolving spiral modes that grow by their \citep{Mark77}
+ mechanism."
+" They invoke shocks in the gas to limit the amplitude of the slowly growing mode, leading to a quasi-steady global spiral pattern."," They invoke shocks in the gas to limit the amplitude of the slowly growing mode, leading to a quasi-steady global spiral pattern."
+ present a model of this kind to account for the spiral structure of M81., \cite{Lowe94} present a model of this kind to account for the spiral structure of M81.
+" The principal objection to their picture is that it is likely that an outer disc with such a low Q will support other, more vigorous, collective responses that will quickly alter the background state by heating the outer disc, as I now show."," The principal objection to their picture is that it is likely that an outer disc with such a low $Q$ will support other, more vigorous, collective responses that will quickly alter the background state by heating the outer disc, as I now show."
+ Here I report simulations designed to test directly whether a galaxy model of the kind invoked by hereafterBLLT) can in fact survive to support the global mode they predict should dominate., Here I report simulations designed to test directly whether a galaxy model of the kind invoked by \cite[][hereafter BLLT]{Bert89} can in fact survive to support the slowly-growing global mode they predict should dominate.
+" The model I adopt here is from the A; survey of BLLT, which is a minor variant of that proposed by (1980)."," The model I adopt here is from the $_1$ survey of BLLT, which is a minor variant of that proposed by \cite{FE80}."
+" The surface density of the disc is with Here R4 and Mz, are respectively the scale length and mass, when A—0, of the unmodified exponential disc."," The surface density of the disc is with Here $R_d$ and $M_d$ are respectively the scale length and mass, when $\Delta=0$, of the unmodified exponential disc."
+ The f(y) factor creates a central dip in the surface density of approximate radius F4/2., The $f(y)$ factor creates a central dip in the surface density of approximate radius $R_d/2$.
+ The rotation curve of their model has the form that is independent of the disc contribution., The rotation curve of their model has the form that is independent of the disc contribution.
+" Because mass removed by the central cutout is effectively replaced by rigid matter, the model can be thought of as including a small central bulge as well as a pseudo-isothermal halo."," Because mass removed by the central cutout is effectively replaced by rigid matter, the model can be thought of as including a small central bulge as well as a pseudo-isothermal halo."
+" Here I adopt units such that G=MaRa1, and also Vo= 1; the orbit period at R=2Rq is 19 time units."," Here I adopt units such that $G=M_d=R_d=1$, and also $V_0=1$ ; the orbit period at $R=2R_d$ is $\sim13$ time units."
+ BLLT also specify the radial variation of Toomre’s local axisymmetric stability parameter Q(R)=Qop(1+/ 1.5e2Ra). Q=2.5Qop atij the centre and decreases rapidly to Q=Φου over most of the outer disc., BLLT also specify the radial variation of Toomre's local axisymmetric stability parameter $Q(R) = Q_{\rm OD}(1 + 1.5e^{2R/R_d})$ ; $Q=2.5Q_{\rm OD}$ at the centre and decreases rapidly to $Q=Q_{\rm OD}$ over most of the outer disc.
+" Because BLLT do not supply stellar dynamical distribution functions, I set the initial velocities to create a rotationally supported disc having the desired Q profile, using the Jeans equations in the epicyclic approximation to determine the azimuthal dispersion and asymmetric drift (BT08, eq."," Because BLLT do not supply stellar dynamical distribution functions, I set the initial velocities to create a rotationally supported disc having the desired $Q$ profile, using the Jeans equations in the epicyclic approximation to determine the azimuthal dispersion and asymmetric drift (BT08, eq."
+ 4.228)., 4.228).
+ I have chosen the case with Qop—1 and A— —0.35., I have chosen the case with $Q_{\rm OD}=1$ and $\Delta = -0.35$ .
+ The upper panel of Fig., The upper panel of Fig.
+" [i] shows the rotation curve, which indicates that this model has a strongly sub-maximal disc tthe disc contributes merely ~25% of the central attraction at R=2Ra4."," \ref{model} shows the rotation curve, which indicates that this model has a strongly sub-maximal disc \citep{Sack97}, the disc contributes merely $\sim25\%$ of the central attraction at $R=2R_d$."
+ The lower panel shows the initial Q-profile., The lower panel shows the initial $Q$ -profile.
+" Because the disc is so light, and the disc surface density dips towards the centre, the initial random velocities are small enough that the epicycle approximationis adequate and the Jeans equationslead to a quite respectable equilibrium model."," Because the disc is so light, and the disc surface density dips towards the centre, the initial random velocities are small enough that the epicycle approximationis adequate and the Jeans equationslead to a quite respectable equilibrium model."
+" BLLT predict that this model has the slowly-growing, tightly-wrapped, bi-symmetric spiral mode illustrated"," BLLT predict that this model has the slowly-growing, tightly-wrapped, bi-symmetric spiral mode illustrated"
+convolved with a gaussian beam with the Full Width Half Maximum (FWHM) of beams (i.e. 33” at 30 GHz and 10’ at 100 GHz) by the pipeline simulator.,convolved with a gaussian beam with the Full Width Half Maximum (FWHM) of beams (i.e. $'$ at 30 GHz and $'$ at 100 GHz) by the pipeline simulator.
+ The code includes the main properties of the payload (e.g. the boresight angle. the scanning strategy) and of the considered receiver (e.g. the beam location on the focal plane. its FWHM. noise properties).," The code includes the main properties of the payload (e.g. the boresight angle, the scanning strategy) and of the considered receiver (e.g. the beam location on the focal plane, its FWHM, noise properties)."
+" The scanning strategy 1s such that every detector observes the same ""ring"" in the sky for an hour before repointing: in our work we average the 60 ]-minute observations of the same sky-ring. considering. |-hour data streams: we refer to these observations as “circles”."," The scanning strategy is such that every detector observes the same “ring"" in the sky for an hour before repointing; in our work we average the 60 1-minute observations of the same sky-ring, considering 1-hour data streams; we refer to these observations as “circles""."
+ The FWHM of the antenna divides every circle (/) in pixels (i)., The FWHM of the antenna divides every circle $j$ ) in pixels $i$ ).
+ The radiometer output is where T° is the observed sky. 2 is a noise term and Go is the true value of the calibration constant.," The radiometer output is, where $T^{\rm obs}$ is the observed sky, $n$ is a noise term and $G_0$ is the true value of the calibration constant."
+" In the noise term we only considered white noise. 1.9. a random Gaussian distribution with rms where Το is the system temperatureof the receiver. Av is the bandwidth (Av/v~ 20%)) and r is the considered integration time: the noise amplitude 07,5, i$ quite constant during the mission."," In the noise term we only considered white noise, i.e. a random Gaussian distribution with rms, where $T_{sys}$ is the system temperatureof the receiver, $\Delta\nu$ is the bandwidth $\Delta\nu/\nu\sim$ ) and $\tau$ is the considered integration time; the noise amplitude $\delta T_{rms}$ is quite constant during the mission."
+ In this case. the instrumental uncertainty on output differences (see Eq. 3))," In this case, the instrumental uncertainty on output differences (see Eq. \ref{eq:sigmagoverg}) )"
+" is: where 07, and 07> are the rms values in the two observed points of the sky."," is:, where $\delta T_1$ and $\delta T_2$ are the rms values in the two observed points of the sky."
+ By averaging 60 rings in one circle we assume that the noise of each ring is uncorrelated with the others. which ts a reasonable assumption for the low 1/f noise expected from the instrument (Mainoetal. 1999)).," By averaging 60 rings in one circle we assume that the noise of each ring is uncorrelated with the others, which is a reasonable assumption for the low $1/f$ noise expected from the instrument \cite{oof}) )."
+ Secondly. we need to determine the pixel-pairs for calibration.," Secondly, we need to determine the pixel-pairs for calibration."
+ If the number of pixels in a circle is N. then N-| is the number of independent pixel-pairs: we must choose them among the N(N—1)/2 possible pairs.," If the number of pixels in a circle is $N$, then $N-1$ is the number of independent pixel-pairs; we must choose them among the $N(N-1)/2$ possible pairs."
+ Different criteria for choosing the N—| independent pixel pairs were considered: the results obtained showed no significant changes., Different criteria for choosing the $N-1$ independent pixel pairs were considered: the results obtained showed no significant changes.
+ The simplest approach was then selected: pairs are formed with opposite pixels in a circle (Fig. 2)):, The simplest approach was then selected: pairs are formed with opposite pixels in a circle (Fig. \ref{fig:pairs}) ):
+ In this way every pixel is considered twice: finally the pair with the lowest signal difference is not considered. thus leaving exactly N—| pairs.," In this way every pixel is considered twice; finally the pair with the lowest signal difference is not considered, thus leaving exactly $N-1$ pairs."
+ The & parameter in the y function (Eq. 4)), The $k$ parameter in the $\chisq$ function (Eq. \ref{eq:chi2}) )
+ becomes in this case k=(m.j). where m=(G.i) refers to pixel-pairs previously defined and / refers to the considered circle.," becomes in this case $k=(m,j)$, where $m=(i,i')$ refers to pixel-pairs previously defined and $j$ refers to the considered circle."
+ In our simulations the G factor is computed over two different time-scales: | hour and | day., In our simulations the $G$ factor is computed over two different time-scales: 1 hour and 1 day.
+ Both are “natural” time scales forPLANcK: the first defines the scan “circles”while the second is the time scale of Earth-satellite communications.," Both are “natural"" time scales for: the first defines the scan “circles”while the second is the time scale of Earth-satellite communications."
+ Thus. the sum in Eq.," Thus, the sum in Eq."
+ 4. is extended to the N—I independent pixel pairs of one circle in the first case (fixed / value). while is extended to the (N—1)-24 circles covered in one day in the second case.," \ref{eq:chi2} is extended to the $N-1$ independent pixel pairs of one circle in the first case (fixed $j$ value), while is extended to the $(N-1)\cdot 24$ circles covered in one day in the second case."
+ We present the analysis carried out at 30 and 100 GHz., We present the analysis carried out at 30 and 100 GHz.
+ The first is the channel where foreground emission is the strongest while the second one has the highest noise level per single radiometer (for 30 GHz radiometers Z44= LOK. for 100 GHz Τις=45K): they represent the worst case for systematic and. statistical error behaviour respectively.," The first is the channel where foreground emission is the strongest while the second one has the highest noise level per single radiometer (for 30 GHz radiometers $T_{sys} \simeq 10$ K, for 100 GHz $T_{sys} \simeq 45$ K): they represent the worst case for systematic and statistical error behaviour respectively."
+ The other frequencies are in intermediate situations., The other frequencies are in intermediate situations.
+ We applied the 2-step procedures described in Sect. 4:: , We applied the 2-step procedures described in Sect. \ref{sect:methodxerr}; ;
+as already noted. the same noise realization has to be used for both steps.," as already noted, the same noise realization has to be used for both steps."
+ In the case of radiometers. since the noise rms value has a slight dependence on the observed sky temperature (see Eq. 17)).," In the case of radiometers, since the noise rms value has a slight dependence on the observed sky temperature (see Eq. \ref{eq:rms_noise}) ),"
+ considering the same realization for both simulation introduces an error of afewon signal terms (Eq. 16))., considering the same realization for both simulation introduces an error of a fewon signal terms (Eq. \ref{eq:signal}) ).
+ Similar results hold for all the radiometers at the same frequency., Similar results hold for all the radiometers at the same frequency.
+" At 30 GHz we have considered the radiometer with the beam position (Gp=5.267. dp= 52.9"") in the focal plane and angular resolution of 33’ that divides every observed circle in 1950pixels?:: at 100 GHz we used the radiometer LFI- with the beam position (G5= 2.937. @p= 0.07) and the angular resolution of 10’ that divides every observed circle in 6498 pixels?."," At 30 GHz we have considered the radiometer LFI-28 with the beam position $\theta_{B}=5.26^{\circ}$, $\phi_{B}=52.9^{\circ}$ ) in the focal plane and angular resolution of $33'$ that divides every observed circle in 1950; at 100 GHz we used the radiometer LFI-01 with the beam position $\theta_{B}=2.93^{\circ}$ , $\phi_{B}=0.0^{\circ}$ ) and the angular resolution of $10'$ that divides every observed circle in 6498 $^6$ ."
+ As deseribed in Sect. 4..," As described in Sect. \ref{sect:methodxerr}, ,"
+ we start by considering the CMB dipole only., we start by considering the CMB dipole only.
+ This allowsus to evaluate the systematic error due, This allowsus to evaluate the systematic error due
+" DI EEa0) den ↥⋅⋜↧↑↕∪∟∖↕∠≻⋀∖↕⋝≨∩∙∩⊇∙⊀≚≼↧↸∖↑⋜⊔↕↸∖≼↧≼∐∖↴∖↴↸⊳↥⋅∏≻⊓∪∐∪↕ tare,∖ in a nebula associated with the NGC 7129 star forming:. region. and is. at a distance. of) about .1250 beyarsecs.",$_\alpha$ 234 $_v$ =11.9mag) is a star embedded in a nebula associated with the NGC 7129 star forming region and is at a distance of about 1250 parsecs.
+ This is oue of the vouusestD kuown Ierbie? Be objects aud has an age ~ 0.1 million vears aud spectral type of D5 (Fucute et al., This is one of the youngest known Herbig Be objects and has an age $\sim$ 0.1 million years and spectral type of B5 (Fuente et al.
+ 2001)., 2001).
+ IIerbie Ac/Be stars are intermediate mass (2 to LOAL..) . ⋅↽⋅ ↴∐⋅↸∖≓⋯⋜⊔∐≓↴∖↴↸∖≺∣⋯∖∐∩∖↴∖↴↑⋜∐⋅↴∖↴⋖↕↕↸∖, Herbig Ae/Be stars are intermediate mass (2 to $_\odot$ ) pre-main-sequence stars (Herbig 1960).
+↥⋅↴⋝↕∶↴∙⊾↓≝⋊⊳↭⋅⊀≚↕∐∐⊔⋝↸∖↥⋅ of Herbie Ac stars have been found to possess circumstellar disks and show both photometric and spectroscopic variability (ITerbst et al., A number of Herbig Ae stars have been found to possess circumstellar disks and show both photometric and spectroscopic variability (Herbst et al.
+ 1999)., 1999).
+" LkIL,231 though much hotter aud more massive also exhibits these properties (Fueute et al.", $_\alpha$ 234 though much hotter and more massive also exhibits these properties (Fuente et al.
+ 2001. Herbst et al.," 2001, Herbst et al."
+ 1999. Polomiski et al.," 1999, Polomski et al."
+ 2002)., 2002).
+ The star shows photometric variability of 1.2 magnitude in the V band and is known posses a circuustellar disk with the disk mass to the stellar iiass ∙ ∙ ∙∙ ∐∖∏⊓−≱⇝↓≼⋯, The star shows photometric variability of 1.2 magnitude in the V band and is known posses a circumstellar disk with the disk mass to the stellar mass ratio $_D$ $_*$ ) $\leq$ 0.02.
+"≀⋅≓⊔⋅∩↕⊔⋞↧⊓⋝↕↴∖⋞↧↴∖↾⋞∐↸∐∏⋝∪↧↸∪↧ LkIL,231 aud its circtuustellar region cau be found in Fuente et al. (", A detailed description of $_\alpha$ 234 and its circumstellar region can be found in Fuente et al. (
+2001. aud references therein).,"2001, and references therein)."
+ Spectroscopic. variatious⋅⋅ in⋅ metallic ⋅⋅lines (like⋅ Na L Ca IL ete.)," Spectroscopic variations in metallic lines (like Na I, Ca II, etc.)"
+ in vouug A type stars like >j Pictoris (20 to 100 million wears) are often ↕∐↑↸∖∏∐⋅↸∖↑↸∖≼↧⋜↧↴∖↴↕∐∖↴∖↴↕∶↴∙⊾∐⋜↕⊓∐⋅↸∖∪↕≯↴∖↴↑⋜∐⋅∶↴∙⊾↥⋅⋜↧∑↕∐∶↴⋁↸⊳∪∐∐∖↑↴∖↴ ∪↥⋅↕⋟⋜↕∐↕∐∶↴∙⊾↸∖↖⇁⋜↧⋯↥⋅⋜↧↑↕∐∶↴⋁↴⋈⋈∐↸∖↴∖↴↖↖⇁↕↑∐↴∖↴↕∑↸∖↴∖↴∪↕⋡↑↸∖∐↴∖↴, in young A type stars like $\beta$ Pictoris (20 to 100 million years) are often interpreted as the signature of star grazing comets or falling evaporating bodies with sizes of tens of kilometers (Grinin et al.
+∪↕⋡ ↨↘↽∐∪∐∐∖↑↸∖↥⋅↴∖↴≺≼∶↥⋅↕∐↕∐↸∖↑⋜↧↕⋅↽⋂⋂⊓⋅≋⋯↕↑∐∙∖↽⊺↸∖↥⋅↥⋅∐↸∖ 1981. Thebault Beust 2001).," 1996, Smith Terrile 1984, Thebault Beust 2001)."
+ The only Ierbig Be star where Sr. fransic:ut events been inferred so far is COMLCTAUENHD. 100516 (Viera et al., The only Herbig Be star where grazing cometary transient events have been inferred so far is HD 100546 (Viera et al.
+ have.1999)., 1999).
+" However. IIDIOO0516 is nach older (2°10 Myrs) aud cooler (spectral type ~B9) compared to LKIL,231."," However, HD100546 is much older $\geq$ 10 Myrs) and cooler (spectral type $\sim$ B9) compared to $_\alpha$ 234."
+ Ilieh resolution spectroscopic monitoring of very vouug (0.1 to 1 Mis) Herbie HBe stars are scauty, High resolution spectroscopic monitoring of very young (0.1 to 1 Myrs) Herbig Be stars are scanty
+in 1816.,in 1876.
+ It is a pleasure to thank Stella Ixafka aud the referee. Michael Friedjung. for helpful commentis.," It is a pleasure to thank Stella Kafka and the referee, Michael Friedjung, for helpful comments."
+ We eratefullv acknowledge the support of this work by NSF grant. 0807892 (o Villanova University and partial summer undergraduate support from the NASA Delaware Space Grant Consortium., We gratefully acknowledge the support of this work by NSF grant 0807892 to Villanova University and partial summer undergraduate support from the NASA Delaware Space Grant Consortium.
+ Bruch. A.. Engel. 1994. AA Supp..," Bruch, A., Engel, 1994, AA Supp.,"
+" 104. 79 Duerbeck. IL1981. PASDP. 93. 165 Katka. ο, et al."," 104, 79 Duerbeck, H.1981, PASP, 93, 165 Kafka, S., et al."
+ 2003. AJ. 126. 1422 Patterson 1984. ÀpJS. 54. 443 Schaeler. D.2010. private communication Selvelli. P.2004. Daltic Astr.13. 93 Waele. RÀ. IInbeny. 1. 1998. ApJ. 509. 350.," 2003, AJ, 126, 1422 Patterson 1984, ApJS, 54, 443 Schaefer, B.2010, private communication Selvelli, P.2004, Baltic Astr.13, 93 Wade, R.A. Hubeny, I. 1998, ApJ, 509, 350."
+is the closest specimen of starbursts with superwinds. at a distance of only (3.9+0.3) MMpe (Sakai&Madore 1999).,"is the closest specimen of starbursts with superwinds, at a distance of only $(3.9 \pm 0.3)$ Mpc \citep{Sakai99}."
+. Its almost edge-on orientation is ideal for the study of the superwind. and more generally the interaction of galaxies with the intergalactic medium and the role of the gas and dust dynamies in the self-regulation of starbursts.," Its almost edge-on orientation is ideal for the study of the superwind, and more generally the interaction of galaxies with the intergalactic medium and the role of the gas and dust dynamics in the self-regulation of starbursts."
+ The exceptional star formation activity of 882 is thought to be caused by its tidal interaction with the M81 galaxy group. redistributing the gas within the disk of 882 toward the central regions where it accumulates. and expelling atomic gas In giant streams connecting the disk to the intragalactic medium (Yunetal.1994).," The exceptional star formation activity of 82 is thought to be caused by its tidal interaction with the 81 galaxy group, redistributing the gas within the disk of 82 toward the central regions where it accumulates, and expelling atomic gas in giant streams connecting the disk to the intragalactic medium \citep{Yun94}."
+. 882 is an evolved starburst. in the sense that its activity has been sustained (episodically) over much of its lifetime.," 82 is an evolved starburst, in the sense that its activity has been sustained (episodically) over much of its lifetime."
+ Coupled stellar population synthesis and photoionization. models reveal two major episodes of star formation over the past MMyr (ForsterSchreiberet4|. 2003)., Coupled stellar population synthesis and photoionization models reveal two major episodes of star formation over the past Myr \citep{Forster03}.
+.. Numerous young supernova remnants have been detected (Muxlowetal.1994)., Numerous young supernova remnants have been detected \citep{Muxlow94}.
+. Another peak of star formation about MMyr ago is inferred from the age-dating of stellar clusters (Konstantopoulosetal.2009). and 200 super star clusters. signposts of starbursts. have been found in 882," Another peak of star formation about Myr ago is inferred from the age-dating of stellar clusters \citep{Konstantopoulos09}, and 200 super star clusters, signposts of starbursts, have been found in 82"
+the Fresnel propagation are compared.,the Fresnel propagation are compared.
+" From the results of these simulations, we can then conclude that the effects of the Fresnel propagation are not too great for our instrument."," From the results of these simulations, we can then conclude that the effects of the Fresnel propagation are not too great for our instrument."
+ This is because of the large conjugated distances of the optics (see Table 2)) and the not very large pupil size., This is because of the large conjugated distances of the optics (see Table \ref{tabfresnelparameter}) ) and the not very large pupil size.
+" Fresnel propagation is much more a concern for extremely large telescopes, like E-ELT (see Antichietal. 2010))."," Fresnel propagation is much more a concern for extremely large telescopes, like E-ELT (see \citealt{Antichi10}) )."
+" Through the simulations described in the previous sections, we have demonstrated the capability of the SPHERE IFS instrument to image extrasolar planets down to a contrast of ~107"" at a separation of a few tenths of arcsec."," Through the simulations described in the previous sections, we have demonstrated the capability of the SPHERE IFS instrument to image extrasolar planets down to a contrast of $\sim 10^{-7}$ at a separation of a few tenths of arcsec."
+" However, to fully characterize the newly discovered planets (temperature, chemical composition of the atmosphere, etc.),"," However, to fully characterize the newly discovered planets (temperature, chemical composition of the atmosphere, etc.),"
+ it is important to be able to reconstruct their spectra at a high-fidelity level., it is important to be able to reconstruct their spectra at a high-fidelity level.
+" To test this capability, we prepared a pipeline for a data analysis of the calibrated datacube resulting from our simulations."," To test this capability, we prepared a pipeline for a data analysis of the calibrated datacube resulting from our simulations."
+ This procedure is composed of five different steps: The simulation was performed using the same PSFs as in the previous simulations., This procedure is composed of five different steps: The simulation was performed using the same PSFs as in the previous simulations.
+ The FOV is rotated by 30? during the observation and the central star is a GOV star at a distance of 10 pc (this corresponds to a magnitude J=3.75)., The FOV is rotated by $30^{\circ}$ during the observation and the central star is a G0V star at a distance of 10 pc (this corresponds to a magnitude J=3.75).
+ Every simulation contains five planets in different positions but at the same separation from the central star (the planets in the same simulation are identical)., Every simulation contains five planets in different positions but at the same separation from the central star (the planets in the same simulation are identical).
+" We perfomed simulations with planets at a separation of 0.3, 0.5 and 1.0 arcsec from the central star."," We perfomed simulations with planets at a separation of 0.3, 0.5 and 1.0 arcsec from the central star."
+" To avoid overlapping of the planets PSFs at the smallest separation, we decided to replace these single simulations with three different ones containing two different planets each (for this reason only for the case at 0.3 arcsec we have six planets for every single case and not five)."," To avoid overlapping of the planets PSFs at the smallest separation, we decided to replace these single simulations with three different ones containing two different planets each (for this reason only for the case at 0.3 arcsec we have six planets for every single case and not five)."
+" We then performed different simulations with different luminosity contrasts between the planets and the central star, and adopted values of 102, 3x1079, 1076, and 3x1077."," We then performed different simulations with different luminosity contrasts between the planets and the central star, and adopted values of $10^{-5}$, $3\times 10^{-6}$, $10^{-6}$, and $3\times 10^{-7}$."
+" Finally, we performed different simulations with different input spectra: we used a late type T-dwarf spectrum (T7), an early type T-dwarf spectrum (T2), a late type L-dwarf spectrum (L8), and an early type L-dwarf spectrum (LO), taken from the spectra libraries described in Section ??.."," Finally, we performed different simulations with different input spectra: we used a late type T-dwarf spectrum (T7), an early type T-dwarf spectrum (T2), a late type L-dwarf spectrum (L8), and an early type L-dwarf spectrum (L0), taken from the spectra libraries described in Section \ref{procdesc}."
+" To test the capability of our procedure to distinguish between a companion object and a background star, we performed simulations usinga flat spectrum (M2) at the low resolution of our instrument as input."," To test the capability of our procedure to distinguish between a companion object and a background star, we performed simulations using a flat spectrum (M2) at the low resolution of our instrument as input."
+" Indeed, at this resolution, all the stellar spectra are expected to be flat."," Indeed, at this resolution, all the stellar spectra are expected to be flat."
+" Moreover, we did not include any faint background galaxies because they are expected to be spatially resolved as extended objects by our While all these spectra come from objects in the solar neighborhood (old objects), our results do not lose generality because, as shown in Section ??, the detectability of the companion objects at a fixed effective temperature is not determined by the gravity effects, which in turn is the main difference between young and older substellar objects."," Moreover, we did not include any faint background galaxies because they are expected to be spatially resolved as extended objects by our While all these spectra come from objects in the solar neighborhood (old objects), our results do not lose generality because, as shown in Section \ref{gravity}, the detectability of the companion objects at a fixed effective temperature is not determined by the gravity effects, which in turn is the main difference between young and older substellar objects."
+ In this section we will describe our reduction procedure in more detail., In this section we will describe our reduction procedure in more detail.
+ The first two steps are performed using the spectral deconvolution method in the same way as described above in Section 4.2.., The first two steps are performed using the spectral deconvolution method in the same way as described above in Section \ref{specdeconv}.
+ The search for companion objects (third step) is composed of three different steps:, The search for companion objects (third step) is composed of three different steps:
+ ~LOO. 104s ~LO°< Cant plancts iu the Solar System have albedos greater than because of the presence of condensed molecules (WhO. ο. NIIS. ete.)," $\sim 100\times$ $\sim10^4 \times$ $\sim10^6\times$ Giant planets in the Solar System have albedos greater than because of the presence of condensed molecules $_{2}$ O, $_{4}$, $_{3}$, etc.)"
+ in their atmospheres., in their atmospheres.
+ Planets with effective temperatures exceeding ~ LOO I& should be cloud free. leading to albedos of 0.050.1 (Marleyetal. 1999).," Planets with effective temperatures exceeding $\sim$ 400 K should be cloud free, leading to albedos of 0.05–0.4 \citep{Marley_1999}."
+". Tf pressure-broadened Na aud [& opacity is portant at optical wavelengths (asitisforbrownανατς,Burrowsetal. 2000).. then the Bond. albedos of hot προς may. be less than (Sudarskyetal. 2000).."," If pressure-broadened Na and K opacity is important at optical wavelengths \citep[as it is for brown dwarfs,][]{Burrows_2000}, then the Bond albedos of hot Jupiters may be less than \citep{Sudarsky_2000}. ."
+ But the very hottest planets. tle so-called class V extrasolar eiat planets (Tig21500 IK). nüght have very hieh albedos due to a high silicate cloud laver (Sudarskyetal. 2000).," But the very hottest planets, the so-called class V extrasolar giant planets $T_{\rm eff} > 1500$ K), might have very high albedos due to a high silicate cloud layer \citep{Sudarsky_2000}."
+. For a planet whose albedo is dominated by clouds (as opposed to Ravleigh scattering) the albedo depends on the composition aud size of cloud particles (Seageretal.2000)., For a planet whose albedo is dominated by clouds (as opposed to Rayleigh scattering) the albedo depends on the composition and size of cloud particles \citep{Seager_2000}.
+. Early attempts to observe reflected Πο from exoplauets (Charbonneauetal.1999:CollierCameroun2010) indicated that they aight not be as reflective as Solar System gas elauts (forareview.seeLanetordetal. 2010).," Early attempts to observe reflected light from exoplanets \citep{Charbonneau_1999, Cameron_2002, Leigh_2003a, Leigh_2003b, Rodler_2008, Rodler_2010} indicated that they might not be as reflective as Solar System gas giants \citep[for a review, see][]{Langford_2010}."
+. Measureimeuts of ΠΟ 209158b taken with the Canadian MOST satellite revealed a very low albedo (<<8%.Roweetal.2008). and it las since been takeu for eranted that all short-periocd planets have albedos ou par with that of charcoal.," Measurements of HD 209458b taken with the Canadian MOST satellite revealed a very low albedo \citep[$< 8$\%,][]{Rowe_2008}, and it has since been taken for granted that all short-period planets have albedos on par with that of charcoal."
+ From the standpoiut ofthe planct’s chinate. the inportant factor is not the albedo at any one waveleneth.," From the standpoint ofthe planet's climate, the important factor is not the albedo at any one wavelength,"
+During the quiescent phase. the optical spectra of LBVs are characterised by effective temperatures ranging from 0000 to 300000 K (B-spectral types) and strong emission lines of hydrogen. Hel. Fell. Call and other singly ionised metals. often with P-Cyeni profiles when observed at sufficiently high resolution (cf.,"During the quiescent phase, the optical spectra of LBVs are characterised by effective temperatures ranging from 000 to 000 K (B-spectral types) and strong emission lines of hydrogen, HeI, FeII, CaII and other singly ionised metals, often with P-Cygni profiles when observed at sufficiently high resolution (cf."
+ Kenyon Gallagher 1985)., Kenyon Gallagher 1985).
+ The mass loss rate in quiescence is smaller and the wind faster., The mass loss rate in quiescence is smaller and the wind faster.
+ While not well-understood an LBV occasionally passes through an “outburst” phase. during which the luminosity remains the same as in quiescence. while in tandem the mass loss rate increases and the wind velocity decreases.," While not well-understood an LBV occasionally passes through an 'outburst' phase, during which the luminosity remains the same as in quiescence, while in tandem the mass loss rate increases and the wind velocity decreases."
+ This pushes outward the pseudo-photosphere. with a consequent icrease in the effective radius and decrease in the surface temperature (whose Wien peak moves from the ultraviolet toward the optical region).," This pushes outward the pseudo-photosphere, with a consequent increase in the effective radius and decrease in the surface temperature (whose Wien peak moves from the ultraviolet toward the optical region)."
+ At outburst maximum. the effective temperature is ~7500-8500 K regardless of the LBV luminosity and temperature in quiescence. and the spectrum turns to that of A-type supergiants. while [Fell]. Hel emission lines. and P-Cygni absorptions weaken or disappear.," At outburst maximum, the effective temperature is $\sim$ 7500-8500 K regardless of the LBV luminosity and temperature in quiescence, and the spectrum turns to that of A-type supergiants, while [FeII], HeI emission lines, and P-Cygni absorptions weaken or disappear."
+ Wolf (1989; see Vink 2008 for an update) pointed out the existence of an luminosity relation for LBV eruptions and called the region of the HR diagram occupied by LBVs in quiescence., Wolf (1989; see Vink 2008 for an update) pointed out the existence of an amplitude-luminosity relation for LBV eruptions and called the region of the HR diagram occupied by LBVs in quiescence.
+ Stothers Chin (1995) suggested the existence of a pertod-luminosity relation of the form My=12.5(40.5)+2.4(40.5)logP where P. expressed in yrs. is the typical interval from one eruption to the next.," Stothers Chin (1995) suggested the existence of a period-luminosity relation of the form $M_{\rm bol} =
+-12.5(\pm 0.5) + 2.4(\pm 0.5) \log P$ where P, expressed in yrs, is the typical interval from one eruption to the next."
+ The spectra that RAVE collected for the LBVs in the LMC are presented in Figs., The spectra that RAVE collected for the LBVs in the LMC are presented in Figs.
+ 1-3. where the ordinate scale is kept constant for an easier inter-comparison.," 1-3, where the ordinate scale is kept constant for an easier inter-comparison."
+. Differences in S/N among the multi-epoch spectra of the same star are mainly due to differences in the throughputs of different fibers and in the sky conditions., Differences in S/N among the multi-epoch spectra of the same star are mainly due to differences in the throughputs of different fibers and in the sky conditions.
+ The radial velocities derived from the absorption and emission lines are given in Table | (for comparison. the heliocentric barycentric velocity of the LMC is +278 km/s. Mateo 1998. although significantly variable across the face of the galaxy).," The radial velocities derived from the absorption and emission lines are given in Table 1 (for comparison, the heliocentric barycentric velocity of the LMC is $+$ 278 km/s, Mateo 1998, although significantly variable across the face of the galaxy)."
+ Table ] also lists the V-magnitude at the time of the RAVE observation as estimated by consulting various databases of variable star amateur observers (AAVSO. VSNET. VSOLJ) and the ASAS Survey database (Pojmanski 2002).," Table 1 also lists the $V$ -magnitude at the time of the RAVE observation as estimated by consulting various databases of variable star amateur observers (AAVSO, VSNET, VSOLJ) and the ASAS Survey database (Pojmanski 2002)."
+ These observations indicate that the RAVE spectra are well suited to detect and study peculiar stars. and LBVs in particular. confirming early investigations in this wavelength interval by Munari (1998. 2002. 2003).," These observations indicate that the RAVE spectra are well suited to detect and study peculiar stars, and LBVs in particular, confirming early investigations in this wavelength interval by Munari (1998, 2002, 2003)."
+ The peculiarities stands out well in terms of absorption features. wind outflow signatures. and emission lines. and will now be considered m turn.," The peculiarities stands out well in terms of absorption features, wind outflow signatures, and emission lines, and will now be considered in turn."
+ The ionised circumstellar gas is easily revealed by the rich emission line spectrum observable over the RAVE wavelength regime., The ionised circumstellar gas is easily revealed by the rich emission line spectrum observable over the RAVE wavelength regime.
+ The principal tons contributing emission lines at these wavelengths are HI. Hel. Call. NI. CIL. CHI. OI. SI. [Fell].," The principal ions contributing emission lines at these wavelengths are HI, HeI, CaII, NI, CII, CIII, OI, SI, [FeII],"
+abundance of individual lines are showed in the file out2.txt (work.dir directory).,abundance of individual lines are showed in the file out2.txt (work.dir directory).
+ There are small differences in the parameters derived using this method by different authors., There are small differences in the parameters derived using this method by different authors.
+ Santosetal.(2000) derived surface gravities systematically higher than the ones obtained by other authors (e.g.Gonzalezetal.2001) by 0.15 dex., \citet{santos00} derived surface gravities systematically higher than the ones obtained by other authors \citep[e.g. ][]{gonzalez01} by 0.15 dex.
+ To solve this problem. Santosetal.(2004) adopted new log ef values for the tron lines.," To solve this problem, \citet{santos04} adopted new log gf values for the iron lines."
+ Then the authors found a small average difference of +25 K in temperature compared to the studies of Gonzalezetal.(2001). and Lawsetal. (57 stars in common).," Then the authors found a small average difference of +25 K in temperature compared to the studies of \citet{gonzalez01} and \citet{laws03}
+ (57 stars in common)."
+ Also their log gs are 40.05 dex (on average) above and the average differences in. metallicity are between -0.10 and +0.10 dex., Also their log gs are +0.05 dex (on average) above and the average differences in metallicity are between -0.10 and +0.10 dex.
+ Ammler-vonΕΙΠetal.(2009) derived the fundamental parameters of host stars with transiting planets. using the 2002 version of the code MOOG and ATLAS model atmospheres.," \citet{ammler09} derived the fundamental parameters of host stars with transiting planets, using the 2002 version of the code MOOG and ATLAS model atmospheres."
+ However for the TrES and HAT objects. the abundance determinations of Sozzettietal.(2007) and Sozzettietal.(2009) are systematically lower than those derived in their work bydex. using essentially the same method.," However for the TrES and HAT objects, the abundance determinations of \citet{sozzetti07}
+ and \citet{sozzetti09} are systematically lower than those derived in their work by, using essentially the same method."
+ The authors also mention differences in temperature (~100 K) and gravity (0.15 dex) due to possible sistematic tendences betweer the works. and they note that the origin of the discrepancies 1 abundanees is commonly unidentified.," The authors also mention differences in temperature $\sim$ 100 K) and gravity $\sim$ 0.15 dex) due to possible sistematic tendences between the works, and they note that the origin of the discrepancies in abundances is commonly unidentified."
+" We compare the parameters derived using and those from literatere (Gonzalez1997,1998;Gonzalezetal.etal.2001:Laws&Gonzalez 2003)."," We compare the parameters derived using and those from literature \citep{gonzalez97,gonzalez98,gonzalez99,santos00,gonzalez-laws00,
+gonzalez01,laws-gonzalez01,laws03}."
+. We selected these works in particular because they present the Fe equivalent widths: we use the same values with FUNDPAR., We selected these works in particular because they present the Fe equivalent widths: we use the same values with FUNDPAR.
+" In the Table 4. we show the T,;;. /ogg. |Fel/H] and & obtained."," In the Table \ref{tabla.pars} we show the $_{eff}$, $log g$, [FeI/H] and $\xi$ obtained."
+ Some stars in the table are listed. twice (identified with a subseript) because we have different equivalent width sources from literature., Some stars in the table are listed twice (identified with a subscript) because we have different equivalent width sources from literature.
+ The sources of the equivalent widths are listed in the Table + as RI.....R8.," The sources of the equivalent widths are listed in the Table \ref{tabla.pars} as R1,...,R8."
+ We show in the Figure 2. an example of abundance vs. excitation potential and abundance vs. (logarithm of the reduced equivalent width) for HD 106252., We show in the Figure \ref{fig2} an example of abundance vs. excitation potential and abundance vs. (logarithm of the reduced equivalent width) for HD 106252.
+ In the Figure 3 we present the difference between FUNDPAR parameters and literature (y axis. ASFUNDPAR-literature) vs. literature (x axis).," In the Figure \ref{fig1} we present the difference between FUNDPAR parameters and literature (y axis, $\Delta$ =FUNDPAR-literature) vs. literature (x axis)."
+ The empty point in the Figure 3. correspond to the, The empty point in the Figure \ref{fig1} correspond to the
+Surveys for DEEP David €. ου Cruz. CA 05064. USA The 21st century promises many faint redshift surveys with the suite of new 8- an class erouud-based optical telescopes.,"Surveys for DEEP David C. Koo } , CA 95064, USA } The 21st century promises many faint redshift surveys with the suite of new 8-10 m class ground-based optical telescopes."
+ Besides complementing data from space aud other wavebanuds with critical redshifts. the high S/N aud high spectral-rvesolution from 8-10 i class telescopes provide three new. powerful diagnostics for the analysis of distant galaxies. namely internal velocities (auc hence masses when size is used). abuudances. and age estimates;," Besides complementing data from space and other wavebands with critical redshifts, the high S/N and high spectral-resolution from 8-10 m class telescopes provide three new, powerful diagnostics for the analysis of distant galaxies, namely internal velocities (and hence masses when size is used), abundances, and age estimates."
+ These paraucters are clearly inportaut. independent probes of galaxies i the early universe witli solid links to theoretical simulations.," These parameters are clearly important, independent probes of galaxies in the early universe with solid links to theoretical simulations."
+ Finallv. because both galaxy. evolution and their large scale patterus are very likely to be couples. survevs with large samples will be needed to address these problems.," Finally, because both galaxy evolution and their large scale patterns are very likely to be complex, surveys with large samples will be needed to address these problems."
+ To mect the challeuge. the Deep Extragalactic Evolutionary Probe (DEEP: more details on participants aud programs at URL: ~deep/home.btinl) was initiated over 6 years ago as a project desigued to gather spectral data for over 10.000 faint ficld galaxies [9]..[3].. [6] using the Keck IT 10-1 telescope and a new spectrograph for Neck IL C(DEIMOS: DEep huagiug Multi-Object Spectrograph: mere information is provided at URL:," To meet the challenge, the Deep Extragalactic Evolutionary Probe (DEEP: more details on participants and programs at URL: ) was initiated over 6 years ago as a project designed to gather spectral data for over 10,000 faint field galaxies \cite{Mou93}, \cite{Koo95a}, \cite{Koo98} using the Keck II 10-m telescope and a new spectrograph for Keck II (DEIMOS: DEep Imaging Multi-Object Spectrograph; more information is provided at URL:"
+for both samples.,for both samples.
+ The far-infrared luminosity is a good measure of the star formation rate (SFR) in dusty svstenms like ULIRGs., The far-infrared luminosity is a good measure of the star formation rate (SFR) in dusty systems like ULIRGs.
+ Comparison of the SFR trom X-ray measurements with the SER from ihe FIR measurements will indicate if there is anv energv contribution from sources other than the starburst (e.g. an AGN)., Comparison of the SFR from X-ray measurements with the SFR from the FIR measurements will indicate if there is any energy contribution from sources other than the starburst (e.g. an AGN).
+ If the galaxy is powered purely by a starburst. then one would expect its “SFR in A-ravs” (o equal its “SFR in FIR’.," If the galaxy is powered purely by a starburst, then one would expect its “SFR in X-rays” to equal its “SFR in FIR”."
+ There are many relerences in the literature that relate the 210 keV hard. X-ray luminosity to SE».40ων., There are many references in the literature that relate the 2–10 keV hard X-ray luminosity to $_{\mathrm {2 - 10~keV}}$.
+ We have chosen to adopt the Ranallietal.(2003) and Persieοἱal.(2004). relations because they appear lo be the best-fits to pure starbursts (ILornschemeieretal.20, We have chosen to adopt the \citet{ranalli} and \citet{persic2} relations because they appear to be the best-fits to pure starbursts \citep[Figure~6b]{horn}.
+"04.Figure6b).. The two Re40jv. values will give approximate upper and lower limits as a function of Ry,""s F", The two $_{\mathrm {2 - 10~keV}}$ values will give approximate upper and lower limits as a function of $_{\rm FIR}$.
+igure 14. relates the SFR from the Ranallietal.(2003) and Persicetal.(2004). relations to (he SER from the FUR luminosity., Figure \ref{fig:sfr} relates the SFR from the \citet{ranalli} and \citet{persic2} relations to the SFR from the FIR luminosity.
+ We have included the Ptalketal.(2003) sample in the plot., We have included the \citet{ptak} sample in the plot.
+ From the figure. it is evident (hat our two Sevlert 1 galaxies have N-ray Iuminosities in excess Of (hat expected from a starburst (as discussed above).," From the figure, it is evident that our two Seyfert 1 galaxies have X-ray luminosities in excess of that expected from a starburst (as discussed above)."
+ The only other source Irom our sample located above the line of equality is (he LINER. F04103-2838., The only other source from our sample located above the line of equality is the LINER F04103-2838.
+ The four galaxies determined to have AGN contributions in their spectra (Mrk 231. Mark 273. IRAS 05139-2524. and NGC 6240) are also located far above the line of equality.," The four galaxies \citet{ptak} determined to have AGN contributions in their spectra (Mrk 231, Mrk 273, IRAS 05189-2524, and NGC 6240) are also located far above the line of equality."
+ The rest of the ULIRGs are most likely powered by starbursts., The rest of the ULIRGs are most likely powered by starbursts.
+ In this section. we explore (he possible emission processes that mav be responsible for the detected. X-ray. emission.," In this section, we explore the possible emission processes that may be responsible for the detected X-ray emission."
+ Persic&Rephaeli(2002) suggested (hat X-ray binaries dominate the 215 keV luminosity in the absence of an AGN., \citet{persic} suggested that X-ray binaries dominate the 2–15 keV luminosity in the absence of an AGN.
+ There are two types of X-ray binaries., There are two types of X-ray binaries.
+ The high mass tvpe (IINXD) produces X-ray. emission from the accretion of wind material of an OB star onto ils neutron star or black hole companion whereas the low mass (vpe (LMXD) produces X-ray, The high mass type (HMXB) produces X-ray emission from the accretion of wind material of an OB star onto its neutron star or black hole companion whereas the low mass type (LMXB) produces X-ray
+parameter. thedeplh. 744. defined in eq. (2)).,"parameter, the, $\tau_{\rm rep}$, defined in eq. \ref{trep}) )."
+ This parameter issize., This parameter is.
+ Our results lor case C are similar but. as shown in section 4.3.. another parameter (I. defined in eq. (18)))," Our results for case C are similar but, as shown in section \ref{casecs}, , another parameter $\Gamma$ , defined in eq. \ref{gamma}) ))"
+ controls the shape of the SED., controls the shape of the SED.
+ The results for ease D differ for single grain size models and models with a size distribution., The results for case B differ for single grain size models and models with a size distribution.
+ For single grain size models. the SEDs are similar only in the IR. while large dillerences are observed in the extinction region. where (he grain opacites depend very much on grain size.," For single grain size models, the SEDs are similar only in the IR, while large differences are observed in the extinction region, where the grain opacities depend very much on grain size."
+ Alodels with a grain size distribution. on the other hand. display a striking similarity in the extinction region. provided (he lower limit of the grain size distribution. (44. is in the geometrical limit in the extinction region.," Models with a grain size distribution, on the other hand, display a striking similarity in the extinction region, provided the lower limit of the grain size distribution, $a_{\rm min}$, is in the geometrical limit in the extinction region."
+" It follows that. if e, is chosen accordingly. the case D SEDs for models with a grain size distribution is controlled solely by (he parameter Trop (Che same result as case A)."," It follows that, if $a_{\rm min}$ is chosen accordingly, the case B SEDs for models with a grain size distribution is controlled solely by the parameter $\tau_{\rm rep}$ (the same result as case A)."
+ The fact that τν (or E lor case C) is the appropriate parameter lor SED classification indicates that the usual approach in the literature - using 7: or 7j - will not reveal the general invariance of the SED cliscussecl in this paper., The fact that $\tau_{rep}$ (or $\Gamma$ for case C) is the appropriate parameter for SED classification indicates that the usual approach in the literature - using $\tau_V$ or $\tau_F$ - will not reveal the general invariance of the SED discussed in this paper.
+ We conclude therefore that the SED is generally avery poor constraint of grain size., We conclude therefore that the SED is generally a very poor constraint of grain size.
+ In most situations. observations must resolve differences of the order of a few percent to extract information about the grain size.," In most situations, observations must resolve differences of the order of a few percent to extract information about the grain size."
+ If. we add to this the fact that other model parameters. such as dust spatial distribution. composition and optical properties are somewhat uncertain. it becomes apparent that the task of extracting inlormation about grain sizes from the SED alone may prove very difficult or even impossible for spherical geometries.," If we add to this the fact that other model parameters, such as dust spatial distribution, composition and optical properties are somewhat uncertain, it becomes apparent that the task of extracting information about grain sizes from the SED alone may prove very difficult or even impossible for spherical geometries."
+" If follows fron: AI that in most circumstances 74, is theobservations.", If follows from AI that in most circumstances $\tau_{\rm rep}$ is the.
+ Using Tw io determine dust column clensity or mass is directlv. subject to the uncertainty in the determination of the grain sizes. which directly affects our ability to measure mass loss rates using the SED.," Using $\tau_{\rm rep}$ to determine dust column density or mass is directly subject to the uncertainty in the determination of the grain sizes, which directly affects our ability to measure mass loss rates using the SED."
+ An important physical effect discussed in section 5.1 is that for models with a grain size distribution. the different grain sizes have different radiative equilibrium temperatures.," An important physical effect discussed in section \ref{sec51} is that for models with a grain size distribution, the different grain sizes have different radiative equilibrium temperatures."
+ This implies (hat a consistent model should include the correct condensation radius for each grain size ol the distribution., This implies that a consistent model should include the correct condensation radius for each grain size of the distribution.
+ However. we show in section 5.1. thal the approximation of using identical condensation radii is reasonable when either all grain sizes have similar condensation temperatures or when a particular range of grain sizes dominate the It emission.," However, we show in section \ref{sec51} that the approximation of using identical condensation radii is reasonable when either all grain sizes have similar condensation temperatures or when a particular range of grain sizes dominate the IR emission."
+ The only situation where the consistent. (treatment (each grain size having its correct. condensation radius) is required is when (wo or more ranges of grain sizes have different. equilibrium temperatures and similar integrated emission., The only situation where the consistent treatment (each grain size having its correct condensation radius) is required is when two or more ranges of grain sizes have different equilibrium temperatures and similar integrated emission.
+ This situation was found only [ον models of hot stars with very small grains (@<0.01 jm) present in the distribution.," This situation was found only for models of hot stars with very small grains $a \lesssim 0.01\,\mu \rm m$ ) present in the distribution."
+ It is importantto notice that the results presented here depend on the parüceular choice ofthe distribution, It is importantto notice that the results presented here depend on the particular choice ofthe distribution
+in a one-dimensional. sell-gravitatiug slab.,"in a one-dimensional, self-gravitating slab."
+ They considered (wo cases: a slab with an embedded inagnelic¢ field oriented parallel to (he slab plane. and one with an internal field in (he normal direction.," They considered two cases: a slab with an embedded magnetic field oriented parallel to the slab plane, and one with an internal field in the normal direction."
+ In the first case. Fatuzzo Adams showed (hat magnetosonic waves provide a normal force.," In the first case, Fatuzzo Adams showed that magnetosonic waves provide a normal force."
+ In the second. it is Alfvénn waves that exert (he force. also in the normal direction.," In the second, it is Alfvénn waves that exert the force, also in the normal direction."
+ Thus. Fattuzzo Aclams verified explicitly that (he waves counteract gravity. even in (he absence of wave dampine.," Thus, Fattuzzo Adams verified explicitly that the waves counteract gravity, even in the absence of wave damping."
+ Melee&Zweibel(1995) extended (his result., \cite{mz95} extended this result.
+ Using the pioneering analvsis of (1970).. thev showed that Alfvénn waves generated bv a (turbulent wave [ield exert. an isolropic pressure. regardless of the background geometry.," Using the pioneering analysis of \citet{d70}, they showed that Alfvénn waves generated by a turbulent wave field exert an isotropic pressure, regardless of the background geometry."
+ Mclxee Zweibel derived a simple dependence of the wave pressure 7? on the local density: If we ignore the relatively small thermal pressure. (hen the cloud can be described as ann polvirope.," McKee Zweibel derived a simple dependence of the wave pressure $P$ on the local density: If we ignore the relatively small thermal pressure, then the cloud can be described as an polytrope."
+ Are the structures of real clouds consistent with this polviropic wave pressure?, Are the structures of real clouds consistent with this polytropic wave pressure?
+ One. indirect. argument indicates they are nol.," One, indirect, argument indicates they are not."
+ Melxee Zweibel also demonstrated Chat J? is proportional to p (mes the square of the (randomly. oriented) velocity [fluctuation 9e., McKee Zweibel also demonstrated that $P$ is proportional to $\rho$ times the square of the (randomly oriented) velocity fluctuation $\delta v$.
+ H follows that in (his model., It follows that in this model.
+" Now the densitv in an polvtrope tends (ο approach a power law outside the central plateau. such that pis proportional to r.ο,"," Now the density in an polytrope tends to approach a power law outside the central plateau, such that $\rho$ is proportional to $r^{-4/3}$."
+ From equation (2). it follows that Ov is proportional to r!oE," From equation (2), it follows that $\delta v$ is proportional to $r^{1/3}$."
+ Consider (he nearly spherical cloud. now gone. that produced the ONC.," Consider the nearly spherical cloud, now gone, that produced the ONC."
+ This cloud was recently driven off by the Trapezitm stars. which themselves have ages of about 10? vr 2001).," This cloud was recently driven off by the Trapezium stars, which themselves have ages of about $10^5$ yr \citep{ps01}."
+. The disruption itself occurred well within the cluster crossing time of about 10° vr., The disruption itself occurred well within the cluster crossing time of about $10^6$ yr.
+ Hence. the present-day velocity dispersion of thes/ars should reflect the prior ovr of thegas.," Hence, the present-day velocity dispersion of the should reflect the prior $\delta v$ of the."
+ But the dispersion of ONC proper motions has negligible variation from (he center to the outskirts of the cluster (Jones&Walker1933)., But the dispersion of ONC proper motions has negligible variation from the center to the outskirts of the cluster \citep{jw88}.
+. These measurements span at least a decade in radius. over which 9e should vary. by a factor of 2.2. according to the polvtropic relation.," These measurements span at least a decade in radius, over which $\delta v$ should vary by a factor of 2.2, according to the polytropic relation."
+ Our conclusion. based on this admittedly limited evidence. is that a more realistic model of the internal turbulence has a spatially constant velocity If we further," Our conclusion, based on this admittedly limited evidence, is that a more realistic model of the internal turbulence has a spatially constant velocity If we further"
+The drift velocity given by equation (10)) has spherical polar components ( ( ,The drift velocity given by equation \ref{EIperp}) ) has spherical polar components ( ( ).
+"The drift velocity is perpendicular to the magnetic field. implving 2B,— 0."," The drift velocity is perpendicular to the magnetic field, implying $\Delta v_rB_r+\Delta v_\theta B_\theta+\Delta v_\phi B_\phi=0$ ."
+ The dift velocity has its extrema αἱ cos(ó—αἱ)=dEL. with values-," The drift velocity has its extrema at $\cos(\phi-\omega t)=\pm1$, with values."
+- The maxinnun drift speed occurs for cos(ax0)=1. and has value jeu;=dersina," The maximum drift speed occurs for $\cos(\alpha\mp\theta)=1$, and has value $\Delta v_{\rm max}=\half\omega r\sin\alpha$."
+ In the case sina«1. the effects of obliquity can be (reated as à perturbation on the aligned moclel. as discussed in (he Appendix.," In the case $\sin\alpha\ll1$, the effects of obliquity can be treated as a perturbation on the aligned model, as discussed in the Appendix."
+ To zeroth order in sina. the magnetosphere is assumed to be rigidly corotating. with the charge density equal to poy.," To zeroth order in $\sin\alpha$, the magnetosphere is assumed to be rigidly corotating, with the charge density equal to $\rho_{\rm GJ}$."
+" The inductive electric Ποια is of first order in sina,"," The inductive electric field is of first order in $\sin\alpha$ ,"
+ , 
+We con]mle a list of 1225 candidate clusters.,We compile a list of 1225 candidate clusters.
+ This ist is certainly neither complete. nor devoid of false detections.," This list is certainly neither complete, nor devoid of false detections."
+ Our alu is to err ou the side of caution. so we visually examine all eiucdidates aud exclude hosethat are either visually nou-distinct or strongly contaminate Ww nearby sources (bright stars or nebulositv).," Our aim is to err on the side of caution, so we visually examine all candidates and exclude those that are either visually non-distinct or strongly contaminated by nearby sources (bright stars or nebulosity)."
+ Even in cases where these rejected candidates represent rea clusters we would be unable to obtain reliable structura fits., Even in cases where these rejected candidates represent real clusters we would be unable to obtain reliable structural fits.
+ We retain 1066 clusters aud their coordinates a1 cross-idenutifications to the Bicaetal.(2008) catalog are resented. i Table 1.. while their images are preseu4( in Figure 1..," We retain 1066 clusters and their coordinates and cross-identifications to the \cite{bica} catalog are presented in Table \ref{tab:Names and Positions}, while their images are presented in Figure \ref{fig:Sample Images}."
+ The Bicaetal.(2008) catalog is an update of fheir earlier work m which they complemet heir collecion of data from the literature with visua exoaninuatioi of photographic survey plates., The \cite{bica} catalog is an update of their earlier work in which they complement their collection of data from the literature with visual examination of photographic survey plates.
+ As such. it is an extensive sunnuurv of the available data. μπιτ ποιονjit. heterogeneous despite best efforts.," As such, it is an extensive summary of the available data, but somewhat heterogeneous despite best efforts."
+ That catalog iucludes all manner of extended objects; whereas we focus ou stellar clusters and associations.," That catalog includes all manner of extended objects, whereas we focus on stellar clusters and associations."
+ Major aud uinor axis 1cheths are inclided. but profiles are not fit.," Major and minor axis lengths are included, but profiles are not fit."
+ The catalog cross-ideutification is done matching to the rearest source within 10 arcsec., The catalog cross-identification is done matching to the nearest source within 10 arcsec.
+ Out of our total of 1066 clusters. we find matching designations in the literature oy 682 )).," Out of our total of 1066 clusters, we find matching designations in the literature for 682 )."
+ Qur suuje. Like alb suuples. nmmst be used with caution 1n evolutionary studies because clusters of outieular ages. masses; aud/or surface brielitucsses nav (0 prefercutially iuissine.," Our sample, like all samples, must be used with caution in evolutionary studies because clusters of particular ages, masses, and/or surface brightnesses may be preferentially missing."
+ Differences du assuniptions τοσαπιο fUs incompleteness have led to appareutly conflicting results in the literature (Claudaretal.2006:Cuelesetal.2007:deGaris&Goodwin 2008).. but correcting fex the incompleteness is. as always. model dependent.," Differences in assumptions regarding this incompleteness have led to apparently conflicting results in the literature \citep{chandar,gieles,degrijs}, but correcting for the incompleteness is, as always, model dependent."
+ Because the source photometric catalog is published. i is possible for subsequent investigators to recover the selection function by placing artificial clusters iuto the catalog aud repeating the cluster selection.," Because the source photometric catalog is published, it is possible for subsequent investigators to recover the selection function by placing artificial clusters into the catalog and repeating the cluster selection."
+ We, We
+shows that the curvature spectrum has a greater E. for a greater oc.,shows that the curvature spectrum has a greater $E_{\rm c}$ for a greater $\varrho_{\rm c}$.
+" Since the pair production mainly takes place between the inward y-rays and the surface X-rays, most of the pairs are produced in the lower altitudes where the X-ray density is large and the collisions take place mostly head-on."," Since the pair production mainly takes place between the inward $\gamma$ -rays and the surface X-rays, most of the pairs are produced in the lower altitudes where the X-ray density is large and the collisions take place mostly head-on."
+" Thus, outward-moving species (e.g., e*'s) run longer distances than inward ones within the gap, leading to a stronger outward 4-ray flux than the inward one."," Thus, outward-moving species (e.g., $e^+$ 's) run longer distances than inward ones within the gap, leading to a stronger outward $\gamma$ -ray flux than the inward one."
+" Therefore, to elucidate the physical mechanism that produces harder spectrum in the trailing peak of a 4-ray light curve, we concentrate on the outward-moving particles and the resultant outward emissions in 2 'To reveal the nature of curvature radiation, we must consider the curvature radius of actual particle motion in the 3-D pulsar magnetosphere."," Therefore, to elucidate the physical mechanism that produces harder spectrum in the trailing peak of a $\gamma$ -ray light curve, we concentrate on the outward-moving particles and the resultant outward emissions in \ref{sec:Ecutoff}
+ To reveal the nature of curvature radiation, we must consider the curvature radius of actual particle motion in the 3-D pulsar magnetosphere."
+" Well inside the light cylinder, the drift motion due to magnetic curvature, magnetic gradient, or time-varying electric field is small compared to the xB drift."," Well inside the light cylinder, the drift motion due to magnetic curvature, magnetic gradient, or time-varying electric field is small compared to the $\mbox{\boldmath$ $}\times\mbox{\boldmath$ $}$ drift."
+" Thus, particles co-rotate with the magnetic field lines."," Thus, particles co-rotate with the magnetic field lines."
+" In the flaring field line geometry of a NS magnetosphere, the path of outward-moving can be illustrated as figure 1.."," In the flaring field line geometry of a NS magnetosphere, the path of outward-moving particles can be illustrated as figure \ref{fig:Rc}."
+" In the leading side, particlesparticle's path is bent towards the rotational direction as the thin solid curve indicates."," In the leading side, particle's path is bent towards the rotational direction as the thin solid curve indicates."
+" in the trailing side, the path is straighter as indicated by However,the thick solid curve."," However, in the trailing side, the path is straighter as indicated by the thick solid curve."
+" Thus, outward-moving particles have greater curvature radii in the trailing side than in the leading side."," Thus, outward-moving particles have greater curvature radii in the trailing side than in the leading side."
+ This leads to a greater Ες in equation (3))., This leads to a greater $E_{\rm c}$ in equation \ref{eq:Ec2}) ).
+ This is the qualitative reason why most of the y-ray pulsars detected with LAT exhibits harder spectrum in the trailing light-curve peak than in the leading peak., This is the qualitative reason why most of the $\gamma$ -ray pulsars detected with LAT exhibits harder spectrum in the trailing light-curve peak than in the leading peak.
+" Next, let us examine the distribution of the curvature radius more quantitatively."," Next, let us examine the distribution of the curvature radius more quantitatively."
+" Due to strong relativistic beaming, curvature photons are emitted along the instantaneous particle velocity measured by the distant static observer (i.e., us)."," Due to strong relativistic beaming, curvature photons are emitted along the instantaneous particle velocity measured by the distant static observer (i.e., us)."
+" In the polar coordinates, the orthonormal components of the particle velocity become (Mestel et al."," In the polar coordinates, the orthonormal components of the particle velocity become (Mestel et al."
+" 1985; Camenzind 1986a,b) where y>1 is assumed and B,=\/(B"")?+(Π6)2, B denotes the strength of the magnetic field, c; the distance from the rotation axis."," 1985; Camenzind 1986a,b) where $\gamma \gg 1$ is assumed and $B_{\rm p}\equiv \sqrt{(B^r)^2+(B^{\hat\theta})^2}$, $B$ denotes the strength of the magnetic field, $\varpi$ the distance from the rotation axis."
+" Both (v"",v’,v®) and (B"",B?,B*) are measured by a distant static observer (i.e., in the inertial observers frame)."," Both $v^r$ $v^{\hat\theta}$ $v^{\hat\phi}$ ) and $B^r$ $B^{\hat\theta}$ $B^{\hat\varphi}$ ) are measured by a distant static observer (i.e., in the inertial observers frame)."
+" The upper sign of f, (eq. [5]])", The upper sign of $f_{\rm v}$ (eq. \ref{eq:v3Df}] ])
+" corresponds to the outward motion along the magnetic field line, whereas the lower sign to the inward motion."," corresponds to the outward motion along the magnetic field line, whereas the lower sign to the inward motion."
+" These are quite general and valid irrespective ofexpressions the force balance on the particles as long as y>1, because they are derived only from the Maxwell equations and the frozen-in condition, the latter of which is justified since we consider <B (or equivalently, consider that the potential dropEj in the gap is much small compared to the electro-motive force exerted on the spinning NS surface)."," These expressions are quite general and valid irrespective of the force balance on the particles as long as $\gamma \gg 1$, because they are derived only from the Maxwell equations and the frozen-in condition, the latter of which is justified since we consider $\Ell \ll B$ (or equivalently, consider that the potential drop in the gap is much small compared to the electro-motive force exerted on the spinning NS surface)."
+" For example, they are valid not only for e~’s being accelerated without dissipation but also for those in the force balance between radiation-reaction and electro-static acceleration."," For example, they are valid not only for $e^\pm$ 's being accelerated without dissipation but also for those in the force balance between radiation-reaction and electro-static acceleration."
+ Note that equations (4)) and (5)) correspond to a generalization of 4.1 of Dyks et al. (, Note that equations \ref{eq:v3D}) ) and \ref{eq:v3Df}) ) correspond to a generalization of 4.1 of Dyks et al. (
+2010) into the higher altitudes.,2010) into the higher altitudes.
+ Note also that equations (4)) and (5)) are applicable not only within the light cylinder but also outside of it and that B?«0 holds in ordinary situations., Note also that equations \ref{eq:v3D}) ) and \ref{eq:v3Df}) ) are applicable not only within the light cylinder but also outside of it and that $B^{\hat\varphi}<0$ holds in ordinary situations.
+" Integrating equation (4)) along particle motion, we obtain the path of the particle, which allows us to compute o. as a function of the distance s along the field line."," Integrating equation \ref{eq:v3D}) ) along particle motion, we obtain the path of the particle, which allows us to compute $\varrho_{\rm c}$ as a function of the distance $s$ along the field line."
+" Since the incorporation of a field-line deformation due to magnetospheric currents is out of the scope of this letter, we adopt a retarded, vacuum dipole field (Cheng et al."," Since the incorporation of a field-line deformation due to magnetospheric currents is out of the scope of this letter, we adopt a retarded, vacuum dipole field (Cheng et al."
+ 2000) as an example., 2000) as an example.
+" In figure 2,, we present θε/σσιις as a function of s for a representative field line in the leading side (LS; y.= 4-60?) andmagnetic that in the side y.= where c. refers to the magnetic trailingazimuthal (TS;angle measured—60°), around the magnetic axis counter-clockwise."," In figure \ref{fig:Rc2}, , we present $\varrho_{\rm c}/\rlc$ as a function of $s$ for a representative magnetic field line in the leading side (LS; $\varphi_\ast=+60^\circ$ ) and that in the trailing side (TS; $\varphi_\ast=-60^\circ$ ), where $\varphi_\ast$ refers to the magnetic azimuthal angle measured around the magnetic axis counter-clockwise."
+" The magnetic inclination angle, a, is assumed to be 60°."," The magnetic inclination angle, $\alpha$, is assumed to be $60^\circ$."
+" In figure 2,, the (red) thick solid and dashed curves denote ϱε/σσιιςο of the actual particle path, whereas the (black) thin solid and dashed curves do oc/cvrc of the last-open magnetic field lines."," In figure \ref{fig:Rc2}, the (red) thick solid and dashed curves denote $\varrho_{\rm c}/\rlc$ of the actual particle path, whereas the (black) thin solid and dashed curves do $\varrho_{\rm c}/\rlc$ of the last-open magnetic field lines."
+" The LS quantities are plotted with dashed curves, while the TS quantities with solid curves."," The LS quantities are plotted with dashed curves, while the TS quantities with solid curves."
+ It follows that ϱε/σσιις of the particle path (labeled with P-TS) becomes greater than that of the field line in the TS., It follows that $\varrho_{\rm c}/\rlc$ of the particle path (labeled with P-TS) becomes greater than that of the field line (B-TS) in the TS.
+" In the LS, on the contrary, θε/σσιις of the (B-TS)particle path (P-LS) becomes less than that of the field line (B-LS)."," In the LS, on the contrary, $\varrho_{\rm c}/\rlc$ of the particle path (P-LS) becomes less than that of the field line (B-LS)."
+" Note that B-TS is corrected to P-TS and B-LS to P-LS, and thatthe magnetic field lines B-LS vs. have greater o. in the LS whereas (seethe particle paths B-TS)(P-LS vs. P-TS) have greater Qc "," Note that B-TS is corrected to P-TS and B-LS to P-LS, and thatthe magnetic field lines (see B-LS vs. B-TS) have greater $\varrho_{\rm c}$ in the LS whereas the particle paths (P-LS vs. P-TS) have greater $\varrho_{\rm c}$ "
+that in modelling the position angle swing the bridge and main pulse emission should. be considered. as belonging to separate. if interlinked. svstenis.,"that in modelling the position angle swing the bridge and main pulse emission should be considered as belonging to separate, if interlinked, systems."
+ In considering this pulsar we adopt the same approach as with PSR 1929|10., In considering this pulsar we adopt the same approach as with PSR 1929+10.
+" The notches appear just 30°zz0.5 radians before the main pulse. but that does not imply that hey originate on leacing-edge fieldlines corresponding to the oesumed. pole at the main pulse phase. since retardation and aberration may have shifted them to earlier phase by considerably more than 30""."," The notches appear just $30^{o}\approx{0.5}$ radians before the main pulse, but that does not imply that they originate on leading-edge fieldlines corresponding to the presumed pole at the main pulse phase, since retardation and aberration may have shifted them to earlier phase by considerably more than $30^{o}$."
+ Reealline from. Section. 3 hat observed. phase shifts are formally. independent. of the unknown sina. and replacing 1.5 in equation (39) with 0.5 we maw attempt to see if the notches could. arise on the »oles last closed Gelcllines.," Recalling from Section 3.3 that observed phase shifts are formally independent of the unknown $\sin{\alpha}$, and replacing 1.5 in equation (39) with 0.5 we may attempt to see if the notches could arise on the pole's last closed fieldlines."
+ However the left. παπά side of (39) has to include the factor to express the mapping of the dipole coordinate onto the pulse phase: Equation (40) remains unchanged. but (41) must also reflect the mapping of c: We can attempt to solve (42). (43) and (40) to see if the notches can possible lie on the last. closecl fieldlines (s=1). allowing for solutions on both the leading and trailing sides of the pole. but we find that for all values of à the leading side solution is unphysical since the height of the emission would be at most 0.06 of the light evlinder radius. less than he separation of the notches. and there is no solution at al on the trailing side (ie retardation and aberration on tha side could never conspire to shift the emission to its observec hase).," However the left hand side of (39) has to include the factor $\frac{1}{\sin{\alpha}}$ to express the mapping of the dipole coordinate onto the pulse phase: Equation (40) remains unchanged, but (41) must also reflect the mapping of $\psi$: We can attempt to solve (42), (43) and (40) to see if the notches can possible lie on the last closed fieldlines (s=1), allowing for solutions on both the leading and trailing sides of the pole, but we find that for all values of $\alpha$ the leading side solution is unphysical since the height of the emission would be at most 0.06 of the light cylinder radius, less than the separation of the notches, and there is no solution at all on the trailing side (ie retardation and aberration on that side could never conspire to shift the emission to its observed phase)."
+ However. physically plausible solutions are possible if the emission lies on the quasi-radial fieldlines surrounding he pole (so that the ellects of curved fieldlines are minimisec and &« 1).," However, physically plausible solutions are possible if the emission lies on the quasi-radial fieldlines surrounding the pole (so that the effects of curved fieldlines are minimised and $k\ll{1}$ )."
+ For example. if we assume that the notches arise rom emission on the polar axis (s=0) then c—0 in (42) and r;&05r. so that the emission is at one quarter of he light evlinder radius.," For example, if we assume that the notches arise from emission on the polar axis (s=0) then $\psi=0$ in (42) and $r_{*}\approx{0.5r_{lc}}$, so that the emission is at one quarter of the light cylinder radius."
+ Again. this result. is independen of o.," Again, this result is independent of $\alpha$."
+" The polar region within which solutions are possible is eracdually reduced to a smaller anc smaller bunele of (trailing) fieldlines (s 1) as the inclination lessens. anc he location of the supposed absorbing region is shifted to 45"" or more. trailing the pole at a distance from the star several times the light evlinder radius."," The polar region within which solutions are possible is gradually reduced to a smaller and smaller bundle of (trailing) fieldlines $s<1$ ) as the inclination lessens, and the location of the supposed absorbing region is shifted to $45^{o}$ or more, trailing the pole at a distance from the star several times the light cylinder radius."
+ Thus we may conclude that the emission obscurecl in the notches is coming from within a few degrees of the polar axis. and at a height. well above the star surface.," Thus we may conclude that the emission obscured in the notches is coming from within a few degrees of the polar axis, and at a height well above the star surface."
+ Within this constraint there are many possible dipole alignments and locations of the absorbing site which can mimic the observations., Within this constraint there are many possible dipole alignments and locations of the absorbing site which can mimic the observations.
+" Establishing the correct configuration requires detailed modelling and comparison with polarisation data. but the suggestion of this model. supported by the very distinctive emission. stretching for some 150"" before the main pulse. is that we are dealing with a raclially-extencded band of weak emission. stretching [from the pole towards the light evlinder. possibly ending at the location of the absorbing region."," Establishing the correct configuration requires detailed modelling and comparison with polarisation data, but the suggestion of this model, supported by the very distinctive emission stretching for some $150^{o}$ before the main pulse, is that we are dealing with a radially-extended band of weak emission stretching from the pole towards the light cylinder, possibly ending at the location of the absorbing region."
+ I£ so. the observed interpulse must. define the upper limit of the band. being the closest point in the magnetosphere to the observer at the moment of emission. and the bridge of emission maps the band down to the polo.," If so, the observed interpulse must define the upper limit of the band, being the closest point in the magnetosphere to the observer at the moment of emission, and the bridge of emission maps the band down to the pole."
+" This configuration has the advantage of naturally explaining the delaved: main pulse to interpulse correlation. (Llankins Cordes 1981) - since the interpulse emission is the upper magnetosphere image of the polar activity - but may require sina to be substantially less than 1 - possibly around. 2. equivalent to an inclination of 30"". to vield the observed long stretch of emission."," This configuration has the advantage of naturally explaining the delayed main pulse to interpulse correlation (Hankins Cordes 1981) - since the interpulse emission is the upper magnetosphere image of the polar activity - but may require $\sin{\alpha}$ to be substantially less than 1 - possibly around $\frac{1}{2}$, equivalent to an inclination of $30^{o}$, to yield the observed long stretch of emission."
+ The notches observed in this pulsar are the most dilfieult to interpret because so little is known about the orientation aud structure of its magnetic field., The notches observed in this pulsar are the most difficult to interpret because so little is known about the orientation and structure of its magnetic field.
+ Is X-ray detection (Zavlin et al 2002) shows a single-hump time curve and is suggestive of a non-dipolar structure. and this picture has theoretical support [rom models of evolving magnetic fields in old neutron stars (Chen ct al 1998). which propose a nearly aligned. but. squeezed magnetic axis.," Its X-ray detection (Zavlin et al 2002) shows a single-hump time curve and is suggestive of a non-dipolar structure, and this picture has theoretical support from models of evolving magnetic fields in old neutron stars (Chen et al 1998), which propose a nearly aligned but ""squeezed"" magnetic axis."
+" Cil & Whrawezvk (1997) have. not cissimilarlyv. modelled the racio emission with an dipole inclined at. 20"". with onlv one pole visible to us."," Gil $\&$ Krawczyk (1997) have, not dissimilarly, modelled the radio emission with an dipole inclined at $20^{o}$, with only one pole visible to us."
+ However the observed: polarisation is complex. an gives no simple interpretation in terms of the rotating vector model (Radhakrishnan & Cooke 1969. Navarro et al 1997).," However the observed polarisation is complex and gives no simple interpretation in terms of the rotating vector model (Radhakrishnan $\&$ Cooke 1969, Navarro et al 1997)."
+ To explain. the notches noted ino ALR withou abanconing our absorption hypothesis we might be temptec o extend our intra-light evlinder assumption., To explain the notches noted in MR without abandoning our absorption hypothesis we might be tempted to extend our intra-light cylinder assumption.
+ Intriguingly e profile of Fig 2 in MEI reveals à number of single notches in addition to the double. notch. feature., Intriguingly the profile of Fig 2 in MR reveals a number of single notches in addition to the double notch feature.
+ Thus. we coulc dlow the absorbing agent to be the star itself ancl assume wt (he radiation is emitted. by downllowing particles in 16 back magnetosphere., Thus we could allow the absorbing agent to be the star itself and assume that the radiation is emitted by downflowing particles in the back magnetosphere.
+ Alternatively we may contrive ibsorbing regions outside the light evlinder which corotate gsuperluminally through which particles flow. counter. to 1 rotation., Alternatively we may contrive absorbing regions outside the light cylinder which corotate superluminally through which particles flow counter to the rotation.
+Without discarding these possibilities. we will nevertheless demonstrate below a self-consistent geometry which can still produce the couble notches within the ramework of our basic model.,"Without discarding these possibilities, we will nevertheless demonstrate below a self-consistent geometry which can still produce the double notches within the framework of our basic model."
+ The positioning of the notches in this pulsar is late (10)., The positioning of the notches in this pulsar is late $70^{o}$ ).
+ So i£ we wish the location of the obscured sources to be on open fieldlines above the visible pole then it is necessary ο assume near alignment in order to interpret the observed ohase as a ereathy expanded projection of the underlving dipole phase (0)., So if we wish the location of the obscured sources to be on open fieldlines above the visible pole then it is necessary to assume near alignment in order to interpret the observed phase as a greatly expanded projection of the underlying dipole phase $\psi$ ).
+" It is convenient to adopt the inclination a=20"" of Gil & Ixrawezyk (1997) to illustrate this point.", It is convenient to adopt the inclination $\alpha=20^{o}$ of Gil $\&$ Krawczyk (1997) to illustrate this point.
+ These authors also suggest. that. the polar cap emission height at 142Mllz is 16% of the light evlinder radius., These authors also suggest that the polar cap emission height at 1.4MHz is $16\%$ of the light cylinder radius.
+" Hence the clearlv-defined emission peak may be ahead of the ""true? pole by 10"" of the observed rotation phase."," Hence the clearly-defined emission peak may be ahead of the ""true"" pole by $10^{o}$ of the observed rotation phase."
+" This means that the notches appear GO""21.0redian after the pole for |20° of clipole angle after applying the factor sin20°= 1).", This means that the notches appear $60^{o}\approx{1.0}radian$ after the pole (or $+20^{o}$ of dipole angle after applying the factor $\sin{20^{o}}\approx{\frac{1}{3}}$ ).
+ ]lence we may replace [0.5 with -1.0 in equation (42). so that," Hence we may replace +0.5 with -1.0 in equation (42), so that"
+that a Coulomb correction and contributions from radiative recombination can be neglected.,that a Coulomb correction and contributions from radiative recombination can be neglected.
+ The diffusive escape of particles from the blast wave scales with the scattering length A as derived in Eq.(??))., The diffusive escape of particles from the blast wave scales with the scattering length $\lambda$ as derived in \ref{65}) ).
+ For the escape timescale we can writeEffame DOODTBd If escape is more efficient than the energy losses via pion production. then the bolometric luminosity of the blast wave will be reduced by a factor Tr/7.—TE|Ainal/>-," For the escape timescale we can write = = If escape is more efficient than the energy losses via pion production, then the bolometric luminosity of the blast wave will be reduced by a factor $T_{\rm E}/\tau_\pi \simeq T_{\rm E}\,\vert
+\dot \gamma_{\rm inel}\vert/\gamma$."
+ The proton spectra evolve differently depending on the relative effect of the losses and the blast wave slow down., The proton spectra evolve differently depending on the relative effect of the losses and the blast wave slow down.
+ If there were no losses. the relativistic mass loading Eq. (79))," If there were no losses, the relativistic mass loading Eq. \ref{massload}) )"
+ would be trivial to calculate., would be trivial to calculate.
+ Combining the mass loading equation with the deceleration equation Eq. (80)), Combining the mass loading equation with the deceleration equation Eq. \ref{decel}) )
+ we obtain after some calculus in the limit D1 where Ty is the initial Lorentz factor of the blast wave., we obtain after some calculus in the limit $\Gamma \gg 1$ = where $\Gamma_0$ is the initial Lorentz factor of the blast wave.
+ By comparison with the simple timescale Eq. (81)), By comparison with the simple timescale Eq. \ref{82}) )
+ we see that the mass loading reduces the deceleration by a factor Τη., we see that the mass loading reduces the deceleration by a factor $\Gamma/\Gamma_0$.
+" The proton continuity equation (84)) can be trivially integrated in the absence of losses to Gamme, So 1f the energy losses can be neglected. the blast wave deceleration alone would cause approximately à ~? spectrum in the relativistic range."," The proton continuity equation \ref{contp}) ) can be trivially integrated in the absence of losses to _0 So if the energy losses can be neglected, the blast wave deceleration alone would cause approximately a $\gamma^{-3}$ spectrum in the relativistic range."
+ Here we discuss a system of rather peculiar geometry. à very thin disk. and a strong angular variation of optical depth effects can be expected.," Here we discuss a system of rather peculiar geometry, a very thin disk, and a strong angular variation of optical depth effects can be expected."
+ While for photons traveling exactly in the blast wave plane the line-of-sight through the system has an average length R. and the system may appear optically thick. the angle-averaged line-of sight has a length ) which is approximately equal to 4.," While for photons traveling exactly in the blast wave plane the line-of-sight through the system has an average length $R$, and the system may appear optically thick, the angle-averaged line-of sight has a length = ) which is approximately equal to $d$."
+ Therefore the total optical depth may be small because the photons would interact only in a very small solid angle element., Therefore the total optical depth may be small because the photons would interact only in a very small solid angle element.
+ In all the cases presented here this is true for the Thompson optical depth τςστηςL. so that Compton scattering will not effectively influence the emission. though στηςδ may be of the order of I.," In all the cases presented here this is true for the Thompson optical depth $\tau =\sigma_{\rm T}\,
+n_{\rm e}\,
+{\bar L}$, so that Compton scattering will not effectively influence the emission, though $\sigma_{\rm T}\,n_{\rm e}\,R$ may be of the order of 1."
+ At photon energies above 511 keV in the blast wave photon-photon pair production may occur and inhibit the escape of from the system., At photon energies above 511 keV in the blast wave frame photon-photon pair production may occur and inhibit the escape of from the system.
+ This process ts not important in our model. because very high apparent luminosities can be produced with moderate photon densities in the source às à result of the highly relativistic motion of the blast wave and the associated high Doppler factors.," This process is not important in our model, because very high apparent luminosities can be produced with moderate photon densities in the source as a result of the highly relativistic motion of the blast wave and the associated high Doppler factors."
+ For example. given a system radius ?=101 em and a Doppler factor D=100. an apparent luminosity £L.=10/5oresec| at MeV energies in the blast wave rest frame corresponds to an optical depth of P.CX104.," For example, given a system radius $R=10^{15}$ cm and a Doppler factor $D=100$, an apparent luminosity $L=10^{48}\ {\rm erg\, sec^{-1}}$ at MeV energies in the blast wave rest frame corresponds to an optical depth of $\tau_{\gamma\gamma}\simeq 10^{-4}$."
+ In Fig., In Fig.
+ 3. we show the reaction and radiation channels of the swept-up particles together with the approximate photon energies for the emission processes., \ref{schema} we show the reaction and radiation channels of the swept-up particles together with the approximate photon energies for the emission processes.
+" Since the source power of secondary electrons and the power emitted in the form of z""- is the same. the bulk of the bolometric luminosity will be emitted in the GeV to TeV energy range. independent of the choice of parameters."," Since the source power of secondary electrons and the power emitted in the form of $\pi^0$ -decay is the same, the bulk of the bolometric luminosity will be emitted in the GeV to TeV energy range, independent of the choice of parameters."
+ Thus the unexpected finding. that emission of blazars generally provides a major part of the bolometric luminosity. would find a natural explanation if as in this model pion production in inelastic collisions were the main source of energetic electrons.," Thus the unexpected finding, that emission of blazars generally provides a major part of the bolometric luminosity, would find a natural explanation if as in this model pion production in inelastic collisions were the main source of energetic electrons."
+" In this paper we concentrate on the high-energy z-ray emission. and especially on the processes of z"" decay. sstrahllung. and pair annihilation. for they scale only with internal parameters like the gas density 5j."," In this paper we concentrate on the high-energy $\gamma$ -ray emission, and especially on the processes of $\pi^0$ decay, lung, and pair annihilation, for they scale only with internal parameters like the gas density $n_{\rm b}$."
+ A discussion of inverse Compton scattering and neutrino emission. will be presented elsewhere., A discussion of inverse Compton scattering and neutrino emission will be presented elsewhere.
+ The inverse Compton process must inevitably occur. but its rate. the spectral and angular distributiondepend on the actual choice of target photon field.," The inverse Compton process must inevitably occur, but its rate, the spectral and angular distributiondepend on the actual choice of target photon field."
+ If this scattering provides significant energy losses for electrons and positrons. the solution of Eq. (85))," If this scattering provides significant energy losses for electrons and positrons, the solution of Eq. \ref{conti}) )"
+ will be particularly arduous., will be particularly arduous.
+" The x"" production spectrum is calculated using the Carlo code DTUNUC described earlier in this paper.", The $\pi^0$ production spectrum is calculated using the Monte-Carlo code DTUNUC described earlier in this paper.
+ The bremsstrahlung spectrum is determined using the relativistic limit of the differential cross section (Blumenthal Gould 1970)) = where and the photon energy e is. in.units. ofC JD ο.," The bremsstrahlung spectrum is determined using the relativistic limit of the differential cross section (Blumenthal Gould \cite{bg70}) ) = + ] where and the photon energy $\epsilon$ is inunits of $m_{\rm e}\,c^2$ ."
+ The annihilation spectrum its obtained from the differential cross section (Jauch Rohrlich 1976)) κ., The annihilation spectrum is obtained from the differential cross section (Jauch Rohrlich \cite{jr76}) ) = - ]
+are available for download from abyss.uoregou.edtu/ js/network.,are available for download from abyss.uoregon.edu/ js/network.
+ E sincerely hope that the reader finds them useful. learus. moclifies them aud. most importantly. seuds me (jschombeeuoregoi.edu) feedback on what works (aud what broke) as well as ideas for [future tools aud research direction.," I sincerely hope that the reader finds them useful, learns, modifies them and, most importantly, sends me (jschombeuoregon.edu) feedback on what works (and what broke) as well as ideas for future tools and research direction."
+ This follows the model of a virtual community for computing knowledge., This follows the model of a virtual community for computing knowledge.
+ Perhaps the most comumon method of file trausfer aud retrieval is sftp (secure shell file transfer srotocol) using some version of OpeuSSH protocols., Perhaps the most common method of file transfer and retrieval is sftp (secure shell file transfer protocol) using some version of OpenSSH protocols.
+ This is certainly the most popular method of rausferring data (rom one [δν system to another (friendly meaning you coutrol both systems: e.g.. transfer of data fromthe telescope to your ollice computer).," This is certainly the most popular method of transferring data from one friendly system to another (friendly meaning you control both systems; e.g., transfer of data from the telescope to your office computer)."
+ A script that uses sftp would take he place of time consuimiug command line typing aud. usiug Pythou’s ability to search and. parse ile directories.," A script that uses sftp would take the place of time consuming command line typing and, using Python's ability to search and parse file directories."
+ Even a simall script eases the laborious typing required to trausler a mixture of file vpes., Even a small script eases the laborious typing required to transfer a mixture of file types.
+ The simplest manner to communicate to the shell [rom Python is the use of the Command the (althoughmuch more complicated module is nowrecommended)., The simplest manner to communicate to the shell from Python is the use of the command (although the much more complicated module is now recommended).
+ The following example pushes a buuch of FITS files onto a remote system by making a temporary file with sftp commands. then uses to send a ‘sftp -b commands to the shell.," The following example pushes a bunch of FITS files onto a remote system by making a temporary file with sftp commands, then uses to send a 'sftp -b' commands to the shell."
+ This type of operation works well as a cron job for moving backup files duriug off hours. or any background task that doesu't require liuunan monitoring.," This type of operation works well as a cron job for moving backup files during off hours, or any background task that doesn't require human monitoring."
+ However. this type of script is clumsy (e.g.. requires the creation and destruction of the temporary file. sftp.cmds).," However, this type of script is clumsy (e.g., requires the creation and destruction of the temporary file, )."
+ A smoother interface uses the module (sourcelorge.uet/projects/pexpect). as in the following example," A smoother interface uses the module (sourceforge.net/projects/pexpect), as in the following example:"
+Angular Dillerential Imaging Direct detections of very faint exoplanets and brown dwarls near bright stars are essential to understand substellar formation and evolution around stars.,Angular Differential Imaging Direct detections of very faint exoplanets and brown dwarfs near bright stars are essential to understand substellar formation and evolution around stars.
+ This endeavor is now one of the major goals for next generation 10-11 telescope instrumentis and future 30- to. telescopes., This endeavor is now one of the major goals for next generation 10-m telescope instruments and future 30- to 100-m telescopes.
+ The task is dauntünelv difficult., The task is dauntingly difficult.
+ The exoplanet or brown dwarf mage is usually much Jainter than the background [rom the brilliant stellar point spread. function (PSF) image., The exoplanet or brown dwarf image is usually much fainter than the background from the brilliant stellar point spread function (PSF) image.
+ Besides the Poisson noise. ground-based telescopes suffer from. atmospheric turbulence that produces random short-livecl speckles that mask faint companions.," Besides the Poisson noise, ground-based telescopes suffer from atmospheric turbulence that produces random short-lived speckles that mask faint companions."
+ II these (wo limitations were (he only ones. a simple solution would be to integrate longer to average (hese random noises ancl gain as the square rool of the integration time.," If these two limitations were the only ones, a simple solution would be to integrate longer to average these random noises and gain as the square root of the integration time."
+ Bul observations have shown that. for integrations longer than a few minutes. the PSF noise converges to a quasi-static noise pattern. thus preventing a gain wilh increasing integration time (Alaroisοἱ2005:Masciadrietal. 2005).," But observations have shown that, for integrations longer than a few minutes, the PSF noise converges to a quasi-static noise pattern, thus preventing a gain with increasing integration time \citep{marois2003,marois2005,masciadri2005}."
+. To achieve better detection limits. itis (hus necessary to sublract the quasi-static noise using a reference PSF.," To achieve better detection limits, it is thus necessary to subtract the quasi-static noise using a reference PSF."
+ Both eround- ancl space-based imaging are plagued with this stellar PSF calibration problem caused by imperfect optics ancl slowly evolving optical aligninents., Both ground- and space-based imaging are plagued with this stellar PSF calibration problem caused by imperfect optics and slowly evolving optical alignments.
+ For ground-based imaging. subtraction of a reference PSF obtained from a star close to the target achieves a [actor ~4 of PSF noise attenuation. leaving residuals (hat are also quasi-static and (hus severely limiting detection of fainter companions (Alaroisοἱal.2005).," For ground-based imaging, subtraction of a reference PSF obtained from a star close to the target achieves a factor $\sim 4$ of PSF noise attenuation, leaving residuals that are also quasi-static and thus severely limiting detection of fainter companions \citep{marois2005}."
+. For space telescopes (hat have a better PSF stability. like IST. a partial solution was found by subtracting (wo stellar images acquired. during the same orbit with a different roll angle.," For space telescopes that have a better PSF stability, like HST, a partial solution was found by subtracting two stellar images acquired during the same orbit with a different roll angle."
+ This technique. called. “roll deconvolution”. successfully subtracts the stellar image by a factor 50 (Schneider&Silverstone2003:Fraquellietal2004) but is also ultimately limited by PSF evolution.," This technique, called “roll deconvolution”, successfully subtracts the stellar image by a factor 50 \citep{schneider2003,fraquelli2004} but is also ultimately limited by PSF evolution."
+ A similar technique. called angular dilferential imaging (ADI). can be used on ground-based. altitude/azimuth telescopes to subtract a signilicant fraction of the stellar quasistatic noise and can potentially achieve detection limits (hat improve as (he square root of the integration (ine.," A similar technique, called angular differential imaging (ADI), can be used on ground-based altitude/azimuth telescopes to subtract a significant fraction of the stellar quasi-static noise and can potentially achieve detection limits that improve as the square root of the integration time."
+source has been added to Table |.,source has been added to Table 1.
+ There were no X-ray sources apparent at the positions of the other 28 objects and X-ray flux upper limits were derived as follows., There were no X-ray sources apparent at the positions of the other 28 objects and X-ray flux upper limits were derived as follows.
+ In the task. we generated images for each detector consisting of smooth models of the X-ray background. to which were added of the significantly detected sources calculated from their count rates and the model point spread function used to detect and parameterise them.," In the task, we generated images for each detector consisting of smooth models of the X-ray background, to which were added of the significantly detected sources calculated from their count rates and the model point spread function used to detect and parameterise them."
+ These images were “noise-free” estimates of the expected X-ray background for any given position., These images were “noise-free” estimates of the expected X-ray background for any given position.
+ The total expected background was summed within circles of radius 20 areseconds surrounding each of the undetected Irwin et al. (, The total expected background was summed within circles of radius 20 arcseconds surrounding each of the undetected Irwin et al. (
+2008) objects.,2008) objects.
+ The number of observed X-ray counts in those areas in the, The number of observed X-ray counts in those areas in the
+based on the data from LB. Denevactal.2009.. Iseaneetal.2010.. Durke-Spolaor&Bailes 2010)). or the burst's properties were nisiuterpreted.,"based on the data from LB, \citealt{deneva}, , \citealt{keane}, \citealt{sbsmb}) ), or the burst's properties were misinterpreted."
+ We announce the detection of 16 pulses that exhibit a frequency-swept signal with simular characteristics to the Loriuer Burst., We announce the detection of 16 pulses that exhibit a frequency-swept signal with similar characteristics to the Lorimer Burst.
+ However. 1[umese new events are clearly of terrestrial origiu with properties unlike anv known sources of broad-xui radio cussion.," However, these new events are clearly of terrestrial origin with properties unlike any known sources of broad-band radio emission."
+ We detail the properties of the new doetectious and scrutinize them iu conmgxuisou to the LB., We detail the properties of the new detections and scrutinize them in comparison to the LB.
+ We also discuss methods hat can be used to distinguish au astrophysical or errestrial origin for future detections of frequeney-swept pulses., We also discuss methods that can be used to distinguish an astrophysical or terrestrial origin for future detections of frequency-swept pulses.
+ Prompted by the ciscovery of the Lorimer Durst. we searched 1078 hours of data for signals exhibiting a dtxv? cold plana dispersion cela.," Prompted by the discovery of the Lorimer Burst, we searched 1078 hours of data for signals exhibiting a $\delta t\propto\nu^{-2}$ cold plasma dispersion delay."
+ The data were archival pulsar surveys aken over the vears 1998 to 2003 with the 20cm uultibeanà receiver installed at Parkes Telescope., The data were archival pulsar surveys taken over the years 1998 to 2003 with the 20cm multibeam receiver installed at Parkes Telescope.
+ Four surveys were searched., Four surveys were searched.
+ For two (Edwards we followed he search aud iuspection process described by Durke-Spolaor&Bailes(2010).. differing only iu interference iitigation filters: we inspected a candidate if it either a) was detected at a signal-to-noise ratio of 26 in less than 9/13 of the receivers. b) had a signal-to-noise ratio of &20. or c) showed a DM of higher than >.," For two \citep{SM,BJ} we followed the search and inspection process described by \citet{sbsmb}, differing only in interference mitigation filters; we inspected a candidate if it either a) was detected at a signal-to-noise ratio of $>6$ in less than 9/13 of the receivers, b) had a signal-to-noise ratio of $\geq20$, or c) showed a DM of higher than $^{-3}$."
+ As these were surveys of düsh Galactic latitude (Jb)> 57). we searched DMs up to 7.," As these were surveys of high Galactic latitude $|b| >
+5^\circ$ ), we searched DMs up to $^{-3}$."
+ For the remaining two surveys (Crawtordetal.2006:Roberts2002.their2aTable1. respectivelv).. we used methods described by (Crawfordetal.2007.2009) applying uo interference excision based on imultiple-beain detections.," For the remaining two surveys \citep[][their observing parameters are detailed in their Sec. 2 and
+Table 1, respectively]{froneydata1,froneydata2}, we used methods described by \citep{froneysearch1,froneysearch2} applying no interference excision based on multiple-beam detections."
+ These were generally surveys of lower ealactic latitude. therefore we searched DAIs up to 1000. (Crawfordetal.2006).. 2000 (for tarect AN JIs26.11300 of the Roberts et al;," These were generally surveys of lower galactic latitude, therefore we searched DMs up to 1000 \citep{froneydata1}, 2000 (for target AX J1826.1Ð1300 of the Roberts et al."
+ survey) aud 2 (for the other tarects of Roberts et al.)., survey) and $^{-3}$ (for the other targets of Roberts et al.).
+" Candidates from these surveys were inspected oue bea at a tine by eve to identify dispersed pulses,", Candidates from these surveys were inspected one beam at a time by eye to identify dispersed pulses.
+ Qur search revealed 16 pulses with two striking features that distineuish them from all others in the data: an apparent fX»7 delay of a maenitude implying an extragalactic origin im the telescopes poiting direction. aud a simultaucous occurrence in all 13 telescope receivers at relative intensitics of less than a factor of four (Fie.," Our search revealed 16 pulses with two striking features that distinguish them from all others in the data: an apparent $\delta t\propto\nu^{-2}$ delay of a magnitude implying an extragalactic origin in the telescope's pointing direction, and a simultaneous occurrence in all 13 telescope receivers at relative intensities of less than a factor of four (Fig."
+ 1)., 1).
+ When a dispersive delay is fit to cach etection. the values cluster about a uct baud delay Af=360 πας. jucdicating a close counectiou with the LB at Af=355 nuus (Fie.," When a dispersive delay is fit to each detection, the values cluster about a net band delay $\Delta t = 360$ ms, indicating a close connection with the LB at $\Delta t =
+355$ ms (Fig."
+ 2)., 2).
+ The LB’s reported sky position was below the horizon for several detections. therefore the pulses could not have conie from the same extragalactic source.," The LB's reported sky position was below the horizon for several detections, therefore the pulses could not have come from the same extragalactic source."
+ Below. we ee evidence that the 16 signals have a terrestrial origin.," Below, we give evidence that the 16 signals have a terrestrial origin."
+ The 29 separation between cach receivers bean position aud >20ddB attcunation bevoud 30° from cach beam center render it inpossible for an on-axis. poiutlike signal to appear iu more than 3 bemus at simnilar intensity (Stavelev-Suuüthetal.1996:ITuut&Wrieht 1992).," The $^{'}$ separation between each receiver's beam position and $>20$ dB attenuation beyond $^{'}$ from each beam center render it impossible for an on-axis, pointlike signal to appear in more than 3 beams at similar intensity \citep{multibeam,feedlegs}. ."
+". Out detections were therefore made through a sidelobe of the Parkes antenna. aud based ou the consistency of sjeuals in the beams. in each case the enütter was positioned sl,295 degrees from the telescope’s »mÓtiue direction."," Out detections were therefore made through a sidelobe of the Parkes antenna, and based on the consistency of signals in the beams, in each case the emitter was positioned $\gg$ 5 degrees from the telescope's pointing direction."
+ Consequeutlv: 1) we lave a uiniuiun-senisitivitv fieldof view of ~20000 dclees. a detectionκ rate aq2.3.«414-*deg2Slyο1 (0.14 | with our observing svsteni). aud poor accuracy or eudtter localisation. 2) the pulses are subject o frequency-depeudent dropouts and scattering roni multi-path propagationvisible iu Fie. 1(b-dj," Consequently: 1) we have a minimum-senisitivity field of view of $\sim$ $^{2}$, a detection rate $2.3\times10^{-7}\,{\rm deg}^{-2}\,{\rm hr}^{-1}$ (0.1 $^{-1}$ with our observing system), and poor accuracy for emitter localisation, 2) the pulses are subject to frequency-dependent dropouts and scattering from multi-path propagation—visible in Fig. 1(b-d)"
+" aud explored belowand 3) the source(s] intrinsic Fux deusity, had we pointed directly at it. is a factor of 2500-850000 ercater than the detected value (0:8πμcz272 kkJv for the xiehtest detection. 0.1<Signs BlissJv for he faimtest: see Table 1)."," and explored below—and 3) the source(s)' intrinsic flux density, had we pointed directly at it, is a factor of 2500-850000 greater than the detected value $0.8 < S_{intrinsic} <
+272$ kJy for the brightest detection, $0.1 < S_{intrinsic} < 34$ kJy for the faintest; see Table 1)."
+ We conclude. a terrestrial origin for these sts based on their extreme brightuess and wo other features., We conclude a terrestrial origin for these bursts based on their extreme brightness and two other features.
+ First. some exhibit deviations roni a model dispersive delav: e.g. the sharp sinks at 1165 MIIz indetection 06 (Fig.," First, some exhibit deviations from a model dispersive delay: e.g. the sharp kink at 1465 MHz indetection 06 (Fig."
+ 1c). aud subtler deviations in otherdetections (see the 4? isting in Table 1).," 1c), and subtler deviations in otherdetections (see the $\chi^2$ listing in Table 1)."
+ Despite a trend mimicking hat expected frou dispersion. such deviations," Despite a trend mimicking that expected from dispersion, such deviations"
+We consider now whether this is the final word on evolving gravitational fields.,We consider now whether this is the final word on evolving gravitational fields.
+ absorbers were the most obvious source of foreground. interference: they. by definition. interseet the line of sight.," absorbers were the most obvious source of foreground interference; they, by definition, intersect the line of sight."
+ Galactic halocs at impact parameters of 120 kpe are thought το produce the rarer absorption elfects. such as LLSs anc DLAs: these have been the extreme cases studied: in this work.," Galactic haloes at impact parameters of 1–20 kpc are thought to produce the rarer absorption effects, such as LLSs and DLAs; these have been the extreme cases studied in this work."
+ Estimates of the frequeney ancl redshift distribution of these objects are hindered by reddening. which can dim quasars and exelude them from lux-limited samples. although the magnification bias can counteract this ellect.," Estimates of the frequency and redshift distribution of these objects are hindered by reddening, which can dim quasars and exclude them from flux-limited samples, although the magnification bias can counteract this effect."
+ Both effects can be expected. particularly if the absorbers are. massive galaxy haloes at low impact parameters.," Both effects can be expected, particularly if the absorbers are massive galaxy haloes at low impact parameters."
+ The resulting evolution of the gravitational field within the filaments and sheets associated with the low-density forest is cillicult ο study analytically., The resulting evolution of the gravitational field within the filaments and sheets associated with the low-density forest is difficult to study analytically.
+ Ehe filaments feed into more massive structures. such as galaxy clusters. with absorption taking ace at impact parameters of ~1 Alpe.," The filaments feed into more massive structures, such as galaxy clusters, with absorption taking place at impact parameters of $\sim$ 1 Mpc."
+" Vhe transverse »eculiar motion of such objects will have a similar ""moving-ens. elfect. the rarity of which has not. been assessed."," The transverse peculiar motion of such objects will have a similar `moving-lens' effect, the rarity of which has not been assessed."
+ Their presence may. not be noticed in the quasar spectra. hough perhaps in direct imaging (at least at low redshilt).," Their presence may not be noticed in the quasar spectra, though perhaps in direct imaging (at least at low redshift)."
+ Further stucies using dark matter simulations will probe he expected motions of all mass scales in the vicinity of the ine of sight to distant quasars. allowing the determination of the full influence of the DG clleet in future redshift drift experiments: this will form the basis of future contributions.," Further studies using dark matter simulations will probe the expected motions of all mass scales in the vicinity of the line of sight to distant quasars, allowing the determination of the full influence of the BG effect in future redshift drift experiments; this will form the basis of future contributions."
+ The authors would like to thank the anonymous referee for the helpful comments ancl suggestions., The authors would like to thank the anonymous referee for the helpful comments and suggestions.
+ acknowledge support from ARC Discovery Project ADIP0665574., They acknowledge support from ARC Discovery Project DP0665574.
+".""M MIN. is also supported. by the University. of Faculty of Science. Postgraduate Award ancl would like to thank Richard Lane for his feedback.", M.K. is also supported by the University of Sydney Faculty of Science Postgraduate Award and would like to thank Richard Lane for his feedback.
+that the ült (or N-S dipole moment) is already. established in (he source region inside the sun and not during the buovant rise of the flix loops.,that the tilt (or N-S dipole moment) is already established in the source region inside the Sun and not during the buoyant rise of the flux loops.
+ The present work supports Chis conclusion by showing. with much enhanced statistics. that the average tilt does not change appreciably as we go to regions with orders of magnitude smaller flux contents.," The present work supports this conclusion by showing, with much enhanced statistics, that the average tilt does not change appreciably as we go to regions with orders of magnitude smaller flux contents."
+ An even much more ingrained. long-term paradigm that is shown to be untenable by the present work is the phenomenological scenario. according to which differential rotation creates Irom a poloidal field a coherent. toroidal flux svstem. from which (he sunspots arise.," An even much more ingrained, long-term paradigm that is shown to be untenable by the present work is the phenomenological scenario, according to which differential rotation creates from a poloidal field a coherent toroidal flux system, from which the sunspots arise."
+ Such a scenario is nol compatible with (he violations of Hale's polarity law that we have presented here., Such a scenario is not compatible with the violations of Hale's polarity law that we have presented here.
+ Our illustrated exemples where well defined mecdiumn-size bipolar magnetic regions occur side by side unambiguouslv show that these oppositely oriented regions caunot be part of the same flux svstem. but that there must be a coexistence of oppositelv oriented toroidal Πας in the same latitude zones al any given time.," Our illustrated examples where well defined medium-size bipolar magnetic regions occur side by side unambiguously show that these oppositely oriented regions cannot be part of the same flux system, but that there must be a coexistence of oppositely oriented toroidal flux in the same latitude zones at any given time."
+ Although the occurrencies of violations of IHale's polarity law represent only a few percent of all cases. thev rule out the possibility of well defined. coherent toroidal flux svslenis as a source of all active regions. even the large ones.," Although the occurrencies of violations of Hale's polarity law represent only a few percent of all cases, they rule out the possibility of well defined, coherent toroidal flux systems as a source of all active regions, even the large ones."
+ Our results make clear thatfluctuations represent an essential inherent physical property of the solar dynanmo. as expected to various degrees [rom turbulent dynamo (theory 2005).," Our results make clear that represent an essential inherent physical property of the solar dynamo, as expected to various degrees from turbulent dynamo theory \citep{stenflo-brandsub05}."
+. The fundamental role of large fluctuations at all scales is further illustrated in Fig., The fundamental role of large fluctuations at all scales is further illustrated in Fig.
+" for four selected. cases of well defined bipolar magnete regions. including large ones. which have orientations (hat differ by 90° - 1007"" from the orientations prescribed by Joys law."," \ref{fig:90deg} for four selected cases of well defined bipolar magnetic regions, including large ones, which have orientations that differ by $90^\circ$ - $100^\circ$ from the orientations prescribed by Joy's law."
+" The existence of many such cases shows (hat we have fluctuations over the whole range of orientation angles. not onlv between. ""Hale and anti-llale (proper ancl reversed) orientations."," The existence of many such cases shows that we have fluctuations over the whole range of orientation angles, not only between “Hale and anti-Hale” (proper and reversed) orientations."
+ Also the largest bipolar regions are subject to these fHuctuations., Also the largest bipolar regions are subject to these fluctuations.
+Therefore. let us reconsider the critical black hole mass such that a neutron star is disrupted outside the ISCO.,"Therefore, let us reconsider the critical black hole mass such that a neutron star is disrupted outside the ISCO."
+ For mpifans=2 aud f=5mxa. Ry=3.35.," For $m_{\rm BH}/m_{\rm NS}=2$ and $R_0=5m_{\rm NS}$, ${\hat R}_0=3.15$ ."
+ At this F6). Wiggeius Lai (2000) fine Pide©3 and heuce r/mpi=7.25.," At this ${\hat R}_0$, Wiggins Lai (2000) find ${\hat r}_{\rm tide}\approx 3$ and hence $r/m_{\rm BH}=7.25$."
+ Therefore. r/Al=ιδ and this orbit is actually inside the ISCO.," Therefore, $r/M=4.83$, and this orbit is actually inside the ISCO."
+ Consider now a miuinal mass black hole of py=2.2M. in a binary with a neutrou star OL nixa=LEM. aud Rp=5mxa.," Consider now a minimal mass black hole of $m_{\rm BH}=2.2\,M_\odot$ in a binary with a neutron star of $m_{\rm NS}=1.4\,M_\odot$ and $R_0=5m_{\rm NS}$."
+ Then Ry23.70. Mide=2.27 according to Wigeins Lai (2000) for n=1. and thus r/impy=8.10.," Then ${\hat R}_0=3.70$, ${\hat r}_{\rm tide}=
+2.27$ according to Wiggins Lai (2000) for $n=1$, and thus $r/m_{\rm BH}=8.40$."
+ This implies r/A/=5.13., This implies $r/M=5.13$.
+ This is technically just outside the ISCO. but as we will see in the next section. augularOm momenttun loss to egravitational radiatiou is sienificaut and the plunge actually starts well outside this radius.," This is technically just outside the ISCO, but as we will see in the next section, angular momentum loss to gravitational radiation is significant and the plunge actually starts well outside this radius."
+ Higher mass ueutron stars are more compact (with Ry« μηνα). aud will also not be tidally disrupted.," Higher mass neutron stars are more compact (with $R_0<5m_{\rm NS}$ ), and will also not be tidally disrupted."
+ Only low mass stars with Ry>Sexy might leave accretion disks. but as we now show. even this is not certain.," Only low mass stars with $R_0>5m_{\rm NS}$ might leave accretion disks, but as we now show, even this is not certain."
+ The ISCO as a sharp dividing liue is only strictly useful in the test particle limit. where uo other elfects could cause inspiral," The ISCO as a sharp dividing line is only strictly useful in the test particle limit, where no other effects could cause inspiral."
+ Any mechanism that leads to inspiral of nearly circular orbits. whether nagnetic angular momeutunm transport iu accretion disks or losses to gravitational radiatiou. will blur this line (see. e.g. Buonanuo Damour 2000 for a detailed discussion ofthis phenomenon).," Any mechanism that leads to inspiral of nearly circular orbits, whether magnetic angular momentum transport in accretion disks or losses to gravitational radiation, will blur this line (see, e.g., Buonanno Damour 2000 for a detailed discussion of this phenomenon)."
+ Near to but outside of the ISCO. the specific angular momentuim of circular orbits is very [lat with radius. meaning that even a stall amount of angular momentum loss can have drastic couseqtuences (see. e.g.. Lombardi. Rasio. Shapiro 1997).," Near to but outside of the ISCO, the specific angular momentum of circular orbits is very flat with radius, meaning that even a small amount of angular momentum loss can have drastic consequences (see, e.g., Lombardi, Rasio, Shapiro 1997)."
+ In particular. even after the neutron star bas reached the tidal radius. the matter that is stripped will take some finite time to separate siguilicantly [rom the star. aud curie this time all the matter is still emitting gravitational radiation.," In particular, even after the neutron star has reached the tidal radius, the matter that is stripped will take some finite time to separate significantly from the star, and during this time all the matter is still emitting gravitational radiation."
+ We suspect that the omission of these effects by Prakash. Ratkovic. Lattimer (2001). who cid a semianalytic treatment of BH-NS mergers uxing the Blanchet (2002) post-Newtouian equations. is the reason that they reached qualitatively different. conclusions from ours.," We suspect that the omission of these effects by Prakash, Ratkovic, Lattimer (2004), who did a semianalytic treatment of BH-NS mergers using the Blanchet (2002) post-Newtonian equations, is the reason that they reached qualitatively different conclusions from ours."
+ Cousider. for example. loss to gravitational radiation.," Consider, for example, loss to gravitational radiation."
+ From Peters Mathews (1963) aud Peters (196D. the lowest-order (euadrupolar) loss rate of angular momeuttun for a nearly. circular orbit is Let the angular momentum at the ISCO be Ljaco.," From Peters Mathews (1963) and Peters (1964), the lowest-order (quadrupolar) loss rate of angular momentum for a nearly circular orbit is Let the angular momentum at the ISCO be $L_{\rm ISCO}$ ."
+ Hf enough augular momentum is lost iu a time short compared to the orbital period so that £L«Lycee. then the black hole aud weutrou star will plunge together.," If enough angular momentum is lost in a time short compared to the orbital period so that $L<L_{\rm ISCO}$, then the black hole and neutron star will plunge together."
+ For the hypothetical 1.LAZ.—2.2M. binary discussed above. an initial circular orbit of radius 5.13A has an angular momentum only 0.00078//M. greater than that of the ISCO.," For the hypothetical $1.4\,M_\odot - 2.2\,M_\odot$ binary discussed above, an initial circular orbit of radius $5.13\,M$ has an angular momentum only $0.00078\mu M$ greater than that of the ISCO."
+ From the formula above for angular momentum loss. this is radiated within a time 0.164. compared to an orbitalperiod of zz75M at that radius.," From the formula above for angular momentum loss, this is radiated within a time $0.16M$, compared to an orbitalperiod of $\approx 75\,M$ at that radius."
+ Clearly. the binary undergoes a plunge well before theneutron star is tidally stripped.," Clearly, the binary undergoes a plunge well before theneutron star is tidally stripped."
+We adopted the modified TAT (1976) value of L.1959787066 «10 Maui (Astronomical Almanac. 1997) for the astronomical unit. and adjusted this value by -L678 Mii to account for the mean displacement between the telescope's: noontine: location: and the Τατκ center. aud by |0.119 Maii for the displacement of the Suus center relative to the barvceuter of the Earth-Sun svstem.,"We adopted the modified IAU (1976) value of 1.4959787066 $\times 10^5$ Mm (Astronomical Almanac, 1997) for the astronomical unit, and adjusted this value by -4.678 Mm to account for the mean displacement between the telescope's noontime location and the Earth's center, and by +0.449 Mm for the displacement of the Sun's center relative to the barycenter of the Earth-Sun system."
+ This distance. combined with Dy from Eq. (," This distance, combined with $D_0$ from Eq. ("
+2). vields the Sts apparent radius.,"2), yields the Sun's apparent radius."
+ Applving the model corrections described in the last section. we obtain We estimate the modeling errors to be 1/2 of the difference between these estimates. or about 0.020 Alu.," Applying the model corrections described in the last section, we obtain We estimate the modeling errors to be $1/\sqrt 2$ of the difference between these estimates, or about 0.020 Mm."
+" Based on the uncertainties in the geometric corrections that were made to the measured radius. we estimate the systematic errors in the measured value to be 0.015 Min. or about twice as large as the random errors,"," Based on the uncertainties in the geometric corrections that were made to the measured radius, we estimate the systematic errors in the measured value to be 0.015 Mm, or about twice as large as the random errors."
+ Averaging our results for Models 1 aud 2. and adding the various error sources in quadrzature. we arrive at our final estimate of The interred solar photospherie radius is πια]: bv about 0.5MBa than the normally used value of 695.99M (Allen 19076).," Averaging our results for Models 1 and 2, and adding the various error sources in quadrature, we arrive at our final estimate of The inferred solar photospheric radius is smaller by about $0.5 \Mm$ than the normally used value of $695.99 \Mm$ (Allen 1976)."
+" À reviewof recent observations was given by Schou (L997). concluding tha these were cousisteut with an angular diameter of 1919.26""£ 0.2"". correspouding to Allen's value of R.."," A review of recent observations was given by Schou (1997), concluding that these were consistent with an angular diameter of $1919.26\arcsec \pm 0.2\arcsec$ , corresponding to Allen's value of $\Rsun$."
+ This is also consistent with the observed value obtained here (ef., This is also consistent with the observed value obtained here (cf.
+ eq., eq.
+ 2)., 2).
+ However. it appears tha the observations considered by Schou refer to the inflection poiut of intensity for. in one ease. to an FFTD determination) aud hence do not contain the correction to photospheric radius.," However, it appears that the observations considered by Schou refer to the inflection point of intensity (or, in one case, to an FFTD determination) and hence do not contain the correction to photospheric radius."
+ Such a correction. iuto account characteristics. Is an talsingessential part of the radius observationaldetermination.," Such a correction, taking into account the observational characteristics, is an essential part of the radius determination."
+ Some . confirmationof. the∙∙∙ reliabilityof. the modeling. comes frou. the coumparisou in Fig., Some confirmation of the reliability of the modeling comes from the comparison in Fig.
+ 2 of computed and observed slopes of tle Iimb position as function of the scan widths., 2 of computed and observed slopes of the limb position as function of the scan widths.
+ Nevertheless. it is striking that. as indicated by the difference between Models 1 aud 2. the major uncertainty in R&R. appears to come roni the modeling.," Nevertheless, it is striking that, as indicated by the difference between Models 1 and 2, the major uncertainty in $\Rsun$ appears to come from the modeling."
+ Iudeed. it is evident that the real solar atmosphere is substantially more complicated han the one-dimensional model resulting from the ATLAS code ov the mean model obtained from the wdrodvuamucal simmlatious.," Indeed, it is evident that the real solar atmosphere is substantially more complicated than the one-dimensional model resulting from the ATLAS code or the mean model obtained from the hydrodynamical simulations."
+ A more accurate determination of the radius correction cau probably be obtained youn a detailed calculation of the limi iuteusity. alius iuto account the inuhomiogeueous nature of the relevant lavers. ou the basis of the simulations.," A more accurate determination of the radius correction can probably be obtained from a detailed calculation of the limb intensity, taking into account the inhomogeneous nature of the relevant layers, on the basis of the simulations."
+ Such an investigation is bevoud the scope of the preseut oper. however.," Such an investigation is beyond the scope of the present paper, however."
+ We find no significant variation in the observed diameter during the observation period (cf., We find no significant variation in the observed diameter during the observation period (cf.
+ Fig., Fig.
+ 2): anuual averages of the radius for the vears 1981 to 1987 all agree within their measurement errors of +#.037 Mau., 2); annual averages of the radius for the years 1981 to 1987 all agree within their measurement errors of $\pm .037$ Mm.
+ These limits are substantially smaller han diuneter changes reported previously for the sae interval of time (ee. Ulrich. Dertello 1995. Laclare 1996). but are in agreement with measurements o» Wittinan (1997).," These limits are substantially smaller than diameter changes reported previously for the same interval of time (e.g. Ulrich Bertello 1995, Laclare 1996), but are in agreement with measurements by Wittman (1997)."
+" On the other hand. the Πο-»oxition slope shows fairly substautial variations,"," On the other hand, the limb-position slope shows fairly substantial variations."
+ We also note that during solu maxima. the dailv slope values tended to be highly variable as well as sinall iu uaenitude: this suggests that the loug-terii variation nav result from localized activity-depeudent features such as faculae.," We also note that during solar maximum, the daily slope values tended to be highly variable as well as small in magnitude; this suggests that the long-term variation may result from localized activity-dependent features such as faculae."
+ Tt is pliραible that the previously inferred variations iu solar diameter with solar activity is in fact a reflection of such variatious in the linib-darkening slope., It is plausible that the previously inferred variations in solar diameter with solar activity is in fact a reflection of such variations in the limb-darkening slope.
+ It is interesting that the value of AR. obtained here is somewhat smaller than that interred: fron the solar fauode frequencies. indicating additional contributions to the differences between the observed aud model values of these frequencies.," It is interesting that the value of $\Rsun$ obtained here is somewhat smaller than that inferred from the solar f-mode frequencies, indicating additional contributions to the differences between the observed and model values of these frequencies."
+ This issue. aud the effects of the reduction of the model radius on the lehoscisunically deteruued structure of the solar interior will be considered elsewhere.," This issue, and the effects of the reduction of the model radius on the helioseismically determined structure of the solar interior will be considered elsewhere."
+ Wenote. however. that Antia (1998) and Schon," Wenote, however, that Antia (1998) and Schou"
+observation to be made from this plot is the fact that the distribution is fitted well for intrinsic redshift values z;X2. which is in excellent aerecluent with the range of intrinsic redshift values observed for QSOs clustered near active ealaxies (Arp1983:Chuetal.1998:Burbidge1999).,"observation to be made from this plot is the fact that the distribution is fitted well for intrinsic redshift values $_{\rm i} \lesssim2$, which is in excellent agreement with the range of intrinsic redshift values observed for QSOs clustered near active galaxies \citep{arp83,chu98,bur99}."
+". Tu other words. fle current dogez - an. distribution of sources can be fitted exactly if the short-lived QSOs have intrinsic redshifts àz2 (identical to those found for sources clustered near active yalarics), together with a distance redshift component defined by the Hubble relation."," In other words, the current logcz - $_{\rm v}$ distribution of sources can be fitted exactly if the short-lived QSOs have intrinsic redshifts $_{\rm i} \lesssim2$ (identical to those found for sources clustered near active galaxies), together with a distance redshift component defined by the Hubble relation."
+ The model of NarlikaraudArp(19050) does provide an alternativo explanation for the IIubble expansion redshift., The model of \citet{nar80a} does provide an alternative explanation for the Hubble expansion redshift.
+ However. if QSO production is ongoing. aud they evolve relatively quickly (105. vrs) iuto something differeut. the distribution iu Fie.," However, if QSO production is ongoing, and they evolve relatively quickly $10^{8}$ yrs) into something different, the distribution in Fig."
+ 9 appears to stil require ali initial Die Bang aud subsequent expausiou., 9 appears to still require an initial Big Bang and subsequent expansion.
+" These local inodel results thus sugeestOO that the QSO phase. in which intrinsic redshifts and absolute lunimosities behave as predicted by the theory of Narlikar and Das. is a transient one and that for most of their lives these objects obey normal physica avs,"," These $local$ model results thus suggest that the QSO phase, in which intrinsic redshifts and absolute luminosities behave as predicted by the theory of Narlikar and Das, is a transient one and that for most of their lives these objects obey normal physical laws."
+" It is interesting o speculate whether after the initial Big Bane uatter firs appeared in clnips as dt does here after the ""nini Banes”."," It is interesting to speculate whether after the initial Big Bang matter first appeared in clumps as it does here after the ""mini Bangs""."
+ Hf so. galaxies would be expected appear within 10? vis of the Dig Bang.," If so, galaxies would be expected to appear within $10^{8}$ yrs of the Big Bang."
+ The fac that they are iuitially sub-bhuuinous bw several magnitudes would prevent much earlier detection., The fact that they are initially sub-luminous by several magnitudes would prevent much earlier detection.
+ It is inuportaut to remember that the high-redshift QSOs detected with za=5 (logeza2 6.17) appear to be born soon after the Big Dang ouly if the cosmological redshift model is assumed., It is important to remember that the high-redshift QSOs detected with $_{\rm c1} > 5$ $_{c1} > 6.17$ ) appear to be born soon after the Big Bang only if the cosmological redshift model is assumed.
+ As can be seen in Fie., As can be seen in Fig.
+ 9. iu the loca? model these objects have much smaller distance redshift Ccomponcuts.," 9, in the $local$ model these objects have much smaller distance redshift components."
+ Although it is impossible to know bow much dispersion nav be preseut in the magnitudes at a eiven redshift. or how iuuch of the measured redshift is Doppler-related. broadly speaking. in Fie.," Although it is impossible to know how much dispersion may be present in the magnitudes at a given redshift, or how much of the measured redshift is Doppler-related, broadly speaking, in Fig."
+ 9 distance decreases horizontally to the left for sources with z;2(0.1 and magnitudes brighter than 18., 9 distance decreases horizontally to the left for sources with $_{\rm i} > 0.1$ and magnitudes brighter than 18.
+ This. then. ecplains the vertical appearance found for the Palomar Bright Quasar distribution (filled dots with ay< l7) from Sehinidtctal.(1983).," This, then, explains the vertical appearance found for the Palomar Bright Quasar distribution (filled dots with $_{\rm v} < 17$ ) from \citet{sch83}."
+. Where this vertical-liug edge at maguitudes near 15.5 is interpreted ο ο edge i the cosmological interpretation. in the local model it is interpreted," Where this vertical-lying edge at magnitudes near 15.5 is interpreted as a high-luminosity edge in the cosmological interpretation, in the local model it is interpreted"
+couut rate depeud substantially ou details of the assumed GRB population.,count rate depend substantially on details of the assumed GRB population.
+ Pessimistic uumbers result if the kinetic energy iun the CRB ejecta is assumed to scale with jet solid augle (Dalal. Ciriest. Pruet 2002).," Pessimistic numbers result if the kinetic energy in the GRB ejecta is assumed to scale with jet solid angle (Dalal, Griest, Pruet 2002)."
+ Fortunately. present evidence suggests that the ust enerey is approximately iudependent of collimation augle (Frail et al 2001: Pauaitescu Ixumar 2002).," Fortunately, present evidence suggests that the burst energy is approximately independent of collimation angle (Frail et al 2001; Panaitescu Kumar 2002)."
+ La this case. au alterglow will be detectable to at oll-axis angle that is iudependent of the initial jet augle. aud the orphan alterglow rate cau greally exceed the GRB rate and. provide a strong diagnostic of collimation (Ciranot et al 2002: Toani Panaitescu 2002: Nakar. Piran. Granot 2002 [NPGO02]).," In this case, an afterglow will be detectable to an off-axis angle that is independent of the initial jet angle, and the orphan afterglow rate can greatly exceed the GRB rate and provide a strong diagnostic of collimation (Granot et al 2002; Totani Panaitescu 2002; Nakar, Piran, Granot 2002 [NPG02])."
+ However. translents 'esembliug CRB alterglows cau also be produced by cosmological explosious hat do not prodice Patina ray bu‘sts.," However, transients resembling GRB afterglows can also be produced by cosmological explosions that do not produce gamma ray bursts."
+" Afterglow emission at any particular frequency reaches peak Intensity at some characteristic tiire aud characteristic Lorentz factor D, of the burst ejecta.", Afterglow emission at any particular frequency reaches peak intensity at some characteristic time and characteristic Lorentz factor $\Gamma_a$ of the burst ejecta.
+ The ybservecl beliavio “is essentially inclependeut of the initial Lorentz [actor Py of the explosion provided hat Py>Py., The observed behavior is essentially independent of the initial Lorentz factor $\Gamma_0$ of the explosion provided that $\Gamma_0 > \Gamma_a$.
+ Optical depth arguments require CgZ100 to produce the observed power law GRB spectra (Paczyussi 1986: Coocinan 19836: Ixrolik Pier 1991: Woods Loeb 1995.), Optical depth arguments require $\Gamma_0 \ga 100$ to produce the observed power law GRB spectra (Paczynski 1986; Goodman 1986; Krolik Pier 1991; Woods Loeb 1995.)
+ Fireballs with 10XPy1060 will not produce delectable gamma ray emission but will still produce detectable alterglows at ταςlo. optical. aud (1£ E is large enough) X-ray. wavelenetls.," Fireballs with $10 \la \Gamma_0 \ll 100$ will not produce detectable gamma ray emission but will still produce detectable afterglows at radio, optical, and (if $\Gamma_0$ is large enough) X-ray wavelengths."
+ Thus. dirty fireballs can prXLce transients with the generic characteristics of GRB afterglows. Le. broken power law spectra aneL light curves produced. by syuchrotron emission in a forward external srock.," Thus, dirty fireballs can produce transients with the generic characteristics of GRB afterglows, i.e., broken power law spectra and light curves produced by synchrotron emission in a forward external shock."
+ Rhoacks (2000) reimarked ou this possibility zud suggested that the cilference iu light curves could be used to cistitetuish the two., Rhoads (2000) remarked on this possibility and suggested that the difference in light curves could be used to distinguish the two.
+" The principle here is that afterglows [rom jetlike bursts see1 off-axis at [requencies :ibove the peak in f, will show a rapid decay fi,x/? (where p is the elecron energy distributiot ;»ower law index).", The principle here is that afterglows from jetlike bursts seen off-axis at frequencies above the peak in $f_\nu$ will show a rapid decay $f_\nu \propto t^{-p}$ (where $p$ is the electron energy distribution power law index).
+" lu contrast. afterglows from dirty fireballs will show a mitch slower decay. f,ix17?U U/Lop 1/2L7P1U (depending""n on whether the coolingn frequency"" LA. lies above or below the observe [ποιον]."," In contrast, afterglows from dirty fireballs will show a much slower decay, $f_\nu
+\propto t^{-3(p-1)/4}$ or $t^{1/2-3p/4}$ (depending on whether the cooling frequency $\nu_c$ lies above or below the observed frequency)."
+" Huaug. Dai. Lu (2002: he'ealter HDLO2) examiued the requiremeuts for distinguishiugjs) orphan alterglows due to ofl-axis jets [rom those due to dirty fireballs (which they teri ""failed eamuaua-ray bursts"" or ""FGBRBS)."," Huang, Dai, Lu (2002; hereafter HDL02) examined the requirements for distinguishing orphan afterglows due to off-axis jets from those due to dirty fireballs (which they term “failed gamma-ray bursts” or “FGRBS”)."
+ They point out that a major difficulty in distinguishiug the two is the unknown triggerMD time., They point out that a major difficulty in distinguishing the two is the unknown trigger time.
+ Iu studies of triggeredMD CRBs. the decay iudex a can be deterinined without ambiguity usine the arrival time of the first gana rays as the time origin fg. (," In studies of triggered GRBs, the decay index $\alpha$ can be determined without ambiguity using the arrival time of the first gamma rays as the time origin $t_0$. ("
+A modest exception ean be made for GRB 9t101209. where the availability of optical ROTSE data during the GRB iself [Akerlof et al 1999] introduces some sensitivity of the early decay index to the precise choice of the time origin.),"A modest exception can be made for GRB 990123, where the availability of optical ROTSE data during the GRB itself [Akerlof et al 1999] introduces some sensitivity of the early decay index to the precise choice of the time origin.)"
+ However. --- Lan orphan afterglow. the trigger time is unknown.," However, in an orphan afterglow, the trigger time is unknown."
+ All we kuow is the discovery. time /y., All we know is the discovery time $t_d$.
+ The age of the burst at discovery cau be defined as /4=ly—ty aud becomes par:uueter of the ligh curve that we may try to fit.," The age of the burst at discovery can be defined as $t_1 \equiv t_d -
+t_0$ and becomes parameter of the light curve that we may try to fit."
+ In general. there is a substantial degeneracy between a anc |/4 unless the observatious span a period of duration M2H.," In general, there is a substantial degeneracy between $\alpha$ and $t_1$ unless the observations span a period of duration $\Delta t \ga 2
+t_1$."
+" HDLO2 suggest a few tests that might. hel» discriminate dirty. fireballs from off-axis GRBs. |ut conclude that the ouly promising one Is the »ossibility (uucoufiruied at present) that most dirty fireballs are accompauled by ""x-ray [laslies"" of he type observed by the BeppoSAX (e.g.. Heise et al 2001) aud HETE-U satellites."," HDL02 suggest a few tests that might help discriminate dirty fireballs from off-axis GRBs, but conclude that the only promising one is the possibility (unconfirmed at present) that most dirty fireballs are accompanied by “x-ray flashes” of the type observed by the BeppoSAX (e.g., Heise et al 2001) and HETE-II satellites."
+strategv of using all of the available time for continuous observing. acquiring a large total umber of observations.,"strategy of using all of the available time for continuous observing, acquiring a large total number of observations."
+ ? showed that the observed Fourier transform of a discrete dataset can be expresse as the convolution of the true Fourier transforma FY(7) with a spectral window 6x(7) that fully describes the interference effects between the sampled frequencies. such tha where F(r) is the complex Fourier transform of a function FG) aud is defined as and the spectral window is given by where Nis the total nuniber of observations and f; is the set of observation times.," \citet{Deeming75UnevenSampling} showed that the observed Fourier transform of a discrete dataset can be expressed as the convolution of the true Fourier transform $F_{\rm N}(\nu)$ with a spectral window $\delta_{\rm N}(\nu)$ that fully describes the interference effects between the sampled frequencies, such that where $F(\nu)$ is the complex Fourier transform of a function $f(t)$ and is defined as and the spectral window is given by where $N$ is the total number of observations and $t_k$ is the set of observation times."
+ The window functiou is even. so onlv positive frequencies are evaluated.," The window function is even, so only positive frequencies are evaluated."
+" The ovsical iuterpretation of the window fiction can be secu w cousicdering the Fourier trausform of an infinite sine :unction (a celta ""unction) convolved with the window πιοΊο, leading to the observed Fourier trausfori."," The physical interpretation of the window function can be seen by considering the Fourier transform of an infinite sine function (a delta function) convolved with the window function, leading to the observed Fourier transform."
+" Adjusting the set o observation tines fj alters the window ""uctiou. and hence the Fourier transform."," Adjusting the set of observation times $t_k$ alters the window function, and hence the Fourier transform."
+ ? showed hat aliasing effects in the observed Fourier trausforiu arising from the window function could be mitigated by introducing irregularities oeiuto the sampling pattern., \citet{Deeming75UnevenSampling} showed that aliasing effects in the observed Fourier transform arising from the window function could be mitigated by introducing irregularities into the sampling pattern.
+ By careful choice of sampling it is possible to reconstruct hnieher frequencies than vy. the highest frequeney that could be accurately reconstructed if he same number of observations were equally spaced.," By careful choice of sampling it is possible to reconstruct higher frequencies than $\nu_{\rm N}$, the highest frequency that could be accurately reconstructed if the same number of observations were equally spaced."
+ ? illustrated the behaviour of the window function hrough the use of a simple single-variable model for the WV525 case. curpirically deriving a form for the spacing sed on visual iuspectiou of the resulting spectral window ποοι.," \citet{Deeming75UnevenSampling} illustrated the behaviour of the window function through the use of a simple single-variable model for the $N=25$ case, empirically deriving a form for the spacing based on visual inspection of the resulting spectral window function."
+ 7. adopted a pseudo-geonmetrie spacing as a wav Oo lnaxinuse the uniforuütv of their phase spaciug. but 10 details of the inplemoenuted optimisation strategv were xovided.," \citet{MouldEtAl00HSTKeyProjectXXI} adopted a pseudo-geometric spacing as a way to maximise the uniformity of their phase spacing, but no details of the implemented optimisation strategy were provided."
+ The Ravleigh resolution criterion may be written as AvAT=1.," The Rayleigh resolution criterion may be written as $\Delta \nu \Delta T
+ = 1$."
+ It means that observations of leneth AT vield the frequeney of au oscillation with an error of Av= 1/AT. at worst.," It means that observations of length $\Delta T$ yield the frequency of an oscillation with an error of $\Delta \nu=1/ \Delta T$ , at worst."
+ After au additional time interval GAT this vields an error in the evcle count AvaAT=06., After an additional time interval $\alpha \Delta T$ this yields an error in the cycle count $\Delta \nu \alpha \Delta T = \alpha$.
+ Thus to keep the evele count eror less than a. subsequeu observations ought to be scheduled no later than (1|ajAT from the start.," Thus to keep the cycle count error less than $\alpha$, subsequent observations ought to be scheduled no later than $(1+ \alpha) \Delta T$ from the start."
+ Recureut aud rigid application of this known principle viclds a geometric series scheduling (woe are eratefu to the referee for sugeesting this line of argument)., Recurrent and rigid application of this known principle yields a geometric series scheduling (we are grateful to the referee for suggesting this line of argument).
+ We shall show that the choice of geometric base is crucial for such a strateey to succeed., We shall show that the choice of geometric base is crucial for such a strategy to succeed.
+ The ecucral question of what form an optimal uneven spacing should take. when the nuuber of observations aud the length of the observing run are free parameters remains. to our l-jowledee. uianswered.," The general question of what form an optimal uneven spacing should take, when the number of observations and the length of the observing run are free parameters remains, to our knowledge, unanswered."
+ Tere we present a detailed analysis of the application of a geometric spacing technique to the uudersiuupliug problem., Here we present a detailed analysis of the application of a geometric spacing technique to the undersampling problem.
+ Although our sampling is not an analytic solution to the general problem of arbitrary placemeut of the dataset fj. if outperforms both the model of ? and pure random distributions. and thus represents the hes current heuristic solution to the sampling xoblei.," Although our sampling is not an analytic solution to the general problem of arbitrary placement of the dataset $t_k$, it outperforms both the model of \citet{Deeming75UnevenSampling} and pure random distributions, and thus represents the best current heuristic solution to the sampling problem."
+ Iu Section 3.. we present the sampling strategy and quality metrics by which it is judeed.," In Section \ref{Section:Method}, we present the sampling strategy and quality metrics by which it is judged."
+ In Section L we compare the ecolmetric samplue to the sampling of ? aud to random suupliung and show that noise has no impact on the salplue strategy., In Section \ref{Section:Results} we compare the geometric sampling to the sampling of \citet{Deeming75UnevenSampling} and to random sampling and show that noise has no impact on the sampling strategy.
+ We also show that the window function retains similar characteristics under random of the spacings. an mniportaut result for practical scheduling scenarios.," We also show that the window function retains similar characteristics under random of the spacings, an important result for practical scheduling scenarios."
+" In Section 5r, we present a erapl indicating optimal geometric samplues for observing rus spanniig the range from 10 to 500 observations."," In Section \ref{Section:Practical
+ application} we present a graph indicating optimal geometric samplings for observing runs spanning the range from 10 to 500 observations."
+ We present our final conclusions im Section 7.., We present our final conclusions in Section \ref{Section:Conclusions}.
+ We generate our observation times with a ecometric distribution using tle scheme where fy ds the (time) value of the point iu the series. ως ds the base of the series. Ais the uuuber of the curreut data point. JV is the total προς of observations aud T is the duratiou of the total observing interval.," We generate our observation times with a geometric distribution using the scheme where $t_k$ is the (time) value of the point in the series, $x$ is the base of the series, $k$ is the number of the current data point, $N$ is the total number of observations and $T$ is the duration of the total observing interval."
+ This produces a econietrically spaced distribution such hat 0xfy«T., This produces a geometrically spaced distribution such that $0 \le t_k \le T$.
+ The parameter ο. which we call thebusc.. wav be arbitrarily chosen.," The parameter $x$, which we call the, may be arbitrarily chosen."
+" Since the window ""unction is in ecneral complex. for comparison with ? we follow the convention of examining the uormatised amplitude AQ). defined as the square of the amplitude of the window function. normalised bv the square of the nuuber of datapoiuts. such that We define thebase.. Jupes as the value for which the overall degree of structure present in (7) jsi nmuuiuised."," Since the window function is in general complex, for comparison with \citet{Deeming75UnevenSampling} we follow the convention of examining the normalised amplitude $A(\nu)$, defined as the square of the amplitude of the window function, normalised by the square of the number of datapoints, such that We define the, $x_{\rm opt}$, as the value for which the overall degree of structure present in $A(\nu)$ is minimised."
+ We use the root inca square (RAIS) of Atv) as a straightforward metric for ideutifiug structure., We use the root mean square (RMS) of $A(\nu)$ as a straightforward metric for identifying structure.
+ \Tinimusing the RAIS thus represeuts the desire ofthe astronomer to achieve relatively even seusitivitv to periods., Minimising the RMS thus represents the desire ofthe astronomer to achieve relatively even sensitivity to periods.
+ Iu order to retain scusitivity to the higher frequency structure we evaluate ο). for frequeucies in the range, In order to retain sensitivity to the higher frequency structure we evaluate $A(\nu)$ for frequencies in the range
+fit.,fit.
+ Although the fit is more meaningful. as it is based on a physically realistic model while having no more free parameters than the Gaussian mocdel. the Gaussian fit confirms and cquantifies the large line widths ancl blue shifts expected from EWS emission.," Although the fit is more meaningful, as it is based on a physically realistic model while having no more free parameters than the Gaussian model, the Gaussian fit confirms and quantifies the large line widths and blue shifts expected from EWS emission."
+ Along similar lines. He-like complexes can be fit with hree Gaussians. providing a direct measure of the f//7 ratio. rom which a single radius of formation can be inferred.," Along similar lines, He-like complexes can be fit with three Gaussians, providing a direct measure of the $f/i$ ratio, from which a single radius of formation can be inferred."
+ We show these results in reffig:hesauss.., We show these results in \\ref{fig:hegauss}.
+ For the complex we find alow ratio of ).40T, For the complex we find alow ratio of $0.40_{-.17}^{+.23}$.
+ We next fit the same mocel. but with the f// ratio ixed at the low density limit that. would. be expected if he X-ray. plasma were very far from the photosphere. where »hotoexcitation cannot alter the ratio.," We next fit the same model, but with the $f/i$ ratio fixed at the “low density limit” that would be expected if the X-ray plasma were very far from the photosphere, where photoexcitation cannot alter the ratio."
+ This fit (which is also shown in reffig:hegauss)) is poor compared to the fs)=0.1 fit. implving that the hot plasma is relatively close to. the photosphere.," This fit (which is also shown in \\ref{fig:hegauss}) ) is poor compared to the $f/i = 0.4$ fit, implying that the hot plasma is relatively close to the photosphere."
+" This conclusion is in good agreement with the result. from fitting the three wind-profile model (hewind). in which the onset radius of a distributed source of X-ray emitting plasma.f2,.. is constrained to be within 1 stellar racius of the photosphere."," This conclusion is in good agreement with the result from fitting the three wind-profile model ), in which the onset radius of a distributed source of X-ray emitting plasma, is constrained to be within 1 stellar radius of the photosphere."
+Porosity. due to optically thick clumps. can allect X- line profiles by reducing the average optical depth of the wind (Oskinovaetal.2006).. and is another elfect we might consider when fitting the emission line data.,"Porosity, due to optically thick clumps, can affect X-ray line profiles by reducing the average optical depth of the wind \citep{ofh2006}, and is another effect we might consider when fitting the emission line data."
+ However. the »orositv. length tlat ds necessary to provide a measurable οσοι is quite large (Owocki&Cohen2006) compared to he small-scale sructure in state-of-the-art 2-D. radiation ivcirocivnamics simulations (Dessart&Owocki2003.5)2005).," However, the porosity length that is necessary to provide a measurable effect is quite large \citep{oc2006} compared to the small-scale structure in state-of-the-art 2-D radiation hydrodynamics simulations \citep{do2003,do2005}."
+. ‘To explore the efects of porosity on the fit quality. and on he other parameers. we fit the line with a model having an clleetive opacity ποιος by »orosity. from spherical clumps.," To explore the effects of porosity on the fit quality and on the other parameters, we fit the line with a model having an effective opacity modified by porosity from spherical clumps."
+ This model is similar to the one described in Owocki&Cohen(2006) using the same »orositv-Iength. (f(r)=(£f. where (is the characteristic chump size and f is the clump filling factor) formalism. rut here employing a racial clump distribution determined ov the wind beta-velocitv law. as used by Oskinovaetal. (2006).," This model is similar to the one described in \citet{oc2006} – using the same porosity-length $h(r) \equiv \ell/f$, where $\ell$ is the characteristic clump size and $f$ is the clump filling factor) formalism – but here employing a radial clump distribution determined by the wind beta-velocity law, as used by \citet{ofh2006}."
+. We fix the characteristic optical depth at τν=4.2. which is the value we would expect at the wavelength of the line assuming a mass-loss rate of 1.8107 (Taresehetal.1997).," We fix the characteristic optical depth at $\taustar = 4.2$, which is the value we would expect at the wavelength of the line assuming a mass-loss rate of $1.8 \times 10^{-5}$ \citep{Taresch1997}."
+. Phe fit we obtain is similar in quality to the fit reported in Table 2.., The fit we obtain is similar in quality to the fit reported in Table \ref{tab:line_fits}.
+ But in order to achieve this &ood fit. a very large terminal porosity length of hy.—5.5 is required. (wi ha lower confidence limit off.=2.5 #.)). where the racially varving porosity length is given by h(r)2hytlRaη)," But in order to achieve this good fit, a very large terminal porosity length of $\hinf = 5.3$ is required (with a lower confidence limit of $\hinf =
+2.5$ ), where the radially varying porosity length is given by $h(r) = \hinf(1-\Rstar/r)^{\beta}$."
+ 1f we use 2.61PO as the stanclareLo smooth-wine miass-loss rate (Itepolustetal. 2004).. then he required. porosity length is even larger.," If we use $2.6 \times 10^{-5}$ as the standard, smooth-wind mass-loss rate \citep{Repolust2004}, then the required porosity length is even larger."
+ Even the minimium porosity length of ha=2.5 js inconsisten with the numerical simulations of the line-clriving instability. (Dessart&Owocki2003. 2005).. requiring. for example. clumps that are individually 0.25r in scale in a wind with a uniform filling [actor of f.— 0.1. Given the charac‘teristic optical depth values (7... listed in," Even the minimum porosity length of $\hinf = 2.5$ is inconsistent with the numerical simulations of the line-driving instability \citep{do2003,do2005}, , requiring, for example, clumps that are individually 0.25 in scale in a wind with a uniform filling factor of $\fv = 0.1$ Given the characteristic optical depth values , listed in"
+disks rotate faster ρω=+.30d).,disks rotate faster $(P_{avg} = 4.30d)$.
+ The period distributions are wide for both (o = 2.44 and 2.57 for disked and clisk-less respectively)., The period distributions are wide for both $\sigma$ = 2.44 and 2.57 for disked and disk-less respectively).
+ A (wo sided K-S test of the period distribution of disked and cisk-Iess stars resulted in D = 0.68., A two sided K-S test of the period distribution of disked and disk-less stars resulted in D = 0.68.
+ Which indicates the probability that the period distributions are drawn from the same parent is 1.6%., Which indicates the probability that the period distributions are drawn from the same parent is $1.6\%$.
+ In the language of gaussian hypothesis (esting. (his is a 240 result.," In the language of gaussian hypothesis testing, this is a $\sigma$ result."
+ So for stars which are believed {ο be single. the period distributions of disked aud disk-less stars. are mareinally different.," So for stars which are believed to be single, the period distributions of disked and disk-less stars, are marginally different."
+ As described in section 2. we obtained M-band and Q-band photometry [or a few stars with strong IR. excesses.," As described in section 2, we obtained M-band and Q-band photometry for a few stars with strong IR excesses."
+ All new M and Q-band photometry are listed with other photometry assembled [rom the literature in Table 3., All new M and Q-band photometry are listed with other photometry assembled from the literature in Table 3.
+ SEDs derived [rom these data are shown in Figure 3., SEDs derived from these data are shown in Figure 3.
+ The data are not simultaneous., The data are not simultaneous.
+ We compared the de-reddened fluxes to emperical stellar photospheres appropriate for each spectral type and models of geometrically-thin. thick disks (Ilillenbraudetal.1992).," We compared the de-reddened fluxes to emperical stellar photospheres appropriate for each spectral type and models of geometrically-thin, optically-thick disks \citep{hillenbrand92}."
+. The model is of a face-0n blackbody disk extending from the stellar surface., The model is of a face-on blackbody disk extending from the stellar surface.
+" Easily visible in all SED plots are signs of accretion as (lie excess fluxes in the U-band (A,jj=0.367).", Easily visible in all SED plots are signs of accretion as the excess fluxes in the U-band $(\lambda_{eff} = 0.36\micron)$.
+ UR excesses are also present and significantly exceed the predictions of the stellar photosphere aud (he flat reprocessing disk model at longer wavelengths., IR excesses are also present and significantly exceed the predictions of the stellar photosphere and the flat reprocessing disk model at longer wavelengths.
+" Most sources lack I excesses al the shortest wavelengths. consistent with clust sublimation and/or inner holes (Muzerolleetal.(2003): οἱ, Meveretal. (1997a)))."," Most sources lack IR excesses at the shortest wavelengths, consistent with dust sublimation and/or inner holes \citet{muzerolle03}; cf. \citet{meyer97a}) )."
+ An inner dust disk hole would be visible as à lack of nem-infrared excess in the SED., An inner dust disk hole would be visible as a lack of near-infrared excess in the SED.
+ Disk regulation theories predict that circumstellar disks have inner holes maintained by pressure balance with the stellar magnetic field (Ostriker&Shu1995)., Disk regulation theories predict that circumstellar disks have inner holes maintained by pressure balance with the stellar magnetic field \citep{OS95}.
+". Measurements of stellar magnetic field strengths are needed along with estimates of AM.M and R, to properly estimate the extent of inner disk clearing."," Measurements of stellar magnetic field strengths are needed along with estimates of $M_{\star},\dot{M}$ and $_{\star}$ to properly estimate the extent of inner disk clearing."
+" The verv large Il excesses at longer wavelengths indicate that the disks are probably not flat but. flared (Chiang&GoldreichL997) with vertical temperature structures commonly referred to as ""disk atimospheres” 1999:Dullemondetal. 2001)."," The very large IR excesses at longer wavelengths indicate that the disks are probably not flat but flared \citep{CG97} with vertical temperature structures commonly referred to as “disk atmospheres” \citep[see][]{dalessio99,dullemond01}."
+. In the following sections we discuss the significance of our results in the context of current theoretical frameworks for rotational evolution., In the following sections we discuss the significance of our results in the context of current theoretical frameworks for rotational evolution.
+ We also compare our observational study to others conducted in Tu-Aur and the ONC., We also compare our observational study to others conducted in Tau-Aur and the ONC.
+ Finally we emphasize the importance of sample selection criteria for studying rotational evolution of low mass PAIS stars., Finally we emphasize the importance of sample selection criteria for studying rotational evolution of low mass PMS stars.
+state.,state.
+ Both objects exhibited radio emission associated with a jet which was detected near (he peak of the secondary maxima., Both objects exhibited radio emission associated with a jet which was detected near the peak of the secondary maxima.
+ Jets are believed to form during the low-hard state (Fender2001)., Jets are believed to form during the low-hard state \citep{fen01}.
+. I the secondary maximum emission does in [act originate [rom jets then one would expect to see it after the object has transitioned into the low-harcl state., If the secondary maximum emission does in fact originate from jets then one would expect to see it after the object has transitioned into the low-hard state.
+ We have presented OIR light curves of the 2002 outburst of 4U 1543-47., We have presented OIR light curves of the 2002 outburst of 4U 1543-47.
+ A strong Ih flare appears soon after (he source transitions into the low-hard state., A strong IR flare appears soon after the source transitions into the low-hard state.
+ Blackbocly ancl broken power-law [its suggest that svnchrotron emission is the most likely source of the secondary maximum flux., Blackbody and broken power-law fits suggest that synchrotron emission is the most likely source of the secondary maximum flux.
+ The detection of a radio source al the peak of this flare supports (his conclusion., The detection of a radio source at the peak of this flare supports this conclusion.
+ Indeed. (he presence of jets during (his X-ray. state would not be surprising eiven that we often observe jets in this accretion phase (Fender2001).," Indeed, the presence of jets during this X-ray state would not be surprising given that we often observe jets in this accretion phase \citep{fen01}."
+. since (he secondary maximunm is so prominent in the IH. it is clear that. I. monitoringD ol SXTs can successfully trigger other simultaneous observations at other wavelengths and eather useful information into the lormation of and flux contribution of jets.," Since the secondary maximum is so prominent in the IR, it is clear that IR monitoring of SXTs can successfully trigger other simultaneous observations at other wavelengths and gather useful information into the formation of and flux contribution of jets."
+ Future observations IE such phenomena can be further improved by obtaining; simultaneous radio observations in two bands so that a confident fit of the power-law gradient may be obtained., Future observations of such phenomena can be further improved by obtaining simultaneous radio observations in two bands so that a confident fit of the power-law gradient may be obtained.
+ UV observations will enable us to ascertain anv upturn in the SED clue to the dominance of an accretion disk blackbody component. although: observations in Chis band are quite difficult. due to strong interstellar extinction.," UV observations will enable us to ascertain any upturn in the SED due to the dominance of an accretion disk blackbody component, although observations in this band are quite difficult due to strong interstellar extinction."
+ X-ray spectra would also provide valuable information on the relationship to the OIR data and if there is any contribution from thermal IC of disk photons., X-ray spectra would also provide valuable information on the relationship to the OIR data and if there is any contribution from thermal IC of disk photons.
+ We would like to thank Emrah lIalemci for his useful comments. Suzanne Tourtellotte and Alfred Chen in assisting with the optical data reduction. and the CTIO observers. David Gonzalez Huerta and Juan Espinoza.," We would like to thank Emrah Kalemci for his useful comments, Suzanne Tourtellotte and Alfred Chen in assisting with the optical data reduction, and the CTIO observers, David Gonzalez Huerta and Juan Espinoza."
+ We also thank Darren Depov for arranging the Stromlo observations., We also thank Darren Depoy for arranging the Stromlo observations.
+ This paper has made use of the Astrophysics Data Svstem Abstract Service and the ASM/BRXTE Data Extraction website., This paper has made use of the Astrophysics Data System Abstract Service and the ASM/RXTE Data Extraction website.
+ MB and CD eratelully acknowledge support from the National Science Foundation through erant. AST-0098421., MB and CB gratefully acknowledge support from the National Science Foundation through grant AST-0098421.
+The average temperature of the active region Is determined to be fogT=6.11 through using the SAT filter ratio technique ou the ALI aud ÀAIMe filters.,The average temperature of the active region is determined to be $log~ T = 6.41$ through using the SXT filter ratio technique on the Al.1 and AlMg filters.
+ We take the volue of the coronal οιτας plasma to be 1075 cu., We take the volume of the coronal emitting plasma to be $10^{29}~$ $^3$.
+ Note hat this is uot the volume of the computation box but rather au upper estimate of the volume of cluitting plasma in the active region., Note that this is not the volume of the computation box but rather an upper estimate of the volume of emitting plasma in the active region.
+ Taking also an upper estimate of the electron. umber deusity to be Lem 7. provides au upper estimate of the coronal racdiative losses as ~los1077 eres lL," Taking also an upper estimate of the electron number density to be $10^9$ $^{-3}$, provides an upper estimate of the coronal radiative losses as $\sim 1 \times 10^{25}$ ergs $^{-1}$."
+ This value can be seen to be of the same order as he euerev build up due to the small scale convecIve motions., This value can be seen to be of the same order as the energy build up due to the small scale convective motions.
+ While the two values are cousistent with one-another. it should be noted that he free magnetic enerev m the simulation is held «thin the coronal field and therefore not muuediaely accessible to heat the corona and account or the radiative losses.," While the two values are consistent with one-another, it should be noted that the free magnetic energy in the simulation is held within the coronal field and therefore not immediately accessible to heat the corona and account for the radiative losses."
+ For such a process to take place au additional energy release mechausin. would have to act.," For such a process to take place an additional energy release mechanism, would have to act."
+ Possible mechauisiis tuchide nanoflares (Parker19588:Browningctal.2008:Wiltuot-Sinithetal.2010) or turbulent recounection(ILevvaert," Possible mechanisms include nanoflares \citep{Parker88,Browning08,Wilmot-Smith10} or turbulent \citep{heyv84}."
+z&Priest105 D...At this point it is unclear if such a iuechanisni may release the cucrey., At this point it is unclear if such a mechanism may release the energy.
+ Also the simple calculation docs not take iuto account the effect ofthermal conduction which is an iuportaut factor iu he enerev balance of hot coronal loops., Also the simple calculation does not take into account the effect of thermal conduction which is an important factor in the energy balance of hot coronal loops.
+ However. taking these limitations into account. the calculation does show that convective motions may play a key role iu the energy build up of the Corolla.," However, taking these limitations into account, the calculation does show that convective motions may play a key role in the energy build up of the corona."
+ Using the rates of energy iujection deteruiued in the simulation and assuiinug that it remains constant over the life-time of the active region. it is possible to estimate how uch free energy. the active region could have ou the L6th December. the start dav of the present siunulations. due to prior evolution as it rotated onto the disk and towards central meridian.," Using the rates of energy injection determined in the simulation and assuming that it remains constant over the life-time of the active region, it is possible to estimate how much free energy the active region could have on the 16th December, the start day of the present simulations, due to prior evolution as it rotated onto the disk and towards central meridian."
+ Asstuning that it elerecd just before disk passage the active region would have evolved for 5 caves., Assuming that it emerged just before disk passage the active region would have evolved for 5 days.
+ As a result the sunall scale convective motions would inject a total of 1.1« 10?!eres of free magnetic enerev iuto the coronal field., As a result the small scale convective motions would inject a total of $1.1 \times 10^{31}$ ergs of free magnetic energy into the coronal field.
+ This vah wis 10:4 of the active region eucrev on the 16th December., This value is $\%$ of the active region energy on the 16th December.
+ While the energv may have been 10:4 hieher. we cannot use this to coustraiu the initial coudition without having vector magnetic field data to constrain the distribution of a on the lower boundary.," While the energy may have been $10\%$ higher, we cannot use this to constrain the initial condition without having vector magnetic field data to constrain the distribution of $\alpha$ on the lower boundary."
+ Iu this paper a new techuique foranodeling nou-linear force-free fields directly fire1u line of sight conrponeut maenetogram observations has been xeseuted., In this paper a new technique for modeling non-linear force-free fields directly from line of sight component magnetogram observations has been presented.
+ The kev new feature of this technique is that sequences of Lue of sight 1naeuetograiis cal )o directly used as lower boundary conditions o dive the evolution of coronal magnetic fields )etweecn successive force-free αςpulibrig, The key new feature of this technique is that sequences of line of sight magnetograms can be directly used as lower boundary conditions to drive the evolution of coronal magnetic fields between successive force-free equilibria.
+ As a result. the lower boundary condition iu the uodel closely follows the observe| iiagnetic field evolution as given by the obseryed. line-ofsiglt uagnetogranms over he period of observations.," As a result, the lower boundary condition in the model closely follows the observed magnetic field evolution as given by the observed line-of-sight magnetograms over the period of observations."
+ The techuique is illustrated bs* applying it to AIDI observations of a decavius active region. NOAA AR s005.," The technique is illustrated by applying it to MDI observations of a decaying active region, NOAA AR 8005."
+ The active region was close to central meridian aud in flux balance to a high degree., The active region was close to central meridian and in flux balance to a high degree.
+ It was uodeled over a 1 day period. where Gl iudividual SoIIO:MDI 96 iin uaenetoeranis where available ο use as lower bouidary couditious.," It was modeled over a 4 day period, where 61 individual SoHO:MDI 96 min magnetograms where available to use as lower boundary conditions."
+ During this time the dispersal of the active region was mainly dominated by random lotions due to sanall scale convective cells aud. fiux cancellation along the internal polarity inversicπι line., During this time the dispersal of the active region was mainly dominated by random motions due to small scale convective cells and flux cancellation along the internal polarity inversion line.
+ Through applying the 61 nxieuetograinis as ower boundary conditions a1 siuulating a continuous evolution of the surface and coronal field between thoi. the build up «ot free magnetic enerev iuto the corona was studied.," Through applying the 61 magnetograms as lower boundary conditions and simulating a continuous evolution of the surface and coronal field between them, the build up of free magnetic energy into the corona was studied."
+ It should ρα noted bere hat the initial conditiou for the coronal field in the siuulation was taken to be a potential field., It should be noted here that the initial condition for the coronal field in the simulation was taken to be a potential field.
+ Due to the availability of oulv inc-of-ieht coluponent maenetograms this was he only unique choice., Due to the availability of only line-of-sight component magnetograms this was the only unique choice.
+ IHowewy atf the time corresponding to the start of tlic siuulation the actual coronal field of the active region need not rave been potential., However at the time corresponding to the start of the simulation the actual coronal field of the active region need not have been potential.
+" Therefore our al was not o reproduce the evolution of the coroual OOPS as seen in coronal observations. but rather to study he trends of euerevOo, input over the perioc of four davs. where this is deduced. frou using observed uagnetoeranus."," Therefore our aim was not to reproduce the evolution of the coronal loops as seen in coronal observations, but rather to study the trends of energy input over the period of four days, where this is deduced from using observed magnetograms."
+ When the initial coronal field is taken to be »otential. the small-scale iuotious iuject around ο2.5«402510°? ores 1 of Dgfee inaetic energy. this enerev is manly stored iu the low corona below 30 Ahn.," When the initial coronal field is taken to be potential, the small-scale motions inject around $2.5 \times 10^{25}$ ergs $^{-1}$ of free magnetic energy, this energy is mainly stored in the low corona below 30 Mm."
+ However. even if the initial state ou the Sun was different. this would not sigmnifemutlv alter the trends of cucreyv injection found within the," However, even if the initial state on the Sun was different, this would not significantly alter the trends of energy injection found within the"
+"(2007a),, to make sure that the results is not just a coincidence because of lack of numerical resolution.",", to make sure that the results is not just a coincidence because of lack of numerical resolution."
+ We have plotted the high resolution data as squares in Figs., We have plotted the high resolution data as squares in Figs.
+ to 6 for comparison., \ref{fig:vmall} to \ref{fig:etam} for comparison.
+ In Figs., In Figs.
+" 7 and 8 we have plotted the ’Via Lactea’ alone, without any cutoffs or re-binning."," \ref{fig:vM} and \ref{fig:vbetadiem} we have plotted the 'Via Lactea' alone, without any cutoffs or re-binning."
+ In these figures we are using the values 7 from Table 1.., In these figures we are using the values $\eta$ from Table \ref{tab:eta}.
+ On both figures we see that the suggested relations are confirmed when comparing with highly resolved data., On both figures we see that the suggested relations are confirmed when comparing with highly resolved data.
+ The ’Via Lactea’ structure did not experience any major mergers since z=1., The 'Via Lactea' structure did not experience any major mergers since $z=1$.
+ All quantities are extracted in spherical bins., All quantities are extracted in spherical bins.
+ Due to numerical softening one can safely trust the radius outside 1 kpc of this galaxy., Due to numerical softening one can safely trust the radius outside 1 kpc of this galaxy.
+" In this simulation the outermost 6-10 points should be considered with care since they are potentially not fully equilibrated yet, as is easily seen when considering the radial derivative of the density profile."," In this simulation the outermost 6-10 points should be considered with care since they are potentially not fully equilibrated yet, as is easily seen when considering the radial derivative of the density profile."
+" We have thus compared the angular momentum of the intermediate resolution structures, the highly resolved Cpung.W3 cluster, and the ’Via Lactea’ simulation, which is one of the best resolved structures published today (Diemandetal. to our suggested relations, and see strong correlations 2007a),,between the rotation, mass and velocity anisotropy of the system."," We have thus compared the angular momentum of the intermediate resolution structures, the highly resolved $_{HR}$ .W3 cluster, and the 'Via Lactea' simulation, which is one of the best resolved structures published today \citep{diemand}, to our suggested relations, and see strong correlations between the rotation, mass and velocity anisotropy of the system."
+" As mentioned in the introduction, people often use the spin parameter A when describing the angular momentum of DM halos."," As mentioned in the introduction, people often use the spin parameter $\lambda$ when describing the angular momentum of DM halos."
+ Combining our relation between the mass and the angular momentum with the spin, Combining our relation between the mass and the angular momentum with the spin
+with the Hevnolds stress teusor Ilere v1N is (he coefficient of turbulent. viscosity and i44j is the coefficient of the A effect (Ixitchatinov&Rüdiger51995).. a non-diffusive angular5 momentum transport caused by turbulence.,"with the Reynolds stress tensor Here $\nu_\mathrm{tv}$ is the coefficient of turbulent viscosity and $\nu_\mathrm{tl}$ is the coefficient of the $\Lambda$ effect \citep{1995A&A...299..446K}, a non-diffusive angular momentum transport caused by turbulence."
+ v1N and mytl are expected to have the same value. since both effects are caused bv turbulence. i.e.. thermal driven convection.," $\nu_\mathrm{tv}$ and $\nu_\mathrm{tl}$ are expected to have the same value, since both effects are caused by turbulence, i.e., thermal driven convection."
+ We discuss this in more detail in 87?.., We discuss this in more detail in \ref{diffusivity}.
+ οth denotes the deformation tensor. which is given in spherical coordinates by An expression lor the A effect. (Aj) is given later.," $E_{ik}$ denotes the deformation tensor, which is given in spherical coordinates by An expression for the $\Lambda$ effect $\Lambda_{ik}$ ) is given later."
+ The amount of energy. that is converted by the Revnolcds stress from kinematic energy (o internal energy is given by, The amount of energy that is converted by the Reynolds stress from kinematic energy to internal energy is given by
+roughly into account bv scaling the viscous dissipation in our modeling (to (he total instead {ο (he gas pressure and (hus including parts which are related to the addition of disk produced magnetic [Iux that raise (he magnetic pressure. ie. lead to a lower beta in the corona.,"roughly into account by scaling the viscous dissipation in our modeling to the total instead to the gas pressure and thus including parts which are related to the addition of disk produced magnetic flux that raise the magnetic pressure, i.e. lead to a lower beta in the corona."
+ Detailed investigations on the formation of a magnetized corona from MILD simulation (e.g. Miller Stone 2000: Lawley Balbus 2002: Machida. Havashi. Matsumoto 2000) show that à strongly magnetized corona can form above an initially weakly magnetizecl disk.," Detailed investigations on the formation of a magnetized corona from MHD simulation (e.g. Miller Stone 2000; Hawley Balbus 2002; Machida, Hayashi, Matsumoto 2000) show that a strongly magnetized corona can form above an initially weakly magnetized disk."
+ Miller Stone (2000) showed that magnetic field is amplified] within 2 disk scale heieht (1120.01) and the energy is mostly cissipated in 3 to 5 scale heieht. (hereby heats up the corona.," Miller Stone (2000) showed that magnetic field is amplified within 2 disk scale height (H=0.01R) and the energy is mostly dissipated in 3 to 5 scale height, thereby heats up the corona."
+ Machida et al. (, Machida et al. (
+2000) also showed that a low-; corona in the form of “patch corona” or active coronal region.,"2000) also showed that a $\beta$ corona in the form of ""patch corona"" or active coronal region."
+ [ere we consider a slab corona in a large vertical extent where the structure is vertically stratified. in contrast (ο an isothermal torus in the ΑΔΗ) simulation.," Here we consider a slab corona in a large vertical extent where the structure is vertically stratified, in contrast to an isothermal torus in the MHD simulation."
+ More importantly. the vertical thermal conduction is (taken into account in our model.," More importantly, the vertical thermal conduction is taken into account in our model."
+ This leads to efficient mass evaporation from the disk to the corona. feeding a corona to higher mass density than that in tlie coronal envelope shown in MIID simulations.," This leads to efficient mass evaporation from the disk to the corona, feeding a corona to higher mass density than that in the coronal envelope shown in MHD simulations."
+ Therefore. we expect relatively higher ;2 value in our case.," Therefore, we expect relatively higher $\beta$ value in our case."
+ We investigate the influence of coronal magnetic fields on the structure of an accretion disk corona sustained by dissipative energy release aud thermal coupling between corona ancl disk., We investigate the influence of coronal magnetic fields on the structure of an accretion disk corona sustained by dissipative energy release and thermal coupling between corona and disk.
+ Numerical calculations show (hat (he relation between mass evaporation rate and radius Qn - r curve) svstematically shilts to smaller radii with both an increase of magnetic pressure (decreasing ;2) and a decrease of heat conductivity (decreasing &)., Numerical calculations show that the relation between mass evaporation rate and radius $\dot m$ - r curve) systematically shifts to smaller radii with both an increase of magnetic pressure (decreasing $\beta$ ) and a decrease of heat conductivity (decreasing $\kappa$ ).
+ The location of maximal evaporation rate lies between 20/24 and 200734 when 9 ranges from 2 to x or & [rom Aa Lo 0μμ., The location of maximal evaporation rate lies between $30R_{\rm S}$ and $200R_{\rm S}$ when $\beta$ ranges from 2 to $\infty$ or $\kappa$ from $\kappa_{\rm Sp}$ to $0.2\kappa_{\rm Sp}$.
+" However. the maximal evaporation rate remains almost the same (~0.03pa. wil[un energv conversion efficiency)= 0.1) If the disk truncation and state transition are indeed caused by an evaporation process. (he transition Iuminosity predicted by our disk evaporation model is L;,=0.03Laq before transition to the soft state."," However, the maximal evaporation rate remains almost the same $\sim 0.03 \dot{M}_{\rm 
+ Edd}$, with energy conversion efficiency $\eta=0.1$ ) If the disk truncation and state transition are indeed caused by an evaporation process, the transition luminosity predicted by our disk evaporation model is $L_{tr}=0.03 L_{\rm Edd}$ before transition to the soft state."
+ The inner edge of the outer thin disk can be several tei (o several hundred Schwarzschild radii depending on the strength of magnetic field and its elfect on the heat conduction.," The inner edge of the outer thin disk can be several ten to several hundred Schwarzschild radii, depending on the strength of magnetic field and its effect on the heat conduction."
+ This alleviates the problem that the previous evaporation models predict too large a disk (runcation radius before (he transition, This alleviates the problem that the previous evaporation models predict too large a disk truncation radius before the transition
+Galaxies with polar rings (PRGs) are an interesting case of,Galaxies with polar rings (PRGs) are an interesting case of
+assuming a planetary magnetic fiekl. the flow of the charged winds will induce a new component of the magnetic field wilh associated currents.,"assuming a planetary magnetic field, the flow of the charged winds will induce a new component of the magnetic field with associated currents."
+ The winds will experience a drag ? show that (his magnetic drag could potentially be strong enough to reduce zonal wind speeds by up to a factor of ~3 and alter the flow structure., The winds will experience a drag – \citet{per10a} show that this magnetic drag could potentially be strong enough to reduce zonal wind speeds by up to a factor of $\sim$ 3 and alter the flow structure.
+ SUL. it ds not straightforward (o. Granslate (hese wind speeds into the (ransniüssion spectral signature of a hot. Jupiter atmosphere.," Still, it is not straightforward to translate these wind speeds into the transmission spectral signature of a hot Jupiter atmosphere."
+ In (his paper we attempt (o do just that., In this paper we attempt to do just that.
+ We couple together an existing 2-D alinospheric dvnamies model (??) with a transmission spectroscopy code (??) such that the 3-D temperature-pressure and wind speed structures are sell-consistentlv. treated. in the transmission radiative transfer.," We couple together an existing 3-D atmospheric dynamics model \citep{rau10, per10a} with a transmission spectroscopy code \citep{mil09, mil10} such that the 3-D temperature-pressure and wind speed structures are self-consistently treated in the transmission radiative transfer."
+ The end result is to produce Doppler-shifted transmission spectra that include the effects of velocity shifts due to winds along with planetary rotation and orbital motion on the planets absorption features., The end result is to produce Doppler-shifted transmission spectra that include the effects of velocity shifts due to winds along with planetary rotation and orbital motion on the planet's absorption features.
+ Previous work bv ? studied the velocity effects of rotation on the transmission spectra of exoplanets but did not include the more dominant. effect of winds., Previous work by \citet{spi07} studied the velocity effects of rotation on the transmission spectra of exoplanets but did not include the more dominant effect of winds.
+ Our work also goes bevond previous work coupling 3-D at(mospheric dvnanmies and radiative transfer models. which have sell-consistenilv ireated the temperature-pressure structure wilh (he radiative transfer. while neelecting the Doppler shift inducing effect of winds (?7?)..," Our work also goes beyond previous work coupling 3-D atmospheric dynamics and radiative transfer models, which have self-consistently treated the temperature-pressure structure with the radiative transfer, while neglecting the Doppler shift inducing effect of winds \citep{for10, bur10}."
+ Our goal is to determine whether measurements of Doppler shifted winds in hot Jupiter abmospheres during (ransit could be used to constrain (he nature of the wind patterns themselves., Our goal is to determine whether measurements of Doppler shifted winds in hot Jupiter atmospheres during transit could be used to constrain the nature of the wind patterns themselves.
+ Furthermore. we attempt to determine whether a 2km + blueshift. as seen in the ? transmission spectrum of WD 209458b. is consistent with the results ol the aliospheric dynamics modeling.," Furthermore, we attempt to determine whether a 2km $^{-1}$ blueshift, as seen in the \citet{sne10} transmission spectrum of HD 209458b, is consistent with the results of the atmospheric dynamics modeling."
+ Specilically. we look at clvuamics models that use clilferent prescriptions for the treatment of drag in the atmospheres of hot Jupiters.," Specifically, we look at dynamics models that use different prescriptions for the treatment of drag in the atmospheres of hot Jupiters."
+ We study. cases with no additional sources of atmospheric drag along with cases where prescriptions for magnetic drag have been included. (ο determine the overall effect on the planetary. wind patterns as well as how these assumptions can alter the transmission signature of such a planet.," We study cases with no additional sources of atmospheric drag along with cases where prescriptions for magnetic drag have been included, to determine the overall effect on the planetary wind patterns as well as how these assumptions can alter the transmission signature of such a planet."
+ We also look into the possibility of recovering the spatial structure of exoplanet winds through (iransmission spectroscopy. both by using individual spectral lines to probe different heights in (he plauet's atmosphere and by measuring Doppler shifts during transit ingress and egress to separately map out wind speeds across the planet's eastern and western terminators.," We also look into the possibility of recovering the spatial structure of exoplanet winds through transmission spectroscopy, both by using individual spectral lines to probe different heights in the planet's atmosphere and by measuring Doppler shifts during transit ingress and egress to separately map out wind speeds across the planet's eastern and western terminators."
+ We present our model and methodology in Section 2.., We present our model and methodology in Section \ref{methods}. .
+ Our transmission spectrum results are presented in Section 3.., Our transmission spectrum results are presented in Section \ref{results}.
+ Finally. we offer up some summarizingthoughts aud conclusions in Section 4..," Finally, we offer up some summarizingthoughts and conclusions in Section \ref{concl}. ."
+in relativistic hyclrodvnamies there are no explicit relations eiving the primitive variables p. 3 and p in terms of the conserved. variables Lp. wl? and wh?p.,"in relativistic hydrodynamics there are no explicit relations giving the primitive variables $\rho$ , $\beta$ and $p$ in terms of the conserved variables $\Gamma \rho$, $w \Gamma^2
+\beta$ and $w\Gamma^2 -p$."
+ Therefore an iterative algorithm must be developed to do this., Therefore an iterative algorithm must be developed to do this.
+ The precise nature of this algorithm can allect how high a Lorentz [actor can be simulated by the code., The precise nature of this algorithm can affect how high a Lorentz factor can be simulated by the code.
+ We have found that the code used here will simulate Lows up to Lorentz factors of at least several hundred. reliably., We have found that the code used here will simulate flows up to Lorentz factors of at least several hundred reliably.
+ telativistic shocks. in the presence of magnetic Ducetuations which enable multiple shock crossings. can accelerate xwticles bv the cillusive shock mechanism (irk&Schnei-cer 1989).," Relativistic shocks, in the presence of magnetic fluctuations which enable multiple shock crossings, can accelerate particles by the diffusive shock mechanism \cite{kirkschneider89}."
+. As described in the introduction it has been shown that at ultra-relativistic shocks this produces a power aw distribution of particles with index 2.23. Le. N(E)poc?5 .so that the energy density does not diverge at high xwlicle energies.," As described in the introduction it has been shown that at ultra-relativistic shocks this produces a power law distribution of particles with index $2.23$, i.e. $N(E)\propto E^{-2.23}$ , so that the energy density does not diverge at high particle energies."
+ The acceleration timescale in relativistic lows will scale roughly. with the particle mean free path in he turbulent magnetic field. divided by the shock speed: although there is no published. derivation of the precise acceleration timescale., The acceleration timescale in relativistic flows will scale roughly with the particle mean free path in the turbulent magnetic field divided by the shock speed; although there is no published derivation of the precise acceleration timescale.
+ In this paper we are interested. in what happens to the power law population produced: by highly. relativistic shocks over the much longer timescale which is the size of the blastwave civided by e., In this paper we are interested in what happens to the power law population produced by highly relativistic shocks over the much longer timescale which is the size of the blastwave divided by $c$.
+ Pherefore. over hvdrocdsynamical timescales. we can treat acceleration as an impulsive injection of energetic particles. with this universal power law up to arbitrarily high energies. into each fluid element which passes through a relativistic shock.," Therefore, over hydrodynamical timescales, we can treat acceleration as an impulsive injection of energetic particles, with this universal power law up to arbitrarily high energies, into each fluid element which passes through a relativistic shock."
+ We will address the role of an intrinsic cut-olf in. the shock accelerated: population. either as a result of a finite acceleration timescale or loss processes. in à future paper.," We will address the role of an intrinsic cut-off in the shock accelerated population, either as a result of a finite acceleration timescale or loss processes, in a future paper."
+ The particles subsequently suller svnchrotron and acdiabatic losses from which we can calculate the emission., The particles subsequently suffer synchrotron and adiabatic losses from which we can calculate the emission.
+ We need to introduce three [ree parameters which. when taken together with the results of our hydrodynamical simulations. will give a complete picture of the blastwave's evolution and the associated emission of energetic particles.," We need to introduce three free parameters which, when taken together with the results of our hydrodynamical simulations, will give a complete picture of the blastwave's evolution and the associated emission of energetic particles."
+ These parameters are Once an electron. is injected at the shock it will. be scattered in the local [uid frame and suller both svachrotron ancl adiabatic losses., These parameters are Once an electron is injected at the shock it will be scattered in the local fluid frame and suffer both synchrotron and adiabatic losses.
+ Ho the length scale. over which the electrons are. isotropised. is much shorter than any other length. scale of interest then the relativistic electrons will respond acdiabatically to the expansion or contraction of the How., If the length scale over which the electrons are isotropised is much shorter than any other length scale of interest then the relativistic electrons will respond adiabatically to the expansion or contraction of the flow.
+Acliabatic losses. or indeed. gains. are then described by the fact that. p/pH is constant where p and. p are the particle momentum and Εις density in the local [uid frame.,"Adiabatic losses, or indeed gains, are then described by the fact that $p/\rho^{1/3}$ is constant where $p$ and $\rho$ are the particle momentum and fluid density in the local fluid frame."
+ With £=pe for ultra-relativistic particles the combined svnehrotron ancl adiabatic losses are described in the comoving frame by where à is à constant and D the magnetic field strength in the local uid. frame., With $E=pc$ for ultra-relativistic particles the combined synchrotron and adiabatic losses are described in the comoving frame by E where $\alpha$ is a constant and $B$ the magnetic field strength in the local fluid frame.
+ Consider now a Iuid element which is shocked at time fo when a power law distribution of energetic particles is injected and. where / is time measured in the comoving frame., Consider now a fluid element which is shocked at time $\that_0$ when a power law distribution of energetic particles is injected and where $\that$ is time measured in the comoving frame.
+" This population then evolves according to the equation E) where NGE.E) is the differential number οἱ particles of energy fat time {,"," This population then evolves according to the equation ) where $N(E,\that)$ is the differential number of particles of energy $E$ at time $\that$."
+" Phe losses. £. are given by equation (6)) while the injection term is (QUE.i=QuatFu)UEιοminds ""(GL being the Leviside function) which describes injection of a power law spectrum. starting at time {ο with a minimum particle energv Luvin."," The losses, $\dot{E}$, are given by equation \ref{energyloss}) ) while the injection term is $Q(E,\that)=Q_0\delta(\that-\that_0)H(E-E_{\rm min})E^{-p}$ $H$ being the Heviside function) which describes injection of a power law spectrum, starting at time $\that_0$ with a minimum particle energy $E_{\rm min}$."
+ We can solve equation (7)) by finding the characteristic curves along which N(LDdEοEs where a particle with energy Ley at fy cools toan energy £ at time /. From equation (6)) we have left(( which can be solved to show that along a characteristic curve AV is conserved so that and the solution to (7)) becomes," We can solve equation \ref{kineticeqn}) ) by finding the characteristic curves along which $N(E,\that)dE=N(E_0,\that_0)dE_0$ where a particle with energy $E_0$ at $\that_0$ cools toan energy $E$ at time $\that$ From equation \ref{energyloss}) ) we have ) which can be solved to show that along a characteristic curve $N$ is conserved so that and the solution to \ref{kineticeqn}) ) becomes"
+between a reddened background star and a voung star.,between a reddened background star and a young star.
+ Therefore. we have used a statistical approach to remove contaminating field stars from the A71/—Iv. CM diagram of the cluster region (we will discuss this further in detail in the following Sect.," Therefore, we have used a statistical approach to remove contaminating field stars from the $K/H-K$ CM diagram of the cluster region (we will discuss this further in detail in the following Sect."
+ 3.2)., 3.2).
+ The cluster member stars are clearly seen in (he statistically cleaned. CM diagram panel)., The cluster member stars are clearly seen in the statistically cleaned CM diagram ).
+ Our ain here is to estimate (he approximate age of the lower mass population by comparing our data with theoretical evolutionary tracks., Our aim here is to estimate the approximate age of the lower mass population by comparing our data with theoretical evolutionary tracks.
+ The low mass evolutionary tracks have been tested on various observational data of voung stellar clusters and are thought to be &ood representations., The low mass evolutionary tracks have been tested on various observational data of young stellar clusters and are thought to be good representations.
+ The evolutionary models for hieh mass stars are more complicated and varied. however. and are not well tested. against observational data sets.," The evolutionary models for high mass stars are more complicated and varied, however, and are not well tested against observational data sets."
+ Therefore. we first use the pre-main-sequence (PAIS) isochrone fitting for the lower mass population.," Therefore, we first use the pre-main-sequence (PMS) isochrone fitting for the lower mass population."
+ Further. it should be noted that a CM diagram that uses only one colour may not be very ellective in describing the less understood massive voung stars.," Further, it should be noted that a CM diagram that uses only one colour may not be very effective in describing the less understood massive young stars."
+ since the stellar photosphere's ave better represented by shorter wavelength data such as (he optical. J and Lf bands. we use these data when comparing source photometry with (he model isochrones.," Since the stellar photosphere's are better represented by shorter wavelength data such as the optical, $J$ and $H$ bands, we use these data when comparing source photometry with the model isochrones."
+ In the region of Fig., In the region of Fig.
+ 1. all optically detected stars in the BVRL bands," 1, all optically detected stars in the $BVRI$ bands"
+being stripped. as well as cliscontinuitics associated with bow shocks (see (see Paper | for a discussion of these features as seen [rom the entropy maps of simulations).,"being stripped, as well as discontinuities associated with bow shocks (see (see Paper I for a discussion of these features as seen from the entropy maps of simulations)."
+ H is tempting to associate such structures to the smallscale features revealed by Chandra observations. which in some cases are associated with bow shocks or cold fronts from merging structures (e.g. Alarkevitch et al.," It is tempting to associate such structures to the small–scale features revealed by Chandra observations, which in some cases are associated with bow shocks or cold fronts from merging structures (e.g., Markevitch et al."
+ 2000. Mazzotta et al.," 2000, Mazzotta et al."
+ 2001. Ettori Fabian 2001. and references therein).," 2001, Ettori Fabian 2001, and references therein)."
+ At 2— Othe ComaLike halo is uncergoing several merging events with massive sub.groups. at dillerent stages of advancement. much like what is seen in the XMM observations of the real Coma cluster (Briel et al.," At $z=0$ the Coma–like halo is undergoing several merging events with massive sub–groups, at different stages of advancement, much like what is seen in the XMM observations of the real Coma cluster (Briel et al."
+ 2001)., 2001).
+" The bolometric luminosity for a set of Δίων gas particles is defined. as (lhwhere m, is the proton mass. maga. ds the mass of à gas particle (see Table 12). p; and 7; are the density and temperature at the position of the 7-th particle. respectively. and j£;=0.6 is the mean molecular weight for a primordial gas composition with mass provided. by hvdrogen."," The bolometric luminosity for a set of $N_{gas}$ gas particles is defined as , where $m_p$ is the proton mass, $m_{gas}$ is the mass of a gas particle (see Table \ref{t:simul}) ), $\rho_i$ and $T_i$ are the density and temperature at the position of the $i$ -th particle, respectively, and $\mu=0.6$ is the mean molecular weight for a primordial gas composition with mass provided by hydrogen."
+ Assumingthat the main contribution to the .X , Assumingthat the main contribution to the $X$ 
+has established in the solar corona there are ubiquitous propagating low amplitude (221 km |) Alfvénnic waves with a wide range of frequencies.,has established in the solar corona there are ubiquitous propagating low amplitude $\approx$ 1 km $^{-1}$ ) Alfvénnic waves with a wide range of frequencies.
+ Realistically interpreting these propagating waves as kink modes. it was predicted thev should exhibit a frequency dependent damping length due to resonant absorption. governed by the TGV relation which shows that naturally occurring transversal plasma inhomogeneily in coronal magnetic [lux (tubes causes them (o act as low-pass filters.," Realistically interpreting these propagating waves as kink modes, it was predicted they should exhibit a frequency dependent damping length due to resonant absorption, governed by the TGV relation which shows that naturally occurring transversal plasma inhomogeneity in coronal magnetic flux tubes causes them to act as low-pass filters."
+ It was found the observed frequency dependence on clamping length can be explained by (he kink wave interpretation. al least up to about. /=8 mllz. and furthermore. (he spatially averaged. equilibrium parameter describing the length scale of transverse plasma density inhomogeneltv over a svstem of coronal loops is consistent with the range of values estimated from TRACE observations of standing kink modes.," It was found the observed frequency dependence on damping length can be explained by the kink wave interpretation, at least up to about $f=8$ mHz, and furthermore, the spatially averaged equilibrium parameter describing the length scale of transverse plasma density inhomogeneity over a system of coronal loops is consistent with the range of values estimated from TRACE observations of standing kink modes."
+ Due to its restricted field of view. full loops observed with TRACE are shorter than the average loop length of the coronal loop svstem observed with CoMP.," Due to its restricted field of view, full loops observed with TRACE are shorter than the average loop length of the coronal loop system observed with CoMP."
+ The similar inferred (ransversal plasma inhomogeneily length scales. suggest this parameter is independent of loop length.," The similar inferred transversal plasma inhomogeneity length scales, suggest this parameter is independent of loop length."
+ Importantly. it was also found that the estimated least squares value of the ratio of initial outward and inward power was approximately unity. suggesting the kink wave power generated at both footpoints in the lower atmosphere was almost equal.," Importantly, it was also found that the estimated least squares value of the ratio of initial outward and inward power was approximately unity, suggesting the kink wave power generated at both footpoints in the lower atmosphere was almost equal."
+ Previously much work has been undertaken (o study the properties of (he well known standing kink modes in post-llare coronal loops., Previously much work has been undertaken to study the properties of the well known standing kink modes in post-flare coronal loops.
+ The results of this Letter now open the wav lo investigating the wave contribution to coronal heating by exploiting the latest hieh spatial/temporal resolution observations of propagating kink waves modes., The results of this Letter now open the way to investigating the wave contribution to coronal heating by exploiting the latest high spatial/temporal resolution observations of propagating kink waves modes.
+ An accurate estimate of coronal kink wave clamping rates due to resonant. absorption is crucial in order to quantify the amount of this wave energv that can contribute to plasma heating., An accurate estimate of coronal kink wave damping rates due to resonant absorption is crucial in order to quantify the amount of this wave energy that can contribute to plasma heating.
+ Droadband frequency disturbances propagating along (ransversally inhomogeneous coronal Εαν tubes will generate many resonances that will phase nix. causing a cascade of energy to smaller scales where dissipation becomes more elficient. (hereby creating the necessary," Broadband frequency disturbances propagating along transversally inhomogeneous coronal flux tubes will generate many resonances that will phase mix, causing a cascade of energy to smaller scales where dissipation becomes more efficient, thereby creating the necessary"
+would have seemed (he strongest candidate source.,would have seemed the strongest candidate source.
+ Indeed. a single temperature WD cloes provide an excellent [it to Lya and fits the widths of the metallic absorption lines. but ihe match to the continuum slope is less (han satisfactory ancl vields a pronounced mocel (hix deficit relative to the data shortware of 1520A.," Indeed, a single temperature WD does provide an excellent fit to $\alpha$ and fits the widths of the metallic absorption lines, but the match to the continuum slope is less than satisfactory and yields a pronounced model flux deficit relative to the data shortward of $\sim 1520$."
+". If E(D-V) = 0.0. the best single temperature WD fit (with log g = 8.3) is T;yy=29.000K. 0,sin/=200xI00km/s with a scale Lactor-clerived distance of 506 pe."," If E(B-V) = 0.0, the best single temperature WD fit (with log g = 8.3) is $T_{eff} = 29,000$ K, $v_{rot} \sin{i}= 200 \pm 100 $ km/s with a scale factor-derived distance of 506 pc."
+ This value of distance is quite close to our adopted distance in Table 1., This value of distance is quite close to our adopted distance in Table 1.
+ The single temperature WD fit is displavecl in figure 4a., The single temperature WD fit is displayed in figure 4a.
+ ILowever. in view of the relatively large distance we favored. il was prudent to examine the effect on the fits if reddening is assumed.," However, in view of the relatively large distance we favored, it was prudent to examine the effect on the fits if reddening is assumed."
+ For (his exercise. we used both the STIS data in quiescence and IUE data of TT Crt in outburst.," For this exercise, we used both the STIS data in quiescence and IUE data of TT Crt in outburst."
+ We assumed different values of E(D-V) (0.00. 0.05. 0.10. 0.15) for both the STIS observations aud the IUE observations.," We assumed different values of E(B-V) (0.00, 0.05, 0.10, 0.15) for both the STIS observations and the IUE observations."
+ We modeled (he STIS data with a WD only and the IUE outburst data with a disk only but the best result. (lor whieh IUE and STIS results are consistent) is for E(B-V)=0.15 which vields a distance of about 220 pe and 7;y=30. 000K for a 1.0... white dwarf.," We modeled the STIS data with a WD only and the IUE outburst data with a disk only but the best result (for which IUE and STIS results are consistent) is for E(B-V)=0.15 which yields a distance of about 220 pc and $T_{eff} = 30,000$ K for a $1.0 M_{\odot}$ white dwarf."
+ This distance is well below our adopted value [rom Table 1., This distance is well below our adopted value from Table 1.
+ For lower values of E(B-V). the fits between IUE and STIS are discordant. e.g. ECB-V)— 0 leads to a distance of 375pe lor STIS ancl LOOpe for IUE (M—1.0).," For lower values of E(B-V), the fits between IUE and STIS are discordant, e.g. E(B-V)= 0 leads to a distance of 375pc for STIS and 100pc for IUE (M=1.0)."
+ In fact. the reddening constrains the distance to be less (han 500 pe because the value 500pc corresponds to E(D-V) = (0.0 for a white dwarf mass of 0.3... or greater.," In fact, the reddening constrains the distance to be less than 500 pc because the value 500pc corresponds to E(B-V) = 0.0 for a white dwarf mass of $0.8 M_{\odot}$ or greater."
+ To get a larger distance. one would have Lo asstume a much smaller mass which is inconsistent with previous work.," To get a larger distance, one would have to assume a much smaller mass which is inconsistent with previous work."
+ In any case. a large distance would be consistent with a larger ECD-V) but as E(D-V) increases. the resulting distance decreases (as the flux increases).," In any case, a large distance would be consistent with a larger E(B-V) but as E(B-V) increases, the resulting distance decreases (as the flux increases)."
+ An optically thick disk structure alone does a poor job of matching (he data because it utterly. fails to fit anv of the sharp metallic lines. mismatches the continuum slope. ancl has a marked flux deficit relative to the data.," An optically thick disk structure alone does a poor job of matching the data because it utterly fails to fit any of the sharp metallic lines, mismatches the continuum slope, and has a marked flux deficit relative to the data."
+ The situation is significantly improved with combinations of a WD plus an optically thick disk ancl a(wo-temperature WD., The situation is significantly improved with combinations of a WD plus an optically thick disk and atwo-temperature WD.
+ The 2-T WD model with Zipp=31. 000Ix and 45.0001. matches the continuum slope and sharp metallic lines as well as Lya but there remains a flux deficit between roughly 1600 aand 1250A.," The 2-T WD model with $T_{eff} = 31,000$ K and 45,000K matches the continuum slope and sharp metallic lines as well as $\alpha$ but there remains a flux deficit between roughly 1600 and 1250."
+. The combination of a WD with 7;;;=30. 000 and an optically thick disk accomplishes everything the 2-T WD does and reveals less of a flux deficit though a deficit of model fluxes relative to the data is still evident shortward of 1400À.," The combination of a WD with $T_{eff} = 30,000$ K and an optically thick disk accomplishes everything the 2-T WD does and reveals less of a flux deficit though a deficit of model fluxes relative to the data is still evident shortward of 1400."
+. In figure th. the best-fitting WD + disk combination is shown for comparison.," In figure 4b, the best-fitting WD + disk combination is shown for comparison."
+ V442 Cen presented a number of challenges., V442 Cen presented a number of challenges.
+ Since it has a large recldeninge and distance. iis very likely that the Lyva feature is mainiv due to the ISM atomic hydrogen opacitv.," Since it has a large reddening and distance, its very likely that the $\alpha$ feature is mainly due to the ISM atomic hydrogen opacity."
+" We therefore decided to include in our modeling the ISM atomic hydrogen by using an ISM moclel with an atomic hydrogen column density N(LLOQ)=5x10?""em 7. a turbulent. velocity of 40"," We therefore decided to include in our modeling the ISM atomic hydrogen by using an ISM model with an atomic hydrogen column density $5 \times 10^{20}$ $^{-2}$ , a turbulent velocity of 40"
+"X,/mqgzxo(0).",$\vec{X}_n=\sqrt{m_n}x_0(0)$.
+" The time evolution of the system is given by: Substituting the initial conditions, and the expression in Eq. (12))"," The time evolution of the system is given by: Substituting the initial conditions, and the expression in Eq. \ref{eigen1}) )"
+" for the eigenvector components, we get The coordinate of the large oscillator is simply given by xo(t)=Xo(t)/V/M."," for the eigenvector components, we get The coordinate of the large oscillator is simply given by $x_0(t)=\vec{X}_0(t)/\sqrt{M}$."
+" For the continuum of small modes, the above expansion breaks down, since the eigenvalue equation has no real solutions inside the continuum range."," For the continuum of small modes, the above expansion breaks down, since the eigenvalue equation has no real solutions inside the continuum range."
+" However, the edge modes are well defined, and they determine the dynamics at late times."," However, the edge modes are well defined, and they determine the dynamics at late times."
+" Therefore, for the continuum case we can still write down the analogous expression which is valid only at late times: The sums of Eq. (14))"," Therefore, for the continuum case we can still write down the analogous expression which is valid only at late times: The sums of Eq. \ref{no10}) )"
+" are replaced with the corresponding integrals, and we have the following expression for the displacement of the large oscillator at late times: This expression is in excellent agreement with the numerical simulations."," are replaced with the corresponding integrals, and we have the following expression for the displacement of the large oscillator at late times: This expression is in excellent agreement with the numerical simulations."
+" In the numerical example of subsection 2.1, the upper edge mode dominates the late-time behavior of the system, and its calculated contribution is plotted in Fig. 3,,"," In the numerical example of subsection $2.1$, the upper edge mode dominates the late-time behavior of the system, and its calculated contribution is plotted in Fig. \ref{Fig2B},"
+ together with the numerically simulated motion., together with the numerically simulated motion.
+ We now add an extra degree of realism by introducing viscous friction into the system., We now add an extra degree of realism by introducing viscous friction into the system.
+" In MHD, continuum modes are spatially localized, and the effect of viscosity is to frictionally couple the neighboring modes (see, e.g., Hollweg 1987)."," In MHD, continuum modes are spatially localized, and the effect of viscosity is to frictionally couple the neighboring modes (see, e.g., Hollweg 1987)."
+" In our simple model we introduce viscosity by adding frictional forces between the neighboring oscillators, where fnn+i is the force from the n’th oscillator acting on the (n+ 1)'th."," In our simple model we introduce viscosity by adding frictional forces between the neighboring oscillators, where $f_{n,n+1}$ is the force from the $n$ 'th oscillator acting on the $(n+1)$ 'th."
+ We now calculate how the system dissipates energy as a function of time., We now calculate how the system dissipates energy as a function of time.
+ We will show that it occurs in two stages (see Fig. 6)): (, We will show that it occurs in two stages (see Fig. \ref{Fig4}) ): (
+"1) Initially, the small oscillators are strongly and simultaneously excited by the ""Landau-damped"" large oscillator, then they become dephased, with the average relative motion between the neighboring oscillators growing linearly in time.","1) Initially, the small oscillators are strongly and simultaneously excited by the ""Landau-damped"" large oscillator, then they become dephased, with the average relative motion between the neighboring oscillators growing linearly in time."
+ This leads to a very rapid dissipation of the bulk of the initial energy. (, This leads to a very rapid dissipation of the bulk of the initial energy. (
+"2) The edge modes persist, since the participating small oscillators move in phase and the energy dissipation is small.","2) The edge modes persist, since the participating small oscillators move in phase and the energy dissipation is small."
+ The energy of the modes is damped exponentially on a timescale much longer than that of the first stage., The energy of the modes is damped exponentially on a timescale much longer than that of the first stage.
+" 'The dissipated energy is given by In the continuum limit, the small oscillators are labeled not by a discrete index n, but by a continuous variable A."," The dissipated energy is given by In the continuum limit, the small oscillators are labeled not by a discrete index $n$, but by a continuous variable $\lambda$."
+ 'The expression for the dissipated energy is then where  is the viscous coefficient., The expression for the dissipated energy is then where $\tilde{\gamma}$ is the viscous coefficient.
+" After the initial exponential damping of the large oscillator and the excitation of the small oscillators, the latter initially move independently, with"," After the initial exponential damping of the large oscillator and the excitation of the small oscillators, the latter initially move independently, with"
+preferred effective radius. nuclear point source contribution. or distribution of residual Iux.,"preferred effective radius, nuclear point source contribution, or distribution of residual flux."
+ The modelled: data and mocel-subtracted residuals are displaved in Fig. I0.. ," The modelled data and model-subtracted residuals are displayed in Fig. \ref{Fig: morph},"
+together with a plot of the variation in reduced. 47 over the parameter space considered., together with a plot of the variation in reduced $\chi^2$ over the parameter space considered.
+ ‘The resulting best-fit parameters give an ellective radius. of rap=1424+0.77 aresec (equivalent to 6.47+3.51kpc in our assumed cosmological moclel). and an unresolved nuclear point source contribution of 21.2-4.9 the residuals are comparable to the background. noise.," The resulting best-fit parameters give an effective radius of $r_{\rm eff} = 1.42 \pm 0.77$ arcsec (equivalent to $6.47 \pm 3.51$ kpc in our assumed cosmological model), and an unresolved nuclear point source contribution of $21.2 
+\pm 4.9$; the residuals are comparable to the background noise."
+" After subtraction of the point source component. the remaining ποσαΕς Lux is comparable to that of a I.4L, galaxy at a redshift of 0.31 (Willott ct al 2003)."," After subtraction of the point source component, the remaining near-IR flux is comparable to that of a $\sim 1.4 \rm L_{\star}$ galaxy at a redshift of $\sim 0.31$ (Willott et al 2003)."
+ Assumingὃν that the point source component is norestellar in nature. and that the bulk old. stellar population of the galaxy. formed at high. redshift (ic. 2=5—20). 1ο galaxy Lux can be mocelled by a 9.5vr. old. stellar population.," Assuming that the point source component is non-stellar in nature, and that the bulk old stellar population of the galaxy formed at high redshift (i.e. $z = 5-20$ ), the galaxy flux can be modelled by a 9.5Gyr old stellar population."
+ We have convolved a 9.56Gvr GISSEL spectral synthesis (Bruzual Charlot 2003) template spectrum with 16 ANs-band. filter profile., We have convolved a 9.5Gyr GISSEL spectral synthesis (Bruzual Charlot 2003) template spectrum with the $K_S$ -band filter profile.
+ Comparing the resulting Dux =μαith that measured. for PINS2250-41. Calter subtraction of the point source contribution measured. via our modelling) sugeests a total stellar mass for the galaxy of 2.1x0.5104M..., Comparing the resulting flux with that measured for PKS2250-41 (after subtraction of the point source contribution measured via our modelling) suggests a total stellar mass for the galaxy of $2.1 \pm 0.5 \times 10^{11} \rm M_{\odot}$ .
+ We have also carried out. aperture photometry of the host galaxy using the E547M WEPC? data of P05. obtaining a host galaxy magnitude in this filter of 20.5440.12.," We have also carried out aperture photometry of the host galaxy using the F547M WFPC2 data of T05, obtaining a host galaxy magnitude in this filter of $20.54 \pm 0.12$."
+ Comparison of the relative (ux densities observed. in. each of the two continuum filters (E547M and Avs) confirms that they are consistent with the spectral energy distribution expected for a 9.5Gvr old stellar population., Comparison of the relative flux densities observed in each of the two continuum filters (F547M and $K_S$ ) confirms that they are consistent with the spectral energy distribution expected for a 9.5Gyr old stellar population.
+ lt is interesting to note that the filamentary line emission displaved in Fig., It is interesting to note that the filamentary line emission displayed in Fig.
+ 6 is matched by an extension in the optical continuum observed. from this galaxy. along the same position angle (Fig. 11))., \ref{oiii_closeup} is matched by an extension in the optical continuum observed from this galaxy along the same position angle (Fig. \ref{overlays}) ).
+ The extension to the host galaxy is clearly visible in the optical. and it is also noteworthy that the lowest surface brightness Ht emission is somewhat disturbed in this region of the galaxy. perhaps rellecting a distortion in the underlving galaxy structure.," The extension to the host galaxy is clearly visible in the optical, and it is also noteworthy that the lowest surface brightness IR emission is somewhat disturbed in this region of the galaxy, perhaps reflecting a distortion in the underlying galaxy structure."
+" Figure 12. clisplavs the FORSIL spectrum for the companion ealaxy to the north east. extracted within a 1"" aperture."," Figure \ref{companion_spec} displays the FORS1 spectrum for the companion galaxy to the north east, extracted within a $1^{\prime\prime}$ aperture."
+ The strongest. feature in the spectrum: is an absorption line consistent with LL? at a redshift of +=0.309., The strongest feature in the spectrum is an absorption line consistent with $\beta$ at a redshift of $z = 0.309$.
+ The continuum shape is also generally consistent with a galaxy at the assumed redshift. and we also see a noisy absorption features at wavelengths appropriate to the C-band. Ες and 19.," The continuum shape is also generally consistent with a galaxy at the assumed redshift, and we also see a noisy absorption features at wavelengths appropriate to the G-band, $\gamma$ and $\delta$."
+ Using the ΝΕΡΟ continuum cata of T05 and the LR ata presented in this paper. we have carried out aperture photometry of this source.," Using the WFPC2 continuum data of T05 and the IR data presented in this paper, we have carried out aperture photometry of this source."
+ In the optical. this companion ealaxy has a magnitude of 20.57+0.12. slightly fainter aan the radio source host galaxy.," In the optical, this companion galaxy has a magnitude of $20.57 \pm 0.12$, slightly fainter than the radio source host galaxy."
+ In the Hit. the observed As band magnitude is 16.48+0.08.," In the IR, the observed $K_S-$ band magnitude is $16.48 \pm 0.08$."
+ The relative [lux ensities of the observed-frame optical and infrared emission rom this galaxy are consistent with a single old stellar xpulation (OSD?) aged between 2-46vr., The relative flux densities of the observed-frame optical and infrared emission from this galaxy are consistent with a single old stellar population (OSP) aged between 2-4Gyr.
+ However. given 1e presence of Balmer absorption features in this galaxy a composite stellar population including both old and. voung," However, given the presence of Balmer absorption features in this galaxy a composite stellar population including both old and young"
+" 1977)). (Carlbere1986.Camu1987.Teruquist.Ill1993)). Veceraraghavan1985.Barnes1988.Lake1989)) (Neeropoute&White19011). (Barnes&Ieruquist1991.SandersI9588b)). (~LO?AL, OY Usormeudy1992)]. rlt -—100. "," \cite{Toomre77}) \cite{Carlberg86,Gunn87,HSH93}) \cite{VW85,Barnes88,Lake89}) \cite{Negroponte83,Noguchi86,BarnesH91}) \cite{BarnesH91,San88b}) $\sim 10^2 M_\odot$ $^{-3}$ \cite{Kormendy92}) $r^{1/4}$ $\sim$ "
+The representative data set for HSB galaxies is taken from the compilation of 3). to which E added data for 6 LSB aud LL dwarls galaxies (PetersonaudCaldwell 1993).. 6 compact ellipticals (CGtubathakturtaanevanderMare 2002).. aud 22 radio galaxies (Bettonietal.2001).,"The representative data set for HSB galaxies is taken from the compilation of \citet{bender93}, to which I added data for 6 LSB and 11 dwarfs galaxies \citep{peterson93}, , 6 compact ellipticals \citep{marel02}, , and 22 radio galaxies \citep{bettoni01}."
+. Throughout the paper. Hy=50 kms th . is used.," Throughout the paper, $_0=50$ km $^{-1}$ $^{-1}$ is used."
+ The acceler‘ation gx at re is fouud comparabe to ag lor HSB galaxies. while being well below this value in LSB.," The acceleration $g_N$ at $r_e$ is found comparable to $a_0$ for HSB galaxies, while being well below this value in LSB."
+ The efect of MOND on objec‘ts characterizes by different. X is shown in Fig. 3.., The effect of MOND on objects characterized by different $\Sigma$ is shown in Fig. \ref{g_sigma}.
+ It is seen that galaxies folOW a contiunuots sectence inplyiug he mass ciiscrepaney. or in other word the dark matter corteu. is uniquely define Nw X.," It is seen that galaxies follow a continuous sequence implying the mass discrepancy, or in other word the dark matter content, is uniquely defined by $\Sigma$."
+ I shall siow that this fact has important consequences for uuderstanding why objects as clierent as elobuar aud galaxy clusters may appear to have the same mass discre)alcv., I shall show that this fact has important consequences for understanding why objects as different as globular and galaxy clusters may appear to have the same mass discrepancy.
+ Ii Fig. 3..," In Fig. \ref{g_sigma},"
+ HSB aud LSB galaxies are no longer istingulslabe because {1e opposite dependence of X «i rs (see Fig. 2)), HSB and LSB galaxies are no longer distinguishable because the opposite dependence of $\Sigma$ on $r_e$ (see Fig. \ref{kormendy}) )
+ disapyears., disappears.
+" The different behavior of HSB aud LSB become evideut when plotti1g g/gy vers wr, orl (Fig.", The different behavior of HSB and LSB become evident when plotting $g/\gn$ versus $r_e$ or $L$ (Fig.
+ E aud 5))., \ref{msul_re} and \ref{msul_lum}) ).
+ It is tjerefore clear tiat the effec oL modified dyiamics is (Oo minc ai iucrease of 7 as a function of L for HSB ealaxles. while doiug the opposite for LSB £alaxies.," It is therefore clear that the effect of modified dynamics is to mimic an increase of $\tau$ as a function of $L$ for HSB galaxies, while doing the opposite for LSB galaxies."
+ Moreove Fie.," Moreover, Fig."
+" 5. shows that with good aj»proxitwliou. auc certaluly within the spread of the data. the depeudence ou luminosity is 7iXLY? Dos[o: HSB galaxies. and ixL"" [or LSB ealanies. COLlslslοιt with the values quoted in the Literatu‘e Pete‘sonandCalcwell1993 )."," \ref{msul_lum} shows that with good approximation, and certainly within the spread of the data, the dependence on luminosity is $\tau \propto L^{0.25}$ for HSB galaxies, and $\tau \propto L^{-0.3}$ for LSB galaxies, consistent with the values quoted in the literature \citep{bender93,peterson93}."
+". The strikiug clierent behavior of HSB aud LSB galaxies can thereore be easily explained by a unique model üsslluiug a moclilication of Newton""s law.", The striking different behavior of HSB and LSB galaxies can therefore be easily explained by a unique model assuming a modification of Newton's law.
+ Under this hypothesis the tilt of the FP is uot due to a systematic increase/decrease ὁ ‘the dark matter couteit of galaxies. but it is the ellec of he variatio 10f tle streugth of their ierial gravitational fiel as comparedl to αμ.," Under this hypothesis the tilt of the FP is not due to a systematic increase/decrease of the dark matter content of galaxies, but it is the effect of the variation of the strength of their internal gravitational field as compared to $a_0$."
+ Anoher wew to investigate the effect of LOD ou galaxies ds shown iu Fig. 6..," Another way to investigate the effect of MOND on galaxies is shown in Fig. \ref{mondExp},"
+ where the acceleralion of gravity from eq., where the acceleration of gravity from eq.
+" L. is compare dto the acceleration derived. using the clyuamica --lass. AL,tyn=.Gon*Jgfru."," 4, is compared to the acceleration derived using the dynamical mass, $M_{dyn}=3G \sigma^2/r_e$."
+ Trough there is signiIcaut spread of tlie data. it is seen that clyuamica --lasses 1iply uo mass cise'epaucy for acceleraion strouger or comparable to ag. while a deviallol is apparent wheu the acceleration goes below it. as prescribed by MOND.," Though there is significant spread of the data, it is seen that dynamical masses imply no mass discrepancy for acceleration stronger or comparable to $a_0$, while a deviation is apparent when the acceleration goes below it, as prescribed by MOND."
+ Note that the stronger iscrepanicy ocurs for LSB galaxies. because tvey are in deep MOND regime.," Note that the stronger discrepancy occurs for LSB galaxies, because they are in deep MOND regime."
+ While observalona ala all ΝΟD predictions agree reasonabv well. to explain this behavior usiug NewtOllal vnanules recti‘es fine tuning of the total amut of dark matter in galaxies of very clierent ype aid morplology.," While observational data and MOND predictions agree reasonably well, to explain this behavior using Newtonian dynamics requires fine tuning of the total amount of dark matter in galaxies of very different type and morphology."
+ These cotSderatious are not limited to ellipical galaxies., These considerations are not limited to elliptical galaxies.
+ In fact if MOND is the unclelviug aw regulatiu[n]0 [n]€yravity inthe universe. then all stellar svstems should be described by ΝOND.," In fact if MOND is the underlying law regulating gravity inthe universe, then all stellar systems should be described by MOND."
+" Thouel a full cliscussion of this issue goes beyou« the scope of the present work. it is worth o add nore data to Fiο,"," Though a full discussion of this issue goes beyond the scope of the present work, it is worth to add more data to Fig."
+ 6 to see whether other stellar systems different from elliptical galaxies follow he same tren., \ref{mondExp} to see whether other stellar systems different from elliptical galaxies follow the same trend.
+ In Fig., In Fig.
+ 7 I plot the data from the data for 1050 objects (51 elobular οἱ151915. κ)+)QU elliptical galaxies. 513 spiral galaxies. 170 galaxy groups aud 16 galaxy clusters) froi the compilation of Bursteinοἱal. (1997).. are plotted.," \ref{burstein1} I plot the data from the data for 1050 objects (51 globular clusters, 300 elliptical galaxies, 513 spiral galaxies, 170 galaxy groups and 16 galaxy clusters) from the compilation of \citet{burstein97}, , are plotted."
+Datacover 7 order of ruaguitudes in iiternal,Datacover 7 order of magnitudes in internal
+The total energv flux. Q(Total)=<Qí(ln)>+<Q(protonic)>.,"The total energy flux, $Q(\mathrm{Total})= <Q(\mathrm{lm})> +
+<Q(\mathrm{protonic})> $."
+ This measure of enerev flux can be distinguished Irom the enerev [lux injected into the plasmoid [rom the central engine., This measure of energy flux can be distinguished from the energy flux injected into the plasmoid from the central engine.
+" In order to assess the origin of the plasmoids. it is much more revealing to look at the power required to initiate the plasmold. Q4,Qnjected). rather than the redistribution of energy. [lux within the plasmoicd. far from the central engine."," In order to assess the origin of the plasmoids, it is much more revealing to look at the power required to initiate the plasmoid, $Q_{\mathrm{protonic}}(\mathrm{injected})$, rather than the redistribution of energy flux within the plasmoid, far from the central engine."
+ The power injected by (the central engine is delined by where {μις is the amount of time [ον the energy to be injected into the plasmoid. the flare rise time.," The power injected by the central engine is defined by where $t_{\mathrm{rise}}$ is the amount of time for the energy to be injected into the plasmoid, the flare rise time."
+ Similarly. lor the lepto-maenetic energy flix that is injected The radio flux densitv during (he rise of the flare was not sampled during the December 1993 flare.," Similarly, for the lepto-magnetic energy flux that is injected The radio flux density during the rise of the flare was not sampled during the December 1993 flare."
+ There is high time resolution radio sampling of four major Lares (hat was found in (he archival literature (hat is plotted in Figure 10., There is high time resolution radio sampling of four major flares that was found in the archival literature that is plotted in Figure 10.
+ The four fIares with rise lime coverage were all monitored at 15 GIz with the Ryle telescope in Rushtonetal(2010)., The four flares with rise time coverage were all monitored at 15 GHz with the Ryle telescope in \citet{rus10}.
+. Two of the flares were actually well sampled at 15 GlIz in Rushtonetal(2010).. the 1997 flare on MJD 50752 and the 1999 flare on MJD 51337.," Two of the flares were actually well sampled at 15 GHz in \citet{rus10}, the 1997 flare on MJD 50752 and the 1999 flare on MJD 51337."
+ The two 1998 flares in Figure 10 were not well sampled on their rise at 15 GIIz in Rushtonetal(2010)., The two 1998 flares in Figure 10 were not well sampled on their rise at 15 GHz in \citet{rus10}.
+. Fortunately. (here is excellent data sampling on (he rise time at 2.3 GlIz and 8.4 GHz for the 1998 flares on MJD 50917 and MJD 50934 in Dhawanetal(2000).," Fortunately, there is excellent data sampling on the rise time at 2.3 GHz and 8.4 GHz for the 1998 flares on MJD 50917 and MJD 50934 in \citet{dha00}."
+". The 8.4 GIIz data is used to estimate /,4 Lor these flares.", The 8.4 GHz data is used to estimate $t_{\mathrm{rise}}$ for these flares.
+ There is also a crude estimate on {μις for the 1.4 GU flux from the March 1994 flare of 1245 hours that was derived in Rodriguezetal(1995) bv extrapolating the motion of VLA components backwards in time combined wilh sparsely sampled radio monitoring., There is also a crude estimate on $t_{\mathrm{rise}}$ for the 1.4 GHz flux from the March 1994 flare of $12 \pm 5$ hours that was derived in \citet{rod95} by extrapolating the motion of VLA components backwards in time combined with sparsely sampled radio monitoring.
+ This point with its large error bars is also added to the plot., This point with its large error bars is also added to the plot.
+ The ligure indicates that stronger flares tend (to have larger μις., The figure indicates that stronger flares tend to have larger $t_{\mathrm{rise}}$.
+ The strongest flare from March 1994 may not follow the linear trend that fits the other 4 flares very well - the error bars are too large to make this determination., The strongest flare from March 1994 may not follow the linear trend that fits the other 4 flares very well - the error bars are too large to make this determination.
+ The linear fit [ον the other four data points is with a coefficient of determination of R?=0.9819., The linear fit for the other four data points is with a coefficient of determination of $R^{2} = 0.9819$.
+ Thus. it seems reasonable to use archival data to estimate the rise time of the December 1993 flare.," Thus, it seems reasonable to use archival data to estimate the rise time of the December 1993 flare."
+ The December 1993 Mares C1 and C2 have 15 GIIz flux densities similar to the MJD 50752 and MJD 50934 flares. so lice&T.5hrs222x10/sec is chosen as a fiducial /4;.," The December 1993 flares C1 and C2 have 15 GHz flux densities similar to the MJD 50752 and MJD 50934 flares, so $t_{\mathrm{rise}} \approx 7.5 \mathrm{hrs}\approx 2 \times 10^{4}
+\mathrm{sec}$ is chosen as a fiducial $t_{\mathrm{rise}}$."
+" The resultant Q values [rom equations (15) and (16) can be simply scaled with different choices of /,;...", The resultant $Q$ values from equations (15) and (16) can be simply scaled with different choices of $t_{rise}$.
+ The uncertainty here. based on monitored Lares in the rising phase seems to be a [actor of 2 at most.," The uncertainty here, based on monitored flares in the rising phase seems to be a factor of 2 at most."
+"For our analysis, we used the all-sky simulated of the IR and radio point sources provided by ? at 30, 90, 148, 219 and 350 GHz.","For our analysis, we used the all-sky simulated of the IR and radio point sources provided by \cite{Sehgal2010} at 30, 90, 148, 219 and 350 GHz."
+ We provide here a brief summary of the map construction., We provide here a brief summary of the map construction.
+" For a detailed description, we refer the reader to the above-cited article."," For a detailed description, we refer the reader to the above-cited article."
+" The maps are based on N-body simulations of the large scale structure, with a volume 1000 h! Mpc on a side, produced using a tree-particle mesh code."," The maps are based on N-body simulations of the large scale structure, with a volume 1000 $\mathrm{h}^{-1} $ Mpc on a side, produced using a tree-particle mesh code."
+ Dark matter (DM) haloes are identified and are then populated with infrared and radio galaxies., Dark matter (DM) haloes are identified and are then populated with infrared and radio galaxies.
+ The model for the radio sources is adapted so that the radio luminosity function matches that of the observed radio sources at 151 MHz., The model for the radio sources is adapted so that the radio luminosity function matches that of the observed radio sources at 151 MHz.
+ The IR-source model was partially based on ?.., The IR-source model was partially based on \cite{Righi2008}.
+ The DM haloes are populated with galaxies of given luminosities taking into account a Poisson term and a correlation term., The DM haloes are populated with galaxies of given luminosities taking into account a Poisson term and a correlation term.
+" The model was constructed so that it is compatible with constraints on the luminosity function, the source counts and the fluctuations from SCUBA, BLAST, Spitzer and ACBAR (see ? for details)."," The model was constructed so that it is compatible with constraints on the luminosity function, the source counts and the fluctuations from SCUBA, BLAST, Spitzer and ACBAR (see \cite{Sehgal2010} for details)."
+" However, the simulation of IR sources, used here, does not account for the most recent observational constraints from ACT, SPT Herschel and Planck results."," However, the simulation of IR sources, used here, does not account for the most recent observational constraints from ACT, SPT Herschel and Planck results."
+" Maps of the different astrophysical contributions, in HEALPix format at Nside=8192, were produced by replication of one simulated octant of the sky."," Maps of the different astrophysical contributions, in HEALPix format at Nside=8192, were produced by replication of one simulated octant of the sky."
+" This procedure does not properly account for the signal at the largest scales especially up to the octopole, (=3 ; it also introduces excess power at €<300 for infrared sources as discussed by ?,, but the power is correct for higher multipoles."," This procedure does not properly account for the signal at the largest scales especially up to the octopole, $\ell=3$ ; it also introduces excess power at $\ell \le 300$ for infrared sources as discussed by \cite{Sehgal2010}, but the power is correct for higher multipoles."
+" The produced maps have half-arcminute resolution but we degraded them to Nside=1024 and used a uniform binning Δέ=64, to speed up computations."," The produced maps have half-arcminute resolution but we degraded them to Nside=1024 and used a uniform binning $\Delta\ell=64$, to speed up computations."
+ We checked that this procedure does not introduce a bias by comparing the binned spectrum in the degraded map to the unbinned spectrum in the original map and find excellent agreement., We checked that this procedure does not introduce a bias by comparing the binned spectrum in the degraded map to the unbinned spectrum in the original map and find excellent agreement.
+ The octant replication in the map-making translates mainly into a lack of power in the first bin (centered around £= 32) which is hence discarded in later plots., The octant replication in the map-making translates mainly into a lack of power in the first bin (centered around $\ell=32$ ) which is hence discarded in later plots.
+ We first investigate the bispectrum dependence on the configurations at a given frequency., We first investigate the bispectrum dependence on the configurations at a given frequency.
+" We plot the angular bispectrum in four commonly used configurations, namely equilateral (6,6,£), isosceles orthogonal (6,/2£), isosceles flat (6,6.26), and squeezed (€min,£) configurations."," We plot the angular bispectrum in four commonly used configurations, namely equilateral $(\ell,\ell,\ell)$, isosceles orthogonal $(\ell,\ell,\sqrt{2}\ell)$, isosceles flat $(\ell,\ell,2\ell)$, and squeezed $(\ell_\mathrm{min},\ell,\ell)$ configurations."
+"£, They furthermore sample rather well the £,configuration space (see Fig. 1)).", They furthermore sample rather well the configuration space (see Fig. \ref{triparam}) ).
+" We see in Fig. 3,,"," We see in Fig. \ref{speconfrad090},"
+" black thick line, that the bispectrum is constant."," black thick line, that the bispectrum is constant."
+ This result is independent of the frequency., This result is independent of the frequency.
+ Moreover the value of the constant is independent of the configuration., Moreover the value of the constant is independent of the configuration.
+ This is what we expect from white noise and it shows that radio sources are indeed randomly distributed over the sky., This is what we expect from white noise and it shows that radio sources are indeed randomly distributed over the sky.
+ We show in Fig., We show in Fig.
+" 4 the dependence with frequency of the bispectrum amplitude averaged over all the configurations, in equivalent temperature unit for the upper panel and flux unit for the bottom panel."," \ref{fig:amplRAD} the dependence with frequency of the bispectrum amplitude averaged over all the configurations, in equivalent temperature unit for the upper panel and flux unit for the bottom panel."
+" The amplitude is maximal at the lowest frequency 30 GHz, it then decreases to become mostly flat above 90 GHz because of free-free emission and inverted spectra sources."," The amplitude is maximal at the lowest frequency 30 GHz, it then decreases to become mostly flat above 90 GHz because of free-free emission and inverted spectra sources."
+" The bispectrum was fitted with the one population prescription described in Sect. ??,,"," The bispectrum was fitted with the one population prescription described in Sect. \ref{sect:onepop},"
+ using the power spectrum extracted from the simulations and the multiplicative constant which minimizes the x? of the observed bispectrum to the prescribed one with error bars from Wick's expansion (see Appendix ??))., using the power spectrum extracted from the simulations and the multiplicative constant which minimizes the $\chi^2$ of the observed bispectrum to the prescribed one with error bars from Wick's expansion (see Appendix \ref{appendix:wngvar}) ).
+ The fit of the bispectrum with the prescription is very good: Fig., The fit of the bispectrum with the prescription is very good: Fig.
+ 5 shows that the relative error (exact to fit) lies in the range between and with a mean relative error always less than 196o., \ref{histerrorfitrad} shows that the relative error (exact to fit) lies in the range between and with a mean relative error always less than $\tcperthousand$.
+ Figure 6 shows the amplitude of the angular bispectrum in the equilateral configuration for the different frequencies., Figure \ref{equibisp-IR} shows the amplitude of the angular bispectrum in the equilateral configuration for the different frequencies.
+" Dusty galaxy emission peaks at IR frequencies, and plummets in the radio domain and so does the amplitude of the bispectrum."," Dusty galaxy emission peaks at IR frequencies, and plummets in the radio domain and so does the amplitude of the bispectrum."
+" The bispectrum decreases with 6, well fitted by a power-law for 100€@<1000, and with a flattening at higher multipoles."," The bispectrum decreases with $\ell$, well fitted by a power-law for $100\leq \ell \leq 1000$, and with a flattening at higher multipoles."
+ The decrease is expected from the clustering of IR galaxies on large scales., The decrease is expected from the clustering of IR galaxies on large scales.
+" We found that the slopes were about the same for equilateral,flat and orthogonal configurations except for the squeezed triangles."," We found that the slopes were about the same for equilateral,flat and orthogonal configurations except for the squeezed triangles."
+" Indeed, for the latter £, is fixed, while for the other configurations £1,£2,/3 are all proportional to 6."," Indeed, for the latter $\ell_1$ is fixed, while for the other configurations $\ell_1,\ell_2,\ell_3$ are all proportional to $\ell$ ."
+In particular. errors in cillerent wavebancl datasets. are probably highly. correlated. so that the evaluation of e.g. colour. phase dillerences on the basis of the formal errors may be misleading.,"In particular, errors in different waveband datasets are probably highly correlated, so that the evaluation of e.g. colour phase differences on the basis of the formal errors may be misleading."
+ We obtained I4 spectra of KD 042215421. starting at LID 2.450.756.8586 (1997 November 4.3583 UT) with the 2.7m telescope at the MeDonald Observatory using the Large Cass Spectrograph. a 600£/mm grating (hlazecl at 4200A)). and the PL S00.500 CCD.," We obtained 14 spectra of KPD 0422+5421 starting at HJD 2,450,756.8586 (1997 November 4.3583 UT) with the 2.7m telescope at the McDonald Observatory using the Large Cass Spectrograph, a $600\ell$ /mm grating (blazed at ), and the TI1 $800\times 800$ CCD."
+ Phe resulting spectral resolution is zz3.5 (ENLIIND with wavelength coverage of 3525-4940A., The resulting spectral resolution is $\approx 3.5$ (FWHM) with wavelength coverage of 3525-4940.
+. The exposure times ranged [rom 10 to 20 minutes. and the signal-to-noise ratios in the final reduced. spectra were z in the continuum near 11.," The exposure times ranged from 10 to 20 minutes, and the signal-to-noise ratios in the final reduced spectra were $\approx 20-40$ in the continuum near $\beta$."
+ “Phere were cirrus clouds present. so the spectra could not be placed on an absolute [ux scale.," There were cirrus clouds present, so the spectra could not be placed on an absolute flux scale."
+ Instead. we used nine observations of five different spectrophotometrie Dux standards from the nights ofNovember 1-3. 1997 (which were photometric) to calibrate out the instrumental response.," Instead, we used nine observations of five different spectrophotometric flux standards from the nights of November 1-3, 1997 (which were photometric) to calibrate out the instrumental response."
+ Since the slit could. not be rotated to follow the parallactic angle. most of our program objects were observed at hour angles near zero.," Since the slit could not be rotated to follow the parallactic angle, most of our program objects were observed at hour angles near zero."
+ The relative flux scale is reasonably accurate between LE3 and the Balmer jumphowever the relative calibration at either end of the covered spectral range is suspect in part because some of the IXPD 042215421 spectra were obtained at. relatively large aipmiass (1.5-2.0)., The relative flux scale is reasonably accurate between $\beta$ and the Balmer jump—however the relative calibration at either end of the covered spectral range is suspect in part because some of the KPD 0422+5421 spectra were obtained at relatively large airmass (1.5-2.0).
+ We used the cross-correlation technique of Tonry. Davis (1979) to determine the radial velocities of the KPD 0422|5421 spectra., We used the cross-correlation technique of Tonry Davis (1979) to determine the radial velocities of the KPD 0422+5421 spectra.
+ The cross-correlation functions (CCTs) were computed over the wavelength interval 3800-4920., The cross-correlation functions (CCFs) were computed over the wavelength interval 3800-4920.
+A.. We used a synthetic spectrum generated. [rom a model atmosphere (sce below) as the template spectrum (although our results were almost. identical when the first observation was used as the template)., We used a synthetic spectrum generated from a model atmosphere (see below) as the template spectrum (although our results were almost identical when the first observation was used as the template).
+ In. cach case the CCEs were strong and well-defined., In each case the CCFs were strong and well-defined.
+ The velocities corresponding to the centroid of the CCE peaks were determined by a parabolic it to the six pixels surrounding the maximum., The velocities corresponding to the centroid of the CCF peaks were determined by a parabolic fit to the six pixels surrounding the maximum.
+ Large racial velocity variations were evident., Large radial velocity variations were evident.
+ We fitted a [our-parameter sinusoid to the 124 racial velocities., We fitted a four-parameter sinusoid to the 14 radial velocities.
+ The est-litting sinusoid had 42=6.5. an indication that the errors on the radial velocities were probably too small.," The best-fitting sinusoid had $\chi^2_{\nu}=6.5$, an indication that the errors on the radial velocities were probably too small."
+ We scaled the errors on the individual velocities to give yo=1 or the fit., We scaled the errors on the individual velocities to give $\chi^2_{\nu}=1$ for the fit.
+ The resulting spectroscopic elements are. given in Table 4., The resulting spectroscopic elements are given in Table 4.
+ Note that the spectroscopic period of P4=09070.0020 clays is roughly twice the photometric period derived. above.," Note that the spectroscopic period of $P_{\rm spect}=0.0907 \pm
+0.0020$ days is roughly twice the photometric period derived above."
+ The spectroscopic period and the velocity semi-amplitude. of ⋅⊳AG=⋅⋅−237x18 km nimply a mass unction of the (unseen) companion star of where Aca is the mass of the unseen companion star. Alam ds the mass of the sdD. star. and 7 is the orbital inclination.," The spectroscopic period and the velocity semi-amplitude of $K_{\rm
+sdB}=237\pm 18$ km $^{-1}$ imply a mass function of the (unseen) companion star of where $M_{\rm comp}$ is the mass of the unseen companion star, $M_{\rm sdB}$ is the mass of the sdB star, and $i$ is the orbital inclination."
+ The above equation implies. Mesa-0.126+0.020AL. .," The above equation implies $M_{\rm comp} > 0.126\pm 0.029\,M_{\odot}$ ."
+" The [act that the photometric period. is half the spectroscopic period suggests that the optical mocdulations we observe are. due. to. the well-known ""ellipsoidal"" variations.", The fact that the photometric period is half the spectroscopic period suggests that the optical modulations we observe are due to the well-known “ellipsoidal” variations.
+ Phe cause of the ellipsoidal modulations is easy to understand., The cause of the ellipsoidal modulations is easy to understand.
+ The sd star we observe is slightly distorted by the unseen companion star., The sdB star we observe is slightly distorted by the unseen companion star.
+ As the sdD star moves around in, As the sdB star moves around in
+We have made a detailed analysis of X-ray point sources detected with a deep ACIS-S observation of the core of the Perseus cluster.,We have made a detailed analysis of X-ray point sources detected with a deep ACIS-S observation of the core of the Perseus cluster.
+ The main observational results and conclusions of our study are:, The main observational results and conclusions of our study are:
+In case of the axial syimunetry of the radiation field. it clearly requires that the plane of polarization be aloug the meridian plane (or. at right angle to it).,"In case of the axial symmetry of the radiation field, it clearly requires that the plane of polarization be along the meridian plane (or, at right angle to it)."
+" Cousequeutly. CEEUU=0 aud the two iuteusities { and J, are sufficient to characterize the radiation field."," Consequently, $U=V=0$ and the two intensities $I_{l}$ and $I_{r}$ are sufficient to characterize the radiation field."
+ With this consideration the phase matrix used to describe the angular distribution of pliotons that uudereo Rayleigh scatteriug cau be written as Iu the present work. plane parallel atinosphiere bas been considered.," With this consideration the phase matrix used to describe the angular distribution of photons that undergo Rayleigh scattering can be written as In the present work, plane parallel atmosphere has been considered."
+ The transfer equation eoverning the inteusities Z; aud { in rest frame with the medium stratified into plane parallel cau be written as where fp=cosOg€[0.1]) and 0 is the angle between the Stokes specific intensity vector and the axis of sviunetry z: b(z) is the internal radiation source. wy is the albedo for single scattering aud 7 is the optical depth.," The transfer equation governing the intensities $I_{l}$ and $I_{r}$ in rest frame with the medium stratified into plane parallel can be written as where $\mu=\cos\theta (\mu\in[0,1])$ and $\theta$ is the angle between the Stokes specific intensity vector and the axis of symmetry $z$; ${\bf b(z)}$ is the internal radiation source, $\omega_0$ is the albedo for single scattering and $\tau$ is the optical depth."
+ In case of spherical geometry the equation of transfer for polarized radiation cau be written as where # is the mass absorption co-efficient for radiation of [requeney v anc p is the density of the material such that between the points s a ids’. For an object having radius as small as hat of Jupiter with geometrically thin atmosphere. the secoucd term in the left hand side of the above equation which incorporates the curvature effect becomes negligible.," In case of spherical geometry the equation of transfer for polarized radiation can be written as where $\kappa$ is the mass absorption co-efficient for radiation of frequency $\nu$ and $\rho$ is the density of the material such that between the points $s$ and $s'$, For an object having radius as small as that of Jupiter with geometrically thin atmosphere, the second term in the left hand side of the above equation which incorporates the curvature effect becomes negligible."
+ Therefore. for solar type ο stars and substellar-imass objects. the plane-parallel stratification is sullicient.," Therefore, for solar type of stars and substellar-mass objects, the plane-parallel stratification is sufficient."
+heat generation in these zones would evantually increase the temperature gradient. and convection would be restored.,"heat generation in these zones would evantually increase the temperature gradient, and convection would be restored."
+ The properties of the internal modes excited in the stable shells are undoubtedly affected by the computational domain and the boundary condition., The properties of the internal modes excited in the stable shells are undoubtedly affected by the computational domain and the boundary condition.
+ Reflective waves are probably the cause for the constant level of velocities we noticed at Figure 11., Reflective waves are probably the cause for the constant level of velocities we noticed at Figure 11.
+ However. the volume below the oxygen shell in these stars is relatively small so a finite value of the velocity is quite plausible.," However, the volume below the oxygen shell in these stars is relatively small so a finite value of the velocity is quite plausible."
+ An interesting point is the validity of the results as we performed 2D simulations (and not 3D)., An interesting point is the validity of the results as we performed 2D simulations (and not 3D).
+ This is not an easy question to answer. however. from various comparisons of 2D and 3D simulations (see for example Kiragaetal. 1999)) it seems that the main features of the interaction of the unstable layer and its neighboring layers are similar. and the overshoot and mixingin in both cases are comparable.," This is not an easy question to answer, however, from various comparisons of 2D and 3D simulations (see for example \cite{krg99}) ) it seems that the main features of the interaction of the unstable layer and its neighboring layers are similar, and the overshoot and mixingin in both cases are comparable."
+ As we do not claim to resolve this interaction quantitatively. we think that our results are valid.," As we do not claim to resolve this interaction quantitatively, we think that our results are valid."
+ Penetration of the convective flow to neighboring stable regions is a common feature that exists in. many multidimensional simulations of convection. in a variety of stellar problems.," Penetration of the convective flow to neighboring stable regions is a common feature that exists in many multidimensional simulations of convection, in a variety of stellar problems."
+ Hurlburt.Toomre.&Massaguer1986. have studied such penetration beyond a shallow convective region in 2D. and found that it generated internal g modes in the stable region below the convection region. all the way to the bottom boundary.," \cite{hrl86} have studied such penetration beyond a shallow convective region in 2D, and found that it generated internal g modes in the stable region below the convection region, all the way to the bottom boundary."
+ Two regions in this stable layer can be identified from their figure 13: close to the unstable layer there is a sharp decrease in the kinetic energy. but it does not go to zero. rather there is an almost constant level of kinetic energy throughout the rest of the stable layer.," Two regions in this stable layer can be identified from their figure 13: close to the unstable layer there is a sharp decrease in the kinetic energy, but it does not go to zero, rather there is an almost constant level of kinetic energy throughout the rest of the stable layer."
+ Kiragaetal.1999 have tested penetration in. both 2D and 3D. they found that the energy fluxes in the stable layer in the 3D are about two thirds of the 2D results and that the constant level is replaced with a continuas decrease when the viscosity coefficient is increased at the bottom to avoid reflection of waves.," \cite{krg99} have tested penetration in both 2D and 3D, they found that the energy fluxes in the stable layer in the 3D are about two thirds of the 2D results and that the constant level is replaced with a continuas decrease when the viscosity coefficient is increased at the bottom to avoid reflection of waves."
+ In both these studies. ideal simplified physics was assumed.," In both these studies, ideal simplified physics was assumed."
+ Freytag.Ludwig.&Steffen(1996) studied the structure and dynamties of shallow stellar surface convection zones. using two dimensional radiation hydrodynamic simulations with more “real” physics.," \cite{fls96} studied the structure and dynamics of shallow stellar surface convection zones, using two dimensional radiation hydrodynamic simulations with more “real” physics."
+" They found convective motions extending ""well beyond the boundary of convectively unstable region. with vertical velocities decaying exponentially with depth in the deeper parts of the lower overshoot region. as expected for linear g modes.”"," They found convective motions extending “well beyond the boundary of convectively unstable region, with vertical velocities decaying exponentially with depth in the deeper parts of the lower overshoot region, as expected for linear $g^-$ modes.”"
+ They suggested to approximate the average velocity at the stable region as the convective velocity at the boundary of the unstable zone multiplied by an exponential decay factor with a lapse rate that is of order of the pressure scale height., They suggested to approximate the average velocity at the stable region as the convective velocity at the boundary of the unstable zone multiplied by an exponential decay factor with a lapse rate that is of order of the pressure scale height.
+ As explamed by Cox(1980.$17.8) g modes are unstable within convection regions. and generally decrease exponentially with increasing distance from the boundaries of the oscillatory region.," As explained by \cite{cox80} $^-$ modes are unstable within convection regions, and generally decrease exponentially with increasing distance from the boundaries of the oscillatory region."
+ As Freytag.Ludwig.&Steffen(1996) point out. Landau&Lifshitz(1959) have a relevant discussion in their $334. where they discuss some properties of potential flow.," As \cite{fls96}
+ point out, \cite{ll59} have a relevant discussion in their 34, where they discuss some properties of potential flow."
+ If compressibility and dissipation can be neglected. a steady state potential flow which is periodic in some plane. must be damped exponentially in a direction perpendicular to that plane.," If compressibility and dissipation can be neglected, a steady state potential flow which is periodic in some plane, must be damped exponentially in a direction perpendicular to that plane."
+ Also. the shortest wavelengths will be damped fastest.," Also, the shortest wavelengths will be damped fastest."
+ Consequently. most of the overshoot is given by the largest scales which are driven at the interface. that is. the largest convective scale.," Consequently, most of the overshoot is given by the largest scales which are driven at the interface, that is, the largest convective scale."
+ We find a striking underlying unity in these results and ours. despite the significantly different astrophysical context involved: our simulations correspond to deep shell burning convection that ts driven by nuclear burning at the bottom and cooled by neutrino radiation and in which spherical properties of the domain are important: while these studies deal with envelope convection driven by photon radiation at the surface in plane parallel geometry.," We find a striking underlying unity in these results and ours, despite the significantly different astrophysical context involved: our simulations correspond to deep shell burning convection that is driven by nuclear burning at the bottom and cooled by neutrino radiation and in which spherical properties of the domain are important; while these studies deal with envelope convection driven by photon radiation at the surface in plane parallel geometry."
+ In both cases The flow extand beyond the formal boundaries of the convection zone. with velocities that decay exponentially in the stable shells producing internal modes.," In both cases The flow extand beyond the formal boundaries of the convection zone, with velocities that decay exponentially in the stable shells producing internal modes."
+ Near the boundaries of the convection zone. where the velocities are still large. quite efficient mixing is taking place.," Near the boundaries of the convection zone, where the velocities are still large, quite efficient mixing is taking place."
+ The amount of mixing beyond that region. is not entirely clear.," The amount of mixing beyond that region, is not entirely clear."
+ Freytag.Ludwig.&Steffen(1996) have used trajectories of test particles to estimate the diffusion coefficient and conclude that the diffusion coefficient corresponds directly to the exponentially decaying velocities., \cite{fls96} have used trajectories of test particles to estimate the diffusion coefficient and conclude that the diffusion coefficient corresponds directly to the exponentially decaying velocities.
+ Some studies of stellar evolution that used these mixing terms for hydrogen core convection (Herwigetal.1997:: Herwig 1998). found that the lapse rate of decay of the diffusion should be much smaller (72 percents) than the pressure scale height.," Some studies of stellar evolution that used these mixing terms for hydrogen core convection \cite{hrw97}; Herwig 1998), found that the lapse rate of decay of the diffusion should be much smaller $\approx$ 2 percents) than the pressure scale height."
+ They speculate that this is due to the large stability of the layers that are near the core convection zone., They speculate that this is due to the large stability of the layers that are near the core convection zone.
+ Yet another explanation might be true. namely the flow characters at the stable shells and the many scales involved in mixing.," Yet another explanation might be true, namely the flow characters at the stable shells and the many scales involved in mixing."
+ Mixing is a complex phenomenon which take place in a very small length scale comparing to the length scale of the star. and since stars have a lower effective viscosity than in the simulations. and since they evolve for longer times. an accurate prediction of mixing is very difficult.," Mixing is a complex phenomenon which take place in a very small length scale comparing to the length scale of the star, and since stars have a lower effective viscosity than in the simulations, and since they evolve for longer times, an accurate prediction of mixing is very difficult."
+ Further. we note that motion does not necessarily imply mixing!," Further, we note that motion does not necessarily imply mixing!"
+" In particular. oscillatory potential flow. which is irrotational.would involve ""stretching"" and ""contracting"" motions that do not give mixing. at least at lowest order."," In particular, oscillatory potential flow, which is irrotational,would involve “stretching” and “contracting” motions that do not give mixing, at least at lowest order."
+ As we can see in Figures 11-14. The flow in the stable shells is different than the flow inside the convection region. even though it is purely potential flow (as revealed when examining the vorticity).," As we can see in Figures 11-14, The flow in the stable shells is different than the flow inside the convection region, even though it is purely potential flow (as revealed when examining the vorticity)."
+ Due to these differences in the flow. it seems plausible that the effective mixing velocity may be less than the hydrodynamic velocity. and its lapse rate different as well.," Due to these differences in the flow, it seems plausible that the effective mixing velocity may be less than the hydrodynamic velocity, and its lapse rate different as well."
+ As the various shells in our model have a different composition. mixing can be directly examined. from the 2D profiles and the 1D average of composition.," As the various shells in our model have a different composition, mixing can be directly examined from the 2D profiles and the 1D average of composition."
+ However. as the scale of mixing is much smaller than numerical resolution. the results must be carefully examined.," However, as the scale of mixing is much smaller than numerical resolution, the results must be carefully examined."
+ One way to examine the results is to compare the mixing in the standard simulation to that 1n the simulation with sub grid mixing model., One way to examine the results is to compare the mixing in the standard simulation to that in the simulation with sub grid mixing model.
+ There are small differences in the average abundances in the oxygen shell between the two simulations. and the main difference is that with SGSM terms. the profile is more uniform because of the extra mixing.," There are small differences in the average abundances in the oxygen shell between the two simulations, and the main difference is that with SGSM terms, the profile is more uniform because of the extra mixing."
+" The difference between the simulations is much more prominent in the lower stable shell: the width of the abundance jump is much wider when SGSM terms were used. and penetrates to deeper zones. in which the ""constant level"" average velocity exists."," The difference between the simulations is much more prominent in the lower stable shell: the width of the abundance jump is much wider when SGSM terms were used, and penetrates to deeper zones, in which the “constant level” average velocity exists."
+ When we used SGSM we implicitly assumed that turbulent flow exists from the scale of the mesh resolution to the scale of molecular diffusion., When we used SGSM we implicitly assumed that turbulent flow exists from the scale of the mesh resolution to the scale of molecular diffusion.
+ Since this mixing occurs in a stable region. in which internal modes have been excited. the validity of this assumption is questionable.," Since this mixing occurs in a stable region, in which internal modes have been excited, the validity of this assumption is questionable."
+ Thus it is fair to say that further work Is needed to resolve this subject of mixing beyond the convective layers., Thus it is fair to say that further work is needed to resolve this subject of mixing beyond the convective layers.
+ Mazzitelli.D'Antona.&Ventura(1999) have examined the, \cite{mdv99} have examined the
+The goal of the smoothing procedure specified in the previous sections is to obtain a (Gaussian ENIM) 27 resolution map of the predieted. cumulative SZ ellect generated. by all cosmic structures within SO. bh+Alpe.,"The goal of the smoothing procedure specified in the previous sections is to obtain a (Gaussian FWHM) $2^{\circ}$ resolution map of the predicted, cumulative SZ effect generated by all cosmic structures within 80 $\hmpc$."
+ Since the SZ distortions are fully specified by the density and temperature of the eas along the line of sight. one max think that reliable SZ maps can be obtained directIy from the constrained hvdrodynamical simulations.," Since the SZ distortions are fully specified by the density and temperature of the gas along the line of sight, one may think that reliable SZ maps can be obtained directly from the constrained hydrodynamical simulations."
+ This is certainly the case for most statistical analyses but [or eross-correlation studies like the one we wish to perform here., This is certainly the case for most statistical analyses but for cross-correlation studies like the one we wish to perform here.
+ Indeed. the constraints in the simulations are only effective above a (Gaussian) resolution scale of 5. bh1Mpe. corresponding to an angular scale >3.57 within our spherical volume.," Indeed, the constraints in the simulations are only effective above a (Gaussian) resolution scale of 5 $\hmpc$, corresponding to an angular scale $>3.5^{\circ}$ within our spherical volume."
+ Since the actual positions of cosmic structures within each resolution clement are essentially random. we cannot clirectly cross-correlate the SZ maps obtained from the hvdrodynamical simulation with the temperature I[uctuations of the CMD with the desired angular resolution of 2.," Since the actual positions of cosmic structures within each resolution element are essentially random, we cannot directly cross-correlate the SZ maps obtained from the hydrodynamical simulation with the temperature fluctuations of the CMB with the desired angular resolution of $2^{\circ}$."
+ Phe PSC: galaxy density maps described in section 2.2. can be used to cireumvent his problem., The $z$ galaxy density maps described in section \ref{sect:pscz} can be used to circumvent this problem.
+" “Phe mass overdensity field. ὃν, in the constrained. simulation has been obtained from the IAS 1.2Jv galaxy density field assuming that WAS galaxies trace the underlving mass distribution."," The mass overdensity field, $\delta_m$ in the constrained simulation has been obtained from the IRAS 1.2Jy galaxy density field assuming that IRAS galaxies trace the underlying mass distribution."
+" On the other hand (""Teodoroctal.2003) rave shown that the LIGAS 1.2Jv and PSC: galaxy density fieles are in very good agreement (apart from à monopole mismatch hat loes not allect our analysis) and thus that a tight relation exists between psc; and ὃν", On the other hand \cite{teodoro} have shown that the IRAS 1.2Jy and $z$ galaxy density fields are in very good agreement (apart from a monopole mismatch that does not affect our analysis) and thus that a tight relation exists between $\delta_\mathrm{PSCz}$ and $\delta_m$.
+" We can therefore specify the gas properties at the correct spatial locations by rank ordering the overdensity field 9psc and 9,, while keeping the original spatial association between ὃν and the gas density and temperature in the simulation.", We can therefore specify the gas properties at the correct spatial locations by rank ordering the overdensity field $\delta_{PSCz}$ and $\delta_m$ while keeping the original spatial association between $\delta_m$ and the gas density and temperature in the simulation.
+" As a result. we are able to specify the emperature 7. and the electron density 7, in the gas associated to the cosmic structures traced by. σος galaxies. hence allowing a cross correlation analysis with a resolution of 2."," As a result, we are able to specify the temperature $T$, and the electron density $n_e$ in the gas associated to the cosmic structures traced by $z$ galaxies, hence allowing a cross correlation analysis with a resolution of $2^{\circ}$."
+ The Sunvaev Zel'dovich. effect. (Sunvaey&Zeldovich1972). arises due to the inverse Compton scatter of CALB photons against the hot and dilfuse electron gas trapped in the potential well of the cosmic web and of each cluster of galaxies., The Sunyaev Zel'dovich effect \cite{sz0} arises due to the inverse Compton scatter of CMB photons against the hot and diffuse electron gas trapped in the potential well of the cosmic web and of each cluster of galaxies.
+ In the latter case these electrons are also responsible for their X-ray emission., In the latter case these electrons are also responsible for their X-ray emission.
+ The CMD temperature change across a single line of sight is given by: where a—hefKdosp., The CMB temperature change across a single line of sight is given by: where $x=h\nu/kT_{\rm CMB}$.
+" Ehe Comptonization parameter is where n, and 1 are the electron density and temperature respectively. ep is the Thomson cross section. and the integral is over a line of sight."," The Comptonization parameter is where $n_e$ and $T $ are the electron density and temperature respectively, $\sigma _T$ is the Thomson cross section, and the integral is over a line of sight."
+ Phe spectral form of this “thermal ellect is described by the function sGre)=«e:coth6e/2)—4]. which is negative (positive) at values of μι smaller (larger) than ο=3.83. corresponding to a critical frequency. vy=211610Iz.," The spectral form of this “thermal effect"" is described by the function $s(x) = [x\cdot \coth(x/2)-4]$, which is negative (positive) at values of $x$ smaller (larger) than $x_0=3.83$, corresponding to a critical frequency $\nu_0=217 \textrm{GHz}$."
+ At the frequencies considered here. the temperature variation is always negative with values sGr)=—1.55 in the Wo channel. αμ) LNOinthe V channel and sir)=191 in the OQ channel.," At the frequencies considered here, the temperature variation is always negative with values $s(x)=-1.55$ in the W channel, $s(x)=-1.80$ in the V channel and $s(x)=-1.91$ in the Q channel."
+ To obtain a map of the SZ ellect. we need to integrate equation (4)) along the line of site in all the shells described in section 2..," To obtain a map of the SZ effect, we need to integrate equation \ref{eq:y}) ) along the line of site in all the shells described in section \ref{sect:data}."
+ Each of the shells are pixelised in a 3D grid of 128.128 pixels., Each of the shells are pixelised in a 3D grid of $128\times128\times128$ pixels.
+ Before making the line of sight integral. we made a 3D linear interpolation to obtain the shells on a 1280.1280.)1280 cubic eric.," Before making the line of sight integral, we made a 3D linear interpolation to obtain the shells on a $1280\times1280\times1280$ cubic grid."
+ From this. we performed the following integration: where / is pixel number on the Lealpix grid. j is grid element in the interpolated 3D erid and the sum is performed over all elements j contained in the solid angle spanned by. Healpix. pixel 7.," From this, we performed the following integration: where $i$ is pixel number on the Healpix grid, $j$ is grid element in the interpolated 3D grid and the sum is performed over all elements $j$ contained in the solid angle spanned by Healpix pixel $i$."
+" The temperature and. electron density at. &ridpoint pare ὃνgiven by Z7 and n7‘ (obtained from the procedure described in section 3.1)). δ0"" and 3/; are the mean solid angles and distances for cach gridpoint j and finally AQ=ἐπΣΑΝ} is the solid angle of each Healpix pixel."," The temperature and electron density at gridpoint $j$ are given by $T^j$ and $n_e^j$ (obtained from the procedure described in section \ref{sect:maps}) ), $\delta\Omega_j$ and $\delta l_j$ are the mean solid angles and distances for each gridpoint $j$ and finally $\Delta\Omega=4\pi/(12N_{side}^2)$ is the solid angle of each Healpix pixel."
+ As all shells were smoothed with a 3D filter corresponding to Gaussian angular beam of ENIM 27. our final map has this resolution.," As all shells were smoothed with a 3D filter corresponding to Gaussian angular beam of FWHM $2^\circ$, our final map has this resolution."
+ 1n ligureὃν 1 owe show the resultingIn map of the Compton ν parameter., In figure \ref{fig:SZmap} we show the resulting map of the Compton y parameter.
+ We can see that the most dominant structure is the Virgo cluster on the upper right side., We can see that the most dominant structure is the Virgo cluster on the upper right side.
+ We have also plotted the power spectrum of this map at two cdilferent. frequencies in figure 2.., We have also plotted the power spectrum of this map at two different frequencies in figure \ref{fig:SZcl}.
+ The dashed line corresponds to the Q 16111) frequeney and the dotted line to the W (9461Iz) frequency., The dashed line corresponds to the Q $41\textrm{GHz}$ ) frequency and the dotted line to the W $94\textrm{GHz}$ ) frequency.
+ Because of the very limited. redshift range considered. the power spectrum starts to fall of at the smallest scales.," Because of the very limited redshift range considered, the power spectrum starts to fall of at the smallest scales."
+ We see that the predicted: effect: is considerably. lower than the €MD. power spectrum (solid line) ancl could. not be distinguished from the cosmic variance at these scales., We see that the predicted effect is considerably lower than the CMB power spectrum (solid line) and could not be distinguished from the cosmic variance at these scales.
+"that the space density develops an infinite cusp if the X;/X, value is large enough.",that the space density develops an infinite cusp if the $\Sigma_d / \Sigma_g$ value is large enough.
+ interpreted (his result as a hint that Che gas distribution can only sustain the weight of a mass of dust. similar to ils own mass. and if surpassed. the cust would precipitate.," \citet{you02} interpreted this result as a hint that the gas distribution can only sustain the weight of a mass of dust similar to its own mass, and if surpassed, the dust would precipitate."
+" By numerically integrating our equation (41)) )) for a given Ri=σουηδ, together with the hydrostatic equilibrium condition in equation (15)). we find that for py>>1 (see also eq. ["," By numerically integrating our equation \ref{muapprox_eq}) ) for a given $\Ri=const.$, together with the hydrostatic equilibrium condition in equation \ref{rhog_eq}) ), we find that for $\mu_0 \gg 1$ (see also eq. ["
+22] in Sekiva 1993).,22] in Sekiya 1998).
+ This means that. for Ri; of order unity. Ma/Mg>Coomar/Cs anplies an exponential growth of the midplane dust abundance.," This means that, for $\Ri_{crit}$ of order unity, $\Sigma_d /
+\Sigma_g > v_{0,max}/c_s$ implies an exponential growth of the midplane dust abundance."
+ The onset of a cusp in densitv would therefore occur (under the condition Ri= cons/.). lor a range of radii ~1—10AU (see eq. [14].," The onset of a cusp in density would therefore occur (under the condition $\Ri = const.$ ), for a range of radii $\sim 1-10 \au$ (see eq. \ref{v0maxR_eq}] ]),"
+" when X;/X, exceeds ~0.1. i. e.. for roughly an order of magnitude dust enhancement."," when $\Sigma_d / \Sigma_g$ exceeds $\sim
+0.1$, i. e., for roughly an order of magnitude dust enhancement."
+ However. this conclusion appears to be an artifact of the Ri=const.condition.," However, this conclusion appears to be an artifact of the $\Ri=const.$condition."
+" If Ri is permitted to vary with height. with hi becoming large near the midplane. then sry varies only linearly rather than exponentially with X,;/X,. and a cusp need not form."," If $\Ri$ is permitted to vary with height, with $\Ri$ becoming large near the midplane, then $\mu_0$ varies only linearly rather than exponentially with $\Sigma_d / \Sigma_g$ , and a cusp need not form."
+using Equatious 5-9.,using Equations 5-9.
+ Figures Sa-f show our ποσα P-L relations. while Figures 92-6 present the optical oues.," Figures 8a-f show our near-IR P-L relations, while Figures 9a-e present the optical ones."
+ Iu cach panel. the solid line represeuts the results of the fitting process described in 8112 while the dashed lines iudicate the uncertaüutv of the fit.," In each panel, the solid line represents the results of the fitting process described in 4.2 while the dashed lines indicate the uncertainty of the fit."
+ Fit results are displaved iu cach panel aud listed in Tables 6 and 7., Fit results are displayed in each panel and listed in Tables 6 and 7.
+ We also tabulate in Table 6 published distance moduli for these galaxies (mostlyfromTable|ofFreecanauctal.2001a).. determined frou substantially larger samples of variables aud using the Stetson(1998) zeropoiuts.," We also tabulate in Table 6 published distance moduli for these galaxies \citep[mostly from Table 4 of][]{fr01a}, determined from substantially larger samples of variables and using the \citet{st98} zeropoints."
+ The optical distance moduli determined from our sinaller salples should take precedence over the above values., The optical distance moduli determined from our smaller samples should take precedence over the above values.
+ The optiuunu combination of the optical and infrared results will appear in Freedmanetal.(2001b)., The optimum combination of the optical and infrared results will appear in \citet{fr01b}.
+. As mentioned iu the introduction. one of the motivations of this project was to further study the ietallicity dependence of the Cepheid Period-Euminositv relation.," As mentioned in the introduction, one of the motivations of this project was to further study the metallicity dependence of the Cepheid Period-Luminosity relation."
+ Our data can contribute to these studies on two wavs: first. a global test of the metallicity dependence can be performed by aualvziug the appareut clistauce moduli of all galaxies: second. a differcutial test of the effect can be performed by analyzing the distance moduli to two regions of the same galaxw. provided they differ siguificantlv in abundance.," Our data can contribute to these studies on two ways: first, a global test of the metallicity dependence can be performed by analyzing the apparent distance moduli of all galaxies; second, a differential test of the effect can be performed by analyzing the distance moduli to two regions of the same galaxy, provided they differ significantly in abundance."
+ The first approach was undertaken bv Wochanek(1997).. while the second one was followed by Freedman&Madore(1990) iu M31 and by Neunicuttetal.(1998) iun MIOI.," The first approach was undertaken by \citet{ko97}, while the second one was followed by \citet{fm90} in M31 and by \citet{ke98} in M101."
+ The differential test is a challenging one. because the Taner field of MIOLT deviates substantially from other Nev Project fields in terms of surface brightucss auc stellar density.," The differential test is a challenging one, because the Inner field of M101 deviates substantially from other Key Project fields in terms of surface brightness and stellar density."
+ Our uear-infared distance moduli to the inner auc outer fields in MIOL exhibit laree differences Age=0.1640.12 mag and Agr;—0.37£0.12 mag.," Our near-infrared distance moduli to the inner and outer fields in M101 exhibit large differences: $\Delta\mu_H = 0.46\pm0.12$ mag and $\Delta\mu_J =
+0.37\pm0.12$ mag."
+ If taken at face value. they imply a very large iietalliity depeudence. of the order of 0.6 1iag/dex.," If taken at face value, they imply a very large metallicity dependence, of the order of 0.6 mag/dex."
+ However. other observational effects could be contributing to the observed. differences iu distance moduli.," However, other observational effects could be contributing to the observed differences in distance moduli."
+" One of them is ""bleudiug. or the contanunation of Cepheid fluxes bv nearby stars. uot physically associated with the variables. that fall within the NICAIOS seeing disk aud. cannot be resolved."," One of them is “blending”, or the contamination of Cepheid fluxes by nearby stars, not physically associated with the variables, that fall within the NICMOS seeing disk and cannot be resolved."
+ This effect has been the subject of a recent iuvestigation in the optical bv Mochejskaetal.(2000)., This effect has been the subject of a recent investigation in the optical by \citet{moc00}.
+.. One way to characterize the effect of blending on our distance determination to MIOI-Duuer is to move nearby. wellbresolved fields to the distance of MIOI. re-observe Cepheids in these fields. aud compare the resulting magnitudes with the ones obtained from the original lmages.," One way to characterize the effect of blending on our distance determination to M101-Inner is to move nearby, well-resolved fields to the distance of M101, re-observe Cepheids in these fields, and compare the resulting magnitudes with the ones obtained from the original images."
+ Our program contains observations of two suitable ealaxies: M21 and M81., Our program contains observations of two suitable galaxies: M31 and M81.
+ The fields observed iu these ealaxies show simular stellar densities and mean nearest-neighbor distances between Cepheids when compared to our MIOLI-Iuner fields., The fields observed in these galaxies show similar stellar densities and mean nearest-neighbor distances between Cepheids when compared to our M101-Inner fields.
+ We used our two ALS1 fields as well as fourteen of our M31 Ποια». which were located in Fields I aud ΤΠ of Daade&Swope(1965).," We used our two M81 fields as well as fourteen of our M31 fields, which were located in Fields I and III of \citet{bs65}."
+. We started by collecting the positions aud magnitudes reported by ALLSTAR for all objects in our input fields., We started by collecting the positions and magnitudes reported by ALLSTAR for all objects in our input fields.
+ The separation between stars were reduced by the ratio of distances between the input galaxies and the MIOT-Outer field (D(ALLOL/AL31)=10.1 and D(MIOL/AMs1)= 2.2)., The separation between stars were reduced by the ratio of distances between the input galaxies and the M101-Outer field $D(M101/M31)=10.1$ and $D(M101/M81)=2.2$ ).
+ The input maguitudes were corrected for the exposure time of the ovieial frames and distance to the iuput ealaxies. and then iiodified to reflect the distance of N101 and the exposure time of the ΑΠΟ frames.," The input magnitudes were corrected for the exposure time of the original frames and distance to the input galaxies, and then modified to reflect the distance of M101 and the exposure time of the M101 frames."
+ The stars were added to blank fines using the ADDSTAR routine found in DAOPUOT. which takes into account the properties of he detector and the PSF. as well as photon statistics.," The stars were added to blank frames using the ADDSTAR routine found in DAOPHOT, which takes into account the properties of the detector and the PSF, as well as photon statistics."
+ The artificial ficlds are shown in Figure 10 and compare avorably with the actual MIO0I-Iuner nuages shown iu Figures Ga-t. Ouce the artificial fields had been generated. they were xXiotoinetered in exactly the same wav as our real data.," The artificial fields are shown in Figure 10 and compare favorably with the actual M101-Inner images shown in Figures 6a-f. Once the artificial fields had been generated, they were photometered in exactly the same way as our real data."
+ To ideutity the variables in our artificial frames. we used he iuput positious as the equivalent of the astrometric information available for the real data. and searched for he objects ucarest to those positions. subject to the same 1.5-pixel rejection criterion from 811.1.," To identify the variables in our artificial frames, we used the input positions as the equivalent of the astrometric information available for the real data, and searched for the objects nearest to those positions, subject to the same 1.5-pixel rejection criterion from 4.1."
+ All Cepheids that were recovered were located at distances smaller than the rejection linüt., All Cepheids that were recovered were located at distances smaller than the rejection limit.
+ Figure Ll shows the P-L relations obtained from the sinulation. compared with the original iuput data.," Figure 11 shows the P-L relations obtained from the simulation, compared with the original input data."
+ Seven long period Cepheids (2=20 d) exhibit changes in magnitude of order 0.1-0.2 mae. most of them beime in the expected direction (.c.. towards brighter magnitudes).," Seven long period Cepheids $P > 20$ d) exhibit changes in magnitude of order 0.1-0.2 mag, most of them being in the expected direction (i.e., towards brighter magnitudes)."
+ Iu the case of the eight Cepheids with short periods (P«20 d). one was not recovered. four exhibited large variatious (which would have resulted iu their rojectiou based ou color-color criteria}. aud three had very modest changes in iiagnuitudoe.," In the case of the eight Cepheids with short periods $P < 20$ d), one was not recovered, four exhibited large variations (which would have resulted in their rejection based on color-color criteria), and three had very modest changes in magnitude."
+ The resulting distance moduli are sinaller than the iuput ones by 0-0.2 mae. depending on the period cutoff applied o these μα] samples.," The resulting distance moduli are smaller than the input ones by 0-0.2 mag, depending on the period cutoff applied to these small samples."
+ Therns scatter of the relations do rot increase significantly. once the short-period outlicrs are removed from the fits.," The scatter of the relations do not increase significantly, once the short-period outliers are removed from the fits."
+ We therefore conclude that a substantial fraction of the difference im distance moduli )etween MIOL-Iuner aud MIOl-Outer could be due to deuding., We therefore conclude that a substantial fraction of the difference in distance moduli between M101-Inner and M101-Outer could be due to blending.
+ This prevents us from performing a differential determination of the metallicity effect: however. a global est can still be performed using the other galaxies we have observed.," This prevents us from performing a differential determination of the metallicity effect; however, a global test can still be performed using the other galaxies we have observed."
+ Another goal of this project is to test whether the mean WF color excess. δεDj. is an appropriate indicator of total extinction and whether it can be used to obtain true distance moduli.," Another goal of this project is to test whether the mean $\vi$ color excess, $\evi$, is an appropriate indicator of total extinction and whether it can be used to obtain true distance moduli."
+ One could claim that the rauge in waveleneth between these two bands is too small to allow a eood extrapolation to A.1=0., One could claim that the range in wavelength between these two bands is too small to allow a good extrapolation to $\lambda^{-1}=0$.
+" There is also no enarantee that a ""standard"" value of Ryyp is applicable to other galaxies.", There is also no guarantee that a “standard” value of $\rvi$ is applicable to other galaxies.
+ If ΕΕ is indeed a good indicator of reddening. aud if a standard reddening law (Cardelliet applies to other galaxies. then one would expect CE(WD aud εςΤΠ to be strongly correlated aud to follow the slope predicted. by the standard reddening law.," If $\evi$ is indeed a good indicator of reddening, and if a standard reddening law \citep{ccm89} applies to other galaxies, then one would expect $\evi$ and $\evh$ to be strongly correlated and to follow the slope predicted by the standard reddening law."
+ The mean Wf and WOJI color excesses of a Cepheid sample are related by, The mean $\vi$ and $\vh$ color excesses of a Cepheid sample are related by
+and following rather different observational procedures. the measurement of line indices in the 1D spectra and their conversion [rom instrumental (ο standard system follows the same well established recipes (hat are only briefly described in this Section.,"and following rather different observational procedures, the measurement of line indices in the 1D spectra and their conversion from instrumental to standard system follows the same well established recipes that are only briefly described in this Section."
+ All line indices discussed in this paper were measured with the code. which is part of the package (Graves&Schiavon2008).," All line indices discussed in this paper were measured with the code, which is part of the package \citep{gs08}."
+. Cluster spectra were first s;moothed from (heir original resolution to the resolution of the Lick index. svstem. given in Table 1 of for each of the indices.," Cluster spectra were first smoothed from their original resolution to the resolution of the Lick index, system, given in Table 1 of \cite{s07} for each of the indices."
+ The original resolution was 5A in the case of the M 31 GCs. and that given in equation (1) of Schiavonetal.(2005).," The original resolution was ${\rm\AA}$ in the case of the M 31 GCs, and that given in equation (1) of \cite{s05}."
+. Instrumental index measurements are based on the definitions provided bv Worthevetal.(1994) aud Wortheyv&Ottaviani (1997)., Instrumental index measurements are based on the definitions provided by \cite{wo94} and \cite{wo97}.
+. Conversion between instrumental ancl standard Lick indices was achieved in the usual fashion. through comparison with measurements taken in spectra of standard stars.," Conversion between instrumental and standard Lick indices was achieved in the usual fashion, through comparison with measurements taken in spectra of standard stars."
+ Observation of Lick standards were obtained using the same instrumental setup used Lor GC observations. aud the spectra reduced following the same procedures as adopted for GC spectra.," Observation of Lick standards were obtained using the same instrumental setup used for GC observations, and the spectra reduced following the same procedures as adopted for GC spectra."
+ Likewise. (he spectra of Lick standards were smoothed to the Lick resolution. and instrumental indices measured adopting the same definitions as above.," Likewise, the spectra of Lick standards were smoothed to the Lick resolution, and instrumental indices measured adopting the same definitions as above."
+ For the M 31 llectospec data. the six standards were observed through single fibers. though each star was observed through a different. fiber.," For the M 31 Hectospec data, the six standards were observed through single fibers, though each star was observed through a different fiber."
+ No significant star-to-star variation is seen in the residuals shown in Figures 3. and 4.. so we conclude that different liber-to-fiber throuehput variations contribute negligible uncertainties to our index measurements.," No significant star-to-star variation is seen in the residuals shown in Figures \ref{zm31a} and \ref{zm31b}, so we conclude that different fiber-to-fiber throughput variations contribute negligible uncertainties to our index measurements."
+ We note (hat the Tololo spectra are allected by a bad CCD column (that runs across the speclra al a wavelength placed within the passband of the Lick CN indices., We note that the Tololo spectra are affected by a bad CCD column that runs across the spectra at a wavelength placed within the passband of the Lick CN indices.
+ We measured the effect of that bad column on our index measurements and found them to be affected at the ~ 0.015 mag level in the worst cases., We measured the effect of that bad column on our index measurements and found them to be affected at the $\sim$ 0.015 mag level in the worst cases.
+ Moreover. the effect. varies in sign and intensity from GC to GC. and thus should not introduce any significant svstematice effect in our," Moreover, the effect varies in sign and intensity from GC to GC, and thus should not introduce any significant systematic effect in our"
+1991)).,).
+ Note that on Lh+ Mpe scales. the importance of cooling plivsies to the gas teniperatuxe is nof inuueciately apparent: in Appendix A. we show this same fieure at the scale of a sinele cell. which shows quite clearly the effects of includiug the radiative plivsies of the gas.," Note that on 1 $h^{-1}$ Mpc scales, the importance of cooling physics to the gas temperature is not immediately apparent; in Appendix A, we show this same figure at the scale of a single cell, which shows quite clearly the effects of including the radiative physics of the gas."
+ We now cousider the equivalent plot for the galaxy formation density field., We now consider the equivalent plot for the galaxy formation density field.
+ Figure 2. shows the conditional mean of the galaxy formation deusity feld 1|6. ax a function of 1|o and T: 1|δαà.T5.," Figure \ref{dTg} shows the conditional mean of the galaxy formation density field $1+\delta_\ast$ as a function of $1+\delta$ and $T$: $\avg{1+\delta_\ast|1+\delta, T}$."
+ At any fixed density. the star formation rate declines as temperature rises.," At any fixed density, the star formation rate declines as temperature rises."
+ It is clear that the densest reeious cannot cool aud collapse at late times because they are too hot., It is clear that the densest regions cannot cool and collapse at late times because they are too hot.
+ To illustrate what Figure 2 inplies aboabout the relationshipatiouship bbetween galaxyzidlaxv formation aukand massass density. FigureFie 33. shows the condition:conditional probability P(1|ο8) as the logaritlinic grey scale and the conditional meau (1|;]1|8) as the solid line.," To illustrate what Figure \ref{dTg}
+ implies about the relationship between galaxy formation and mass density, Figure \ref{dendenstar} shows the conditional probability $P(1+\delta_\ast|1+\delta)$ as the logarithmic grey scale and the conditional mean $\avg{1+\delta_\ast|1+\delta}$ as the solid line."
+ The dashed lines are lo limits around the mean., The dashed lines are $\sigma$ limits around the mean.
+ At hieh redshift. the hieh density regions are still cool enough to have chumps of eas in them cooling quickly: in this seuse. the temperature on these scales is not particularly inportaut. making the relatiouship between galaxy formation aud mass deusity close to deterministic.," At high redshift, the high density regions are still cool enough to have clumps of gas in them cooling quickly; in this sense, the temperature on these scales is not particularly important, making the relationship between galaxy formation and mass density close to deterministic."
+ At low redshift. many regious become too hot to form new galaxies: in particular. galaxy formation is completely quenched in the deusest regions. causing a turnover in (1|.]15.," At low redshift, many regions become too hot to form new galaxies; in particular, galaxy formation is completely quenched in the densest regions, causing a turnover in $\avg{1+\delta_\ast|1+\delta}$."
+" In addition. at low redshifts. the ealaxy formation rate depends stronely on temperature. adding scatter to the relationship between galaxy formation aud mass (Blantonefαἱ, 1999))."," In addition, at low redshifts, the galaxy formation rate depends strongly on temperature, adding scatter to the relationship between galaxy formation and mass \cite{blanton98ap}) )."
+ The most obvious consequence of this evolution isthat the bias of galaxy formation is a strong function of time., The most obvious consequence of this evolution isthat the bias of galaxy formation is a strong function of time.
+ Figure Lo shows 5. aud re {efined in Section 13) as a functiou of redshift for top hat smoothing radii of R=1 aud R—8 h! Mpce., Figure \ref{brstartime} shows $b_\ast$ and $r_\ast$ (defined in Section \ref{motiv}) ) as a function of redshift for top hat smoothing radii of $R=1$ and $R=8$ $h^{-1}$ Mpc.
+ The bias of Oogalaxy. formation P. is clearly a stroneC» fiction of redshift. especially at smiall scales.," The bias of galaxy formation $b_\ast$ is clearly a strong function of redshift, especially at small scales."
+ The decrease in the bias is due to the fact that the ealaxy formation has moved to lower σ peaks and out of the hottest (and densest) regions ofthe universe. as described iu Section 1..," The decrease in the bias is due to the fact that the galaxy formation has moved to lower $\sigma$ peaks and out of the hottest (and densest) regions of the universe, as described in Section \ref{motiv}."
+ We can compare these results to those of the peaks-bias formalisin for the bias of dark matter haloes., We can compare these results to those of the peaks-bias formalism for the bias of dark matter haloes.
+" If we assume that the distribution of galaxy formation is traced by AL>1072 AZ. halos (that is. halos with typical bright galaxy masses} which have just collapsed. then according to the peaks-bias formalis (Doroshkevich1970: kaiser198 E: Bardeeneta£ 19856: Boudetαἱ. 1991: Mo&White 1996)). the bias of mass in such halos with respect to the general mass deusitv field is: where à,~~1.7 is the liner overdensity of a spliere which collapses exactly at the redshift of observation. and a(Af.2) is the xis mass fluctuation on scales corresponding to mass M. again at the redslüft of the observations."," If we assume that the distribution of galaxy formation is traced by $M > 10^{12}$ $M_\odot$ halos (that is, halos with typical bright galaxy masses) which have just collapsed, then according to the peaks-bias formalism \cite{doroshkevich70ap}; \cite{kaiser84ap}; \cite{bardeen86ap}; \cite{bond91ap}; \cite{mo96ap}) ), the bias of mass in such halos with respect to the general mass density field is: where $\delta_c\sim 1.7$ is the linear overdensity of a sphere which collapses exactly at the redshift of observation, and $\sigma(M,z)$ is the rms mass fluctuation on scales corresponding to mass $M$, again at the redshift of the observations."
+ Note that this formula differs from the standard one for the bias of all halos which have collapsed before redshift +., Note that this formula differs from the standard one for the bias of all halos which have collapsed before redshift $z$.
+ The dotted line iu the upper paucl of Figure 1. shows the resulting bias., The dotted line in the upper panel of Figure \ref{brstartime} shows the resulting bias.
+ Thus. the prediction for b. of the pure halo bias model is not much differeut from that of our lydrodvuamic model on large scales. even though the former does not include all the plivsical effects of the latter.," Thus, the prediction for $b_\ast$ of the pure halo bias model is not much different from that of our hydrodynamic model on large scales, even though the former does not include all the physical effects of the latter."
+ That is. measting /5.(:) alone cannot casily distinguish the two models.," That is, measuring $b_\ast(z)$ alone cannot easily distinguish the two models."
+ More interesting is the evolution of the correlation coefficient. à as a function of redshift. shown iu the lower panel of Figure L.," More interesting is the evolution of the correlation coefficient $r_\ast$ as a function of redshift, shown in the lower panel of Figure \ref{brstartime}."
+ At all scales. the correlation coefficient falls precipitously for 2<<1. when the densest regions become too hot to form new galaxies.," At all scales, the correlation coefficient falls precipitously for $z<1$, when the densest regions become too hot to form new galaxies."
+ Ou the other haud. in the case of peaks-biasing. find a scaling between the halo aud mass cross- aud autocorrelatious that implies ο~1 ou large scales at all redshitts!..," On the other hand, in the case of peaks-biasing, \cite{mo96a} find a scaling between the halo and mass cross- and autocorrelations that implies $r_\ast\sim 1$ on large scales at all ."
+ This is an important difference between the predictions of the models of, This is an important difference between the predictions of the models of
+separation of only 2z2 [rom the Galactic center.,"separation of only $\approx
+2\degr$ from the Galactic center."
+ The lime at which the jets were produced is constrained to Z5=ALJD 52767.641.1 days. independent of the prior.," The time at which the jets were produced is constrained to $T_0 = $ MJD $52767.6 \pm 1.1$ days, independent of the prior."
+ This timing supports a connection between I1743's X-ray [lare and the production of its jets., This timing supports a connection between H1743's X-ray flare and the production of its jets.
+ The speed of the jets. Ty. has a relatively low maximum a posteriori estimate. Py1.4. but is poorly constrained al high values and tracks its prior.," The speed of the jets, $\Gamma_0$, has a relatively low maximum a posteriori estimate, $\Gamma_0 \sim 1.4$, but is poorly constrained at high values and tracks its prior."
+ For the kinematic energv of each jet. we obtain a large uncertainty of 20.5 dex centered around Ex10% ere.," For the kinematic energy of each jet, we obtain a large uncertainty of $\approx 0.5$ dex centered around $\tilde{E}
+\approx 10^{45}$ erg."
+ This implies that 111743's jets are only about a tenth as energetic as those produced in the 1998 outburst of NTE J1550564 or. alternativelv. for I11743 either (1) the density of the surrounding medium is much lower or (2) the jet opening anele is substantially smaller.," This implies that H1743's jets are only about a tenth as energetic as those produced in the 1998 outburst of XTE J1550–564 or, alternatively, for H1743 either (1) the density of the surrounding medium is much lower or (2) the jet opening angle is substantially smaller."
+ We now estimate the spin of 11743 by fitting its X-ray spectra., We now estimate the spin of H1743 by fitting its X-ray spectra.
+ For the three crucial inpul parameters. we use the values of D and / derived in the preceding section and the distribution of black hole masses discussed below.," For the three crucial input parameters, we use the values of $D$ and $i$ derived in the preceding section and the distribution of black hole masses discussed below."
+ All of our analvsis is performed. using NSPEC v12.7.0 CArnaud1996).., All of our analysis is performed using XSPEC v12.7.0 \citep{XSPEC}.
+ Following Steineretal.(2009) and making minor adjustmentV. our spectral model has the form TBABS(SIMPLCKERRBDB2)). where and are. respectively. (he low-enerev-absorplion and accretion-disk components.," Following \cite{Steiner_2009} and making minor adjustments, our spectral model has the form $\otimes$ ), where and are, respectively, the low-energy-absorption and accretion-disk components."
+ The component scatters a fraction of the thermal disk photons into a Compton power law., The component scatters a fraction of the thermal disk photons into a Compton power law.
+ For 111743. this simple convolution model describes only the broad continuum components ancl is unaffected by Che inclusion of weaker features due. e.g... lo warm absorbers or spectral reflection.," For H1743, this simple convolution model describes only the broad continuum components and is unaffected by the inclusion of weaker features due, e.g., to warm absorbers or spectral reflection."
+" The four free parameters of (he spectral are the (1) Traction of thermal photons [sc sealtered into the Compton power law: (2) power-law index D: (3) spin parameter a,: and mass accretion rate M.", The four free parameters of the spectral are the (1) fraction of thermal photons $\fsc$ scattered into the Compton power law; (2) power-law index $\Gamma$; (3) spin parameter $a_*$; and mass accretion rate $\dot M$.
+ Mass. inclination. and distance are varied in 5x10* Monte-Carlo saniples. and all 110 spectra are fitted for each setting.," Mass, inclination, and distance are varied in $5\times10^3$ Monte-Carlo samples, and all 170 spectra are fitted for each setting."
+ Uncertainty in (he absolute calibration ol the X-ray fIux is accounted for by randomly varving the overall flix normalization by for each triplet selling of AZ. 7. aud D (e.g.. Steineretal. 20111).," Uncertainty in the absolute calibration of the X-ray flux is accounted for by randomly varying the overall flux normalization by for each triplet setting of $M$ , $i$ , and $D$ (e.g., \citealt{Steiner_j1550spin_2011}) )."
+ We similarly mareinalize over uncertaintv in the viscosity parameter a bv randomly assigning either a=0.01 or a=0.1 (e.g.. Ningetal.2007:Pessal 2007)). which are (wo representative values available to our model.," We similarly marginalize over uncertainty in the viscosity parameter $\alpha$ by randomly assigning either $\alpha = 0.01$ or $\alpha = 0.1$ (e.g., \citealt{King_2007, Pessah_2007}) ), which are two representative values available to our model."
+" Both of these uncertainties have asmall effect on a, compared to", Both of these uncertainties have asmall effect on $a_*$ compared to
+right elfective temperature. it may actually be impossible to have p-mocle instability without e-mode instability.,"right effective temperature, it may actually be impossible to have p-mode instability without g-mode instability."
+ Indeed. the stars Balloon 090100001. and LES 070216043. lie on the boundary. between the observed EC14026 anc PGITIO variables and exhibit both types of pulsational behaviour (Schuhetal.2006:Oreiroct2005).," Indeed, the stars Balloon 090100001 and HS 0702+6043 lie on the boundary between the observed EC14026 and PG1716 variables and exhibit both types of pulsational behaviour \citep{Sch06,Ore05}."
+. Llowever. there is a problem.," However, there is a problem."
+ The high temperature limit of the unstable e-mocles occurs around. τω20000 Ix for /=1 modes. increasing to ~25000 Ix [or /=4.," The high temperature limit of the unstable g-modes occurs around $T_{\rm eff} \sim 20\,000$ K for $l=1$ modes, increasing to $\sim 25\,000$ K for $l=4$."
+ The observed high temperature limit for POGITI6 variables is 29000 Ix. and it is unlikely that /72.," The observed high temperature limit for PG1716 variables is $\sim
+29\,000$ K and it is unlikely that $l>2$."
+ We have investigated e-model stabilitv for ZALB models having a range of chemical composition., We have investigated g-model stability for ZAHB models having a range of chemical composition.
+ Figs., Figs.
+ 11 and 12. demonstrate the cllects of varving the overall metallicity Z or the overall hydrogen concentration .X whilst maintaining a high iron concentration (f£.= 10).," \ref{nrad_z0102}
+ and \ref{nrad_x9050} demonstrate the effects of varying the overall metallicity $Z$ or the overall hydrogen concentration $X$ whilst maintaining a high iron concentration $f=10$ )."
+ Increasing X (equivalent to reducing the helium fraction) has the cllect of a modest. increase in the ellective temperature of the e-mocde blue edge., Increasing $X$ (equivalent to reducing the helium fraction) has the effect of a modest increase in the effective temperature of the g-mode blue edge.
+ Similarly. reducing Z while conserving the iron abundance (Fig. 11))," Similarly, reducing $Z$ while conserving the iron abundance (Fig. \ref{nrad_z0102}) )"
+ shifts the blue edge οἱ &-modes blue-ward a little., shifts the blue edge of g-modes blue-ward a little.
+ VPhese ellects can be understood from equation (4)) However. in neither case is à substantial change in the envelope composition capable of shifting the e-mode bluc-edge to its observed location.," These effects can be understood from equation \ref{eq_tauth}) ) However, in neither case is a substantial change in the envelope composition capable of shifting the g-mode blue-edge to its observed location."
+ Figs., Figs.
+ 9 to 12. exclude results for /=4., \ref{nrad_f0110} to \ref{nrad_x9050} exclude results for $l=4$.
+ The /=+ blue edge is always bluer than the /=3 blue edge by ~0.01 dex., The $l=4$ blue edge is always bluer than the $l=3$ blue edge by $\sim 0.01$ dex.
+ In addition to the lower likelihood of observing /=4 modes. this is again insullicient to account [or the observed blue edge.," In addition to the lower likelihood of observing $l=4$ modes, this is again insufficient to account for the observed blue edge."
+ These results closely mirror those obtained: previously (Fontaineetal.2003)... where radiative levitation provides the necessary enhancement in the iron abundance in the driving region.," These results closely mirror those obtained previously \citep{Fon03}, where radiative levitation provides the necessary enhancement in the iron abundance in the driving region."
+ These authors found g-mocde instability with f=3.4 21 ellective temperatures up to 240000 Ix (logμι 1.38). or with /S up to logZi1.43. still too cool o account for the observed. blue edge.," These authors found g-mode instability with $l=3,4$ at effective temperatures up to 000 K $\log T_{\rm eff} \leq 4.38$ ), or with $l\leq8$ up to $\log T_{\rm eff} \leq 4.43$, still too cool to account for the observed blue edge."
+ The main dillerence is that our models show unstable g-mocdes at /=1 and 2 where those of Fontaine ct al., The main difference is that our models show unstable g-modes at $l=1$ and 2 where those of Fontaine et al.
+ did not., did not.
+ One reason may » that our model envelopes have a homogeneous chemical composition with an artificial iron. enhancement. whereas he earlier. models have a stratified: composition with iron enhanced in critical lavers by a well-understood mechanism.," One reason may be that our model envelopes have a homogeneous chemical composition with an artificial iron enhancement, whereas the earlier models have a stratified composition with iron enhanced in critical layers by a well-understood mechanism."
+ We have found that the elective temperature at the Xue edge of e-modes hardly changes when the degree. of iron enrichment increases as long as fc10., We have found that the effective temperature at the blue edge of g-modes hardly changes when the degree of iron enrichment increases as long as $f\ge 10$.
+ Εις indicates hat to have a blue edge of e-moces consistent with the observation. we have to change something in the structure of models.," This indicates that to have a blue edge of g-modes consistent with the observation, we have to change something in the structure of models."
+ One solution to the bluc-cdge question might be to consider the evolution of these stars., One solution to the blue-edge question might be to consider the evolution of these stars.
+ We have computed the evolution of iron-enriched: extendec-horizontal branch models: with N=0.75.4002.f10 and Al=0476M. through the completion of core helium-burning ancl bevond (Fig. 13).," We have computed the evolution of iron-enriched extended-horizontal branch models with $X=0.75, Z_0=0.02, f=10$ and $M_{\rm c}=0.476\Msolar$ through the completion of core helium-burning and beyond (Fig. \ref{nrad_evol}) )."
+ Note the “breathing pulses” (small loops in the UR diagram) towards the end of core helium-burning., Note the “breathing pulses” (small loops in the HR diagram) towards the end of core helium-burning.
+ Phese are encountered. in our models with artificially high. iron abundances., These are encountered in our models with artificially high iron abundances.
+ They are due to the increased core opacity and are not significant in models with a normal core composition., They are due to the increased core opacity and are not significant in models with a normal core composition.
+ Seniconvection is included. following the description. by Spruit(1992) We have carried oul a stability analysis for. selected models along cach evolutionary sequence.," Semiconvection is included following the description by \citet{Spr92}
+ We have carried out a stability analysis for selected models along each evolutionary sequence."
+ Fig., Fig.
+ 13. shows those mocels where either. p-nioces and/or. e-modes are, \ref{nrad_evol} shows those models where either p-modes and/or g-modes are
+"to the gas as [ar as the initial field strength satislies D.<10""(ngj/10*cm7)?G.",to the gas as far as the initial field strength satisfies $ B\lesssim 10^{-5}(n_{\rm H}/10^3~{\rm cm^{-3}})^{0.55}~{\rm G}$.
+ This condition holds when the magnetic field does not allect the dvnaimics of the gravitational contraction., This condition holds when the magnetic field does not affect the dynamics of the gravitational contraction.
+ It is also found that seed. magnetic field induced by radiation foree is strong enough to activate MB in the accretion disk surrounding Poplll stars., It is also found that seed magnetic field induced by radiation force is strong enough to activate MRI in the accretion disk surrounding PopIII stars.
+ The MBI induced turbulence might play an important role in transporting the angular momentum in Poplll accretion disk., The MRI induced turbulence might play an important role in transporting the angular momentum in PopIII accretion disk.
+ We are grateful to J. Tan. who made important comments and suggestions on this paper as a referee.," We are grateful to J. Tan, who made important comments and suggestions on this paper as a referee."
+ We (hank R. Nishi. Kk. Omukai. 5. Inutsuka. N. Shibazaki. M. Umeniura and A. Ferrara for stimulating discussion.," We thank R. Nishi, K. Omukai, S. Inutsuka, N. Shibazaki, M. Umemura and A. Ferrara for stimulating discussion."
+ We also thank F. Nakamura who provided the data im his caleulations., We also thank F. Nakamura who provided the data in his calculations.
+ The analvsis has been made with computational facilities at Rikkvo University., The analysis has been made with computational facilities at Rikkyo University.
+ We acknowledge Research Grant from Japan Society [or the Promotion of Science (IS 15740122). Rikkvo University Special Fund [ον Research (IM IS).," We acknowledge Research Grant from Japan Society for the Promotion of Science (HS 15740122), Rikkyo University Special Fund for Research (HM HS)."
+v Do,$\chi^2$.
+cThe process was repeated. (a72=a;n|ον ete.)," The process was repeated $a^2_i=a^1_i+\delta$, etc.)"
+ until 47 differed substantially from xz., until $\chi^2$ differed substantially from $\chi^2_{\rm min}$ .
+ Figure 6 shows the X7 iandi hex? qeurecs, Figure \ref{chicut1} shows the $\chi^2$ $i$ and the $\chi^2$ $q$ curves.
+ Achangeintheinclinalioniof zz0.5% is required. to force. 47 to change by 1., A change in the inclination $i$ of $\approx 0.5^{\circ}$ is required to force $\chi^2$ to change by 1.
+" We adopt σι= 0.507. although we note that this is a rough approximation since the curve is not parabolic ancl xi, is much larger than the number of data points fit."," We adopt $\sigma_i=0.50^{\circ}$ , although we note that this is a rough approximation since the curve is not parabolic and $\chi^2_{\rm
+min}$ is much larger than the number of data points fit."
+" Based. on the \7 οσο, weadopla,=0.15 for the sake of discussion. (and with the same caveats discussed for σε)."," Based on the $\chi^2$ $q$ curve, we adopt $\sigma_q=0.15$ for the sake of discussion (and with the same caveats discussed for $\sigma_i$ )."
+ The quoted errors on the other parameters given in Table 5 were derived assuming Caussian errors., The quoted errors on the other parameters given in Table 5 were derived assuming Gaussian errors.
+" Eggleton's (1983) formula was used to compute the elfective radii of the Roche lobes The W-D code computes the radii of the components using the values of οι, O». and the orbital separation e."," Eggleton's (1983) formula was used to compute the effective radii of the Roche lobes The W-D code computes the radii of the components using the values of $\Omega_1$, $\Omega_2$, and the orbital separation $a$."
+ The sdB star is well within its Roche lobe., The sdB star is well within its Roche lobe.
+ We tried model fits with the temperature of the sdD star set to 25.000 Ix and to 27.000 Ix. Phe best-fit values of7 and q were nearly identical to those values given in Table 5.," We tried model fits with the temperature of the sdB star set to 25,000 K and to 27,000 K. The best-fit values of $i$ and $q$ were nearly identical to those values given in Table 5."
+ In that same vein. we also tried mocdel fits where the limb darkening coellicients were set to 0.25 and 0.35 for both filters.," In that same vein, we also tried model fits where the limb darkening coefficients were set to 0.25 and 0.35 for both filters."
+ Again. the best-fit values of ; and q did not change significantly.," Again, the best-fit values of $i$ and $q$ did not change significantly."
+ Finally. we found that the best-fit values of / and q did not change significantly when the elfective wavelengths of the U and D filters were changed by 100A.," Finally, we found that the best-fit values of $i$ and $q$ did not change significantly when the effective wavelengths of the $U$ and $B$ filters were changed by 100."
+ We conclude that our results presented in Table 5 are insensitive to the exact values chosen for the sdB parameters., We conclude that our results presented in Table 5 are insensitive to the exact values chosen for the sdB parameters.
+ We emphasize. that we have only estimated. the statistical errors and not any svstematic errors., We emphasize that we have only estimated the statistical errors and not any systematic errors.
+ There are no doubt systematic errors present. given the small light curve amplitudes and the nature of the observations hhigh speed photometry).," There are no doubt systematic errors present, given the small light curve amplitudes and the nature of the observations high speed photometry)."
+ However. the white light curve from the 0.9m and the 2 light curve [rom the 2.1m are very similar in amplitude. and their relative phases are as expected.," However, the white light curve from the 0.9m and the $B$ light curve from the 2.1m are very similar in amplitude, and their relative phases are as expected."
+ There is also nothing unusual about the derived component masses., There is also nothing unusual about the derived component masses.
+" The mass of the κα star (0.72£0.26 A.) is consistent with the ""canonical"" extended horizontal branch mass of 0.5Al. (Caloi 1989: Dorman. Rood. O'Connell 1993: Salfer οἱ 11994)."," The mass of the sdB star $0.72 \pm 0.26\,M_{\odot}$ ) is consistent with the “canonical” extended horizontal branch mass of $0.5\,M_{\odot}$ (Caloi 1989; Dorman, Rood, O'Connell 1993; Saffer et 1994)."
+ Webbink (1990. see also Warner 1995) finds a mean white chwarf mass of 0.742E0.04AZ. in S4 cataclysmic variable binaries. and the mean mass of white chwarls in the Ποια is near 0.56AL. (Bergeron. Saller. Lichert 1992).," Webbink (1990, see also Warner 1995) finds a mean white dwarf mass of $0.74\pm 0.04\,M_{\odot}$ in 84 cataclysmic variable binaries, and the mean mass of white dwarfs in the field is near $0.56\,M_{\odot}$ (Bergeron, Saffer, Liebert 1992)."
+ Phe mass of the white chwarlin IXPD 0422|5421 (0.6240.18AL. ) is consistent with either of these.," The mass of the white dwarf in KPD 0422+5421 $0.62 \pm 0.18\,M_{\odot}$ ) is consistent with either of these."
+ These considerations sugeest that the svstematic errors are reasonably small ancl that ou derived parameters are trustworthy., These considerations suggest that the systematic errors are reasonably small and that ou derived parameters are trustworthy.
+ A close examination of the model light curve near the photometric phase 0.5 (when the sclB star isbehind the companion) shows that the white dwarf companion passes in [ront of the sdB star. causing a clip of z0.005 magnitudes which lasts zz0.05 in phase (6.5 minutes).," A close examination of the model light curve near the photometric phase 0.5 (when the sdB star is the companion) shows that the white dwarf companion passes in front of the sdB star, causing a dip of $\approx 0.005$ magnitudes which lasts $\approx 0.05$ in phase (6.5 minutes)."
+ There is also a very slight distortion at phase 1.0 where the nearly. invisible white cbwaurf is totally eclipsed by the sdB star., There is also a very slight distortion at phase 1.0 where the nearly invisible white dwarf is totally eclipsed by the sdB star.
+ Interestingly enough. the white-light curve from the 09m shows a possible indication of a small dip near the spectroscopic phase 0.25 tthe photometric phase 0.5).," Interestingly enough, the white-light curve from the 0.9m shows a possible indication of a small dip near the spectroscopic phase 0.25 the photometric phase 0.5)."
+ Caution suggests that the observation of a transit be considered only tentative., Caution suggests that the observation of a transit be considered only tentative.
+ Since the depth of the depression caused by the transit is only z0.005 magnitudes. its exact shape is quite sensitive to slight drifts in the transparency. ete.," Since the depth of the depression caused by the transit is only $\approx 0.005$ magnitudes, its exact shape is quite sensitive to slight drifts in the transparency, etc."
+ Extended observations with a CCD on a largetelescope (where one can use reasonably short exposures) should. be done to confirm the existence of the transit., Extended observations with a CCD on a largetelescope (where one can use reasonably short exposures) should be done to confirm the existence of the transit.
+ The use of a CCD enables one to do precise dilferential photometry over many orbital eveles., The use of a CCD enables one to do precise differential photometry over many orbital cycles.
+aas determined from the 35 anc 3N field. pair: although we could not measure the velocity dispersion from the spectrum in 3N. we were able to determine the positions of the absorption lines in the spectrum and therefore the velocity ollset. with respect to the 38 field. opposite the nucleus.,"as determined from the 3S and 3N field pair: although we could not measure the velocity dispersion from the spectrum in 3N, we were able to determine the positions of the absorption lines in the spectrum and therefore the velocity offset with respect to the 3S field, opposite the nucleus."
+ By averaging the two velocities from these field we find == 938 c 16 km/s. However. repeating this exercise for the field. pair at 3.5 {οι which are 4S and + owe find == S92 + 23 km/s. which seems mareinally deviant [rom our other two measurements. although this is just. outside the 1-0 interval.," By averaging the two velocities from these field we find = 938 $\pm$ 16 km/s. However, repeating this exercise for the field pair at 3.5 $R_e$, which are 4S and 4N, we find = 892 $\pm$ 23 km/s, which seems marginally deviant from our other two measurements, although this is just outside the $\sigma$ interval."
+ We therefore decided to subtract for each field. pair their own corresponding systemic velocity., We therefore decided to subtract for each field pair their own corresponding systemic velocity.
+ In Figure 4. we compare our results with the long-slit data of Statler Smoecker-Hane (1999)..., In Figure \ref{fig:kin_ngc3379} we compare our results with the long-slit data of Statler Smecker-Hane \nocite{1999AJ....117..839S}.
+ Phe radial extent of the absorption line kinematics has been increased. by a factor two compared to the lone-slit data. and our data show a smooth continuation of kinematics out to larger radii.," The radial extent of the absorption line kinematics has been increased by a factor two compared to the long-slit data, and our data show a smooth continuation of kinematics out to larger radii."
+" We see little rotation at large radii. and the dispersion. profile remains Lat out to 3.5. H,. although there is a hint that the profile is slightly dropping."," We see little rotation at large radii, and the dispersion profile remains flat out to 3.5 $R_e$, although there is a hint that the profile is slightly dropping."
+ We also plot the kinematic profiles of the planetary. nebulae. as presented in. C'occato et al. (2009)...," We also plot the kinematic profiles of the planetary nebulae, as presented in Coccato et al. \nocite{2009MNRAS.394.1249C}."
+ We folded these profiles to negative racii. to allow comparison with our data at large radii also in these ranges.," We folded these profiles to negative radii, to allow comparison with our data at large radii also in these ranges."
+ We find a good agreement between the several datasets. although. the rotational velocities found. by the planetary. nebulae seem to be a bit higher (~30 km/s) than our findings.," We find a good agreement between the several datasets, although the rotational velocities found by the planetary nebulae seem to be a bit higher $\sim$ 30 km/s) than our findings."
+ The spectra and results for our fields inNGC 821 can be found in Figure 5. and Table 4.., The spectra and results for our fields inNGC 821 can be found in Figure \ref{fig:ppxf_821} and Table \ref{tab:kin821}. .
+ Unfortunately. weather," Unfortunately, weather"
+he extreme properties of NLSIs is that these are active ealactic nuclei. (AGN) containing black holes of relatively niodest mass which. nevertheless. are accreting at a high rate (Pounds. Done Osborne 1995).,"the extreme properties of NLS1s is that these are active galactic nuclei (AGN) containing black holes of relatively modest mass which, nevertheless, are accreting at a high rate (Pounds, Done Osborne 1995)."
+ In this mocdel. the lower mass black hole gives rise to smaller Ixeplerian. velocities of he broad-line region. clouds while the high accretion rate vields comparable luminosity to that of larger mass black roles and leads to increased disc emission and an enhance soft excess (Ross. Fabian Mineshige 1992).," In this model, the lower mass black hole gives rise to smaller Keplerian velocities of the broad-line region clouds while the high accretion rate yields comparable luminosity to that of larger mass black holes and leads to increased disc emission and an enhanced soft excess (Ross, Fabian Mineshige 1992)."
+ In some NLS1s here is tentative evidence for reflection from a highly ionizec disc (e.g. Comastri 11998). which fits naturally with the high accretion rate model. as the disc surface is expected to be highly ionize: (Matt. Fabian loss 1903).," In some NLS1s there is tentative evidence for reflection from a highly ionized disc (e.g. Comastri 1998), which fits naturally with the high accretion rate model, as the disc surface is expected to be highly ionized (Matt, Fabian Ross 1993)."
+ Reeent studies suggest. tha NLSIs. in addition to their extreme soft. X-ray spectra. also exhibit steeper intrinsic hard X-ray continua than their broad-line counterparts (e.g. Brandt. Mathur Elvis 1997). an οσο which could. arise. for example. if the strong soft excess ellectivelv cools the accretion disk corona in which the uncerlving power-law is formed. (Pounds. Done Osborne 1995: Maraschi Hardt. 1997) Our objective in this paper is to investigate the X-ray Es»ect ral.properties of a sample of NLS1s.," Recent studies suggest that NLS1s, in addition to their extreme soft X-ray spectra, also exhibit steeper intrinsic hard X-ray continua than their broad-line counterparts (e.g. Brandt, Mathur Elvis 1997), an effect which could arise, for example, if the strong soft excess effectively cools the accretion disk corona in which the underlying power-law is formed (Pounds, Done Osborne 1995; Maraschi Haardt 1997) Our objective in this paper is to investigate the X-ray spectral properties of a sample of NLS1s."
+ Specifically we are interested in whether there are any further X-ray spectral characteristics of NLSIs. over and above those noted earlier. which might be described. as distinctive of this class. of object.," Specifically we are interested in whether there are any further X-ray spectral characteristics of NLS1s, over and above those noted earlier, which might be described as distinctive of this class of object."
+ In. practical terms we have carried out à systematic analvsis of 24 observations of 22 NLSIs taken from the ppublic data archive., In practical terms we have carried out a systematic analysis of 24 observations of 22 NLS1s taken from the public data archive.
+ “Phe remainder of this paper is organised as follows., The remainder of this paper is organised as follows.
+ In the next section we celine the source sample and detail the observations and data reduction techniques., In the next section we define the source sample and detail the observations and data reduction techniques.
+ In Section 8 we go on to describe the results of our X-ray spectral analvsis and in Section +) present a discussion. of these results., In Section 3 we go on to describe the results of our X-ray spectral analysis and in Section 4 present a discussion of these results.
+ Finally we summarise our conclusions in Section 5., Finally we summarise our conclusions in Section 5.
+ We have selected 22 objects. identified in the literature as Sevfert 1 galaxies or quasars with H5 ENIM <2000 kms. for which data were available at the time of the analysis (late 1998).," We have selected 22 objects, identified in the literature as Seyfert 1 galaxies or quasars with $\beta$ FWHM $\leq$ 2000 km/s, for which data were available at the time of the analysis (late 1998)."
+ ‘Lhe sample. given in Table 1. is not complete in any sense except that it consists of all suitable spectra of NLSIs in the ppublic archive.," The sample, given in Table 1, is not complete in any sense except that it consists of all suitable spectra of NLS1s in the public archive."
+ The redshift’ distribution. of the sample is shown in Fie., The redshift distribution of the sample is shown in Fig.
+ 1., 1.
+ The NLSI galaxy Mkn 957 was also observed by bbut these data have not been analysed due to the very short exposure time of the observation (3 ks)., The NLS1 galaxy Mkn 957 was also observed by but these data have not been analysed due to the very short exposure time of the observation $\sim$ 3 ks).
+ In addition data on Ixaz 163 have been excluded: Kaz 163 is in the same CUS field as Mkn 507 but its image falls at the extreme edge of the SIS detector. making the data unsuitable for detailed spectral analysis.," In addition data on Kaz 163 have been excluded; Kaz 163 is in the same GIS field as Mkn 507 but its image falls at the extreme edge of the SIS detector, making the data unsuitable for detailed spectral analysis."
+ ccarries four instruments which are operated simultaneously namely two solid-state imaging spectrometers (SIS-0. and SIS-1: Burke 1991) ancl two gas imaging spectrometers (GIS-2 and. GIS-3: IXohmura 1993)., carries four instruments which are operated simultaneously namely two solid-state imaging spectrometers (SIS-0 and SIS-1; Burke 1991) and two gas imaging spectrometers (GIS-2 and GIS-3; Kohmura 1993).
+ In the current analysis. stancard data screening criteria were applied to all the data.," In the current analysis, standard data screening criteria were applied to all the data."
+ Table 1 lists the resulting exposure times (specifically for SIS-0. although data from all four instruments are utilised here).," Table 1 lists the resulting exposure times (specifically for SIS-0, although data from all four instruments are utilised here)."
+ Since we are concerned only with the spectral properties of the sources. time-average spectra have been extracted from each. observation.," Since we are concerned only with the spectral properties of the sources, time-averaged spectra have been extracted from each observation."
+ Counts were accumulated. where possible. in circular apertures of 34 aremin radius centred on the source.," Counts were accumulated, where possible, in circular apertures of 3–4 arcmin radius centred on the source."
+ Phe backerounc was estimated using source free regions from the same observation and. in the case of the GIS observations. a a similar olf-axis angle to the source.," The background was estimated using source free regions from the same observation and, in the case of the GIS observations, at a similar off-axis angle to the source."
+ After backgrouim subtraction the pulse height spectra were binned to give at least 20 counts per spectral channel., After background subtraction the pulse height spectra were binned to give at least 20 counts per spectral channel.
+ In the subsequen analysis the SIS clata were restricted to the 0.610. keV. encreyv range and CIS data to the 0.810 keV band so as to reduce the possible impact of calibration uncertainties., In the subsequent analysis the SIS data were restricted to the 0.6–10 keV energy range and GIS data to the 0.8–10 keV band so as to reduce the possible impact of calibration uncertainties.
+ As a check of the possible impact of spectral variability on the time-averaged source spectra. we extracted the 2.10 keV and 0.52.0 keV light curves of cach source (SIS-0 and SIS-1 co-added) and binned these on the satellite orbital timescale (5760 s).," As a check of the possible impact of spectral variability on the time-averaged source spectra, we extracted the 2–10 keV and 0.5–2.0 keV light curves of each source (SIS-0 and SIS-1 co-added) and binned these on the satellite orbital timescale (5760 s)."
+ We then used a us test to check. the constancy of the spectral hardness (defined as ratio of the 10 keV count rate to the 0.52.0 keV. count rate) during the observation., We then used a $\chi^{2}$ test to check the constancy of the spectral hardness (defined as ratio of the 2--10 keV count rate to the 0.5–2.0 keV count rate) during the observation.
+ Five out of 22 sources gave unacceptable values of v at the level (NAB 10205|024. RN 0439.45. NCC 4051. Mkn 478. Ark 564) indicative of underlving spectral variability.," Five out of 22 sources gave unacceptable values of $\chi^{2}$ at the level (NAB 10205+024, RX J0439–45, NGC 4051, Mkn 478, Ark 564) indicative of underlying spectral variability."
+ However. with the exception of NGC 4051 there was little evidence that these variations corresponded to any clear spectral trends. or to a systematic ellect (c.g. a spectral hardness versus Hux correlation).," However, with the exception of NGC 4051 there was little evidence that these variations corresponded to any clear spectral trends or to a systematic effect (e.g. a spectral hardness versus flux correlation)."
+ We conclude that the implicit assumption of the analysis which follows. namely that the observation-averageck spectra. provide a," We conclude that the implicit assumption of the analysis which follows, namely that the observation-averaged spectra provide a"
+etter by au order of maguitucle.,better by an order of magnitude.
+ Sky spectra were obtained by poiuting ~20 of the available fibers at empty fields duriug the orograum star exposures., Sky spectra were obtained by pointing 20 of the available fibers at empty fields during the program star exposures.
+ After extraction aud calibration. the skies were averaged together to build a clean single background spectrum. which was subtracted from the object spectra.," After extraction and calibration, the skies were averaged together to build a clean single background spectrum, which was subtracted from the object spectra."
+ Finally. the orograim star spectra were flattened with a high order spline that was fitted through the coutinuuu.," Finally, the program star spectra were flattened with a high order spline that was fitted through the continuum."
+ o tellurie line corrections were made since there were no telluric lines in our spectra during the secoud epoch aud the oues appearing in the earlier data are outside our wiudow of interest., No telluric line corrections were made since there were no telluric lines in our spectra during the second epoch and the ones appearing in the earlier data are outside our window of interest.
+ Alter extracting and calibrating the spectra. we measured radial velocities in order to determine cluster membership.," After extracting and calibrating the spectra, we measured radial velocities in order to determine cluster membership."
+ While the etalou lamps gave very high precision waveleugth calibratious. they are uncertain in absolute wavelength.," While the etalon lamps gave very high precision wavelength calibrations, they are uncertain in absolute wavelength."
+ Raw radial velocities were determined [or each star usiug strong lines in our highest S/N ratio spectra., Raw radial velocities were determined for each star using strong lines in our highest S/N ratio spectra.
+ We calibrated the velocity scale for each cluster observation by setting the mode of the velocity distribution equal to the published racial velocity. using 45.2+1.5 (lo) for M&O aud —27.9+0.8 (1o) for NGC 6752 (Harris1996).," We calibrated the velocity scale for each cluster observation by setting the mode of the velocity distribution equal to the published radial velocity, using $+8.2~{\pm}~1.5$ ) for M80 and $-27.9~{\pm}~0.8$ ) for NGC 6752 \citep{Harris96}."
+ We accepted for further study all stars that that had. radial velocities within 22 aud 15 ol the cluster meats for MSO aud NGC 6752. respectively. aud which have sufficient S/N ratio for reliable abundance deteriniuations (typically above 100 for epoch 1 aud 60 lor epoch 2).," We accepted for further study all stars that that had radial velocities within 22 and 15 of the cluster means for M80 and NGC 6752, respectively, and which have sufficient S/N ratio for reliable abundance determinations (typically above 100 for epoch 1 and 60 for epoch 2)."
+ We determined a set of lines to measure by examining the spectra with the highest S/N [rom each set of cluster giauts., We determined a set of lines to measure by examining the spectra with the highest S/N from each set of cluster giants.
+ The lines were identified aud compared with the solar spectrum ol Wallace. Hinkle. Livinestou (1993). and. with Table IL of Thévenin(1990).. from which the oscillator streugths were chosen. with the following exceptious: the waveleugtlis aud excitation potentials of the Fe I lines come from Naveetal.(LO91).. aud we use the atomic data for the Eu II liue at AGGLSA that is consistentM with Lyperline splitting. as described in sec 5.3..," The lines were identified and compared with the solar spectrum of Wallace, Hinkle, Livingston (1993), and with Table II of \citet{Thev90}, from which the oscillator strengths were chosen, with the following exceptions: the wavelengths and excitation potentials of the Fe I lines come from \citet{Nave94}, and we use the atomic data for the Eu II line at 6645 that is consistent with hyperfine splitting, as described in sec \ref{sec:results}."
+ This produced au Eu oscillator streneth that is 0.22 dex strouger than the value listed by Shetrone(1996a) auc is 40.51 dex stronger thau the Thévenin(1990) value., This produced an Eu oscillator strength that is +0.22 dex stronger than the value listed by \citet{S96a} and is $+0.51$ dex stronger than the \cite{Thev90} value.
+ For the other lines. we used only those that Thévenin(1990). cites as having au uncertainty in of less than 0.05 dex. and stayed away [rom stroug bleuds where the separate cores of each line could uot be incivicually detected iu our spectra.," For the other lines, we used only those that \citet{Thev90} cites as having an uncertainty in of less than 0.05 dex, and stayed away from strong blends where the separate cores of each line could not be individually detected in our spectra."
+ We measured equivalent widths using Gaussian fits with the IRAF routine SPLOT aud mace use of the SPLOT deblencing algorithun for lines that overlapped. but appeared. visually distinct.," We measured equivalent widths using Gaussian fits with the IRAF routine SPLOT and made use of the SPLOT deblending algorithm for lines that overlapped, but appeared visually distinct."
+ The line list aud equivalent widths are giveu in Table L (or NGC 6752. aud Table 5. [or ALSO.," The line list and equivalent widths are given in Table \ref{tab:t04}
+ for NGC 6752, and Table \ref{tab:t05} for M80."
+Graphical Astronomy and Image Analysis Tool (GAIA).,Graphical Astronomy and Image Analysis Tool (GAIA).
+" Flux from the galaxy could be recovered out to a radius of 8"", reaching a surface brightness of 23.5 mag arcsec~?, after which stellar contamination became significant."," Flux from the galaxy could be recovered out to a radius of $8^{\prime\prime}$, reaching a surface brightness of 23.5 mag $^{-2}$, after which stellar contamination became significant."
+ The recovered flux within this radius was mg = 16.9--0.1 magnitudes., The recovered flux within this radius was $_B$ = $16.9 \pm 0.1$ magnitudes.
+" Optical follow-up observations with the SARA telescope are ongoing. B, V,"," Optical follow-up observations with the SARA telescope are ongoing. $B$, $V$,"
+" and R band and H-alpha observations are planned over the coming months. B,"," and $R$ band and H-alpha observations are planned over the coming months. $B$,"
+" V, and R band observations are being taken to a lo surface brightness of 26.5 mag arcsec~?."," $V$ , and $R$ band observations are being taken to a $\sigma$ surface brightness of 26.5 mag $^{-2}$."
+ This should be sufficient to see low surface brightness features that correspond to the faint outer parts of a normal spiral and will allow us to measure the galaxy’s diameter at the standard 25 mag arcsec” level., This should be sufficient to see low surface brightness features that correspond to the faint outer parts of a normal spiral and will allow us to measure the galaxy's diameter at the standard 25 mag $^{-2}$ level.
+ ALFA ZOA J1952+1428 has a heliocentric velocity of Όλοι = +279 kms-!., ALFA ZOA J1952+1428 has a heliocentric velocity of $v_{hel}$ = +279 $\mathrm{km \; s^{-1}}$.
+ Solving for its Local Group centered velocity using derivations of the solar motion with respect to the Local Group byCourteau & van den Bergh (1999) gives upg = 491 kms!., Solving for its Local Group centered velocity using derivations of the solar motion with respect to the Local Group byCourteau $\&$ van den Bergh (1999) gives $v_{LG}$ = 491 $\mathrm{km \; s^{-1}}$.
+ Using Hubble’s Law with Πο = 70 kms-!Mpc! puts this source at a distance of 7 Mpc., Using Hubble's Law with $H_{0}$ = 70 $\; \mathrm{km \; s^{-1} \; Mpc^{-1}}$ puts this source at a distance of 7 Mpc.
+" However, Hubble’s Law is not a reliable distance indicator here because the dispersion of peculiar velocities in the local universe {υμει«10,000 kms-1) is c=298+34 km s! (Masters 2008)."," However, Hubble's Law is not a reliable distance indicator here because the dispersion of peculiar velocities in the local universe $v_{hel} < 10,000$ $\mathrm{km \; s^{-1}}$ ) is $\sigma = 298 \pm 34$ km $^{-1}$ (Masters 2008)."
+ The galaxy is probably not closer than 3 Mpc as the linear size at this distance would be smaller than most compact galaxies containing (Huchtmeier et al., The galaxy is probably not closer than 3 Mpc as the linear size at this distance would be smaller than most compact galaxies containing (Huchtmeier et al.
+ 2007)., 2007).
+" For the following analysis, we take the distance to be 7 Mpc, although future observations may well revise this number."," For the following analysis, we take the distance to be 7 Mpc, although future observations may well revise this number."
+" As can be seen in the EVLA and SARA images, the peak is slightly offset (A0=8.7"") from the optical emission, indicating either a false counterpart or a disturbed distribution."," As can be seen in the EVLA and SARA images, the peak is slightly offset $\Delta\theta = 8.7^{\prime\prime}$ ) from the optical emission, indicating either a false counterpart or a disturbed distribution."
+" The offset is ~300 pc at 7 Mpc, which is not uncommon even for isolated galaxies (c.f."," The offset is $\sim$ 300 pc at 7 Mpc, which is not uncommon even for isolated galaxies (c.f."
+" 400 pc offset in VV 124, Bellazzini et al."," $\sim$ 400 pc offset in VV 124, Bellazzini et al."
+ 2011)., 2011).
+ This could conceivably be a pair of low-surface brightness dwarf galaxies (c.f., This could conceivably be a pair of low-surface brightness dwarf galaxies (c.f.
+" HIZSS 3 with separation of ~900 pc, Begum et al."," HIZSS 3 with separation of $\sim$ 900 pc, Begum et al."
+" 2005), but there is no evidence for a second peak in the high signal-to-noise spectrum shown in Figure 1."," 2005), but there is no evidence for a second peak in the high signal-to-noise spectrum shown in Figure 1."
+" Further, ALFA ZOA J1952+1428 has half the velocity width that the pair in HIZSS 3 appeared to have; Ws9 = 55 km s-! for HIZSS 3 (Henning et al."," Further, ALFA ZOA J1952+1428 has half the velocity width that the pair in HIZSS 3 appeared to have; $_{50}$ = 55 km $^{-1}$ for HIZSS 3 (Henning et al."
+ 2000) compared to Wso = 28 km s! here., 2000) compared to $_{50}$ = 28 km $^{-1}$ here.
+ Any second galaxy would have to be much closer both spatially and in velocity than the pair in HIZSS in order to escape detection., Any second galaxy would have to be much closer both spatially and in velocity than the pair in HIZSS 3 in order to escape detection.
+ Deeper interferometric observations3 would be needed to be entirely conclusive., Deeper interferometric observations would be needed to be entirely conclusive.
+ It is possible that ALFA ZOA J1952--1428 is a high velocity cloud (HVC) co-incident with an optical source (c.f., It is possible that ALFA ZOA J1952+1428 is a high velocity cloud (HVC) co-incident with an optical source (c.f.
+" HIPASS J1328-30, Grossi et al."," HIPASS J1328-30, Grossi et al."
+" 2007), though this is unlikely as there is strong evidence that ALFA ZOA J1952--1428 is not an HVC."," 2007), though this is unlikely as there is strong evidence that ALFA ZOA J1952+1428 is not an HVC."
+ Its recessional velocity does not lie near HVCs in this part of the sky (c.f., Its recessional velocity does not lie near HVCs in this part of the sky (c.f.
+ Figure 3a in Morras et al., Figure 3a in Morras et al.
+ 2000)., 2000).
+ The nearest population of HVCs is the Smith Cloud which lies 10? and 170 kms! away (Lockman et al., The nearest population of HVCs is the Smith Cloud which lies $10^{\circ}$ and 170 $\mathrm{km \; s^{-1}}$ away (Lockman et al.
+ 2008) at its nearest point., 2008) at its nearest point.
+" If ALFA ZOA J1952--1428 were an HVC, it would be a remarkable outlier."," If ALFA ZOA J1952+1428 were an HVC, it would be a remarkable outlier."
+" Furthermore, the velocity field of ALFA ZOA J19524-1428 shows a gradient ten times larger than those of HVCs (Begum et al."," Furthermore, the velocity field of ALFA ZOA J1952+1428 shows a gradient ten times larger than those of HVCs (Begum et al."
+ 2010)., 2010).
+ ALFA ZOA J1952--1428 appears to be a dwarf galaxy judging from its Gaussian profile and low mass., ALFA ZOA J1952+1428 appears to be a dwarf galaxy judging from its Gaussian profile and low mass.
+" At adistance of 7 Mpc, My; = 10'? Mo, which is significantly lower than the gaseous content of spiral- galaxies (Roberts and Haynes 1994)."," At adistance of 7 Mpc, $M_{HI}$ = $10^{7.0}$ $_\odot$ , which is significantly lower than the gaseous content of spiral-type galaxies (Roberts and Haynes 1994)."
+" Also, its low"," Also, its low"
+Whiteetal.(2007) applied au innovative technique to measure radio core Chussion from the quasars in the SDSS DR3 quasar catalog (Schneideretal.2005).,\citet{White07} applied an innovative technique to measure radio core emission from the quasars in the SDSS DR3 quasar catalog \citep{Schneider05}.
+. By stacking nuages from the FIRST survey (Beckeretal.1995).. they fouud radio core huuinosities iu the rauge of 1073. — 102WITZtsr1 consistent with huuinosities measured ou millizarcsecoud scales (BlIuudelletal.1996:etal. 2005).," By stacking images from the FIRST survey \citep{Becker95}, they found radio core luminosities in the range of $10^{21}$ – $10^{24} {\rm W
+Hz^{-1} sr^{-1}}$, consistent with luminosities measured on milli-arcsecond scales \citep{Blundell96,Falcke96,BlundBeas98,Ulvestad05}."
+. Of course. one cannot be sure that the core Dhuuinosities nieasured by Whiteetal.(2007). are not elevated because of contamination by svuchrotron Cluission from moving jet knots.," Of course, one cannot be sure that the core luminosities measured by \citet{White07} are not elevated because of contamination by synchrotron emission from moving jet knots."
+ Indeed. πιο].Beasley&Bicknell(2003) have direct. observational evidence for a highly relativistic jet in the bona fide raclio-quiet quasar 107|263.," Indeed, \citet*{Blundell03} have direct observational evidence for a highly relativistic jet in the bona fide radio-quiet quasar 1407+263."
+" A comparison of the radio wind Iguniuosities predicted bv equí((6)) and observed radio core hJuuinuosities indicates that disk winds can make a non-neegligible contribution to observed radio core hpuuinosities oulv if there is significant nass loss. with A‘. exceedius the Eddington Πιτ, Mgq23494vr.1. where M=LORS is the black hole mass M&4Mand y=0.d1go4 ds the radiative efficiency."," A comparison of the radio wind luminosities predicted by \ref{e:Lnu}) ) and observed radio core luminosities indicates that disk winds can make a non-negligible contribution to observed radio core luminosities only if there is significant mass loss, with $\Mdotw$ exceeding the Eddington limit, $\dot M_{\rm Edd} \simeq 3\,
+\eta_{0.1} M_8 M_\odot \, {\rm yr}^{-1}$, where $M = 10^8 M_8$ is the black hole mass and $\eta = 0.1
+\eta_{0.1}$ is the radiative efficiency."
+" This in turn maplies accretion rates AL,o>Afmaa."," This in turn implies accretion rates $\dot M_{\rm a} \gg \dot
+M_{\rm Edd}$."
+ Tle main role of the disk wind is then to relmove excess angular momentum. thereby reducing the accretion rate at smaller radii.," The main role of the disk wind is then to remove excess angular momentum, thereby reducing the accretion rate at smaller radii."
+ Arve super-Eddineton accretion and outflow rates ausible?, Are super-Eddington accretion and outflow rates plausible?
+ We remark that the Eddington lait is owed on smiplifius (aud incorrect) assuniptious of spherical svunuetry and a fiducial 106 radiative eficiency., We remark that the Eddington limit is based on simplifying (and incorrect) assumptions of spherical symmetry and a fiducial $10$ radiative efficiency.
+" The brightest radio galaxies aud quasars have volometric luminosities E-10!Mores1. requiring nass accretion rates of Af,21LOAL.vyHi."," The brightest radio galaxies and quasars have bolometric luminosities $L \simeq 10^{46-47}\,{\rm erg \, s}^{-1}$, requiring mass accretion rates of $\Mdota \gtapprox 1 - 10 \, M_\odot \, {\rm
+yr}^{-1}$."
+ This is a lower huit because the observed radiative output is almost certainly oulv a small fraction of the total oer extracted from accretion. much of which can be converted iuto low radiative efficiency. phenomena such as magnetized jets and/or a corona (seeI&uncic&Bick-well2001. 20010).," This is a lower limit because the observed radiative output is almost certainly only a small fraction of the total power extracted from accretion, much of which can be converted into low radiative efficiency phenomena such as magnetized jets and/or a corona \citep[see][]{KunBick04,KunBick07b}."
+" Similarly. mass outflow rates for warm. deuse iouized eas inferred from X-ray absorption spectra of quasars and Sevferts ave typically 0,1.LOAL.vr+ (ChaatasetPounds&Page 2006)..."," Similarly, mass outflow rates for warm, dense ionized gas inferred from X-ray absorption spectra of quasars and Seyferts are typically $0.1 - 10 \, M_\odot \, {\rm yr}^{-1}$ \citep{Chartas02,Pounds03a,Pounds03b,obrien05,PoundsPage06}."
+ Again. these values imply super-Eddington accretion.," Again, these values imply super-Eddington accretion."
+ In the case of SS133. the mass outflow rate inferred for the observed radio wiud is —E10JA3 (Bluudelletal.20013.," In the case of SS433, the mass outflow rate inferred for the observed radio wind is $\simeq
+4 \times 10^{-4} M_\odot \, {\rm yr}^{-1}$ \citep{Blundell01}."
+ If this Galactic nücroquasar isa LOAL.. black hole. then this tuples au astonishing iiss loss rate of AL.~10°Mg.," If this Galactic microquasar is a $10 M_\odot$ black hole, then this implies an astonishing mass loss rate of $\Mdotw \sim 10^3 \dot M_{\rm Edd}$."
+ If radio core cinission arises at least im part from an optically-thin disk wind. then we may expect to observe a correlation between the radio wind cussion and the optical/UWV blackbody disk enmüssion.," If radio core emission arises at least in part from an optically-thin disk wind, then we may expect to observe a correlation between the radio wind emission and the optical/UV blackbody disk emission."
+ This is expected because the disk wind cau modifv the local radial structure of the accretiou flow. resulting im a disk enission spectrun that is redder aud dinuuer than that of a disk without outflows (I&uucic&Bickuell2007a).," This is expected because the disk wind can modify the local radial structure of the accretion flow, resulting in a disk emission spectrum that is redder and dimmer than that of a disk without outflows \citep{KunBick07a}."
+. Such a correlation has indeed been reported by11).. who find that the SDSS DR3 quasars with redder colours teud to have higher mean radio flux deusities and are radio-louder than quasars with standardcolourst.," Such a correlation has indeed been reported by, who find that the SDSS DR3 quasars with redder colours tend to have higher mean radio flux densities and are radio-louder than quasars with standard."
+. Additionally. Whiteetal.(2007.theirfie.9). find a remarkably tight correlation between the 5GIIz median radio luminosity and the absolute maguitude at 2500 resthiune wavelength. with the specific radio aud optical Iuninosities related via Ly~Ln.," Additionally, \citet[their fig.\ 9]{White07} find a remarkably tight correlation between the $5\,{\rm GHz}$ median radio luminosity and the absolute magnitude at $2500$ restframe wavelength, with the specific radio and optical luminosities related via $L_R
+\sim L_{\rm opt}^{0.85}$."
+" The disk. wind model predicts that Ly2Lo, aud that optical bremsstrahline wind cussion should begin to contribute significantly to the observed (Av-corrected) mean optical quasar luminosity at 2500 (restirame). £a.c«dUeresIEj2 when Af210AL.xr+ (see refe:Lun))."," The disk wind model predicts that $L_R \approx
+L_{\rm opt}$ and that optical bremsstrahlung wind emission should begin to contribute significantly to the observed $K$ -corrected) mean optical quasar luminosity at $2500$ (restframe), $L_{\rm
+ opt} \simeq 4 \times 10^{30} \, {\rm erg \, s^{-1} \, {\rm
+ Hz}^{-1}}$, when $\Mdotw \gtapprox 10 \, M_\odot \, {\rm
+ yr}^{-1}$ (see \\ref{e:Lnu}) )."
+ Thus. unless the mass loss rate is extraordinarily bhüeh (that is. higher than the rates inferred from other iudepeudeut observations — c.f.," Thus, unless the mass loss rate is extraordinarily high (that is, higher than the rates inferred from other independent observations – c.f."
+ 33). it is likely that the optically-thick disk. cuissiou. rather than the optically-thin disk wind ciission. is primarily responsible for the observed. 2500 cluission.," 3), it is likely that the optically-thick disk emission, rather than the optically-thin disk wind emission, is primarily responsible for the observed $2500$ emission."
+" Remarkably, when Wlüteetal.(2007) calculate the radio loudness parameter, RY=Le/Lop. adjusted to remove the stroug radio/optical correlation. they fud logR*~0 (their 111)."," Remarkably, when \citet{White07} calculate the radio loudness parameter, $R^\star = L_R /L_{\rm opt}$, adjusted to remove the strong radio/optical correlation, they find $\log R^\star \sim 0$ (their 11)."
+ While it is not clear that this result is definitive. because of their asstuption that all radio cores at all redshifts lave a straight spectral iudex of a=0.5 (for which there are many counterexamples. some of them sccminely svsteniatic with redshift}. we note that our disk wind model provides a natural explanation for logR* (corrected for blackbody disk emissiou) being cousisteutlv close to zero.," While it is not clear that this result is definitive, because of their assumption that all radio cores at all redshifts have a straight spectral index of $\alpha = -0.5$ (for which there are many counterexamples, some of them seemingly systematic with redshift), we note that our disk wind model provides a natural explanation for $\log R^\star$ (corrected for blackbody disk emission) being consistently close to zero."
+ The parsec-scale cores of. radio-quiet quasars are comparable in luminosity. size and brightucss— teiiperature to those in racdio-loud quasars.," The parsec-scale cores of radio-quiet quasars are comparable in luminosity, size and brightness temperature to those in radio-loud quasars."
+ We sugecst that the model explored in this paper has applicability to the stationary. unresolved. flat spectruu compoucuts revealed with milli-arcsecoud VLBI observations. now an established technique for revealing moving aud evolving svuchrotron-enüttiung jet knots (e.g.Gómezetal.2000:Jorstad2005:Ly 2007).. and even revealed with elobal VLBI techniques with 108 of nülcro-arcesecond resolution (Bautakvroetal.," We suggest that the model explored in this paper has applicability to the stationary, unresolved, flat spectrum components revealed with milli-arcsecond VLBI observations, now an established technique for revealing moving and evolving synchrotron-emitting jet knots \citep[e.g.\,][]{Gomez00,Jorstad05,Ly07}, and even revealed with global VLBI techniques with 10s of micro-arcsecond resolution \citep{Rantakyro98}."
+1998).. Dluudell&Rawlings(2000.1999) have preseuted evidence (four physical arguineuts and from the observed near-absence of any relie radio lobes) that obe chussion mst disappear onu relatively rapicl nuescales.," \citet{BlundellRawlings2000,BlundellRawlings1999} have presented evidence (from physical arguments and from the observed near-absence of any relic radio lobes) that lobe emission must disappear on relatively rapid timescales."
+ Frou the sinularity between the dutv- of iutermittent jet activity in microquasars (their daring mode) and the fraction of quasars that are observed to be radio-oud. NipotiBlundell&Bin- posit that radio-louduess is a function of he epoch at which a quasar is observed.," From the similarity between the duty-cycle of intermittent jet activity in microquasars (their flaring mode) and the fraction of quasars that are observed to be radio-loud, \citet*{Nipoti05} posit that radio-loudness is a function of the epoch at which a quasar is observed."
+ The prescut eer sugeests that when there is no evidence for, The present paper suggests that when there is no evidence for
+"in RCS 2318+0034 and fitting the spectrum to an absorbed model yields a temperature of 7,,,.—6.32 (5.05-8.18) keV (16 uncertainty).",in RCS 2318+0034 and fitting the spectrum to an absorbed model yields a temperature of $T_{spec}$ =6.32 (5.05-8.18) keV $1\sigma$ uncertainty).
+ The resulting gas mass. total gravitating mass and gas mass fraction at σου and rag) assuming the gas Is isothermal are shown in Table 4.," The resulting gas mass, total gravitating mass and gas mass fraction at $r_{2500}$ and $r_{500}$ assuming the gas is isothermal are shown in Table 4."
+ The values for οσο and rsog are calculated from the total gravitating mass profile., The values for $r_{2500}$ and $r_{500}$ are calculated from the total gravitating mass profile.
+ Since the resulting values for Γκορ in Table 4 are slightly less than IMpe. we also extracted spectra from several annuli with inner radii of 7Okpe and outer radi of 700. 800 and 900kpc.," Since the resulting values for $r_{500}$ in Table 4 are slightly less than 1Mpc, we also extracted spectra from several annuli with inner radii of 70kpc and outer radii of 700, 800 and 900kpc."
+ Due to the limited photon statistics in the observation. the best-fit temperatures are all consistent within the uncertainties.," Due to the limited photon statistics in the observation, the best-fit temperatures are all consistent within the uncertainties."
+ Table 4 shows that assuming an isothermal D. model for RCS 231840034 produces gas mass fractions of 0.07+.02 and 0.08+.02 at rosgg and rasoo. respectively.," Table 4 shows that assuming an isothermal $\beta$ model for RCS 2318+0034 produces gas mass fractions of $0.07 \pm .02$ and $0.08 \pm .02$ at $r_{2500}$ and $r_{500}$, respectively."
+ The error bars take into account the uncertainty in the gas mass noted above and the uncertainty in Jy..., The error bars take into account the uncertainty in the gas mass noted above and the uncertainty in $T_{spec}$.
+ The gas mass fraction is essentially independent of radius at large radit since. for an isothermal B—0.7 model. Pegs~r and Pro~177.," The gas mass fraction is essentially independent of radius at large radii since, for an isothermal $\beta=0.7$ model, $\rho_{gas} \sim r^{-2.1}$ and $\rho_{tot} \sim r^{-2}$."
+ Based on the analysis of observations of a sample of 13 rich clusters of galaxies. Vikhlinin et al. (," Based on the analysis of observations of a sample of 13 rich clusters of galaxies, Vikhlinin et al. ("
+"2005) found that the temperature profile of these clusters could be fit with a broken power-law model. given by: where «Τὸ is the emission-weighted temperature excluding the central 70kpe (essentially 7,54).","2005) found that the temperature profile of these clusters could be fit with a broken power-law model, given by: where $\langle T\rangle$ is the emission-weighted temperature excluding the central 70kpc (essentially $T_{spec}$ )."
+ For clusters. που~0.6710.," For clusters, $r_{500} \sim 0.6r_{180}$."
+ With this substitution. the temperature profile in eq. (," With this substitution, the temperature profile in eq. ("
+5) is very similar to that derived for a sample of groups by Sun et al. (,5) is very similar to that derived for a sample of groups by Sun et al. (
+2009).,2009).
+ Gas masses. total gravitating masses and gas mass fractions within r»soo and rao) are shown in Table 5 using our standard f model and the broken power-law temperature profile in eq. (," Gas masses, total gravitating masses and gas mass fractions within $r_{2500}$ and $r_{500}$ are shown in Table 5 using our standard $\beta$ model and the broken power-law temperature profile in eq. ("
+5) normalized by Typ...,5) normalized by $T_{spec}$.
+ Since eq. (, Since eq. (
+5) depends on rsoo. we iterated In που until a self-consistent solution was found.,"5) depends on $r_{500}$, we iterated in $r_{500}$ until a self-consistent solution was found."
+ The self- temperature profiles are shown in Figure 4., The self-consistent temperature profiles are shown in Figure 4.
+ Table 5 shows that assuming a broken power-law temperature profile for RCS 2318+0034 produces gas mass fractions of 0.06+.02 and O.10-+.02 at rasog and rsoo. respectively.," Table 5 shows that assuming a broken power-law temperature profile for RCS 2318+0034 produces gas mass fractions of $0.06 \pm .02$ and $0.10 \pm .02$ at $r_{2500}$ and $r_{500}$, respectively."
+" The error bars take into account the uncertainty in the gas mass and the uncertainty in. 7,4.", The error bars take into account the uncertainty in the gas mass and the uncertainty in $T_{spec}$.
+ AS opposed to the isothermal model. the gas mass fraction increases with radius due to the declining gas temperature.," As opposed to the isothermal model, the gas mass fraction increases with radius due to the declining gas temperature."
+ While RCS 2318-0034 1s presently undergoing a merger. Nagai et al. (," While RCS 2318+0034 is presently undergoing a merger, Nagai et al. ("
+2007) have shown by performing mock observations of a sample of simulated clusters that X-ray derived gas masses 1n both relaxed and unrelaxed clusters are accurate to within6%.,2007) have shown by performing mock observations of a sample of simulated clusters that X-ray derived gas masses in both relaxed and unrelaxed clusters are accurate to within.
+. Including this additional uncertainty in our derived gas mass would not a have a significant affect on our estimate of the gas mass fraction., Including this additional uncertainty in our derived gas mass would not a have a significant affect on our estimate of the gas mass fraction.
+ Nagai et al. (, Nagai et al. (
+2007) also found that X-ray derived gravitating masses are biased low by about with no trend in cluster mass or redshift.,2007) also found that X-ray derived gravitating masses are biased low by about with no trend in cluster mass or redshift.
+ This suggests that all X-ray derived gas mass fraction estimates are biased high by this amount., This suggests that all X-ray derived gas mass fraction estimates are biased high by this amount.
+ Our estimated gas mass fractions using the combined 200 ksee ofChandra data are approximately greater at rosgg and significantly greater at Áo) compared with the values derived from the earlier 50 ksec observation (Hicks et al., Our estimated gas mass fractions using the combined 200 ksec of data are approximately greater at $r_{2500}$ and significantly greater at $r_{500}$ compared with the values derived from the earlier 50 ksec observation (Hicks et al.
+ 2008)., 2008).
+ Comparing the derived parameters for RCS 23184-0034. in Hicks et al., Comparing the derived parameters for RCS 2318+0034 in Hicks et al.
+ with our results. we find that both studies produced similar values for the gas temperature (ντ5.05.8.18 keV compared to Typee=5.207.40 keV) and core radius (7.= kpe compared to —r.=171+6 kpe) for the surface brightness profile.," with our results, we find that both studies produced similar values for the gas temperature $T_{spec}=5.05-8.18$ keV compared to $T_{spec}=5.20-7.40$ keV) and core radius $r_c=142 \pm 23$ kpc compared to $r_c=171 \pm 6$ kpc) for the surface brightness profile."
+ However. we obtain a value for the outer slope of the surface brightness profile corresponding to," However, we obtain a value for the outer slope of the surface brightness profile corresponding to"
+due to the following reasons: (7) it is very uulikelv that iron abundance chanecs by a factor of about two im just six mouths. and (77) it is dificult to explain the variation iu the shape of the non-thermal X-ray contin by changing the iron abundance.,"due to the following reasons: $i$ ) it is very unlikely that iron abundance changes by a factor of about two in just six months, and $ii$ ) it is difficult to explain the variation in the shape of the non-thermal X-ray continuum by changing the iron abundance."
+ Increasing the reflection cussion frou a ueutral disk may fatten the continua., Increasing the reflection emission from a neutral disk may flatten the continuum.
+" This will require a change in the ecometry of the accretion disk-coroua svsteni (οιο, Increasing the covering factor at the corona will result in more reflection cuussion).", This will require a change in the geometry of the accretion disk-corona system (e.g. increasing the covering factor at the corona will result in more reflection emission).
+ However. this will iucrease he equivalent width of the iron Ίνα liue with flattening of the contiuuuu.," However, this will increase the equivalent width of the iron $\alpha$ line with flattening of the continuum."
+ Large equivalent widths of Fel&o can )o produced if the accretion disk is ionized (Weaverctal., Large equivalent widths of $\alpha$ can be produced if the accretion disk is ionized \citep{Weaveretal97}.
+ 1997).. Dewangan(2002) reported a correlation between he shape of the N-ray spectrum and the αισον of the Felxa line among radio-quiet type 1 ACNs aud suggested hat the accretion disks of AGNs with steeper coutinua are nore ionized., \citet{Dewangan02} reported a correlation between the shape of the X-ray spectrum and the energy of the $\alpha$ line among radio-quiet type 1 AGNs and suggested that the accretion disks of AGNs with steeper continua are more ionized.
+ The treud in Dy adl the equivalent width of iron Wa cussion from ALCC-5-23-16 could be uuderstood if the shape of the N-vayv coutimmiun is coupled with the ionization stage of the disk material in MCC-5-25-16 ina similar way to that observed among the racdio-quict ACUNSs., The trend in $\Gamma_X$ and the equivalent width of iron $\alpha$ emission from MCG-5-23-16 could be understood if the shape of the X-ray continuum is coupled with the ionization stage of the disk material in MCG-5-23-16 in a similar way to that observed among the radio-quiet AGNs.
+ As shown in section 3.. the broad feature at 6.1keV. can be inodeled as a broad Caussiau: (PFWIIM--10000kins1 7) or a line profile arising frou au accretion disk around a Schwarzschild or a Nery black hole.," As shown in section \ref{analysis}, the broad feature at $6.4\kev$ can be modeled as a broad Gaussian $\sim40000\kms$ ) or a line profile arising from an accretion disk around a Schwarzschild or a Kerr black hole."
+ Alternatively. most of the fux iu the broad component can also be accounted bv reflection enussion.," Alternatively, most of the flux in the broad component can also be accounted by reflection emission."
+ However. the cerived reflection fraction (R~ 3) ds much lareer than that inferred frou the observations of AICC-5-23-16 (R~0.5: Risaliti (200211) and the average value for Sevtert 1 ealaxics (2~1: Perolaetal. 20023).," However, the derived reflection fraction $R\sim3$ ) is much larger than that inferred from the observations of MCG-5-23-16 $R\sim 0.5$; \citet{Risaliti02}) ) and the average value for Seyfert 1 galaxies $R \sim 1$; \citealt{Perolaetal02}) )."
+" This is not entirely impossible if the reflector is a putative torus far away from the central contin, source.", This is not entirely impossible if the reflector is a putative torus far away from the central continuum source.
+ The reflection fraction can be large if the central source was brighter before a time period which is equal to the time delay between the reflected aud primary X-ray , The reflection fraction can be large if the central source was brighter before a time period which is equal to the time delay between the reflected and primary X-ray emission.
+However. such a large reflection fraction for AICC-5-2:-16Cuisson. nationis uulikely due to the following reasons — (1) the calil of the EPIC PN above 10keV is reliable at a level of ~20% (Suowdeu. private communication).," However, such a large reflection fraction for MCG-5-23-16 is unlikely due to the following reasons – (i) the calibration of the EPIC PN above $10\kev$ is reliable at a level of $\sim20\%$ (Snowden, private communication)."
+ MOS data do not show clear excess cuission above 10keV as is the case for PN. G, MOS data do not show clear excess emission above $10\kev$ as is the case for PN. (
+i) The broad feature at 6.1keV weakeus with flattening of the continua.,ii) The broad feature at $\sim6.4\kev$ weakens with flattening of the continuum.
+ This is not expected if the broad feature were the reflection wission., This is not expected if the broad feature were the reflection emission.
+ Siuultauneous observations with id.μου would be crucial to. characterize both the Felxa emission aud reflection compoucut., Simultaneous observations with and would be crucial to characterize both the $\alpha$ emission and reflection component.
+ observations of December 2001. have revealed the presence of au Trou Is-shell absorption edec near ~7.1keV., observations of December 2001 have revealed the presence of an Iron K-shell absorption edge near $\sim 7.1\kev$.
+ However. the edge is not well constrained aud is detected oulv at a siguificauce level of 93%.," However, the edge is not well constrained and is detected only at a significance level of $93\%$."
+ The edge energy is uot well determinedbut it appears to be ~Tl |eFie. 2)), The edge energy is not well determined but it appears to be $\sim 7.1\kev$ (seeFig. \ref{f2}) ).
+ The optical depth is found to be 0.03Kew. sorption{x, The optical depth is found to be $0.03_{-0.02}^{+0.04}$.
+cAsstunine that the edge arises due to the Is-shell al of ueutral iron aud the corresponding photoionization cross section 2000)).-(2.6«10.78αμ7: Domuelly.Bell.Scott.&IWeenan we determine the col deusity of neutral rou to be Ng.z1.1«1075em2.," Assuming that the edge arises due to the K-shell absorption of neutral iron and the corresponding photoionization cross section $2.6\times10^{-20}{\rm~cm^{-2}}$; \citealt{Donnellyetal00}) ), we determine the column density of neutral iron to be $N_{\rm Fe} \approx 1.1\times10^{18}{\rm~cm^{-2}}$."
+ This corresponds to an effective Ibdrogeu. column density of Agmodos1072cm7. which is a factor of two higher hau the hydrogen coluun density derived bv fitting au absorbed power-law model to the observed data.," This corresponds to an effective Hydrogen column density of $\nh \approx 4\times 10^{22}{\rm~cm^{-2}}$, which is a factor of two higher than the hydrogen column density derived by fitting an absorbed power-law model to the observed data."
+ Given he uncertainty in the derived value of the optical depth. it cannot be ruled out that the Fe I&-edee arises eutirely in he absorbing material along our line of sight iu the form of a putative torus.," Given the uncertainty in the derived value of the optical depth, it cannot be ruled out that the Fe K-edge arises entirely in the absorbing material along our line of sight in the form of a putative torus."
+ It is interesting to Investigate whetler th. the absorption {οσο aud the narrow coniponenut of the Ίνα line due to ron could be produced im the the mitative torus around au accretion disk., It is interesting to investigate whether both the absorption K-edge and the narrow component of the $\alpha$ line due to iron could be produced in the the putative torus around an accretion disk.
+ Iu the case of isotropic enussion. the equivalent width of the emission iue should be ~σιμεs1077.," In the case of isotropic emission, the equivalent width of the emission line should be $\sim 64f_c \nh A_{Fe} \times 10^{-23}$."
+ For a typical covering factor f.~0.7 for the torus aud solar abundances Aye=1). the absorption column required to produce an equivalent width of LOeV for the narrow ion Ίνα line is ~LO?«n2.," For a typical covering factor $f_c\sim0.7$ for the torus and solar abundances $A_{fe} = 1$ ), the absorption column required to produce an equivalent width of $40\ev$ for the narrow iron $\alpha$ line is $ \sim 10^{23}{\rm~cm^{-2}}$."
+ This column is larger by a factor of five than the absorption column derived. from the power-law fit., This column is larger by a factor of five than the absorption column derived from the power-law fit.
+ It is also larger by about a factor of three than the absorption column derived from the iron I-edee. since the absorption features are produced only by a column along our line of sight while the ciissiou features are spatially integrated in the case of a source which is unresolved.," It is also larger by about a factor of three than the absorption column derived from the iron K-edge, since the absorption features are produced only by a column along our line of sight while the emission features are spatially integrated in the case of a source which is unresolved."
+ The above results Πρίν that we are viewing the outer οσο of a putative torus where the column density is lower than the iuner regions of the torus., The above results imply that we are viewing the outer edge of a putative torus where the column density is lower than the inner regions of the torus.
+ This is consistent with that expected for a Sevtert 1.9 ealaxy from the unification scheme (CAntonucci 1993).., This is consistent with that expected for a Seyfert 1.9 galaxy from the unification scheme \citep{Antonucci93}. .
+ We studied the 2.5.10keV spectu) of MCC-5-23-16 which was observed twice with iu May aud December 2001., We studied the $2.5-10\kev$ spectrum of MCG-5-23-16 which was observed twice with in May and December 2001.
+ Our main results are as follows., Our main results are as follows.
+tested in future observations.,tested in future observations.
+ The critical density of the ttransition is high (about 44=10° em (2))) and the low reduced mass of the mmolecule makes its excitation very sensitive to radiative pumping by dust., The critical density of the transition is high (about $n_{\rm crit} = 10^6$ $\cmmd$ \citep{phillips92}) ) and the low reduced mass of the molecule makes its excitation very sensitive to radiative pumping by dust.
+ We therefore assume that the coupling between the colour temperature of the IR emission and the excitation temperature of hholds as long as the dust colour temperature is above 30 K (see discussion in ?))., We therefore assume that the coupling between the colour temperature of the IR emission and the excitation temperature of holds as long as the dust colour temperature is above 30 K (see discussion in \citet{vandertak08}) ).
+ For each galaxy we ran a RADEX model (?) with temperatures ranging from 7210-500 K. In these calculations. the volume density Is set to an arbitrary large value. so that the excitation of iis thermalized.," For each galaxy we ran a RADEX model \citep{vandertak07} with temperatures ranging from $T$ =10-500 K. In these calculations, the volume density is set to an arbitrary large value, so that the excitation of is thermalized."
+ The adopted excitation temperature Is essentially the radiation temperature and a radiative excitation with one excitation temperature is simulated., The adopted excitation temperature is essentially the radiation temperature and a radiative excitation with one excitation temperature is simulated.
+ We therefore do not use the full non-LTE capacity of RADEX., We therefore do not use the full non-LTE capacity of RADEX.
+ The results are shown in Fig. 2.., The results are shown in Fig. \ref{f:radex}.
+ Each plot has temperature (kinetic or radiation) on the x axis and ccolumn on the y axis., Each plot has temperature (kinetic or radiation) on the x axis and column on the y axis.
+ The contour line corresponds to the observed brightness temperature., The contour line corresponds to the observed brightness temperature.
+ This depends on our assumption for the source size., This depends on our assumption for the source size.
+ aabundances relative to Hs (X¥(H307))) are listed in table 3.., abundances relative to $_2$ $X$ )) are listed in table \ref{t:column}.
+" We have used CO 2-] spectra from the literature and used a Galactic conversion factor from CO luminosity to H+ mass (2.5x10°"" em? K! km! s) to estimate NCH>).", We have used CO 2–1 spectra from the literature and used a Galactic conversion factor from CO luminosity to $_2$ mass $2.5 \times 10^{20}$ $\cmmt$ $^{-1}$ $^{-1}$ s) to estimate $N$ $_2$ ).
+ Note that for, Note that for
+"and the displacement of the maximum towards larger distance in comparison with (9°,8°).","and the displacement of the maximum towards larger distance in comparison with $(9^\circ,8^\circ)$."
+" In the last panel we show one of the central lines-of-sight (3°,6°) for the boxy bulge."," In the last panel we show one of the central lines-of-sight $(3^\circ,6^\circ)$ for the boxy bulge."
+" Repeating the identification of maxima for all fields in —6?<|€30° and —8?€b8°, we produce a similar plot as in C07."," Repeating the identification of maxima for all fields in $-6^\circ \le l \le 30^\circ$ and $-8^\circ \le b \le 8^\circ$, we produce a similar plot as in C07."
+ The distance modulus for each maximum is converted to distance in kpc from the observer's position and plotted in Figure 4 in galactocentric coordinates., The distance modulus for each maximum is converted to distance in kpc from the observer's position and plotted in Figure \ref{fig:posdiagram} in galactocentric coordinates.
+" We show such a plot both for the snapshot with leading bar ends (Figure 4a, corresponding to Figure and for the snapshot with straight ends (Figure 4b, corresponding1)), to Figure 2))."," We show such a plot both for the snapshot with leading bar ends (Figure \ref{fig:posdiagram}{, corresponding to Figure \ref{fig:snapshot1}) ), and for the snapshot with straight ends (Figure \ref{fig:posdiagram}{, corresponding to Figure \ref{fig:snapshot2}) )."
+" In both models, the points computed from star counts above the plane follow the heavy line indicating the true orientation(pink crosses)angle α=25? fixed by us for the model'sbar, except at the very center and at negative | where the maxima are closer to the observer due to tangent point effects."," In both models, the points computed from star counts above the plane (pink crosses) follow the heavy line indicating the true orientation angle $\alpha=25^\circ$ fixed by us for the model's, except at the very center and at negative $l$ where the maxima are closer to the observer due to tangent point effects."
+" For lower latitudes (black crosses), the points move further back."," For lower latitudes (black crosses), the points move further back."
+" In both models, for longitudes 0?«|9? in the region of the boxy bulge, points in the plane now approximately follow an imaginary line with a’=43°."," In both models, for longitudes $0^\circ<l<9^\circ$ in the region of the boxy bulge, points in the plane now approximately follow an imaginary line with $\alpha'=43^\circ$."
+ 'This is due to a volume effect combined with the shallow disk density distribution: N(u)οςn(D)D?dD/d., This is due to a volume effect combined with the shallow disk density distribution: $N(\mu)\propto n(D)D^2 dD/d\mu$.
+" There is an intermediate region with 9?<|20°, where the boxy bulge becomes thinner and transits into the planar bar, and then there is a region, 20?<|«30? (~3 to ~5 kpc), where we can see just particles in the plane."," There is an intermediate region with $9^\circ<l<20^\circ$, where the boxy bulge becomes thinner and transits into the planar bar, and then there is a region, $20^\circ<l<30^\circ$ $\sim 3$ to $\sim 5$ kpc), where we can see just particles in the plane."
+" In this region, both models differ: the volume effect alone, as illustrated for the bar with the straight ends in Figure 40, accounts for only half the shift from α=25° to a’=43°."," In this region, both models differ: the volume effect alone, as illustrated for the bar with the straight ends in Figure \ref{fig:posdiagram}{, accounts for only half the shift from $\alpha=25^\circ$ to $\alpha'=43^\circ$."
+ Whereas for the snapshot with the leading bar ends the in-plane points now approximately follow the (Figureimaginary4a) o/—43? line right until the end of the bar is reached., Whereas for the snapshot with the leading bar ends (Figure \ref{fig:posdiagram}{ ) the in-plane points now approximately follow the imaginary $\alpha'=43^\circ$ line right until the end of the bar is reached.
+" If we were not aware of the original structure of the model, the results shown in Figure 4 could easily be interpreted as evidence for a model with two structures, one thicker and shorter at a=25? and another thinner and longer at o/—43?."," If we were not aware of the original structure of the model, the results shown in Figure \ref{fig:posdiagram} could easily be interpreted as evidence for a model with two structures, one thicker and shorter at $\alpha = 25^\circ$ and another thinner and longer at $\alpha'\sim 43^\circ$."
+" Clearly, the case discussed shows that a separatebar is not necessarily implied by the observed star count distributions and that a singlebar structure is a valid and dynamically simpler interpretation of these data."," Clearly, the case discussed shows that a separate is not necessarily implied by the observed star count distributions and that a single structure is a valid and dynamically simpler interpretation of these data."
+" Several large radial velocity surveys of Galactic bulge and disk stars are currently on-going or planned, and we thus provide some predictions for our model."," Several large radial velocity surveys of Galactic bulge and disk stars are currently on-going or planned, and we thus provide some predictions for our model."
+" We scale the average circular velocity curve of the model to the data of Clemens(1985) for radii 3-4.5 kpc, and then compute the radial velocities at different longitudes as seen from the Sun (assumingvg@=250km/s;Reidetal.2009)."," We scale the average circular velocity curve of the model to the data of \citet{Clemens85} for radii 3-4.5 kpc, and then compute the radial velocities at different longitudes as seen from the Sun \citep[assuming $v_\sun$=250 km/s;][]{Reid+09}."
+". We show the mean velocity, dispersion, and particle number ddistance modulus in Figure 5 for two in-plane positions near the end of the planar bar."," We show the mean velocity, dispersion, and particle number distance modulus in Figure \ref{fig:velhisto} for two in-plane positions near the end of the planar bar."
+ The non-circular motions in the barred models leave a clear signature at 14c13.0—13.6 compared with the initial axisymmetric disk., The non-circular motions in the barred models leave a clear signature at $\mu\simeq 13.0-13.6$ compared with the initial axisymmetric disk.
+ In this range the bar’s velocity profile is flat., In this range the bar's velocity profile is flat.
+ For the leading-ends bar it has a break at ys~13.7 and then increases towards an axisymmetric type of profile., For the leading-ends bar it has a break at $\mu\simeq13.7$ and then increases towards an axisymmetric type of profile.
+ For the straight-ends bar the break happens earlier yu~13.5 and the radial velocities are slightly smaller., For the straight-ends bar the break happens earlier $\mu\simeq13.5$ and the radial velocities are slightly smaller.
+ The difference between axisymmetric and bar case is ~30—40 km/s. 'The velocity dispersions of both bar models in the two, The difference between axisymmetric and bar case is $\simeq 30-40$ km/s. The velocity dispersions of both bar models in the two
+"1n Model 2 of Fükugita&Peebles(2006).. the density. of the corona decreases with radius as r and at LOkpe has electron density n,=2.610""em.","In Model 2 of \cite{FukugitaP06}, the density of the corona decreases with radius as $r^{-3/2}$ and at $10\kpc$ has electron density $n_\e=2.6\times10^{-3}\cm^{-3}$."
+" The corona is assumed to be isothermal with Zi=1.810""Ix.", The corona is assumed to be isothermal with $T_\h=1.8\times10^6\K$.
+ shows the cooling time of gas with metallicitics Fe/1l]=1 and 0.5 that is in pressure equilibrium with plasma of this emperature= and the mean density of this corona between 8 and 12kpe., shows the cooling time of gas with metallicities $\hbox{[Fe/H]}=-1$ and $-0.5$ that is in pressure equilibrium with plasma of this temperature and the mean density of this corona between 8 and $12\kpc$.
+" At 2—10""I. the cooling time is ~300Myr."," At $T\sim2\times10^6\K$, the cooling time is $\sim300\Myr$."
+" Once the temperature has fallen to 510""KK. the cooling ime has dropped by nearly. forty to ~SMyr. which is ess than the dynamical time at the solar galactocentric radius. fy."," Once the temperature has fallen to $5\times10^5\K$, the cooling time has dropped by nearly forty to $\sim8\Myr$, which is less than the dynamical time at the solar galactocentric radius, $R_0$."
+ By the time the temperature has dropped. by a further factor of two to 2.510WK. the cooling time ias fallen by more than a further. order of magnitude and is a mere 0.4Myr.," By the time the temperature has dropped by a further factor of two to $2.5\times10^5\K$, the cooling time has fallen by more than a further order of magnitude and is a mere $0.4\Myr$."
+ Thus although the cooling time of ambient gas in the lower corona is long. any diminution in temperature will dramatically shorten the cooling time.," Thus although the cooling time of ambient gas in the lower corona is long, any diminution in temperature will dramatically shorten the cooling time."
+ 1n the model of Fükugita&Peebles(2006). the racdiating plasma [lows inwards at a rate LAL.ve, In the model of \cite{FukugitaP06} the radiating plasma flows inwards at a rate $\sim1\msun\yr^{-1}$.
+" The work done by compression olfsets radiative losses. so the plasma emperature remains L8«10""Ix."," The work done by compression offsets radiative losses, so the plasma temperature remains $1.8\times10^6\K$."
+ The mass of an thalo is rather ill-defined: in an external galaxy it is. not clear at what value of |z] we should. place the boundary »etween disc ancl haloHELLE. and in the Galaxy. there is a similar ambiguity in the line-of-sight. velocity that clivides 1alo [from elise L," The mass of an halo is rather ill-defined: in an external galaxy it is not clear at what value of $|z|$ we should place the boundary between disc and halo, and in the Galaxy there is a similar ambiguity in the line-of-sight velocity that divides halo from disc ."
+LL. Win NGC 891 the disc-halo boundary is conservatively place at 1.3kpe. the mass of the ihalo is ~6LO”AZ. (FPBOG).," If in NGC 891 the disc-halo boundary is conservatively place at $1.3\kpc$, the mass of the halo is $\sim6\times10^8\msun$ (FB06)."
+ Ixalberla&Dedes(2008) conclude that ~10 per cent of the Galaxys iis halo gas. and that the total nunass within 12kpec (the region within which most star formation occurs) is 4.5«10°AZ.. so the mass of the hhalo within 12kpe is 4.510AJ..," \cite{KalberlaD08} conclude that $\sim10$ per cent of the Galaxy's is halo gas, and that the total mass within $12\kpc$ (the region within which most star formation occurs) is $4.5\times10^9\msun$, so the mass of the halo within $12\kpc$ is $\sim4.5\times10^8\msun$."
+ The hhalo is largely confined to the region |z|<Skpe. so we compare the mass of the hhalo at Ro«Ikpe to the mass of coronal gas in the cylindrical annulus 4Rikpe12. ];|«Skpe.," The halo is largely confined to the region $|z|<5\kpc$, so we compare the mass of the halo at $R<12\kpc$ to the mass of coronal gas in the cylindrical annulus $4\le R/\!\kpc\le12$, $|z|<5\kpc$."
+ The coronal mass is 2«107A4... so within this volume there is over twice as much aas coronal gas. although the coronal gas will occupy nearly all the space.," The coronal mass is $\sim2\times10^8\msun$, so within this volume there is over twice as much as coronal gas, although the coronal gas will occupy nearly all the space."
+ The clouds that make up the hhalo typically take 100Myr to travel from their launch point in the plane through the halo and back to the plane., The clouds that make up the halo typically take $100\Myr$ to travel from their launch point in the plane through the halo and back to the plane.
+" Consequently. in one Cyr 4.510"" of iis passed through the 2«107AL. of gas in the lower corona."," Consequently, in one Gyr $\sim4.5\times10^9$ of is passed through the $2\times10^8\msun$ of gas in the lower corona."
+ Lence the temperature of the coronal gas would be halved ifjust 4.5 per cent of the gas that passes through the corona in à Gyr were to mix with the coronal gas., Hence the temperature of the coronal gas would be halved if just 4.5 per cent of the gas that passes through the corona in a Gyr were to mix with the coronal gas.
+ This drop in temperature would bring the cooling time of the coronal gas down to 0.1Gyr. the local orbital time.," This drop in temperature would bring the cooling time of the coronal gas down to $\sim 0.1\Gyr$, the local orbital time."
+ The eclouds of the halo plough through the corona at. speeds à that are less than but comparable to the sound speed of the coronal gas., The clouds of the halo plough through the corona at speeds $v$ that are less than but comparable to the sound speed of the coronal gas.
+ Consequently the Low around these clouds is likely to be in a high Itevnolds number regime. and. cach cloud must be decelerated by ram. pressure of order puc. where py is the coronal density.," Consequently the flow around these clouds is likely to be in a high Reynolds number regime, and each cloud must be decelerated by ram pressure of order $\rho_\h v^2$, where $\rho_\h$ is the coronal density."
+ In these circumstances the cloud. loses its momentum on a timescale equal to. pc/pn times the time ἐς for the cloud to move its own length (EDOS)., In these circumstances the cloud loses its momentum on a timescale equal to $\rho_\c/\rho_\h$ times the time $t_\c$ for the cloud to move its own length (FB08).
+ Clouds have diameters of a few tens of. parsecs and travel a few κρο while pofpuc200.," Clouds have diameters of a few tens of parsecs and travel a few kpc, while $\rho_\c/\rho_\h\simeq200$."
+ Hence clouds must surrender a significant fraction of their momentum to coronal gas (Benjamin2002)., Hence clouds must surrender a significant fraction of their momentum to coronal gas \citep{ben02}.
+. Above we showed that in a Gyr 4.5«10A4. of ppasses through the 2;107AZ. of the lower corona., Above we showed that in a Gyr $\sim4.5\times10^9\msun$ of passes through the $2\times10^8\msun $ of the lower corona.
+ Clearly- the coronal gas cannot absorb a significant [raction of the momentum of more than 20 times its mass of ((FBOS)., Clearly the coronal gas cannot absorb a significant fraction of the momentum of more than 20 times its mass of (FB08).
+ Moreover. if the elise were losing angular momentum to the corona at the rate this calculation implies. it would be contracting on a Car timescale.," Moreover, if the disc were losing angular momentum to the corona at the rate this calculation implies, it would be contracting on a Gyr timescale."
+ Phe natural resolution of these problems is that the disc most of the coronal gas that cclouds encounter., The natural resolution of these problems is that the disc most of the coronal gas that clouds encounter.
+ Phen the momentum lost by ecloucs would be returned to the disc. and it would not build up in the corona.," Then the momentum lost by clouds would be returned to the disc, and it would not build up in the corona."
+ As was mentioned in Section 1. coronal gas in dark-matter halos more massive than those of spiral galaxies has been extensively studied for four decades.," As was mentioned in Section 1, coronal gas in dark-matter halos more massive than those of spiral galaxies has been extensively studied for four decades."
+ Cooling within these halos shows no tendeney to produce a cold. stellar. disc the colclest gas is near the centre. where the cooling time is shortest. and whatever gas cools out of the corona feeds the central black hole rather than forming a star-forming disc.," Cooling within these halos shows no tendency to produce a cold stellar disc – the coldest gas is near the centre, where the cooling time is shortest, and whatever gas cools out of the corona feeds the central black hole rather than forming a star-forming disc."
+ Phe argument of the last subsection may help us to understand why in less massive dark halos cooling coronal eas Lows into the disc rather than onto a central black hole: so long as the halo has a star-forming clisc. that disc sustains its star formation by reaching up ancl grabbing coronal gas.," The argument of the last subsection may help us to understand why in less massive dark halos cooling coronal gas flows into the disc rather than onto a central black hole: so long as the halo has a star-forming disc, that disc sustains its star formation by reaching up and grabbing coronal gas."
+ Dises are disrupted. by major mergers. which occur rather frequently.," Discs are disrupted by major mergers, which occur rather frequently."
+ Analytic arguments and hyelrodynanical simulations ofcosmological clustering suggest that when the dise of a relatively low-mass galaxy is disrupted by a major merger. it will quickly re-form from filaments of cold," Analytic arguments and hydrodynamical simulations ofcosmological clustering suggest that when the disc of a relatively low-mass galaxy is disrupted by a major merger, it will quickly re-form from filaments of cold"
+in this redshift range.,in this redshift range.
+ However. 904 and Dahlenetal.(2004) have shown that ος} is a reflection of the SFRy:(2) which has been delaved by the assembly time function (or delay lime function). i. e. the time for SN Ia progenitors to go from formation to explosion.," However, S04 and \citet{Dahlen2003} have shown that $\mathcal R_{Ia}(z)$ is a reflection of the $_U$ $z$ ) which has been delayed by the assembly time function (or delay time function), i. e. the time for SN Ia progenitors to go from formation to explosion."
+" S04 find that the dillerence in the evolution of the SFRy:(2) and the R,,,(2) ean be well modeled as a Gaussian delav time function with a mean delay of 4 Gyr and a dispersion (hat is a small fraction of the delay.", S04 find that the difference in the evolution of the $_U$ $z$ ) and the $\mathcal{R}_{Ia}(z)$ can be well modeled as a Gaussian delay time function with a mean delay of $\sim 4$ Gyr and a dispersion that is a small fraction of the delay.
+ Subsequently Dallenetal.(2004). find that based on the model delay. (he fraction of stellar mass in (he range of 3 (o 8 solar masses which will explode as SNe la is about 5%.," Subsequently \citet{Dahlen2003} find that based on the model delay, the fraction of stellar mass in the range of 3 to 8 solar masses which will explode as SNe Ia is about $5\%$."
+ Past searches for SNe la in high-redshift clusters with HST by show rates that are similar to field SN Ta rates. indicating some similarity in their rate evolution.," Past searches for SNe Ia in high-redshift clusters with HST by \citet{2002MNRAS.332...37G} show rates that are similar to field SN Ia rates, indicating some similarity in their rate evolution."
+ Unfortunately. these cluster rate measurement al zS>0.1 are not precise enough. nor sufficiently well saniplec in redshift space. to definitivelv compare their SN Ia rate history io those measured in normal (field) environments.," Unfortunately, these cluster rate measurement at $z\gg0.1$ are not precise enough, nor sufficiently well sampled in redshift space, to definitively compare their SN Ia rate history to those measured in normal (field) environments."
+" An anecdotal anomolv in the consideration of the decline in j,(:) al z>1 is the serendipitous discovery of SN 19971E. found in only one re-observation of a pointing withHST and WEPC? (CGililandetal.1999:Riess2001).."," An anecdotal anomoly in the consideration of the decline in $\mathcal R_{Ia}(z)$ at $z>1$ is the serendipitous discovery of SN 1997ff, found in only one re-observation of a pointing with and WFPC2 \citep{Gilliland1999,2001ApJ...560...49R}."
+" The combined depth of the template-searehi pair lor the SN 19971E discovery image was only mareinally deeper than for a sinele tile of template-search epoch of the GOODS,", The combined depth of the template-search pair for the SN 1997ff discovery image was only marginally deeper than for a single tile of template-search epoch of the GOODS.
+ The SN itself. although likely magnilied by 0.3 mag (Benitezetal.2002).. was discovered well above the detection threshold for the tiny survey.," The SN itself, although likely magnified by 0.3 mag \citep{2002ApJ...577L...1B}, was discovered well above the detection threshold for the tiny survey."
+ The 5 month durration of the UDEF--UDEP search and the factor of 2 increase in area of ACS over WEPC? suggest that the number of such SNe Ia at z>1 found in the UDF should. exceed those found by Gilliland et al. (, The 5 month durration of the UDF+UDFP search and the factor of 2 increase in area of ACS over WFPC2 suggest that the number of such SNe Ia at $z>1$ found in the UDF should exceed those found by Gilliland et al. (
+1999) by a factor of ~5.,1999) by a factor of $\sim$ 5.
+ Yet. when dealing with such small. pencil-beam surveys. the survey vields can be subject more to chance (han expectations.," Yet, when dealing with such small, pencil-beam surveys, the survey yields can be subject more to chance than expectations."
+ The low munber of expected SNe serve more as a probability [or [finding SNe. governed by Poisson statistics.," The low number of expected SNe serve more as a probability for finding SNe, governed by Poisson statistics."
+ In (he previous section we concluded (hat our searches of (he UDF. UDFP. and IRUDE had failed to vield any SNe Ia ad z>1.4.," In the previous section we concluded that our searches of the UDF, UDFP, and IRUDF had failed to yield any SNe Ia at $z>1.4$."
+ To test the significance of this null result. we have simulated the vield expected in the multiple campaigns (image stacks) of the survey based on the best-fit δις) determined in SOL.," To test the significance of this null result, we have simulated the yield expected in the multiple campaigns (image stacks) of the survey based on the best-fit $\mathcal R_{Ia}(z)$ determined in S04."
+" As in SOL. the number of expected SNe per redshift interval was determined by: where /. is the ""control time’ probability function. which takes into account the survey efficienev with difference magnitude (described in 82). O is the area surveved. and AW(2) "," As in S04, the number of expected SNe per redshift interval was determined by: where $t_c$ is the “control time” probability function, which takes into account the survey efficiency with difference magnitude (described in 2), $\Theta$ is the area surveyed, and $\Delta V(z)$ "
+hat it should. be considered when discussing the C-chwarl oblem.,that it should be considered when discussing the G-dwarf problem.
+ This raises the question of the exact contribution of the thick disc at Fe/HI]«-0.4., This raises the question of the exact contribution of the thick disc at $<$ -0.4.
+ Wyse Gilmore (195 nave contended that the important metallicity tail below his limit (20% )) is likely due to contamination by the thin disc., Wyse Gilmore (1995) have contended that the important metallicity tail below this limit $>$ ) is likely due to contamination by the thin disc.
+ “Phis problem is then connected to another., This problem is then connected to another.
+ ΙΓ the thin disc contribution at Fe/L]«-0.4 is as important as envisaged ov Wvse Cilmore (1995). it leads to the conclusion that he SER must have been very efficient at early times in the disc.," If the thin disc contribution at $<$ -0.4 is as important as envisaged by Wyse Gilmore (1995), it leads to the conclusion that the SFR must have been very efficient at early times in the disc."
+ “Phis conclusion is reinforced by the apparent rapid rise of metallicity in the age-metallicity relation (Scullyοἱal.1907) - leaving even less time to produce the material at Fefl]«e-04., This conclusion is reinforced by the apparent rapid rise of metallicity in the age-metallicity relation \shortcite{97SCU_EA} - leaving even less time to produce the material at $<$ -0.4.
+ As a matter of fact. it is a feature common to most models with infall that they use a decreasing SER. in accordance with the infell decay rate (Losi.1996)...," As a matter of fact, it is a feature common to most models with infall that they use a decreasing SFR, in accordance with the infall decay rate \shortcite{96TOS}."
+ In contrast. all recent determinations of the SER history of the disc point to a constant. or even slightly increasing SER (Ιανwoodetal..1997:Binneyct2000).," In contrast, all recent determinations of the SFR history of the disc point to a constant, or even slightly increasing SFR \cite{97HAY_EA2,00BIN_EA}."
+. We show here that these apparent contradictions are mostly an οσοι of pre-Llipparcos data. and that the metallicity tailat Fe/L]«- 0.4 is drastically reclucecl with Llipparcos parallaxes.," We show here that these apparent contradictions are mostly an effect of pre-Hipparcos data, and that the metallicity tail at $<$ -0.4 is drastically reduced with Hipparcos parallaxes."
+ Third. while chemical evolution models have increase considerably in complexity. over. the last decade. the sophisticated: solutions cnvisagecl to solve. the ανα problem! (such as infall. variable initial mass function. ete). are still essentially adhoc. ancl direct. observationa evidence favoring one or another alternative have remain extremely elusive.," Third, while chemical evolution models have increased considerably in complexity over the last decade, the sophisticated solutions envisaged to solve the `G-dwarf problem' (such as infall, variable initial mass function, etc), are still essentially adhoc, and direct observational evidence favoring one or another alternative have remain extremely elusive."
+ Infall models have focused. most. of the ellorts of galactic chemical evolution studies in recent vears. while relatively little attention has been devote to other solutions.," Infall models have focused most of the efforts of galactic chemical evolution studies in recent years, while relatively little attention has been devoted to other solutions."
+ Since the Ce-dwarf problem remains the main (indirect) support. for long time-scale gas acerction hy the Milkv Way disc. it is important to check the failure (or success) of no-infall moclels.," Since the G-dwarf problem remains the main (indirect) support for long time-scale gas accretion by the Milky Way disc, it is important to check the failure (or success) of no-infall models."
+ Also. when considering scenarios where the thick disc is envisaged as a genuine galactic population (as is the case here). it must be born in mind that no-infall models may come as a more natural alternative.," Also, when considering scenarios where the thick disc is envisaged as a genuine galactic population (as is the case here), it must be born in mind that no-infall models may come as a more natural alternative."
+ The existence of the thick disc. if considered. as cogenetic. implies that a substantial stellar disc must have been in place at early times (Vvse. 2000 astro-ph/0012270).," The existence of the thick disc, if considered as cogenetic, implies that a substantial stellar disc must have been in place at early times (Wyse, 2000 astro-ph/0012270)."
+ As remarked by Wyse. this suggestion. could. linc some support or invalidation in a more general context. from studies of high redshift ealactic discs (Brinchmann&Ellis. 2000).," As remarked by Wyse, this suggestion could find some support or invalidation in a more general context, from studies of high redshift galactic discs \shortcite{00BRI_EA}."
+. Finally. another incentive for this work has been the (naive) realisation by the author that while models. of chemical evolution predict. distributions of stellar as a function of metallicity. previous studies of the dwarf metallicity distributions have only been able {ο give gaalistics of as a function of metallicity.," Finally, another incentive for this work has been the (naive) realisation by the author that while models of chemical evolution predict distributions of stellar as a function of metallicity, previous studies of the dwarf metallicity distributions have only been able to give statistics of as a function of metallicity."
+ The accurate positioning of the stars in the LR diagram allowed: by Hipparcos parallaxes permits the conversion from (colour. magnitude) to masses.," The accurate positioning of the stars in the HR diagram allowed by $Hipparcos$ parallaxes permits the conversion from (colour, magnitude) to masses."
+ Ht will be demonstrated that. while the conversion itself has a relatively minor elfect. the colour limits of the sample bias the distribution bv underestimating the contribution of solar-metallicity and metal-rich stars.," It will be demonstrated that, while the conversion itself has a relatively minor effect, the colour limits of the sample bias the distribution by underestimating the contribution of solar-metallicity and metal-rich stars."
+ The remainder of the paper is divided into 5 parts., The remainder of the paper is divided into 5 parts.
+ In the following section. we present the calibrations used to estimate metallicity from Geneva. and Strómmgren photometry.," In the following section, we present the calibrations used to estimate metallicity from Geneva and Strömmgren photometry."
+ After selecting a basic sample from considerations of completeness and. available photometry. we select long-lived stars ancl discuss the biases introcucee by our selection. procedure.," After selecting a basic sample from considerations of completeness and available photometry, we select long-lived stars and discuss the biases introduced by our selection procedure."
+ A first estimate of a correcte metallicity. cistribution of star frequency. distribution is given., A first estimate of a corrected metallicity distribution of star frequency distribution is given.
+ This distribution. is then compared το previous works in section 3. where we also discuss the problem of biases in other samples.," This distribution is then compared to previous works in section 3, where we also discuss the problem of biases in other samples."
+ We evaluate the proportion of Merayv cmitter stars in the sample., We evaluate the proportion of X-ray emitter stars in the sample.
+ Since X-ray corona activity is usually considered as a tracer of voung stars. we cliscuss the significance of the high percentage of voung star candidates in the sample.," Since X-ray coronal activity is usually considered as a tracer of young stars, we discuss the significance of the high percentage of young star candidates in the sample."
+ In section ?? we conver our να metallicity distribution to a mass metallicity distribution., In section \ref{finaldis} we convert our dwarf metallicity distribution to a mass metallicity distribution.
+ A proper correction for the mass bias ancl a final metallicity. distribution are egiven. for both iron. ane oxvgen abundances., A proper correction for the mass bias and a final metallicity distribution are given for both iron and oxygen abundances.
+ In. section 27... we brielly review the observed: parameters of the dise that enter the SCB model. and. following Occanis Razor. we explore how the Simple model fits other local constraints of chemical evolution.," In section \ref{models}, we briefly review the observed parameters of the disc that enter the SCB model, and following Occam's Razor, we explore how the Simple model fits other local constraints of chemical evolution."
+ We conclude in the last section., We conclude in the last section.
+ The solar neighbourhood. as sampled. by the. Hipparcos Catalogue. is. considered. to be essentially. complete to Al=8.5 [or stars with w—ο mas (Jahreiss&\Wie-len. 1997).. and contains 959 entries within these limits.," The solar neighbourhood, as sampled by the Hipparcos Catalogue, is considered to be essentially complete to $M_v$ =8.5 for stars with $\pi>$ mas \cite{97JAH_EA}, and contains 959 entries within these limits."
+ At Af.=8.5. main sequence stars have. DV colour between 1.3 anc 1.5.," At $M_v$ =8.5, main sequence stars have $B-V$ colour between 1.3 and 1.5."
+ Present day. available photometric metallicity calibrations are limited to. blace colour., Present day available photometric metallicity calibrations are limited to bluer colour.
+ The Geneva photometric system. is probably the best. suited system in that regard. since the calibration is available down to Ds 1120.65 (Cirenon..1978).. corresponding to D V=1.05 approximately.," The Geneva photometric system is probably the best suited system in that regard, since the calibration is available down to $B_2-V_1$ =0.65 \cite{78GRE}, corresponding to $B-V$ =1.05 approximately."
+ Xecording to the available data. the main sequence of the ealactic globular cluster M92 (Fefll]z-2) passes through the point D-V—1.00.. M —8.0 (Stetson&Παν».1988).," According to the available data, the main sequence of the galactic globular cluster M92 $\approx$ -2) passes through the point B-V=1.00, $_v$ =8.0 \shortcite{88STE_EA}."
+. This means that the limit at AL.=8.5 is unlikely to produce any significant bias against low metallicity stars in the sanple., This means that the limit at $_v$ =8.5 is unlikely to produce any significant bias against low metallicity stars in the sample.
+ Within these limits. the saniple contains 681 entries.," Within these limits, the sample contains 681 entries."
+" If we further select stars with M, 23.5 and οV 0.25 and which are not Dagued as either G. O. S. or V suspected binaries. we find 475 entries. of which 82 are [lagued. 7€ component solutions in the main catalogue."," If we further select stars with $_v>$ 3.5 and $B-V>$ 0.25 and which are not flagued as either G, O, S, or V suspected binaries, we find 475 entries, of which 82 are flagued “C” component solutions in the main catalogue."
+ Lor 27 of these S2 systems. no Geneva photometry is available [rom the General Catalogue of Photometric Data CMermilliodetab. 1997).," For 27 of these 82 systems, no Geneva photometry is available from the General Catalogue of Photometric Data \cite{97MER_EA1}."
+. Those πο component solutions for which photometry is available are usually not separated or simply not detected as binaries in the GCPD data base (except for 3 systems. for which photometry for the primary is available).," Those “C” component solutions for which photometry is available are usually not separated or simply not detected as binaries in the GCPD data base (except for 3 systems, for which photometry for the primary is available)."
+ We decided therefore to exclude entries with the μας 7€7 from further consideration. keeping in mind that we are," We decided therefore to exclude entries with the flag “C” from further consideration, keeping in mind that we are"
+ , 
+Thus. our static inodel iuherits iuauv observational xoperties of the full dvuauic simulation.,"Thus, our static model inherits many observational properties of the full dynamic simulation."
+ The model reproduces the C-band brightening iu the MBP ceutre., The model reproduces the G-band brightening in the MBP centre.
+ Stokes-I and -V profiles for the 6302A Fel Ine are comparable to solar observatious in both the naguctised region aud in the region corresponding to the 101-1122cuoetic solar granulation., Stokes-I and -V profiles for the $6302\mathrm{\AA}$ FeI line are comparable to solar observations in both the magnetised region and in the region corresponding to the non-magnetic solar granulation.
+ The continua formation evel in the centre of the MBP is locates iu the regiou of stronelv uaenetised plasma. where the Alfvéóun speed is ereater than that of sound.," The continuum formation level in the centre of the MBP is located in the region of strongly magnetised plasma, where the Alfvénn speed is greater than that of sound."
+ We have used the model to examine the observational consequences of sound wave propagation though the uaenetic field concentration of the AIBP., We have used the model to examine the observational consequences of sound wave propagation though the magnetic field concentration of the MBP.
+ This preliminary oevestigation shows that the variation of continu oeitensity js more xonounced in the ABP compared to je average eranule., This preliminary investigation shows that the variation of continuum intensity is more pronounced in the MBP compared to the average granule.
+ Using the radiative diagnostics with 1ο Stokes-V profile of the G302A FeI absorption line. we demoustrate the detectability of the maguetic field variatiou in the bright poiut aud show the appearance of 1e Stokes-V profile asvuuuetiy caused by the oscillations.," Using the radiative diagnostics with the Stokes-V profile of the $6302\mathrm{\AA}$ Fe I absorption line, we demonstrate the detectability of the magnetic field variation in the bright point and show the appearance of the Stokes-V profile asymmetry caused by the oscillations."
+ The radiative heating terii is uot included iuto the system of MIID equations used to perform the wave dynamics study., The radiative heating term is not included into the system of MHD equations used to perform the wave dynamics study.
+" The absence of the heating term, may result in reduced vertical radiative flux in the magnetic flux tube.", The absence of the heating term may result in reduced vertical radiative flux in the magnetic flux tube.
+ As a result. we have somewhat smaller continua aud €-baud intensity in the MBPs. compared to the full simulations aud the real Sun.," As a result, we have somewhat smaller continuum and G-band intensity in the MBPs, compared to the full simulations and the real Sun."
+ Ilowever. for wave dvuaimies the effect of the radiative term dis expected to be small.," However, for wave dynamics the effect of the radiative term is expected to be small."
+ As it has recently been shown by YellesChaoucheetal. (2009).. the magnetic flux. tubes in the solar photosphere are “reasonably well reproduced” by flux tube approsimation with no radiative teri. included.," As it has recently been shown by \citet{yelles}, the magnetic flux tubes in the solar photosphere are ""reasonably well reproduced"" by a thin flux tube approximation with no radiative term included."
+ Future imwvestiseatious8 will examine the respouse to sources at differen locations and with a range of frequencies., Future investigations will examine the response to sources at different locations and with a range of frequencies.
+ The behaviour aud observational signatures of Alfvénn waves in πα maeuetic clemenuts is particularly Huportant when used as a diagnostic for solar plasma paraueters. and for understaundius the enerewv transport to the corona.," The behaviour and observational signatures of Alfvénn waves in small magnetic elements is particularly important when used as a diagnostic for solar plasma parameters, and for understanding the energy transport to the corona."
+ Thus we plan to exteud the model to three dimensions., Thus we plan to extend the model to three dimensions.
+ This work has been supported bv the UK Science aud Technology Facilities Council (STFC)., This work has been supported by the UK Science and Technology Facilities Council (STFC).
+ EPK is exateful to AWE Aldermaston for the award of a William Penney Fellowship., FPK is grateful to AWE Aldermaston for the award of a William Penney Fellowship.
+"would become enereelically important only [or injection flares significantly contributing al SEP<102LZ/(GD/G)T,] pe. where Lj; is the accretion disk luminosity in units of 10? eves ! and Ty=T/10.","would become energetically important only for injection flares significantly contributing at $x_0^{\rm ECD} \lesssim 10^{-2}
+\, L_{45}^{1/2} / [(B / {\rm G}) \, \Gamma_1]$ pc, where $L_{45}$ is the accretion disk luminosity in units of $10^{45}$ ergs $^{-1}$ and $\Gamma_1 = \Gamma / 10$."
+ For a recent discussion of this radiation field on the light curves see. e.g.. Dermer&Schlickeiser(2002).," For a recent discussion of this radiation field on the high-energy light curves see, e.g., \cite{ds02}."
+". The (me-averaged photon spectra Irom our simulations with varving BLA. Thomson depth are shown in Fig. ὃν,"," The time-averaged photon spectra from our simulations with varying BLR Thomson depth are shown in Fig. \ref{ext_intspectra},"
+ which clearly shows theemergence of the external. Compton (EC!) component al GeV 5-ray energies., which clearly shows theemergence of the external Compton (EC) component at GeV $\gamma$ -ray energies.
+" The impact of this additional emission component on the lower-frequency. emission is small as long; as its bolometric energy output is smaller than or comparable to the svnchrotron vf, flux.", The impact of this additional emission component on the lower-frequency emission is small as long as its bolometric energy output is smaller than or comparable to the synchrotron $\nu F_{\nu}$ flux.
+ Only when EC cooling becomes dominant over svnchrotron cooling. are (he effects on the svuchrotvon + SSC dominated portion of the spectrum (radio frequencies — MeV. 2-ravs) noticeable.," Only when EC cooling becomes dominant over synchrotron cooling, are the effects on the synchrotron + SSC dominated portion of the spectrum (radio frequencies – MeV $\gamma$ -rays) noticeable."
+" This is the case when the comoving energy densities of the magnetic field and the external photons. wi, and uL. respectively, become comparable."," This is the case when the comoving energy densities of the magnetic field and the external photons, $u'_{\rm B}$ and $u'_{\rm ext}$, respectively, become comparable."
+ From where ry» is the radius of the inner boundary of the DLR. in units of 0.2 pc we can estimate that this happens at 7jpigc0.1.," From — where $r_{0.2}$ is the radius of the inner boundary of the BLR in units of 0.2 pc --- we can estimate that this happens at $\tau_{\rm T, BLR} 
+\sim 0.1$."
+ Specifically. for Typir20.1. the effect of dominant. external-Couplon cooling results in a reduction of the time-averaged enission around the svnchrotron peak — leading to a spectral hardening of the svnchrotron emission —. and a suppression of the SSC emission.," Specifically, for $\tau_{\rm T, BLR} \gtrsim 0.1$, the effect of dominant external-Compton cooling results in a reduction of the time-averaged emission around the synchrotron peak — leading to a spectral hardening of the synchrotron emission —, and a suppression of the SSC emission."
+ The suppression of the low-frecquency svuchrotron emission can be explained as the combined effect of (wo causes: First. electrons al energies above the low-energy. eut-off are radiatively cooling on a time scale much shorter than the dvnamical one (see Eq.," The suppression of the low-frequency synchrotron emission can be explained as the combined effect of two causes: First, electrons at energies above the low-energy cut-off are radiatively cooling on a time scale much shorter than the dynamical one (see Eq."
+ 7 below)., \ref{B_decay} below).
+ Consequently. (he particle spectrum of electrons injected duriug the flare is rapidly depleted within less (han one dynamical lime scale [rom (he end of the flaring episode.," Consequently, the particle spectrum of high-energy electrons injected during the flare is rapidly depleted within less than one dynamical time scale from the end of the flaring episode."
+ However. alter this episode. we are still injecling electrons (although at à much smaller rate corresponding to Lyr(qu) into the blob. which continue to present a high-energv electron population.," However, after this episode, we are still injecting electrons (although at a much smaller rate corresponding to $L_{\rm inj}^{\rm qu}$ ) into the blob, which continue to present a high-energy electron population."
+ In the case of the extremely [ast cooling rate al Tppig1. this additional high-energv electron population leads to a significant. additional contribution of svnchrotron emission at intermediate energies bevond the down-shilted maximum energy of electrons injected during the flare ancl. consequently. to a hardening of the svnchrotron spectrum.," In the case of the extremely fast cooling rate at $\tau_{\rm T, BLR}
+\sim 1$, this additional high-energy electron population leads to a significant additional contribution of synchrotron emission at intermediate energies beyond the down-shifted maximum energy of electrons injected during the flare and, consequently, to a hardening of the synchrotron spectrum."
+ The second cause of spectral hardening of the low-frequency svnchrotron spectrum is related to our parametrization of (he magnetic [ield in terms of an equipartition parameter ej: For electron injection spectral indices q7 2. most of the energy in the electron population is carried by electrons near the low-energv cutoff.," The second cause of spectral hardening of the low-frequency synchrotron spectrum is related to our parametrization of the magnetic field in terms of an equipartition parameter $\epsilon_B$: For electron injection spectral indices $q > 2$ , most of the energy in the electron population is carried by electrons near the low-energy cutoff."
+ Consequently. the co-moving magnetic," Consequently, the co-moving magnetic"
+component at. | LO to | 13 km tlre ~ 3100 kms 1). with much weaker absorption extending from about 50 o |55 km ,"component at $+$ 10 to $+$ 13 km $^{-1}$ $v_{\rm LSR}$ $\sim$ $-$ 3 to 0 km $^{-1}$ ), with much weaker absorption extending from about $-$ 20 to $+$ 55 km $^{-1}$."
+Toward ID 72127. the absorption rom dominant ions of typically mildly. depleted: elements (c.g.. 11)) is concentrated in several components between about | 10 anc |20 km with hints of much weaker absorption at. lower ancl higher velocities.," Toward HD 72127A, the absorption from dominant ions of typically mildly depleted elements (e.g., ) is concentrated in several components between about $+$ 10 and $+$ 20 km $^{-1}$, with hints of much weaker absorption at lower and higher velocities."
+ The absorption rom dominant ions of typically more severely depleted elements (e.g.SilL. bel1)) also is strongest between | 10 and 120 km +. but. becomes increasingly. prominent. at roth lower and higher velocities for the species most alfected w depletion. ," The absorption from dominant ions of typically more severely depleted elements (e.g., ) also is strongest between $+$ 10 and $+$ 20 km $^{-1}$, but becomes increasingly prominent at both lower and higher velocities for the species most affected by depletion. –"
+which can be allected by both ionization and. depletion (e.g... Welty ct al.," which can be affected by both ionization and depletion (e.g., Welty et al."
+ 1996) shows moderately strong absorption over the entire velocity range., 1996) – shows moderately strong absorption over the entire velocity range.
+ Alulticomponent fits to observed. complex absorption-line profiles enable blended. lines ancl components to be disentangled. account explicitly for saturation cllects in the stronger components. anc allow spectra obtained at (somewhat) clillercnt resolutions to be. compared.," Multicomponent fits to observed complex absorption-line profiles enable blended lines and components to be disentangled, account explicitly for saturation effects in the stronger components, and allow spectra obtained at (somewhat) different resolutions to be compared."
+ The method of profile fitting was therefore used. to estimate column densities CN). line widths (6 ~ EWLIIM/1.665). and velocities (0) for the individual components cdiscernible in the spectra (e.g.. Hobbs et al.," The method of profile fitting was therefore used to estimate column densities $N$ ), line widths $b$ $\sim$ FWHM/1.665), and velocities $v$ ) for the individual components discernible in the spectra (e.g., Hobbs et al."
+ 1991: Welty. Hobbs. Morton 2003).," 1991; Welty, Hobbs, Morton 2003)."
+ For each species. the minimum number of components needed to achieve ‘satisfactory’ fits to the detected lines was adopted (given the resolution and S/N ratios characterizing he spectra and assuming that each component may be represented by a svnimetric Voigt. profile).," For each species, the minimum number of components needed to achieve `satisfactory' fits to the detected lines was adopted (given the resolution and S/N ratios characterizing the spectra and assuming that each component may be represented by a symmetric Voigt profile)."
+ Multiple lines or à given species (ee. the two lines of Call)) were fitted simultaneously.," Multiple lines for a given species (e.g., the two lines of ) were fitted simultaneously."
+ The structure determined lor was used ο estimate individual component column densities for rom the weak A3383 lines., The structure determined for was used to estimate individual component column densities for from the weak $\lambda$ 3383 lines.
+ The lines fromSill.Cril..Fert. ill. and seen in the GURS ECL-B spectra. were ittecl simultaneously — using a common velocity structure rut adjusting the component column densities (ancl some b-values) for each species.," The lines from, and seen in the GHRS ECH-B spectra were fitted simultaneously – using a common velocity structure but adjusting the component column densities (and some $b$ -values) for each species."
+ The adopted interstellar component parameters derived rom the UVES spectra. of LD 72127. and. HD. 72127D ave Listed in Table 2: the component parameters obtained rom the simultaneous fit to the GIRS ECII-B spectra of HD. 72127 are given in Table 3..., The adopted interstellar component parameters derived from the UVES spectra of HD 72127A and HD 72127B are listed in Table \ref{tab:comps}; the component parameters obtained from the simultaneous fit to the GHRS ECH-B spectra of HD 72127A are given in Table \ref{tab:compuv}.
+ For cach component in cach sight line. Table 2. gives the component number. he heliocentric velocity anc b-value (lor ancl Cail)). and the column densities ofNal.Cal.Call. ancl Tii.," For each component in each sight line, Table \ref{tab:comps} gives the component number, the heliocentric velocity and $b$ -value (for and ), and the column densities of, and ."
+ Component column densities or b-valucs in square brackets were fixed. but relatively well determined in the fits: values in parentheses also were fixed. but are less well determined.," Component column densities or $b$ -values in square brackets were fixed, but relatively well determined in the fits; values in parentheses also were fixed, but are less well determined."
+ ForNat. the column density of the strongest. component in each sight line was fixed at a value consistent with that derived. from the weak A8302 doublet: the parameters for the other components were derived. [rom fits to the much stronger D lines.," For, the column density of the strongest component in each sight line was fixed at a value consistent with that derived from the weak $\lambda$ 3302 doublet; the parameters for the other components were derived from fits to the much stronger D lines."
+ Independent fits to the profiles of the LD lines and the LH and Ex lines vielded rather similar component velocities (in both sight lines). and components seen in the two species are considered to be associated if the velocities agree within about 12 km," Independent fits to the profiles of the D lines and the H and K lines yielded rather similar component velocities (in both sight lines), and components seen in the two species are considered to be associated if the velocities agree within about 1–2 km $^{-1}$."
+ The Loo uncertainties given for JN. b. and v in ‘Tables 2 and 3. rellect only photon noise. and thus should be considered. as lower limits to the true uncertainties.," The 1 $\sigma$ uncertainties given for $N$, $b$, and $v$ in Tables \ref{tab:comps} and \ref{tab:compuv} reflect only photon noise, and thus should be considered as lower limits to the true uncertainties."
+ Consideration of the elfects of uncertainties in continuum placement (most significant for the stronger lines) ancl of allowing fixed parameters (usually b) to vary would increase the uncertainties in the various parameters (primarily IN)., Consideration of the effects of uncertainties in continuum placement (most significant for the stronger lines) and of allowing fixed parameters (usually $b$ ) to vary would increase the uncertainties in the various parameters (primarily $N$ ).
+ In view of the statistics of component width and separation obtained from higher resolution. higher S/N spectra of Galactic targets (Welty et al.," In view of the statistics of component width and separation obtained from higher resolution, higher S/N spectra of Galactic targets (Welty et al."
+" 1994. 1996: Welty Llobbs 2001). however. it ds likely that the ""true component structures are even more complex than those listed in the tables and that uncertainty may well be even more significant for the individual component. parameters."," 1994, 1996; Welty Hobbs 2001), however, it is likely that the `true' component structures are even more complex than those listed in the tables – and that uncertainty may well be even more significant for the individual component parameters."
+ Dilferences in the component. velocities derived for andCal... for example. may be due to component-to-component dillerences in relative abundances in an underlying more complex structure.," Differences in the component velocities derived for and, for example, may be due to component-to-component differences in relative abundances in an underlying more complex structure."
+ Lligher resolution spectra of toward HD 72127. clo exhibit. more complex structure (ig. Iz , Higher resolution spectra of toward HD 72127A do exhibit more complex structure (Fig. \ref{fig:naca}; ;
+Welty et al., Welty et al.
+ 1996)., 1996).
+ Dillerences between the component velocities derived from optical and. UV spectra of LD 72127A maw rellect. (in addition) possible temporal variations between 1992 and 2003., Differences between the component velocities derived from optical and UV spectra of HD 72127A may reflect (in addition) possible temporal variations between 1992 and 2003.
+ The total sight. line column. densities are given in Table 4.. with values for the main neutral components toward 283 Ori (Welty ct al.," The total sight line column densities are given in Table \ref{tab:coldens}, with values for the main neutral components toward 23 Ori (Welty et al."
+ L999b) and ς Oph (Welty οἱ al., 1999b) and $\zeta$ Oph (Welty et al.
+ 2003 and references therein) given for comparison., 2003 and references therein) given for comparison.
+" In all cases. the values derived from the profile fits are consistent with estimates orlimits based on integrating the ""apparent optical depth over the absorption lines: the total sight line column densities are in most cases fairly well determined."," In all cases, the values derived from the profile fits are consistent with estimates orlimits based on integrating the `apparent' optical depth over the absorption lines; the total sight line column densities are in most cases fairly well determined."
+ In view of the spectral types of LLD 72127X and HD 721273. it is likely that the Lyman-a absorption observed: toward the two stars (in archivalZU spectra) is an inseparable blend of stellar and interstellar components.," In view of the spectral types of HD 72127A and HD 72127B, it is likely that the $\alpha$ absorption observed toward the two stars (in archival spectra) is an inseparable blend of stellar and interstellar components."
+ Several indirect methods thus have been used to estimate the interstellar hvdrogen column densities toward the two stars: (1) The first makes use of the strong. nearly linear correlation observed: between the equivalent width of the cülfuse interstellar band (DIB) at 5780 aand ΑΟ) (eg. Herbie 1993. L995: Welty ct al.," Several indirect methods thus have been used to estimate the interstellar hydrogen column densities toward the two stars: (1) The first makes use of the strong, nearly linear correlation observed between the equivalent width of the diffuse interstellar band (DIB) at 5780 and $N$ ) (e.g., Herbig 1993, 1995; Welty et al."
+ 2006: D. €. York et abl.," 2006; D. G. York et al.,"
+ in preparation)., in preparation).
+ The equivalent widths estimated from our UVES spectra — M (5780) = 15+6 flor LID 72127X and M (5780) = 1255 {for LLD 12127B.— suggest. logIN (111))] of about 20.2. [or LID 72127 and about 20.1 for LID 72127DB. The observed Η (5780) are somewhat low. relative to the corresponding N(Nal)). however. suggesting that the 5780 DDID max be weaker than usual toward LIL) 7212TAB as has been observed. for example. toward some Galactic and Magellanic Clouds targets (Herbig 1993: Welty et al.," The equivalent widths estimated from our UVES spectra – $W$ (5780) = $\pm$ 6 for HD 72127A and $W$ (5780) = $\pm$ 5 for HD 72127B – suggest $N$ )] of about 20.2 for HD 72127A and about 20.1 for HD 72127B. The observed $W$ (5780) are somewhat low, relative to the corresponding $N$ ), however, suggesting that the 5780 DIB may be weaker than usual toward HD 72127AB – as has been observed, for example, toward some Galactic and Magellanic Clouds targets (Herbig 1993; Welty et al."
+ 2006)., 2006).
+ The ΔΕΟ} estimated. from W (5780) might thus be best viewed as lower limits. (, The $N$ ) estimated from $W$ (5780) might thus be best viewed as lower limits. (
+2) An estimate for ΔΙ) toward LD 72127. mav be obtained from the observed column clensity of. IL. as zinc is generally only very. mildly depleted: in. diffuse interstellar clouds.,"2) An estimate for $N$ (H) toward HD 72127A may be obtained from the observed column density of , as zinc is generally only very mildly depleted in diffuse interstellar clouds."
+ H£ the depletion of zinc D(Zn) = 0.2, If the depletion of zinc D(Zn) =$-$ 0.2
+question arises: are flare-induced variations compatible with our observations?,question arises: are flare-induced variations compatible with our observations?
+" ?? carried out a high temporal resolution spectroscopic monitoring of the stars V1054 Oph and AD Leo, respectively."," \cite{2004Ap&SS.292..697C,2006A&A...452..987C} carried out a high temporal resolution spectroscopic monitoring of the stars V1054 Oph and AD Leo, respectively."
+" Studying the temporal evolution of the Balmer series and H K lines, they detected frequent (>0.71 flares/hour) short and weak non white-light flares."," Studying the temporal evolution of the Balmer series and H K lines, they detected frequent $> 0.71$ flares/hour) short and weak non white-light flares."
+ Some of these events lasted as little as 14 min., Some of these events lasted as little as 14 min.
+ Additional observations by ? suggest that this kind of frequent short flare is of common occurence on UV Ceti-type flare stars., Additional observations by \cite{2005ESASP.560..491C} suggest that this kind of frequent short flare is of common occurence on UV Ceti-type flare stars.
+" In the present work, the observed spectra were obtained using typical exposure times ~25 min."," In the present work, the observed spectra were obtained using typical exposure times $\sim 25$ min."
+" Thus, one or two of this kind of flare events probably occurred on the flare stars during their observation."," Thus, one or two of this kind of flare events probably occurred on the flare stars during their observation."
+" However, even if this were the case, it would not explain the different behaviour found in the flux—flux relationships between the stars in the upper branch and the other ones."," However, even if this were the case, it would not explain the different behaviour found in the flux–flux relationships between the stars in the upper branch and the other ones."
+" In fact, ? derived (for AD Leo) a value of the order of 1030 eerg for the energy released in the Ha line by flare event of the above type with a duration of 30 min."," In fact, \cite{2006A&A...452..987C} derived (for AD Leo) a value of the order of $10^{30}$ erg for the energy released in the $H\alpha$ line by a flare event of the above type with a duration of 30 min."
+ This impliesa that flares of this type would contribute with a mean Ha surface flux of the order of 10° eerg cm? s~!., This implies that flares of this type would contribute with a mean $H\alpha$ surface flux of the order of $10^5$ erg $^{-2}$ $^{-1}$.
+" Thus, one, or even two, of these events taking place on some of the target stars during our observations would only introduce a small increase in their measured line fluxes, which would not explain the whole observed ‘deviation’ of the stars in the upper branch from the general trend followed by the remainder."," Thus, one, or even two, of these events taking place on some of the target stars during our observations would only introduce a small increase in their measured line fluxes, which would not explain the whole observed `deviation' of the stars in the upper branch from the general trend followed by the remainder."
+" Moreover, the clear trend followed by the stars of the upper branch in the flux—flux relationship cannot be the result of events of this"," Moreover, the clear trend followed by the stars of the upper branch in the flux–flux relationship cannot be the result of events of this"
+It has long been known that the spectral energy distributions of bright early-type galaxies have a pronounced upturn at wavelengths shortward of (Code Welch 1979: Brown et al.,It has long been known that the spectral energy distributions of bright early-type galaxies have a pronounced upturn at wavelengths shortward of (Code Welch 1979; Brown et al.
+ 1997: O'Connell et al., 1997; O'Connell et al.
+ 1992: O'Connell 1999)., 1992; O'Connell 1999).
+ This far-UV excess was unexpected as early-type galaxies were believed to contain only old and cold stellar populations., This far-UV excess was unexpected as early-type galaxies were believed to contain only old and cold stellar populations.
+ From the shape of the spectral energy distribution in the UV. it appears that the temperature of the stars which give rise to the excess has a narrow range around 20.00043000K (Brown et al.," From the shape of the spectral energy distribution in the UV, it appears that the temperature of the stars which give rise to the excess has a narrow range around $\pm$ 3000K (Brown et al."
+ 1997) which rules out upper main sequence stars as an origin., 1997) which rules out upper main sequence stars as an origin.
+ It is believed to be caused instead by hot subdwarfs. also known as extreme horizontal branch CEHB) stars (see Yi 2008 for a review).," It is believed to be caused instead by hot subdwarfs, also known as extreme horizontal branch (EHB) stars (see Yi 2008 for a review)."
+ A feature of the UV upturn is that it is strongest in the centres of bright elliptical galaxies (O'Connell et al., A feature of the UV upturn is that it is strongest in the centres of bright elliptical galaxies (O'Connell et al.
+ 1992: 1999)., 1992; 1999).
+ The positive UV colour gradients have been interpreted as abundance gradients. but it is still unclear which abundance or abundances drive these gradients.," The positive UV colour gradients have been interpreted as abundance gradients, but it is still unclear which abundance or abundances drive these gradients."
+ Such stars are found in globular clusters. and are generally considered to be helium-core burning stars with extremely thin hydrogen. envelopes (« 0.02A7.).," Such stars are found in globular clusters, and are generally considered to be helium-core burning stars with extremely thin hydrogen envelopes $<0.02M_\odot$ )."
+ There are a number of competing models for their origin., There are a number of competing models for their origin.
+ Han et (2002. 2003. 2007) and Han (2008). propose a binary model. in which a star loses all of its envelope near the tip of the red-giant branch (RGB) by mass ransfer to a companion star or ejection in a common-envelope shase. or where two helium white dwarfs merge with a combined mass larger than ~0.35A7..," Han et (2002, 2003, 2007) and Han (2008), propose a binary model, in which a star loses all of its envelope near the tip of the red-giant branch (RGB) by mass transfer to a companion star or ejection in a common-envelope phase, or where two helium white dwarfs merge with a combined mass larger than $\sim 0.35\,M_{\odot}$."
+ In each ease. the remnant star ignites yelium and becomes a hot subdwarf.," In each case, the remnant star ignites helium and becomes a hot subdwarf."
+ A number of single-star mechanisms have been proposed to explain the EHB populations of ellipticals., A number of single-star mechanisms have been proposed to explain the EHB populations of ellipticals.
+ Park Lee (1997) tind that ~20% of the populations are extremely old and metal-»oor., Park Lee (1997) find that $\sim$ of the populations are extremely old and metal-poor.
+ This appears to contradict the observed positive correlation between FUV excess and A/g» line strength (Burstein et al., This appears to contradict the observed positive correlation between FUV excess and $Mg_2$ line strength (Burstein et al.
+ 1988)., 1988).
+" lost attempts to explain the FUV excess. using single stars. are based upon the idea of Greggio Renzini (1990) that extremely ow-mass metal-rich HB stars might skip the AGB phase altogether CAGB-Manquee stars"") and move directly to the EHB."," Most attempts to explain the FUV excess, using single stars, are based upon the idea of Greggio Renzini (1990) that extremely low-mass metal-rich HB stars might skip the AGB phase altogether (“AGB-Manqúee stars”) and move directly to the EHB."
+ Brown et al. (, Brown et al. (
+1997). using Hopkins Ultraviolet Telescope (HUT) spectra found that the spectral energy distributions of the UV excess in the cores of bright ellipticals required both enhanced,"1997), using Hopkins Ultraviolet Telescope (HUT) spectra found that the spectral energy distributions of the UV excess in the cores of bright ellipticals required both enhanced"
+Iu summary. applying the FLD approximation involves computing the flux limiter from the value of Hin each zone. aud then using equations (10)) to (11)) to obtain the radiatiou flux ancl radiation pressure from tle radiation energy. deusity.,"In summary, applying the FLD approximation involves computing the flux limiter from the value of $R$ in each zone, and then using equations \ref{eqn:fick}) ) to \ref{eqn:flambda}) ) to obtain the radiation flux and radiation pressure from the radiation energy density."
+ As with the MHD equations in ZEUS-2D. the dynamical equations for the radiation are operator-split into source and trausport terms.," As with the MHD equations in ZEUS-2D, the dynamical equations for the radiation are operator-split into source and transport terms."
+ In the source step. the radiation aud material euergy deusities are updated using finite difference approximations to aud separately. and Iu the transport step. an integral formulation is used to generate a conservative dillereuciug scheme for the advection terms. where vg is au arbitrary coordinate velocity which allows grid motion.," In the source step, the radiation and material energy densities are updated using finite difference approximations to and separately, and In the transport step, an integral formulation is used to generate a conservative differencing scheme for the advection terms, where ${\bf v_g}$ is an arbitrary coordinate velocity which allows grid motion."
+ Equations (20)) to (23)) are evolved in the radiation nodule. as is the source terim due to radiative acceleration iu the eas momentum equation (2)).," Equations \ref{eqn:erdivflux}) ) to \ref{eqn:ertransport}) ) are evolved in the radiation module, as is the source term due to radiative acceleration in the gas momentum equation \ref{eqn:gasmomentum}) )."
+ The rest of tlie set of coupled equatious is solved using the existing ivclrodsuzunic algorithuus., The rest of the set of coupled equations is solved using the existing hydrodynamic algorithms.
+ Centering of the radiatiou variables on the hydrodyuaimie grid is described iu 'e[sec:centering.. update of the radiatiou flux term iu 'e[sec:delflux: and the remaining source terms in 'e[sec:ieeriut.. treatineut of the compression terius in the case of strictly scattering opacity in 'e[sec:ieerdv.. aud the transport step in 'e[secitrausport..," Centering of the radiation variables on the hydrodynamic grid is described in \\ref{sec:centering}, update of the radiation flux term in \\ref{sec:delflux} and the remaining source terms in \\ref{sec:eerint}, treatment of the compression terms in the case of strictly scattering opacity in \\ref{sec:eerdv}, and the transport step in \\ref{sec:transport}."
+ Stability of the method aud the choice of timestep are discussed iu 'e[secxlt.., Stability of the method and the choice of timestep are discussed in \\ref{sec:dt}.
+ The depeucent variables for the radiation are cliscretized outo the computational mesh as iu ZEUS-2D (Stone Norman 1992a. herealter SN).," The dependent variables for the radiation are discretized onto the computational mesh as in ZEUS-2D (Stone Norman 1992a, hereafter SN)."
+ Scalars. and teusors of even rank. are while tensors of odd rank are [ace-centered.," Scalars, and tensors of even rank, are zone-centered, while tensors of odd rank are face-centered."
+ The mesh is labeled by coordinate vectors ory , The mesh is labeled by coordinate vectors $x_1$ 
+2 (GSC-2: Lasker et 22008) are saturated in the stacked images. we used shorter exposures (1015 s) of the fields taken with the same imstrumeut configurations as those in Tab.,"2 (GSC-2; Lasker et 2008) are saturated in the stacked images, we used shorter exposures (10–15 s) of the fields taken with the same instrument configurations as those in Tab."
+ 1 and available in he aarchive., 1 and available in the archive.
+ We measured the star centroids through Caussian fitting using the Cawaphical Astronomy and huage Analysis (GALA) aud used the code to compute the pixclto-sky coordinate transformation through an high-order polvnonuüal. which accounts for the CCD distortions.," We measured the star centroids through Gaussian fitting using the Graphical Astronomy and Image Analysis (GAIA) and used the code to compute the pixel-to-sky coordinate transformation through an high-order polynomial, which accounts for the CCD distortions."
+ For both pulsar fields. he rms of the astrometric fits was σενO71 in the radia direction.," For both pulsar fields, the rms of the astrometric fits was $\sigma_{\rm r} \sim 0\farcs1$ in the radial direction."
+ To this value we added in quadrature the τουΤατν Or=OL of the registration of the liuagee on the CGSC2 reference frames. σι=VIE Aueqses. where dose=073 is the Weal positional error of the GSC2 coordinates and Wy is the uunuber of stars used to compute the astrometric solution (Lattanzi ct 11997).," To this value we added in quadrature the uncertainty $\sigma_{\rm tr}=0\farcs1$ of the registration of the image on the GSC2 reference frames, $\sigma_{\rm tr}= \sqrt(3/N_{s}) \sigma_{\rm GSC2}$ , where $\sigma_{\rm GSC2}=0\farcs3$ is the mean positional error of the GSC2 coordinates and $N_{s}$ is the number of stars used to compute the astrometric solution (Lattanzi et 1997)."
+" After accounting for the ~0715 accuracy of the luk of the CSC2 coordinates to the Tuternational Celestial Reference Frame (ICRF). we thus estimate that the overall (10) uncertainty of our astrometry is à,~072,"," After accounting for the $\sim 0\farcs15$ accuracy of the link of the GSC2 coordinates to the International Celestial Reference Frame (ICRF), we thus estimate that the overall $1 \sigma$ ) uncertainty of our astrometry is $\delta_{r}\sim0\farcs2$."
+ Ax a reference for our astrometry. we started from the most receuth published radio positious of thle two pulsars.," As a reference for our astrometry, we started from the most recently published radio positions of the two pulsars."
+ Moreover. they both have precise X-ray positions obtained withChandra.," Moreover, they both have precise X-ray positions obtained with."
+.. The published radio aud N-ray coordinates of 6129 and 5832 are sumunarised in Table 2., The published radio and X-ray coordinates of $-$ 6429 and $-$ 5832 are summarised in Table 2.
+ We note that the pposition of 6129 (Zavlin 2007) is quite different from that derived from the racdio-iterferometry observations performed with the Australia Telescope Compact Arrav (ATCA) by Camilo et (2001)., We note that the position of $-$ 6429 (Zavlin 2007) is quite different from that derived from the radio-interferometry observations performed with the Australia Telescope Compact Array (ATCA) by Camilo et (2004).
+ Since this apparent incousistency laupers the correct pulsar localisation ou the images. heuce the search for its optical counterpart. we independently checked all. the available sets of coordinates.," Since this apparent inconsistency hampers the correct pulsar localisation on the images, hence the search for its optical counterpart, we independently checked all the available sets of coordinates."
+ For 6129. multiple oobservations are available.," For $-$ 6429, multiple observations are available."
+ Thus. we could use them to check the internal cousisteney of the astrometry.," Thus, we could use them to check the internal consistency of the astrometry."
+ Firstly. we re-coniputed tlhe N-ray position from the analysis of the two. independent IRC observations (Esposito et 22007: 22007) and we found values consistent with that published in ((2007).," Firstly, we re-computed the X-ray position from the analysis of the two, independent /HRC observations (Esposito et 2007; 2007) and we found values consistent with that published in (2007)."
+ The pulsar has been receutlv observed also with the ACIS instriuncut (ODS-ID-108580: ALID—55112.28)., The pulsar has been recently observed also with the ACIS instrument (OBS-ID=10880; MJD=55112.28).
+ We downloaded the data from the public aarchive aud we fone that the pulsu coordinates are also consistent. within the nominal ppositiou uncertainty of 0769 confidence level). with those obtained with the WRC.," We downloaded the data from the public archive and we found that the pulsar coordinates are also consistent, within the nominal position uncertainty of $0\farcs6$ confidence level), with those obtained with the HRC."
+" We then averaged the three sets of coordinates, which vields: a=13557902260 and 6=6172929780 with a radial error of 07338."," We then averaged the three sets of coordinates, which yields: $\alpha =13^{\rm h} 57^{\rm m} 02\fs60$ and $\delta = -64^\circ 29\arcmin 29\farcs80$ with a radial error of $\sim$ 38."
+ This iniplies a difference of 172340%1 between the aand racdio-interterometry position., This implies a difference of $1\farcs23 \pm 0\farcs4$ between the and radio-interferometry position.
+ Then. we verified the absolute astrometry by matching the positions of serendipitous X-ray sources detected in the {ACTS image with those of their putative counterparts detected in the 2\TASS catalogue (Skrutskie et 22006).," Then, we verified the absolute astrometry by matching the positions of serendipitous X-ray sources detected in the /ACIS image with those of their putative counterparts detected in the 2MASS catalogue (Skrutskie et 2006)."
+ Since the accuracy of the aastrometrv rapidly degrades at laree off-axis angles. we restricted our search to within ffroii the centre of the ACTS fieldofview.," Since the accuracy of the astrometry rapidly degrades at large off-axis angles, we restricted our search to within from the centre of the ACIS field–of–view."
+ In particular. we found 2MASS matches for { N-rav sources in the 6129 field.," In particular, we found 2MASS matches for 4 X-ray sources in the $-$ 6429 field."
+ Although the unecrtainty ou the computed bore-sight correction does uot vield any significant improvement ou the absolute, Although the uncertainty on the computed bore-sight correction does not yield any significant improvement on the absolute
+Although location on the CMD was not a criterion. all of our variable stars were found near the expected location of short-period Cepheids — brighter and bluer (han the red chup.,"Although location on the CMD was not a criterion, all of our variable stars were found near the expected location of short-period Cepheids – brighter and bluer than the red clump."
+ Figure 4 shows the Sextans A CMD. with the variable stars highlighted.," Figure \ref{figVCMD} shows the Sextans A CMD, with the variable stars highlighted."
+ The apparent horizontal extensions of the Cepheid population at /2223.9 are artifacts of poor phase coverage in F555W.De.Ln explained in the next section.," The apparent horizontal extensions of the Cepheid population at $I \approx 23.9$ are artifacts of poor phase coverage in F555W, explained in the next section."
+ lt should be emphasized that our detection efficiency was not10056., It should be emphasized that our detection efficiency was not.
+. Thus we cannot rule out the existence of Population I] Cepheids. nor are we confident that the non-variable stars [alling within the instabilitv stvip are not. in fact. variables Chat were not detecte.," Thus we cannot rule out the existence of Population II Cepheids, nor are we confident that the “non-variable” stars falling within the instability strip are not, in fact, variables that were not detected."
+ Additionally. while the CMD from combined photometry is complete to £=27. the sinele epoch signal-to-noise of stars below £=25 is such that detection of fainter periodic variables. such as RR Lyraes. would have been extremely diflicull.," Additionally, while the CMD from combined photometry is complete to $I = 27$, the single epoch signal-to-noise of stars below $I=25$ is such that detection of fainter periodic variables, such as RR Lyraes, would have been extremely difficult."
+ For example. the of our individual epoch photometry at the expected location of RR Lyraes is ~5 in Zi thus a variable at that magnitude would have to have an rms variability of ~0.5 magnitudes in / (o pass all of our variability tests.," For example, the signal-to-noise of our individual epoch photometry at the expected location of RR Lyraes is $\sim 5$ in $I$; thus a variable at that magnitude would have to have an rms variability of $\sim 0.5$ magnitudes in $I$ to pass all of our variability tests."
+ since (he primary goal of these observations was the obtaining of deep photometry. the program was nol carried out in a wav (hat would maximize the phase coverage.," Since the primary goal of these observations was the obtaining of deep photometry, the program was not carried out in a way that would maximize the phase coverage."
+ Thus. a variable star wilh one period max have excellent coverage of its light curve while one with another period may have poor coverage.," Thus, a variable star with one period may have excellent coverage of its light curve while one with another period may have poor coverage."
+" Figure 5 shows the maxinunm gap in our phase coverage as a function of period. assuming that the star was measured at every epoch,"," Figure \ref{figdphase} shows the maximum gap in our phase coverage as a function of period, assuming that the star was measured at every epoch."
+ While the ES14W observations consist of 16 epochs and thus give the better coverage. we will have poor phase coverage of anv variable whose period is between 0.5 and 1.1 days.," While the F814W observations consist of 16 epochs and thus give the better coverage, we will have poor phase coverage of any variable whose period is between 0.8 and 1.1 days."
+ For such a star. all 16 FSI4W observations would cover only about of the light curve. while the 8 F555W observations cover about50%.," For such a star, all 16 F814W observations would cover only about of the light curve, while the 8 F555W observations cover about."
+.. Thus our detection efficiency will suffer. since such a star may not show sufficient variability over the phase range imcluded in our observations {ο (rigger our V? flag.," Thus our detection efficiency will suffer, since such a star may not show sufficient variability over the phase range included in our observations to trigger our $\chi^2$ flag."
+ Additionally. if such a variable is recovered. the absence of half the light curve prevenis an accurate mean magnitude determination and creates ambiguity. between periods of a day and half a day.," Additionally, if such a variable is recovered, the absence of half the light curve prevents an accurate mean magnitude determination and creates ambiguity between periods of a day and half a day."
+ Fortunately. the color of such a star should be measured acceplably since both light curves would have their gaps at close to the same phases.," Fortunately, the color of such a star should be measured acceptably since both light curves would have their gaps at close to the same phases."
+ In addition to detection efficiency. poor phase coverage affects our period and mean magnitude determinations.," In addition to detection efficiency, poor phase coverage affects our period and mean magnitude determinations."
+ Eight stars could have been fit with either one-day. or hall-cay light curves: fortunately all such stars could be identified using their locations in the diagram., Eight stars could have been fit with either one-day or half-day light curves; fortunately all such stars could be identified using their locations in the period-luminosity diagram.
+ The mean magnitude errors are more problematic. since the lack of half the light curve can bias the ealeulation in either direction if the star is observed only near," The mean magnitude errors are more problematic, since the lack of half the light curve can bias the calculation in either direction if the star is observed only near"
+is still greater (han the threshold. active deblending continues untill (he program reaches the user defined limit on Νε.,"is still greater than the threshold, active deblending continues untill the program reaches the user defined limit on $N_p$."
+ We performed simulations to test (he limits of passive and active deblending., We performed simulations to test the limits of passive and active deblending.
+ To test passive deblending a set of point sources wilh the same flux was added (ο a smooth background image., To test passive deblending a set of point sources with the same flux was added to a smooth background image.
+ The average separation between (he adjacent sources was approximately (he EWIIM (Pull width half maxinum) of the PRE., The average separation between the adjacent sources was approximately the FWHM (full width half maximum) of the PRF.
+ Two examples of such a cluster with 10 point source is shown in Figure 5.., Two examples of such a cluster with 10 point source is shown in Figure \ref{MIPS_Sim}.
+ These sources were separated al (he detection stage. ie. {μον were detected as separate sources.," These sources were separated at the detection stage, i.e. they were detected as separate sources."
+ Thev were marked as belonging to one initial cluster and were fit simultaneouslv., They were marked as belonging to one initial cluster and were fit simultaneously.
+ The simulation was repeated several (mes for a number of sources ranging from 2 to 20., The simulation was repeated several times for a number of sources ranging from 2 to 20.
+ In Table 20 we show the processing (ime and (he accuracy. of determining positions and fluxes., In Table \ref{MIPS_Sim_P} we show the processing time and the accuracy of determining positions and fluxes.
+ The positional error ὃμ is defined as the average distance between the true RZ and extracted RY positions of the point sources n: The relative flux error 9; is delined as the average ratio of the absolute difference between the rue /7 and extracted flux. /F to the true this: The processing time was measured on a 1 Gllz Sparc Sun workstation with the fitting area of the size of 7 by 7 pixels., The positional error $\delta_R$ is defined as the average distance between the true ${\mathbf R}_n^{T}$ and extracted ${\mathbf R}_n^{E}$ positions of the point sources $n$: The relative flux error $\delta_f$ is defined as the average ratio of the absolute difference between the true $f_n^T$ and extracted flux $f_n^E$ to the true flux: The processing time was measured on a 1 GHz Sparc Sun workstation with the fitting area of the size of 7 by 7 pixels.
+ The processing time depends on a variety of other factors. so (he numbers quoted in Table 2. can be used as a general guide only.," The processing time depends on a variety of other factors, so the numbers quoted in Table \ref{MIPS_Sim_P}
+ can be used as a general guide only."
+ To test active deblending we perlormed similar simulations., To test active deblending we performed similar simulations.
+ The goal of (hese simulations was to test the limit in terms of the algorithms ability to de-blend sources that cannot be separated at the detection stage., The goal of these simulations was to test the limit in terms of the algorithm's ability to de-blend sources that cannot be separated at the detection stage.
+ We simulated images with the clusters of 2 and 3 pointsources., We simulated images with the clusters of 2 and 3 pointsources.
+ The distance between (he sources in (he 2-source clusters was ~ 1/2 of the EWIIM. the distance between (the sources in (he 3-source clusters was ~ 2/3 of the EWIIM.," The distance between the sources in the 2-source clusters was $\sim$ 1/2 of the FWHM, the distance between the sources in the 3-source clusters was $\sim$ 2/3 of the FWHM."
+ In the bottow row of Figure 5. we show an example of a 2-source cluster and a 3-source cluster., In the bottow row of Figure \ref{MIPS_Sim} we show an example of a 2-source cluster and a 3-source cluster.
+ The positional and f[Iux errors were 05;=0.03 pixel and 0;=0.01 for the 2-source clusters and 05=0.01 pixel and 0;=0.004 [or the 3-source clusters., The positional and flux errors were $\delta_R = 0.03$ pixel and $\delta_f = 0.01$ for the 2-source clusters and $\delta_R = 0.01$ pixel and $\delta_f = 0.004$ for the 3-source clusters.
+ The numbers are slightly lower for the 3-source clusters., The numbers are slightly lower for the 3-source clusters.
+ This is explained by the fact that the separation is greater for (hese clusters., This is explained by the fact that the separation is greater for these clusters.
+ The important figure of merit is the failure rate 2. which is the fraction of the cases for which the algorithm could not succesfully deblencd the sources.," The important figure of merit is the failure rate $ER$, which is the fraction of the cases for which the algorithm could not succesfully deblend the sources."
+ For the 2-source clusters LA~ 0.06., For the 2-source clusters $ER \sim 0.06$ .
+ For the 3-source clusters. even though the separation is greater. (he," For the 3-source clusters, even though the separation is greater, the"
+operated bv the Johns Πορ Universitv,operated by the Johns Hopkins University.
+ Financial support to U. S. participants has been provided by NASA NÀS5-32085. Penton., Financial support to U. S. participants has been provided by NASA contract NAS5-32985.
+" 5. uw V. 72. S=0.5. S>(0.5 ο<0.5, Q-paraincter). he central mass couceutration.", $S$ $\omega$ $V$ $R$ $S=0.5$ $S>0.5$ $S<0.5$ $Q$ the central mass concentration.
+ Galaxies with higher ceutral mass concentrations. re. higher overall densities Gucluding dark matter). are predicted to have more ightly wound spiral anus.," Galaxies with higher central mass concentrations, i.e. higher overall densities (including dark matter), are predicted to have more tightly wound spiral arms."
+ The models of Fuchs (1991. 2000) result in disks with rigidly rotating spiral niodes. wherein bulges act as Πο reflectors of waves or modes induced by the swing amplification method. thus leading to modal spiral waves. which form as a result of Tooinre swing amplification (Toomre 1981).," The models of Fuchs (1991, 2000) result in disks with rigidly rotating spiral modes, wherein bulges act as inner reflectors of waves or modes induced by the swing amplification method, thus leading to modal spiral waves, which form as a result of Toomre swing amplification (Toomre 1981)."
+ Fuchs (2000) adopts a stellar-dvuamical analogue (Julian Toomre 1966) of the CGoldreich Lyucdeu-Bell (1965). sheet. which describes the local dwuamics of an uubounded patch of thin. differcutially rotating stellar disk.," Fuchs (2000) adopts a stellar-dynamical analogue (Julian Toomre 1966) of the Goldreich Lynden-Bell (1965) sheet, which describes the local dynamics of an unbounded patch of thin, differentially rotating stellar disk."
+ Ue then increases Q in the Goldreich Lyxudeu-Dell (1965) sheet. which has the affect that the bulge acts as au inner boundary.," He then increases $Q$ in the Goldreich Lynden-Bell (1965) sheet, which has the affect that the bulge acts as an inner boundary."
+ The result is that. instead of shearing deusity waves (Cas in the unbounded sheet). spiral modes appear.," The result is that, instead of shearing density waves (as in the unbounded sheet), spiral modes appear."
+ These models best show how central mass concentration correlates with spiral arii pitch anele., These models best show how central mass concentration correlates with spiral arm pitch angle.
+ If the disk is very ight (low c where o is the disk density) the mode cau be very tight. aud we are in the domain of small epicveles.," If the disk is very light (low $\sigma$ where $\sigma$ is the disk density) the mode can be very tight, and we are in the domain of small epicycles."
+ Formally. if the stability paramcter Q=eKzGao is close o unity. the value of e ust also be πα. where e dis the radial velocity dispersion aud & is the epieyclie frequency.," Formally, if the stability parameter $Q=c\kappa / \pi G \sigma$ is close to unity, the value of $c$ must also be small, where $c$ is the radial velocity dispersion and $\kappa$ is the epicyclic frequency."
+ If we increase the relative mass of the disk. we find a rend towards more open structures.," If we increase the relative mass of the disk, we find a trend towards more open structures."
+ Recent observational studies of spiral structure have Hehlighted a difference iu the morphologies secu iu he optical compared with the uear-infrared., Recent observational studies of spiral structure have highlighted a difference in the morphologies seen in the optical compared with the near-infrared.
+ The classification of ealaxies by IIubble type (ITubble 1926) is performed iu the optical regime. where dust extinction still has a large effect. and where the helt is dominated by the vouug Population I stars.," The classification of galaxies by Hubble type (Hubble 1926) is performed in the optical regime, where dust extinction still has a large effect, and where the light is dominated by the young Population I stars."
+Inthenearufrareddust extinction isminimized aud the light is dominated by,Inthenear-infrareddust extinction isminimized and the light is dominated by
+mas) and proper motion measurement. j=—57.80+O.SS mas land yrs=155.920.54 mas vr.|. were recently obtained trough very long baseline interferometry (VLBI) observations at the Australian Long Baseline Array (Delleretal.2009).,"mas) and proper motion measurement, $\mu_{\alpha}
+= -57.89 \pm 0.88$ mas $^{-1}$ and $\mu_{\delta} = -155.9 \pm 0.54$ mas $^{-1}$, were recently obtained trough very long baseline interferometry (VLBI) observations at the Australian Long Baseline Array \citep{deller09}."
+.. At the time of writing. according to the continuously updated on-line ATNF Pulsar Catalogue (Manchesteretal.200513 iis the closest isolated radio pulsar.," At the time of writing, according to the continuously updated on-line ATNF Pulsar Catalogue \citep{manchester05}, is the closest isolated radio pulsar."
+ Despite its small distance. hhas never been detected outside of the radio waveband.," Despite its small distance, has never been detected outside of the radio waveband."
+ Here we present the first deep optical and X-ray observations of3933.. performed in 2009 with aandXMM-Newton.," Here we present the first deep optical and X-ray observations of, performed in 2009 with and."
+. oobserved ffor about 40 ks on 2009 October 24., observed for about 40 ks on 2009 October 24.
+ Since the aim of the observation was the search for X-ray emission from the pulsar. only the data collected by the EPIC instrument. composed by three imaging detectors sensitive in the 0.2—15 keV energy range. were The EPIC PN camera (Strüderetal.2001) was operated in full frame mode. while the two EPIC MOS units (Turneretal.2001). were operated in large window mode.," Since the aim of the observation was the search for X-ray emission from the pulsar, only the data collected by the EPIC instrument, composed by three imaging detectors sensitive in the 0.2–15 keV energy range, were The EPIC PN camera \citep{struder01short} was operated in full frame mode, while the two EPIC MOS units \citep{turner01short} were operated in large window mode."
+ The thin optical blocking filter was used for the three EPIC cameras., The thin optical blocking filter was used for the three EPIC cameras.
+ All the data were processed using the SScience Analysis Software version 9.0.0)., All the data were processed using the Science Analysis Software version 9.0.0).
+ To improve the sensitivity at the lowest energy. where thermal emission from iis more likely to be detectable. the PN data were further processed using the tool which substantially reduces the low energy instrumental background due to electronic noise.," To improve the sensitivity at the lowest energy, where thermal emission from is more likely to be detectable, the PN data were further processed using the tool which substantially reduces the low energy instrumental background due to electronic noise."
+ The standard pattern. selection criteria for the EPIC X-ray events (patterns 0—4 for PN and 0-12 for MOS) were adopted., The standard pattern selection criteria for the EPIC X–ray events (patterns 0–4 for PN and 0–12 for MOS) were adopted.
+ After filtering out the time intervals affected by a high level of particle background. the net exposure times were 22 ks for the PN and 26 ks for the MOS.," After filtering out the time intervals affected by a high level of particle background, the net exposure times were 22 ks for the PN and 26 ks for the MOS."
+ A blind search for X-ray emission from wwas performed using different source detection algorithms in many energy bands. but no source compatible with the pulsar position was detected.," A blind search for X-ray emission from was performed using different source detection algorithms in many energy bands, but no source compatible with the pulsar position was detected."
+" The closest X-ray source is 25.6£0.5 arcsec away from the proper-motion-corrected position of (Delleretal.2009.. see Figure 1). while the absolute astrometry of the EPIC images is l aresec (root mean square. In order to set an upper limit on the pulsar X-ray emission."" we selected the events from a 160 arcsec radius circular region centred at the position ofJ2144—3933.. obtaining. in the 10 keV energy range. 40 and 21 counts in the PN and in the two MOS cameras summed together. respectively (see Table | for other energy bands)."," The closest X-ray source is $25.6\pm0.5$ arcsec away from the proper-motion-corrected position of \citealt{deller09}, see Figure \ref{ximage}) ), while the absolute astrometry of the EPIC images is $\sim$ 1 arcsec (root mean square, In order to set an upper limit on the pulsar X-ray emission, we selected the events from a 10 arcsec radius circular region centred at the position of, obtaining, in the 0.2--10 keV energy range, 40 and 21 counts in the PN and in the two MOS cameras summed together, respectively (see Table \ref{xlimits} for other energy bands)."
+ To better characterise the background intensity in different energy bands. we extracted for each camera the background events from a larger source-free region in the same chip. deriving the estimates reported in Table |. for the number of background events expected in the source region.," To better characterise the background intensity in different energy bands, we extracted for each camera the background events from a larger source-free region in the same chip, deriving the estimates reported in Table \ref{xlimits} for the number of background events expected in the source region."
+ Taking into account these background estimates and the Poissonian fluctuations (Gehrels1986) of the counts extracted from the region ofJ2144—3933.. we obtain the 30 upper limits on the pulsar count rate in different energy bands reported in Table |..," Taking into account these background estimates and the Poissonian fluctuations \citep{gehrels86} of the counts extracted from the region of, we obtain the $\sigma$ upper limits on the pulsar count rate in different energy bands reported in Table \ref{xlimits}."
+ The conversion from count rates to fluxes for each is detectorperformed by assuming different spectral models (see Section 4) within the (version 11.3.1) spectral fitting software and using the response matrices and ancillary files produced by the software fora 10 aresee circular region centred at the pulsar position., The conversion from count rates to fluxes for each detector is performed by assuming different spectral models (see Section 4) within the (version 11.3.1) spectral fitting software and using the response matrices and ancillary files produced by the software for a 10 arcsec circular region centred at the pulsar position.
+ We observed with the aat the ESO Paranal observatory on 2009 August 21 in visitor mode with (FORS2)) at the Cassegrain focus of the Antu UTI telescope., We observed with the at the ESO Paranal observatory on 2009 August 21 in visitor mode with ) at the Cassegrain focus of the Antu UT1 telescope.
+ To enhance the sensitivity in the blue part of the spectrum. we used the," To enhance the sensitivity in the blue part of the spectrum, we used the"
+statistical approach is that a «οπλο description for 8 py is valid for all pulsars.,statistical approach is that a common description for $R$ $\rho_{\nu}$ is valid for all pulsars.
+" The couuuon description should. however, account for relevant depeudences on quautitics. such as frequency. period. a. etc."," The common description should, however, account for relevant dependences on quantities, such as frequency, period, $\alpha$, etc."
+ properly., properly.
+ The radio cussion at different frequencies is expected to originate at different altitudes above the stellar surface. with the higher frequency radiation associated with regions of lower altitude.," The radio emission at different frequencies is expected to originate at different altitudes above the stellar surface, with the higher frequency radiation associated with regions of lower altitude."
+ This phenomenon known asmapping. fads overwhelming support from observations.," This phenomenon known as, finds overwhelming support from observations."
+" '""Thorsett (1991) has sugeested aun enipircal relation or the observed pulse width as a function of frequency. which seems to provide adequate description of the observed behaviour."," Thorsett (1991) has suggested an empirical relation for the observed pulse width as a function of frequency, which seems to provide adequate description of the observed behaviour."
+" We adopt a similar relation for the frequeney evolution of the beam radius p, as follows where pis the value of p, at infinite frequency. ¢ the spectral iudex. and Jv a coustaut."," We adopt a similar relation for the frequency evolution of the beam radius $\rho_{\nu}$ as follows where $\hat{\rho}$ is the value of $\rho_{\nu}$ at infinite frequency, $\zeta$ the spectral index, and $K$ a constant."
+" Note that both ¢ Iv are expected to have positive values. so that the ΙΙΙ value of p, is p. which should correspoud to the angular size of the polar cap."," Note that both $\zeta$ $K$ are expected to have positive values, so that the minimum value of $\rho_{\nu}$ is $\hat{\rho}$, which should correspond to the angular size of the polar cap."
+ Raukin (1993a) has demonstrated (see also Cul. Iijak Seiradakis 1993: Wramer et al.," Rankin (1993a) has demonstrated (see also Gil, Kijak Seiradakis 1993; Kramer et al."
+ 1991) that the beam radius p varies as P9 (sehere P is the period of the pulsar). a result which is in excellent aereeiiceut with that expected ποια a dipole geometry (Cal 1981).," 1994) that the beam radius $\hat{\rho}$ varies as $P^{-0.5}$ (where $P$ is the period of the pulsar), a result which is in excellent agreement with that expected from a dipole geometry (Gil 1981)."
+ Eq 3. thus takes the form where pa is the minim beam radius for P=1 sec., Eq \ref{eq:rho} thus takes the form where $\rho_{\circ}$ is the minimum beam radius for $P = 1$ sec.
+ Biges (1990) has sugeested that 72? should be a function of a. such that the beam shape is circular for a=0 ancl is Iucreasinely compressed in the latitudinal direction as o inercases to 90°., Biggs (1990) has suggested that $R$ should be a function of $\alpha$ such that the beam shape is circular for $\alpha=0$ and is increasingly compressed in the latitudinal direction as $\alpha$ increases to $90^{\circ}$.
+ We therefore model the functional dependence of Ron a as R= Rar. where BR.> is the axial ratio of the beam at a=0. and 7 is a fiction ofa.," We therefore model the functional dependence of $R$ on $\alpha$ as $R = R_{\circ}\tau$ , where $R_{\circ}$ is the axial ratio of the beam at $\alpha = 0$, and $\tau$ is a function of$\alpha$."
+ AccordingoO to BieesCoco» (1990). R.=1 and 7 is Oogiven by where Au. Ao are constants aud o is in degrees.," According to Biggs (1990), $R_{\circ} = 1$ and $\tau$ is given by where $K_{1}$, $K_{2}$ are constants and $\alpha$ is in degrees."
+ Diges finds that Avy and A» are 3.5 and LL. respectively.," Biggs finds that $K_{1}$ and $K_{2}$ are 3.3 and 4.4, respectively."
+ We. however. treat Avy1.2 as free parameters in our model.," We, however, treat $K_{1,2}$ as free parameters in our model."
+" BBased ou the study of conal componcuts. Raukiu (1993a) has argued for two nested hollow cones of ciission namely, the outer and the inner cone."," Based on the study of conal components, Rankin (1993a) has argued for two nested hollow cones of emission-- namely, the outer and the inner cone."
+" Assuming the bemus to be cieular in shape. opening half angles of the two cones at 1 GIIz were found to be 1.3"" aud 5.77 respectively."," Assuming the beams to be circular in shape, opening half angles of the two cones at 1 GHz were found to be $4.3^{\circ}$ and $5.7^{\circ}$, respectively."
+ During our preliminary examination of the present suuple. we noticed a need to allow for three cones of chussion.," During our preliminary examination of the present sample, we noticed a need to allow for three cones of emission."
+" To incorporate this feature iu our model. we introduce two ratios. rl<l audr2>1. to define the size scaling of the iunerauost and the outeranost cone. respectively, with reference to a uuiddle cone. for which the detailed shape is defined."," To incorporate this feature in our model, we introduce two ratios, $r1 < 1$ and$r2 > 1$ , to define the size scaling of the inner-most and the outer-most cone, respectively, with reference to a `middle' cone, for which the detailed shape is defined."
+ Using the model here defined. we need to solve for Rk... ον NK. ps. Wy. fo. rl and c2 in this three-conal-ring model.," Using the model here defined, we need to solve for $R_{\circ}$, $\zeta$, $K$, $\rho_{\circ}$, $K_{1}$, $K_{2}$ , $r1$ and $r2$ in this three-conal-ring model."
+ The parameter set thus represents an average! description of the beam., The parameter set thus represents an `average' description of the beam.
+ AAu optimized erid search was performed for suitable ranges of the parameter values aud in fine enough steps., An optimized grid search was performed for suitable ranges of the parameter values and in fine enough steps.
+ For ¢. the search rauge allowed for both |ve aud -ve values.," For $\zeta$, the search range allowed for both +ve and -ve values."
+ By definition. rjx1 and οzc1.," By definition, $r_{1}\leq 1$ and $r_{2}\geq 1$."
+" The best fit was obtained by 1ininiziug the standard deviation σ., defined by where D; is the deviationof the i” data point from the nearest coual ring in the model aud γω denotes the degrees of freedom.", The best fit was obtained by minimizing the standard deviation $\sigma_{\circ}$ defined by where $D_{i}$ is the deviationof the $i^{th}$ data point from the nearest conal ring in the model and $N_{dof}$ denotes the degrees of freedom.
+" The factor 180/75 gives σε, iu units of degrees under the sxiidl-angle approximation.", The factor $180/\pi$ gives $\sigma_{\circ}$ in units of degrees under the small-angle approximation.
+ Table 2 lists the parameter values which correspond to the best ft for the cutire sample set., Table 2 lists the parameter values which correspond to the best fit for the entire sample set.
+ With these values. the eq.," With these values, the eq."
+" f£. can now be rewritten as where p, is in deerees.", \ref{eq:rho1} can now be rewritten as where $\rho_{\nu}$ is in degrees.
+" This average description for the ""uuiddle cone applies also to the other two cones when Pp Is scaler bv the ratio rl=0.8 or r2=1.3 (for the inner and the outermost. respectively)."," This average description for the `middle' cone applies also to the other two cones when $\rho_{\nu}$ is scaled by the ratio $r1=0.8$ or $r2=1.3$ (for the inner and the outermost, respectively)."
+ Fig 1. shows the data (plotted to a common scale) for one quadrant ofthe beam and the three solid curves correspouding to the best fit cones., Fig \ref{fig:fig4} shows the data (plotted to a common scale) for one quadrant of the beam and the three solid curves corresponding to the best fit cones.
+ The poiuts iu the fleure. though corresponding to different pulsars and frequencies.are translated to a colon reference scale appropriate for P—1sec. a=0 and v= x.," The points in the figure, though corresponding to different pulsars and frequencies,are translated to a common reference scale appropriate for $P=1$sec, $\alpha = 0$ and $\nu =\infty $ ."
+ We have assumed the period dependence of py as P0.5 whereasLyne ando Manchester(1088) found a dependence of P ," We have assumed the period dependence of $\rho_{\nu}$ as $P^{-0.5}$ , whereasLyne and Manchester(1988) found a dependence of $P^{-\frac{1}{3}}$ ."
+We have exanuned the latter possibility and fouud that the difference in the staudard, We have examined the latter possibility and found that the difference in the standard
+"(PA=0"") we acquired spectra in a direction perpendicular to the known outflows.",$^\circ$ ) we acquired spectra in a direction perpendicular to the known outflows.
+ Thus. our non-detections can be reconciled with the Spitzer observations if most of the enission arises in jets/outflows. which can be verified with future high-resolution spectra covering the region of (he known outflows.," Thus, our non-detections can be reconciled with the Spitzer observations if most of the emission arises in jets/outflows, which can be verified with future high-resolution spectra covering the region of the known outflows."
+ Viscous accretion is thought to be the dominant disk dispersal mechanism but alone cannot explain the relatively short (S10 MMwyr) disk dispersal timescales inferred from observations (e.g. review by Dullemondοἱal. 2007))., Viscous accretion is thought to be the dominant disk dispersal mechanism but alone cannot explain the relatively short $\lesssim 10$ Myr) disk dispersal timescales inferred from observations (e.g. review by \citealt{dullemond07}) ).
+ Lieh-enerey stellar photons can heat ancl ionize the disk surface and induce photoevaporation (ie. disk mass loss) from outer disk regions., High-energy stellar photons can heat and ionize the disk surface and induce photoevaporation (i.e. disk mass loss) from outer disk regions.
+ Exactly when pholoevaporation becomes an important disk dispersal mechanism and how much disk mass it can remove is still a matter of debate., Exactly when photoevaporation becomes an important disk dispersal mechanism and how much disk mass it can remove is still a matter of debate.
+ FUV. and X-ray photons can penetrate much larger column densities of gas than stellar EUV photons., FUV and X-ray photons can penetrate much larger column densities of gas than stellar EUV photons.
+ As a consequence. thev are expected to drive efficient photoevaporation Irom the early (mes of disk evolution at rates that may exceed (he EUV pholoevaporation rate by one or even (wo orders of magnitudes (Gorli&IHollenbach2009:Ercolanoetal.2009).," As a consequence, they are expected to drive efficient photoevaporation from the early times of disk evolution at rates that may exceed the EUV photoevaporation rate by one or even two orders of magnitudes \citep{gorti09,ercolano09}."
+. stellar EVV photons start to penetrate (he disk wind. ancl (thus reach the disk surface. only when accretion rates fall below ~LO “AL. /vr. corresponding to hydrogen column density screens of Z10P ? (Hollenbach&Gorti2009).," Stellar EUV photons start to penetrate the disk wind, and thus reach the disk surface, only when accretion rates fall below $\sim 10^{-8}$ $_{\sun}$ /yr, corresponding to hydrogen column density screens of $\lesssim$ $^{17}$ $^{-2}$ \citep{hollenbach09}."
+. Although EUV photons likely produce little mass loss at all stages (Gorli&IHollenbach2009).. they can help to quickly (~LO? vvr) disperseremaining disk gas once stellar accretion rates [all below a few 10. MAL. {ντ CAlexanderetal. 2006a)..," Although EUV photons likely produce little mass loss at all stages \citep{gorti09}, they can help to quickly $\sim 10^5$ yr) disperseremaining disk gas once stellar accretion rates fall below a few $10^{-10}$ $_{\sun}$ /yr \citep{alexander06a}. ."
+We analyzed the new transit times in conjunction with previously published midtransit times. to seek evidence for significant discrepancies from strict periodicity and refine the ephemerides to allow for accurate prediction of future events.,"We analyzed the new transit times in conjunction with previously published midtransit times, to seek evidence for significant discrepancies from strict periodicity and refine the ephemerides to allow for accurate prediction of future events."
+ The new transit times are given in Tables 8.. 9. and 10.. with uncertainties determined by the MCMC analysis described in Section 3.," The new transit times are given in Tables \ref{tbl:midtransits-tres4}, \ref{tbl:midtransits-hat3}~ and \ref{tbl:midtransits-wasp12}, with uncertainties determined by the MCMC analysis described in Section 3."
+ Previously published midtransit times were taken from Gibsonetal.(2010).. Hebbetal. (2009).. Sozzettietal. (2009).. Torresetal. (2007).. and Mandushevetal.(2007)..," Previously published midtransit times were taken from \citet{2010MNRAS.401.1917G}, \citet{2009ApJ...693.1920H}, \citet{2009ApJ...691.1145S}, \citet{2007ApJ...666L.121T}, and \citet{2007ApJ...667L.195M}."
+ The transit times for each system were fitted with a linear function. where E ts an integer (the epoch). P is the period. and 75 is à particular reference time.," The transit times for each system were fitted with a linear function, where $E$ is an integer (the epoch), $P$ is the period, and $T_0$ is a particular reference time."
+ The best-fitting values of P and Ti were determined by linear regression., The best-fitting values of $P$ and $T_0$ were determined by linear regression.
+ Figures 5-7 show the residuals to the fits. and the captions specify the minimum 4 and the number ofdegrees of freedom.," Figures 5-7 show the residuals to the fits, and the captions specify the minimum $\chi^2$ and the number of degrees of freedom."
+ In no case tis there clear evidence of timing anomalies., In no case is there clear evidence of timing anomalies.
+ For TrES-4. there is formally only a chance of obtaining such a large X with random Gaussian errors. but we do not deem this significant enough to warrant special attention.," For TrES-4, there is formally only a chance of obtaining such a large $\chi^2$ with random Gaussian errors, but we do not deem this significant enough to warrant special attention."
+ The case of WASP-12 required somewhat special treatment because Hebbetal.(2009) did not report individual midtransit times. but rather a consensus reference time based on observations of multiple transits.," The case of WASP-12 required somewhat special treatment because \citet{2009ApJ...693.1920H} did not report individual midtransit times, but rather a consensus reference time based on observations of multiple transits."
+ We included their quoted reference time as a single data point., We included their quoted reference time as a single data point.
+ Campoetal.(2010) provided many transit times obtained over several years. but most of the data were from amateur observers and were not presented in detail or evaluated critically.," \citet{2010arXiv1003.2763C} provided many transit times obtained over several years, but most of the data were from amateur observers and were not presented in detail or evaluated critically."
+ For this reasons we did not include them: however. as a consequence. there are not many points in our fit.," For this reasons we did not include them; however, as a consequence, there are not many points in our fit."
+ The results for P and Το are given in Tables 5-7., The results for $P$ and $T_0$ are given in Tables 5-7.
+ They are based on a fit in which all of the uncertainties of the transit times were rescaled by à common factor to give V(Nor= 1., They are based on a fit in which all of the uncertainties of the transit times were rescaled by a common factor to give $\chi^2/N_{\rm dof} = 1$ .
+ Our intention is to provide conservative error estimates to allow for planning of future observations., Our intention is to provide conservative error estimates to allow for planning of future observations.
+ The uncertainties on the individual transit times given in Tables 8-10 werenor rescaled in this way. nor were the error bars that are plotted in Figures 5-7.," The uncertainties on the individual transit times given in Tables 8-10 were rescaled in this way, nor were the error bars that are plotted in Figures 5-7."
+ We have presented new photometry and new analyses of the transiting exoplanets TrES-4+b. HAT-P-3b. and WASP-12b.," We have presented new photometry and new analyses of the transiting exoplanets TrES-4b, HAT-P-3b, and WASP-12b."
+ Whereas the discovery papers reporting TrES-4b and HAT-P-3b included only a few high-precision light curves. our analyses are based on 5-6 such datasets.," Whereas the discovery papers reporting TrES-4b and HAT-P-3b included only a few high-precision light curves, our analyses are based on 5-6 such datasets."
+ Likewise. the WASP-I2b discovery paper featured only one high- light curve.to which we have added two.," Likewise, the WASP-12b discovery paper featured only one high-precision light curve,to which we have added two."
+ We have applied consistent and conservative procedures for parameter, We have applied consistent and conservative procedures for parameter
+SSED by scaling the Netzer courposite spectrum by an arbitrary factor. we fid a 4?=573 for 32 DOF. which is orlmally a very poor fit. but the discrepancy is dominated bv the disagreement iu the cutoff at short waveleneths and he bump near 7 mnücrons.,"SED by scaling the Netzer composite spectrum by an arbitrary factor, we find a $\chi^2 
+= 573$ for 32 DOF, which is formally a very poor fit, but the discrepancy is dominated by the disagreement in the cutoff at short wavelengths and the bump near 7 microns."
+ The infrared spectrmm of aalso looks ταν similar to low redshift Sevtert galaxies taken from a sample of 23 galaxies (Corian ot al.," The infrared spectrum of also looks very similar to low redshift Seyfert galaxies taken from a sample of 23 galaxies (Gorjian et al.,"
+ iu xeparation)., in preparation).
+ The spectrum of iis plotted with the two most similar Sevtert spectra in Figure 5.. scaled to match the ux of0305.," The spectrum of is plotted with the two most similar Seyfert spectra in Figure \ref{fig:seyferts}, scaled to match the flux of."
+.. The qualitative shape of the SEDs matches well from 1-10. pau. although Myk 509 shows strouger emissiou features. oesmuablv due to silicates.," The qualitative shape of the SEDs matches well from 4-10 $\mu$ m, although Mrk 509 shows stronger emission features, presumably due to silicates."
+ The similarity of the infrared SED of to other Sevferts and quasars indicates that our microleusimg studies of this object will broadly apply to racio-quict active ealaxies., The similarity of the infrared SED of to other Seyferts and quasars indicates that our microlensing studies of this object will broadly apply to radio-quiet active galaxies.
+ To improve our plivsical understanding of the cmission from we fitted the SED with the models of ?.., To improve our physical understanding of the emission from we fitted the SED with the models of \citet{Fritz2006}.
+ The models utilize the Mathis-Rumpl-Noxrdsieck (MBRN) dust size distribution (?) with scatteriug aud absorption opacities from ?.., The models utilize the Mathis-Rumpl-Nordsieck (MRN) dust size distribution \citep{Mathis1977} with scattering and absorption opacities from \citet{Laor1993}.
+ The model fixes the geometry as a torus with au opening angle that is independent of radius aud au iuner radius of the torus. which is set by a dust sublimation temperature of 1500 Ik. where Lig is the AGN luminosity in units of 1015," The model fixes the geometry as a torus with an opening angle that is independent of radius and an inner radius of the torus, which is set by a dust sublimation temperature of 1500 K, where $L_{46}$ is the AGN luminosity in units of $10^{46}$ erg $^{-1}$."
+ cre : ∙∙ ⋅ ∣∣∣∣⋯⋯⋅↽∕∏∐∖∶↴↜⊾↥⋅↕≺⇂∪↕⋟↕⊔≺⋉∐∖↕↴∖↴↸⊳∪↖↽↸∖↥⋅↴∖↴⋜↧↥⋅⋜⋯∶↴∙⊾↸∖∪↕≯↻⋜∐⋅⋜∐⊔↸∖↑↸∖↥⋅↴∖↴↕⋟∪↥⋅↑↕∐∖≼⊔↕↴∖↴↑↖↖⇁↕↑∐∶∐⋟↑∐↸∖↥⋅⋜↧↑↕∪∪↕⋟↑∐↸∖∪∏↑↸∖↥⋅↑∪↕∐↕∐∖↥⋅↥⋅⋜∥∐∏↴∖↴∙⋅ ↴↜∖↗," The dust density in the model is described by $\rho(r,\theta) \propto
+r^\beta e^{-\gamma |\cos{(\theta)}|}$, with a dust-free cone within polar-angle $\theta < \theta_{cone}$."
+ ⋅ : ⋅⋅ ↴∖↴↓∙↽∕∏∐∖≺↧∏↴∖↴↑≼∐∖∐↴∖↴↕↑⋅↖↽↕∐↑↕∐∖⋯∪≼∐∖↕↕↴∖↴≼∐∖↴∖↴≼⊳↥⋅↕↴⋝↸∖≺↴⋝∙↖↽∣≻⋖∣⋮⋅∣∣⋝∖∣↶⋅↙↖⋯⊔∣⋝↖∙↖∏↑∐⋜↧≼↧∏↴∖↴↑≓↕↥⋅↸∖↸∖⋯∐↸∖↖∏↑↕∐∐↻∪↕⋜∐⋅≓⋜⋯∶↴∙↕↸∖ ∶≩∩∫↿≽∩⋯⊽↗⋅∫↿⋟⋯∣⊽⊓∶≩∩∩∶⋖⊇⋝↑∐↸∖↖⇁⋜∐⋅↕⋜↧↑↕∪∐∪↕⋟≼↧∏↴∖↴↑≺∐∖∐↴∖↴↕↑⋅↖⇁↖↖⇁↕↑∐↥⋅⋜∥∐↿↴∖↴∙↕↗∣↕⊇∶⋖∶≩⋟↑∐↸∖↸∖≺∣∏⋜↧↑∪∏⋜↧↕∪↻↑↕⋯↕ ≺∐∖↻↑∐⋜↧↑∩∙⊤∕∣↕⊔∙∩∙↕∣−≝≱⋡⊤∐∣∶≺⊔↕∐↸⊳∐∐⋜↕⊓∪↕⋜⋯∶↴∙⊾↕↸∖↴∖↴∙∕∙↥⋅⋜↧∐∶↴∙⊾↕∐∶↴∙⊾↕⋟↥⋅∪⊔∩∙∩↕≺↸∖≺∟≼↴∙⊾↸∖⊣≻∐⋝↑∪≺∖∖∩≺∐∖∶↴⋁↥⋅↸∖↸∖↴∖↴⋖↕⋟⋜↧↸⊳↸∖⊣≻∐⋝∙⋜⋯≺↧ ∐∙⊇↕∙∶∐∙∐∙⋅↱⊐↕∙⊓↕∙⊤↕⋅≺∖∖↕↴↕∐↴⋝↸∖↑↖↖⇁↸∖↸∖∐∶⋖⋅↱⊐⋝≼↧∏↴∖↴↑≓↕≯↥⋅↸∖↸∖↸⊳∪↕∐∖↖↖↽↕↑∐↴∖↴↕∠↸∖⊇∩↴∣∣⋯⋯∢∔∪↴∶⋜⋯≼⊔⊓⋟⋜⋯⋜↧∐∶↴∙⊾∏↕⋜∐⋅↸⊳∏↑∪∐∎↕∐ ≼∐∖∐↴∖↴↕↑⋅↖↽↖↖↽↕↑∐∩↰⊓⋅," The grid of models covers a range of parameters for the dust with: (1) the ratio of the outer to inner radius, $30 < R_{max}/R_{min} < 300$; (2) the variation of dust density with radius, $-1 < \beta < 1/2$ ; (3) the equatorial optical depth at 9.7 $\mu$ m, $0.1 < \tau_{9.7} < 10$; (4) inclination angles, $i$, ranging from 0.01 (edge-on) to 89 degrees (face-on), and $11,21,31,41,51,61,71,81^\circ$ in between; (5) dust-free cone with size $20^\circ < \theta_{cone} < 60^\circ$; and (6) an angular cutoff in density with $0 < \gamma < 6$."
+ We have scaled the infrared portion of the models by an arbitrary coustant to optimize the match with our Spitzer SED., We have scaled the infrared portion of the models by an arbitrary constant to optimize the match with our Spitzer SED.
+" The best-fit model is shown in Figure 6.. which has parameters ,,.//8,,;,=LOO. =1/2. r210. =τα"". 0,=207. and >=3."," The best-fit model is shown in Figure \ref{fig:q2237_fritz}, , which has parameters $R_{max}/R_{min} = 100$, $\beta = -1/2$, $\tau_{9.7} = 10$, $i = 71^\circ$, $\theta_{cone} = 20^\circ$, and $\gamma=3$."
+ The best fit implies an AGN huninosity of Lis10! eres L| remarkably close to the luuinosity found from the SSED (£=LO« L0Merg/s). despite the fact that we have ouly fit the Fritz models to the infrared data.," The best fit implies an AGN luminosity of $4.4 \times 10^{46}$ erg $^{-1}$, remarkably close to the luminosity found from the SED $L=4.0 \times 10^{46} erg/s$ ), despite the fact that we have only fit the Fritz models to the infrared data."
+ This may Πρίν that the dust acts as a fairly good calorimeter of the total ACN flux., This may imply that the dust acts as a fairly good calorimeter of the total AGN flux.
+ The qualitative fit to the data is fair: the model shows apeak near 2.5 jan in the rest frame. aud gradualdecline towards longer wavelongths. aud an emission," The qualitative fit to the data is fair: the model shows apeak near 2.5 $\mu$ m in the rest frame, and gradualdecline towards longer wavelengths, and an emission"
+The braking iudex » can be obtained differcutiating Eq.(13)). aud using Eqxs.(.11)): where we have introduced This quantity is the ratio between the gravitational aud the electromagnetic contributions to the pulsu spin-down rate and we will see in the next section that its preseut values are always less than 1.,"The braking index $ n$ can be obtained differentiating \ref{omp}) ), and using \ref{ompem}, \ref{ompgw}) ): where we have introduced This quantity is the ratio between the gravitational and the electromagnetic contributions to the pulsar spin-down rate and we will see in the next section that its present values are always less than $1$."
+" Note that. in terms of O and its derivative. which iniplies 0,4,30<4,"," Note that, in terms of $\Omega$ and its derivative, Inverting \ref{n1}) ), which implies $n_{em}\le n <5$."
+" Iu (1)) we have plotted the functiou Y(0) for different values of »,,, in the rauge [1.3]. As we expect. for cach ian. Y(0)=ο ναι i=no (i.c. no gravitational torque) while Y(0)>|x τω »5 (i.c. negligible electromagnetic torque)."," In \ref{fig1}) ) we have plotted the function $ Y(n)$ for different values of $ n_{em}$ in the range $ [1,3].$ As we expect, for each $n_{em}$, $Y(n)=0$ when $n=n_{em}$ (i.e. no gravitational torque) while $Y(n)\rightarrow +\infty $ for $n\rightarrow 5$ (i.e. negligible electromagnetic torque)."
+ We see that a braking index less than 3 cau be obtained combining the effects of the electromagnetic aud eravitational torques., We see that a braking index less than 3 can be obtained combining the effects of the electromagnetic and gravitational torques.
+ For iustance. for the Crab pulsar we have 9=2.51 aud if we assume. Say. 04)=2.1. then it iunediatelv follows Y (LOL that is the gravitational radiation coutributiou to the spin-down rate is about I4 of the electromagnetic onec.," For instance, for the Crab pulsar we have $ n=2.51$ and if we assume, say, $ n_{em}=2.4$, then it immediately follows $Y\simeq 0.04$ , that is the gravitational radiation contribution to the spin-down rate is about $4\% $ of the electromagnetic one."
+" On the other haud. if 5.4,=1. then Yz 0.6."," On the other hand, if $n_{em}=1$, then $Y\simeq 0.6$ ."
+" The spin-down of a pulsar in which both the electromagnetic and gravitational torques are acting. Eq.(13)). cau be expressed as where both the ""effective I,yy auc are functiou of he augulu velocity O. As a consequence of its dependence ou O. the braking index 0 is not constant i time."," The spin-down of a pulsar in which both the electromagnetic and gravitational torques are acting, \ref{omp}) ), can be expressed as where both the “effective” $ K_{eff}$ and $ n$ are function of the angular velocity $ \Omega .$ As a consequence of its dependence on $ \Omega $, the braking index $ n$ is not constant in time."
+ In he next section we will derive the time evolution of n., In the next section we will derive the time evolution of $n$ .
+ Iu all subsequent calculations we will assiune the pulsar Ποιοι of inertia beie equal to the canonical value P=10%η”., In all subsequent calculations we will assume the pulsar momentum of inertia being equal to the canonical value $I=10^{38}~kg~m^2$ .
+ From Eq.(11)) we cau express the ellipticity of the oulsar as Combining Eqx(ll.17)) we can re-write LEq.(21)) as a function of the observed pulsar period and its derivative. aud of the ratio Y between the gravitational and electromagnetic spin-down rates: Eq.(22)) plays a basic role in the determination of the upper limits of the pulsars ellipticity.," From \ref{ompgw}) ) we can express the ellipticity of the pulsar as Combining \ref{ompgw}, \ref{ipsi}) ) we can re-write \ref{ep}) ) as a function of the observed pulsar period and its derivative, and of the ratio $Y$ between the gravitational and electromagnetic spin-down rates: \ref{epy}) ) plays a basic role in the determination of the upper limits of the pulsars ellipticity."
+ Differentiating Eq.(17)) we easily fiud where Jy=S44 being ng the observed braking index (0. its present Havalue).," Differentiating \ref{ipsi}) ) we easily find and then where $Y_0={{n_0-n_{em}}\over{5-n_0}}$, being $n_0$ the observed braking index (i.e. its present value)."
+ We use this solution to solve the equation for the time evolution of the pulsar augular velocity., We use this solution to solve the equation for the time evolution of the pulsar angular velocity.
+" To this purpose. let us write Eq.(13)) iu the form We nuuercallv solve this differeutial equation. fixing the pulsar age tog. aud finding a solution depeudiug ou three wukuowns: the ellipticity € (contained i A,). Dou, Con Which (Q) depends) audthe initialaneular velocity QO; (0;= Otage))."," To this purpose, let us write \ref{omp}) ) in the form We numerically solve this differential equation, fixing the pulsar age $t_{age}$ and finding a solution depending on three unknowns: the ellipticity $\epsilon$ (contained in $K_{gw}$ ), $n_{em}$ (on which $Y(\Omega)$ depends) andthe initialangular velocity $\Omega_i$ $\Omega_i=\Omega(t_{age})$ )."
+ These quantitiesmust be chosen in order to verifv two conditions: those elven bv O(f—0)=Oy aud by Eq.(22)). caleulated for thevalues of the parameters at f£= 01:," These quantitiesmust be chosen in order to verify two conditions: those given by $\Omega(t=0)=
+\Omega_0$ and by \ref{epy}) ), calculated for thevalues of the parameters at $t=0$ :"
+"ound. one with D,=0.7+0.2 ΚΚΟ and the other with £5,=15+ 2kkG. Our high resolution spectra clearly rule out a polar tield strength as large as KKG. since that would produce clear Zeeman splitting. which is not observed. leaving onlythe model with B, 0.ΤΚΚΟ. Por the latter model. agreement with the observed pulsation amplitude and phase variations requires that ~=366° and =62"".","found, one with $B_p = 0.7 \pm 0.2$ kG and the other with $B_p = 15 \pm 2$ kG. Our high resolution spectra clearly rule out a polar field strength as large as kG, since that would produce clear Zeeman splitting, which is not observed, leaving onlythe model with $B_p = 0.7$ kG. For the latter model, agreement with the observed pulsation amplitude and phase variations requires that $\gamma = \beta = 
+66^{\circ}$ and $i=62^{\circ}$."
+ reftig:modulation shows the model amplitude and phase variations as a function of rotation phase for an =16 dipole mode.in good agreement with the observations shown in reftig:phampl..," \\ref{fig:modulation} shows the model amplitude and phase variations as a function of rotation phase for an $n = 16$ dipole mode,in good agreement with the observations shown in \\ref{fig:phamp1}."
+ reftig:multip compares the observed amplitudes of the multiplet structure with the observations., \\ref{fig:multip} compares the observed amplitudes of the multiplet structure with the observations.
+ À higher than normal value of the limb-darkening coefficient of sp=0.5 was needed to get good agreement for the -L42 and «ανω components., A higher than normal value of the limb-darkening coefficient of $\mu = 0.8$ was needed to get good agreement for the $A_{\pm 2}$ and $A_{\pm 3}$ components.
+" The eigenfunction shown in the bottom panel is significantly. deformed from the Legendre function £,(cos6). even though νο and 2345 are much smaller than 24) and cy."," The eigenfunction shown in the bottom panel is significantly deformed from the Legendre function $P_1(\cos\theta)$, even though $A_{\pm2}$ and $A_{\pm3}$ are much smaller than $A_1$ and $A_0$."
+ This is because for an odd mode such as a dipole mode. /&3 components (which produce the sla. and els: sidelobes) are visible only through the limb-darkening effect: without limb- darkening these components would not be visible due to cancellation on the surface.," This is because for an odd mode such as a dipole mode, $l \ge 3$ components (which produce the $A_{\pm2}$ and $A_{\pm3}$ sidelobes) are visible only through the limb-darkening effect; without limb- darkening these components would not be visible due to cancellation on the surface."
+ This increased limb-darkening coefficient is consistent with the atmospheric temperature gradient in roAp stars. which is steeper than in normal stars. as is usually shown by the core-wing anomaly in the Balmer lines (2). for most roAp stars.," This increased limb-darkening coefficient is consistent with the atmospheric temperature gradient in roAp stars, which is steeper than in normal stars, as is usually shown by the core-wing anomaly in the Balmer lines \citep{cowleyetal01} for most roAp stars."
+ The magnetic field derived in the previous section is rather weak for roAp stars and its effect on the mode inclination can be balanced by the centrifugal force., The magnetic field derived in the previous section is rather weak for roAp stars and its effect on the mode inclination can be balanced by the centrifugal force.
+ In this case. as shown by ?.. there is no reason to have an axis of pulsation aligned along the magnetic axis or the rotation axis.," In this case, as shown by \citet{bigot02}, there is no reason to have an axis of pulsation aligned along the magnetic axis or the rotation axis."
+ In order to calculate the mode inclination  and its polarization &. we assume a value of [j//75|=1 tratio of magnetic to centrifugal shift of frequency).," In order to calculate the mode inclination $\gamma$ and its polarization $\psi$ , we assume a value of $\left | \mu^{\rm mag} \right | = 1$ (ratio of magnetic to centrifugal shift of frequency)."
+ The ratio 4 between the Coriolis and centrifugal shift is proportional to —».0. +. where, The ratio $\chi$ between the Coriolis and centrifugal shift is proportional to $\sim n^{-3}\Omega ^{-1}$ where
+"to overlap at +20°«b with the RAVE survey area, though we do not detect such a population in the data.","to overlap at $+20 ^ \circ<b$ with the RAVE survey area, though we do not detect such a population in the data."
+" Future releases of RAVE data with more observations in this area, combined with more careful modelling of the stream, will enable a better understanding of whether this area is indeed populated by Aquarius stream stars."," Future releases of RAVE data with more observations in this area, combined with more careful modelling of the stream, will enable a better understanding of whether this area is indeed populated by Aquarius stream stars."
+" The above scenario of a dynamically young stream is not inconsistent with an age of 10Gyr for the Aquarius candidates, as estimated from the isochrone fit in Section 4.1:: the stream can be seen as a remnant of an old satellite that has been recently disrupted."," The above scenario of a dynamically young stream is not inconsistent with an age of $10\,\gyr$ for the Aquarius candidates, as estimated from the isochrone fit in Section \ref{subsec:mdfcmd}: the stream can be seen as a remnant of an old satellite that has been recently disrupted."
+" As to the nature of the progenitor of the Aquarius stream, it could either be a dwarf galaxy or a globular cluster."," As to the nature of the progenitor of the Aquarius stream, it could either be a dwarf galaxy or a globular cluster."
+ The survival of this progenitor would depend on its concentration: it could either have been tidally stripped or have undergone complete disruption., The survival of this progenitor would depend on its concentration: it could either have been tidally stripped or have undergone complete disruption.
+" To search for possible globular clusters that the Aquarius stream could have been tidally stripped from, we performed a search of the Harris catalog of known globular clusters, selecting those (1996)with —1.5«[Fe/H]« —0.5, 1.5«Rac<9 kpc and Z«4 kpc, where the latter limits are taken from the model satellite orbit in Table 4.."," To search for possible globular clusters that the Aquarius stream could have been tidally stripped from, we performed a search of the \citet{Harris1996} catalog of known globular clusters, selecting those with $-1.5<\mathrm{[Fe/H]}<-0.5$ , $1.5<R_\mathrm{GC}<9$ kpc and $Z<4$ kpc, where the latter limits are taken from the model satellite orbit in Table \ref{tab4}."
+" We then compared the distribution of the clusters in 1,b,Vios to that of the model satellite stream (t,=—700 Myr), and found no globular clusters that match the simulation."," We then compared the distribution of the clusters in $l,\ b, \vrad$ to that of the model satellite stream $t_s=-700\,\myr$ ), and found no globular clusters that match the simulation."
+" Also, w Cen, with an l,b,Vio, of (309°,+15°,233kms!) and L,=—413kpckms'!, is not consistent with not-yet-phased-mixed scenario in Figure 8 and Table 4.."," Also, $\omega$ Cen, with an $l,\ b,\ \vrad$ of $(309^\circ,\ +15^\circ,\ 233\,\kms)$ and $L_z=-413\,\kpc\,\kms$, is not consistent with not-yet-phased-mixed scenario in Figure \ref{f8} and Table \ref{tab4}."
+ With our uncalibrated metallicities it is difficult to compare the MDF to that of w Cen., With our uncalibrated metallicities it is difficult to compare the MDF to that of $\omega$ Cen.
+" A high-resolution spectroscopic abundance study, such in Wylie-deBoeretal. is required well as further modelling of Aquarius) to (2010),,definitively understand(as if Aquarius is related to w Cen."," A high-resolution spectroscopic abundance study, such in \citet{Wylie2010}, is required (as well as further modelling of Aquarius) to definitively understand if Aquarius is related to $\omega$ Cen."
+ The progenitor of the Aquarius stream therefore is currently unknown., The progenitor of the Aquarius stream therefore is currently unknown.
+ In this paper we report the detection of a new halo stream found as an overdensity of stars with large heliocentric radial velocities in the RAVE data-set., In this paper we report the detection of a new halo stream found as an overdensity of stars with large heliocentric radial velocities in the RAVE data-set.
+ The detection is enabled by RAVE’s selection criteria creating no kinematic biases., The detection is enabled by RAVE's selection criteria creating no kinematic biases.
+" The fifteen member stars detected have Vi,=—199+27kms‘! and lie between"," The fifteen member stars detected have $\RV=-199\pm27\,\kms$ and lie between"
+ol molecules such as INCO from the nuclear region to the starburst ring )00)..,of molecules such as HNCO from the nuclear region to the starburst ring \citep{Martin09}.
+ The tentative detection of IICO in the eireunnuelear disk (COND) of 11063 (Useroοἱal.2004).. suggest a rough Η0Ο to HCO line intensity ratio of ~2—3. which turns into an abundance ratio in the range of 0.09—0.13. closer to the values derived in SD dominated galaxies.," The tentative detection of HCO in the circunnuclear disk (CND) of 1068 \citep{Usero04}, suggest a rough $\rm H^{13}CO^+$ to HCO line intensity ratio of $\sim 2-3$, which turns into an abundance ratio in the range of $\sim 0.09-0.13$, closer to the values derived in SB dominated galaxies."
+ This value at the CND can be significantly biased bv the emission [rom the star forming ring covered at hall-power by the 28” beam., This value at the CND can be significantly biased by the emission from the star forming ring covered at half-power by the $28''$ beam.
+" Therefore. it is clear that this ratio does not significantly decrease towards the nuclear AGN in thisgalaxy. with the (vpical angular resolution of 20""—30""."," Therefore, it is clear that this ratio does not significantly decrease towards the nuclear AGN in thisgalaxy with the typical angular resolution of $20''-30''$."
+ On the other hand. the ratio HCO/MOC is only à [actor of 2 lower in the starbursts galaxies than towards the nuclear AGN in 110683.," On the other hand, the ratio $\rm HCO^+/HOC^+$ is only a factor of 2 lower in the starbursts galaxies than towards the nuclear AGN in 1068."
+ From these observations. it is unclear whether photodissociation does play a major role in the AGN dominated center of 11068.," From these observations, it is unclear whether photodissociation does play a major role in the AGN dominated center of 1068."
+ Unfortunately. no observations of are available towards the SF ring in this galaxy.," Unfortunately, no observations of $^+$ are available towards the SF ring in this galaxy."
+ Similar to 11068. the obseured Seyfert 2 nucleus in 44945 is surrounded by a starburst ring more prominent (han in 11068 (Genzeletal.19983).," Similar to 1068, the obscured Seyfert 2 nucleus in 4945 is surrounded by a starburst ring more prominent than in 1068 \citep{Genzel98}."
+. The “μου ratio found in 44945 is even lower than that found in the other galaxies., The $^+$ /HCO ratio found in 4945 is even lower than that found in the other galaxies.
+ However. the fit to these lines was claimed to be very uncertain by Wangetal.(2004).," However, the fit to these lines was claimed to be very uncertain by \citet{Wang04}."
+. We have compared our observed Iractional abundances and abundance ratios with those predicted by the PPDHB model (Belletal.2006)., We have compared our observed fractional abundances and abundance ratios with those predicted by the PDR model \citep{Bell06}.
+. The PPDR code is a lime and depth dependent one dimensional PDR model that simultaneously. solve the chemistry. thermal balance and radiative transfer within a cloud more details)..," The PDR code is a time and depth dependent one dimensional PDR model that simultaneously solve the chemistry, thermal balance and radiative transfer within a cloud \citep[see][for more details]{Bell06}."
+ We have adopted a hydrogen density of LO?em.7. a radiation field C~5000 in units of Habing field. and a cosmic radiation rate of 10. '.," We have adopted a hydrogen density of $\rm 10^5\,cm^{-3}$, a radiation field $G_0\sim5000$ in units of Habing field, and a cosmic radiation rate of $10^{-16}$ ."
+ The high density of 10?em is derived [rom the multiline analvsis of CS and HCSN," The high density of $\rm 10^5\,cm^{-3}$ is derived from the multiline analysis of CS and $_3$ N"
+resolution calculations.,resolution calculations.
+ The resolution dependence of the high k turnover is also consistent with the expected Re3/4 dependence of the dissipation scale., The resolution dependence of the high $k$ turnover is also consistent with the expected $\mathcal{R}_{\rm e}^{-3/4}$ dependence of the dissipation scale.
+" With mean motions that are around 1/3 of the sound speed and transsonic fluctuations, one cannot simply reduce the SPH artificial viscosity parameters arbitrarily (such as the a=0.1 anda=0 calculations attempted by ?)), since this term is necessary to capture the physical dissipation that occurs due to the non-linear steepening of waves."," With mean motions that are around $1/3$ of the sound speed and transsonic fluctuations, one cannot simply reduce the SPH artificial viscosity parameters arbitrarily (such as the $\alpha = 0.1$ and $\alpha = 0$ calculations attempted by \citealt{bs11}) ), since this term is necessary to capture the physical dissipation that occurs due to the non-linear steepening of waves."
+" Such an approach may be adequate for very low Mach number (i.e., incompressible) calculations but it provides no easy answer at ~0.3c,."," Such an approach may be adequate for very low Mach number (i.e., incompressible) calculations but it provides no easy answer at $\sim 0.3 c_{\rm s}$."
+ Instead it is clear that to achieve similar results with SPH an improved viscosity switch is necessary in order to both capture non-linear steepening and shocks as well as reducing the viscosity to very low values where it is not needed., Instead it is clear that to achieve similar results with SPH an improved viscosity switch is necessary in order to both capture non-linear steepening and shocks as well as reducing the viscosity to very low values where it is not needed.
+" The switch proposed recently by ? represents the current state-of-the-art in this regard, essentially a thoroughly enhanced and improved version of the ? approach."," The switch proposed recently by \citet{cd10} represents the current state-of-the-art in this regard, essentially a thoroughly enhanced and improved version of the \citet{mm97} approach."
+" In particular, they show that they are able to simulate linear waves for over 50 periods with essentially no numerical dissipation, using the same parameters as would be applied in shock problems."," In particular, they show that they are able to simulate linear waves for over 50 periods with essentially no numerical dissipation, using the same parameters as would be applied in shock problems."
+" Thus, with an implementation of the ? switch it may be expected that significantly higher Reynolds numbers are achievable in SPH."," Thus, with an implementation of the \citet{cd10} switch it may be expected that significantly higher Reynolds numbers are achievable in SPH."
+" ? argue that “large errors in SPH's gradient estimate"" are responsible for the failure to reproduce a Kolmogorov-like turbulent cascade in their SPH calculations.", \citet{bs11} argue that “large errors in SPH's gradient estimate” are responsible for the failure to reproduce a Kolmogorov-like turbulent cascade in their SPH calculations.
+" Figure 2 demonstrates that this argument is incorrect, since we are able to obtain a k-5/? spectrum using only standard SPH gradient terms and a very similar SPH neighbour number to that employed in their preprint."," Figure \ref{fig:pspec} demonstrates that this argument is incorrect, since we are able to obtain a $k^{-5/3}$ spectrum using only standard SPH gradient terms and a very similar SPH neighbour number to that employed in their preprint."
+" However, we find that the appearance of a power-law inertial range in the power spectrum strongly depends on the Reynolds number employed in the calculations, requiring at least Re21500."," However, we find that the appearance of a power-law inertial range in the power spectrum strongly depends on the Reynolds number employed in the calculations, requiring at least $\mathcal{R}_{e} \gtrsim 1500$."
+" This explains the failure to produce a turbulent cascade in their SPH results, since the maximum Reynolds numbers they achieve are &600."," This explains the failure to produce a turbulent cascade in their SPH results, since the maximum Reynolds numbers they achieve are $\approx 600$."
+" With the ? viscosity switch employed in this Letter we estimate that we are able to achieve R.&6,000 at 256? particles which already brings the SPH results into much better agreement with the grid-based results shown in their work."," With the \citet{mm97} viscosity switch employed in this Letter we estimate that we are able to achieve $\mathcal{R}_{e} \approx 6,000$ at $256^{3}$ particles which already brings the SPH results into much better agreement with the grid-based results shown in their work."
+" Indeed, both ? and ? have already pointed out that using this switch could substantially improve SPH simulations of turbulence in galaxy clusters."," Indeed, both \citet{dolagetal05} and \citet{valdarnini11} have already pointed out that using this switch could substantially improve SPH simulations of turbulence in galaxy clusters."
+" It should be noted that ? do experiment with reduced viscosity parameters in their preprint, using either a fixed a=0.1 or the ? switch and also a run with zero viscosity (as we have already discussed in Sec. 3,,"," It should be noted that \citet{bs11} do experiment with reduced viscosity parameters in their preprint, using either a fixed $\alpha = 0.1$ or the \citet{balsara95} switch and also a run with zero viscosity (as we have already discussed in Sec. \ref{sec:av},"
+" it is not clear that one can simply reduce the parameters arbitrarily, so this approach is questionable — particularly the a=0 calculation)."," it is not clear that one can simply reduce the parameters arbitrarily, so this approach is questionable — particularly the $\alpha=0$ calculation)."
+" Indeed, both the a=0.1 and Balsara-switch calculations show a dramatic improvement in the power spectrum at large scales."," Indeed, both the $\alpha = 0.1$ and Balsara-switch calculations show a dramatic improvement in the power spectrum at large scales."
+ The authors dismiss this result because of a corresponding increase in power at k=100., The authors dismiss this result because of a corresponding increase in power at $k \gtrsim 100$.
+" However, the power at these scales is low amplitude (~10) and thus sensitive to all manner of numerical artefacts (e.g. the interpolation procedure as demonstrated in Fig."," However, the power at these scales is low amplitude $\sim 10^{-4}$ ) and thus sensitive to all manner of numerical artefacts (e.g. the interpolation procedure as demonstrated in Fig."
+ 4 of their preprint)., 4 of their preprint).
+" Indeed, we do not find the upturn in power at large k seen in their results (c.f."," Indeed, we do not find the upturn in power at large $k$ seen in their results (c.f."
+" Fig. 2)),"," Fig. \ref{fig:pspec}) ),"
+" most likely due to our improved power spectrum estimation — here computed by interpolating the SPH data to a 3D grid using the kernel and employing a Fast Fourier Transform, rather than the “nearest neighbour sampling"" procedure employed in their preprint."," most likely due to our improved power spectrum estimation — here computed by interpolating the SPH data to a 3D grid using the kernel and employing a Fast Fourier Transform, rather than the “nearest neighbour sampling” procedure employed in their preprint."
+" Finally, it is important to compare the Reynolds numbers achievable in numerical simulations to the physical Reynolds numbers in the problems of interest (c.f."," Finally, it is important to compare the Reynolds numbers achievable in numerical simulations to the physical Reynolds numbers in the problems of interest (c.f."
+ Sec. 1))., Sec. \ref{sec:intro}) ).
+" For the cold ISM, Re~ 105—107 which, though high, is not as high as is often assumed, and is certainly within reach of being resolved with currently achievable resolutions (c.f."," For the cold ISM, $\mathcal{R}_{\rm e}\sim 10^{5}$ $10^{7}$ which, though high, is not as high as is often assumed, and is certainly within reach of being resolved with currently achievable resolutions (c.f."
+ Sec. 2))., Sec. \ref{sec:reynolds}) ).
+ This implies at the very least that physical viscosity should be introduced into ISM turbulence simulations in the near future., This implies at the very least that physical viscosity should be introduced into ISM turbulence simulations in the near future.
+" In the ICM/IGM, Re~52 would seem to imply laminar flow, though very high estimates for Re (> 10:19) apply in the presence of magnetic fields."," In the ICM/IGM, $\mathcal{R}_{\rm e}\sim 52$ would seem to imply laminar flow, though very high estimates for $\mathcal{R}_{\rm e}$ $\gtrsim 10^{10}$ ) apply in the presence of magnetic fields."
+ Reaching such Reynolds numbers is not presently achievable with any numerical code., Reaching such Reynolds numbers is not presently achievable with any numerical code.
+" However, this may also imply that ultimately it is quite incorrect to try to simulate purely hydrodynamic ICM/IGM turbulence at high Reynolds number without taking into account more detailed physics, such as magnetic fields."," However, this may also imply that ultimately it is quite incorrect to try to simulate purely hydrodynamic ICM/IGM turbulence at high Reynolds number without taking into account more detailed physics, such as magnetic fields."
+" In this Letter we have emphasized the importance of accounting for the Reynolds number in numerical turbulence simulations, particularly in the subsonic regime where it is the main parameter controlling not only whether the flow is turbulent but also the character of such turbulence."," In this Letter we have emphasized the importance of accounting for the Reynolds number in numerical turbulence simulations, particularly in the subsonic regime where it is the main parameter controlling not only whether the flow is turbulent but also the character of such turbulence."
+" In particular differences in the Reynolds number can produce strong differences between calculations, independent of any numerical influences."," In particular differences in the Reynolds number can produce strong differences between calculations, independent of any numerical influences."
+" We have shown that use of viscosity switches is the key to higher Reynolds numbers in SPH turbulence calculations, producing spectra consistent with Kolmogorov theory."," We have shown that use of viscosity switches is the key to higher Reynolds numbers in SPH turbulence calculations, producing spectra consistent with Kolmogorov theory."
+" We are very grateful to the referee, Walter Dehnen, both for constructive feedback and prompt reviewing of this Letter."," We are very grateful to the referee, Walter Dehnen, both for constructive feedback and prompt reviewing of this Letter."
+ We, We
+"power spectrum of the power spectrum — windowed with a narrow filter gives the variation of the large separation with frequency, Av(v).","power spectrum of the power spectrum – windowed with a narrow filter gives the variation of the large separation with frequency, $\Delta\nu(\nu)$ ."
+" ? has shown that, with a dedicated comb filter for analyzing the power spectrum, it is possible to obtain independently the values of the large separation for the odd (J= 1) and even (/2 0,2) ridges."," \citet{Mosser10}
+ has shown that, with a dedicated comb filter for analyzing the power spectrum, it is possible to obtain independently the values of the large separation for the odd $l\!=\!1$ ) and even $l\!=\!0,2$ ) ridges."
+ We denote them by Avoqq and AYeven-, We denote them by $\Delta\nu_{\rm{odd}}$ and $\Delta\nu_{\rm{even}}$.
+ Proxies of the mode frequencies can then be integrated from the Av(v) frequency pattern., Proxies of the mode frequencies can then be integrated from the $\Delta\nu(\nu)$ frequency pattern.
+" In practice, they are derived from the correlation between the observed PDS and a synthetic spectrum based on the Av(v) pattern."," In practice, they are derived from the correlation between the observed PDS and a synthetic spectrum based on the $\Delta\nu(\nu)$ pattern."
+ The values of Avoag and Aveven for KIC 10273246 are given in Fig. [2]., The values of $\Delta\nu_{\rm{odd}}$ and $\Delta\nu_{\rm{even}}$ for KIC 10273246 are given in Fig. \ref{fig_autodeltanuridge_10273246}.
+ Notice the agreement between the proxies of the frequencies derived with the EACF and the frequencies of the (Table ))., Notice the agreement between the proxies of the frequencies derived with the EACF and the frequencies of the (Table \ref{Freq_Mulder}) ).
+ The EACF emphasizes the low values of Avoga as well as its large gradient at the low- end., The EACF emphasizes the low values of $\Delta\nu_{\rm{odd}}$ as well as its large gradient at the low-frequency end.