Correlated radial velocity and X-ray variations in \srco/\srcx\\ % Draft of \today
Duncan K. Galloway
We present evidence for approximately 400-d variations in % both the X-ray flux of \srcx\, and the radial velocity of \srco\ (V934 Her), the suggested optical counterpart of \srcx. The variations are correlated with the previously reported $\approx400$~d variations in the X-ray flux of \srcx, which supports the association of these two objects, % the red giant \srco\ with the X-ray source \srcx, as well as the identification of this system as the second known X-ray binary in which a neutron star accretes from the wind of a red giant. %wind-fed neutron star symbiotic-like binary. The \srco\ radial velocity variations %of order $1\ \kms$, can be fit with an eccentric orbit with period $404\pm3$~d, amplitude $K=0.75\pm0.12\ \kms$ and eccentricity $e=0.26\pm0.15$. There are also indications of variations on longer time scales $\ga2000$~d. We have re-examined all available ASM data following an unusually large X-ray outburst in 1997--98, and confirm that the 1-d averaged 2--10~keV X-ray flux from \srcx\ is modulated with a period of $400 \pm 20$~d. %$392\pm14$~d. The mean profile of the persistent X-ray variations was approximately sinusoidal, with an amplitude of $0.108 \pm 0.012$~ASM~$\cts$ (corresponding to 31\% rms). The epoch of X-ray maximum was approximately 40~d after the time of periastron according to the eccentric orbital fit. % The 1997--98 outburst, as well as the weaker X-ray % maxima which followed, all occurred 40--60~d after one of the predicted % times of periastron passage of the 404~d orbit. If the 400-d oscillations from \srco/\srcx\ are due to orbital motion, then the system parameters are probably close to those of %bears a strong resemblance to the only other neutron-star symbiotic-like binary, GX~1+4. We discuss the similarities and differences between these two systems.
preprint.tex
%% $Id: ms.tex,v 1.11 2002/07/03 19:03:16 duncan Exp duncan $
%%
%% Beginning of file 'preprint.tex'
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%% This is the ApJ letters paper on the ASM/optical observations of 1700+24
%%
%% Copied from ms.tex (resubmitted version) by D.Galloway Jul 3 2002
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\newcommand{\vdag}{(v)^\dagger}
\newcommand{\myemail}{duncan@space.mit.edu}
\newcommand{\xte}{{\it RXTE}}
\newcommand{\eps}{{\rm erg\,s^{-1}}}
\newcommand{\epcs}{{\rm erg\,cm^{-2}\,s^{-1}}}
\newcommand{\cts}{{\rm count\,s^{-1}}}
\newcommand{\kms}{{\rm km\,s^{-1}}}
\newcommand{\srco}{HD~154791}
\newcommand{\srcx}{4U~1700+24}
\newcommand{\asmend}{2002 April 18}
\newcommand{\asmendmjd}{52382}
\newcommand{\msun}{{M}_{\odot}}
\newcommand{\msyr}{\msun \ {\rm yr^{-1}}}
\newcommand{\mdot}{\dot{M}}

% asmfold results for data as of 2002 April 18
% IDL> asmfold,'x1704+240_current.1day',/plot,hifac=.1,ofac=32,perr=1000,/binerr
%   Peak  35.45 at P= 382.52 +/-   3.98 d, sig. 8.36e-13
%   Amplitude =    0.14312 +/-   0.01286 count/s (36.50% rms)
%   Epoch of maximum = (50374.23955 + n * 382.52) +/-  5.64 d
%   Mean period = 382.25 +/-  5.72 d
\newcommand{\pall}{383}
\newcommand{\pwrall}{35.45}
\newcommand{\sigall}{$8\times10^{-13}$}
\newcommand{\errall}{6}
%

% Results for data as of 2002 March 20
% asmfold results with ofac=16,perr=1000,/binerr:
% Peak power 34.52 @ 381.87, sig 2.1e-13; ampl. 0.12 count/s
% Epoch of maximum 50373.8+n*381.87
% 1000 trials with perr method gives 381\pm6 d
%             with pboot and ofac=32 method gives +/- 5.15 d
% With errthresh=0.5 we get
% peak power 51.05 @ 384.75, sig 1.2e-20; ampl. 0.14\pm0.01
% 1000 trials with perr method gives 383\pm4 d
%
% Excluding the outburst (>800 d) gives
% peak power 13.11 @ 378.56 d, sig 2.72e-4; ampl 0.11 count/s
% Epoch of maximum 51128.4+n*378.56.
% 1000 trials gives 370 \pm 30 d
% For the same time interval with errthresh=0.5 we get
% peak power 20.92 @ P= 391.18, sig 9.42e-8, ampl. 0.094 \pm 0.012
% epoch of peak at (51114 + n*391.18) +/- 8 d
% 1000 trials gives 390 \pm 20 d(!)
% 33000 bootstrap trials gives a lower bound to the significance of 4 sigma
%
% Excluding around the outburst [550,800] with errthresh=0.5 gives
% peak 21.66
%% You can insert a short comment on the title page using the command below.

\slugcomment{To appear in the Astrophysical Journal (2002)}

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\shortauthors{Galloway et al.}
\shorttitle{}


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\begin{document}

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\title{
Correlated radial velocity and X-ray variations in \srco/\srcx\\
% Draft of \today
}


%% Use \author, \affil, and the \and command to format
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\author{Duncan K. Galloway}
%            Jennifer L. Sokoloski\altaffilmark{1}, and
%            Scott J. Kenyon\altaffilmark{1}}
%, \and            Phil Uttley\altaffilmark{2}}
\affil{Center for Space Research,
  Massachusetts Institute of Technology, \\Cambridge,MA 02139}
\email{duncan@space.mit.edu}
\and
\author{J. L. Sokoloski and Scott J. Kenyon}
\affil{Smithsonian Astrophysical Observatory, 60 Garden St., Cambridge, MA 02138}

%% Notice that each of these authors has alternate affiliations, which
%% are identified by the \altaffilmark after each name.  Specify alternate
%% affiliation information with \altaffiltext, with one command per each
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% \altaffiltext{1}{}
%\altaffiltext{1}{\em Harvard/Smithsonian Center for Astrophysics}
%\altaffiltext{2}{\em University of Southampton, UK}


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\begin{abstract}
We present evidence for approximately 400-d variations in
% both the X-ray flux of \srcx\, and
the radial velocity of \srco\ (V934 Her), the suggested optical
counterpart of \srcx.  The variations are correlated with the previously
reported $\approx400$~d variations in the X-ray flux of \srcx, which
supports the association of these two objects,
% the red giant \srco\ with the X-ray source \srcx,
as well as the identification of this system as the second known X-ray
binary in which a neutron star accretes from the wind of a red giant.
%wind-fed neutron star symbiotic-like binary.  
The \srco\ radial velocity variations 
%of order $1\ \kms$, 
can be fit with an eccentric orbit with period $404\pm3$~d, amplitude
$K=0.75\pm0.12\ \kms$ and eccentricity $e=0.26\pm0.15$. There are also
indications of variations on longer time scales $\ga2000$~d.  We have
re-examined all available ASM data following an unusually large X-ray
outburst in 1997--98, and confirm that the 1-d averaged 2--10~keV X-ray
flux from \srcx\ is modulated with a period of $400 \pm 20$~d.
%$392\pm14$~d.  
The mean profile of the persistent X-ray variations was approximately
sinusoidal, with an amplitude of $0.108 \pm 0.012$~ASM~$\cts$ (corresponding to
31\% rms).
The epoch of X-ray maximum was approximately 40~d after the time of
periastron according to the eccentric orbital fit.
% The 1997--98 outburst, as well as the weaker X-ray
% maxima which followed, all occurred 40--60~d after one of the predicted
% times of periastron passage of the 404~d orbit.
If the 400-d oscillations
from \srco/\srcx\ are due to orbital motion, then the system parameters
are probably close to those of
%bears a strong resemblance to 
the only other neutron-star symbiotic-like binary, GX~1+4. We discuss the
similarities and differences between these two systems.
\end{abstract}

%% Keywords should appear after the \end{abstract} command. The uncommented
%% example has been keyed in ApJ style. See the instructions to authors
%% for the journal to which you are submitting your paper to determine
%% what keyword punctuation is appropriate.

\keywords{binaries: symbiotic --- X-rays: binaries --- stars: neutron}


%% From the front matter, we move on to the body of the paper.
%% In the first two sections, notice the use of the natbib \citep
%% and \citet commands to identify citations.  The citations are
%% tied to the reference list via symbolic KEYs. The KEY corresponds
%% to the KEY in the \bibitem in the reference list below. We have
%% chosen the first three characters of the first author's name plus
%% the last two numeral of the year of publication as our KEY for
%% each reference.

\section{Introduction}

The X-ray source \srcx\ (also known as 2A 1704+241; $l=45\fdg15$,
$b=32\fdg99$) was first discovered in {\it Ariel V}\/ scans for
high-latitude X-ray sources \cite[]{2acat}, and by the $Uhuru$ (SAS A)
X-ray observatory \cite[]{4ucat}.  It
%(as measured by a range of X-ray missions) 
has
%exhibited multiple-year periods of quiescence, with 
a typical flux of (1--$10)\times10^{-11}\ {\rm erg\,cm^{-2}\,s^{-1}}$
\cite[2--10~keV, equivalent to 0.5--5~mCrab as measured by a range of
X-ray missions;][]{masetti01}, and experiences occasional outbursts.
During a 100-d X-ray high state in 1997--98,
\srcx\, was observed by the {\it Rossi X-ray Timing Explorer}\/
\cite[\xte;][]{xte96}.   This high state was the larger of
two known outbursts of this system, and the X-ray flux reached
a peak of 40~mCrab.
% Combined Proportional Counter Array (PCA, 2--60~keV) and High-Energy X-ray
% Timing Experiment (HEXTE, 20--250~keV) spectra accumulated during the
% outburst were reasonably well fit by a combination blackbody and power
% law, with blackbody $kT=1.3$--2.0~keV and spectral index $\alpha=2.4$--2.0
% in quiescence and outburst respectively.
The X-ray spectrum of \srcx\, is similar to that of other accreting
neutron stars, and generally requires a blackbody with a temperature
of 0.9--1.3~keV plus a hard component
%to account for excess emission
at higher energies for an acceptable fit.  
The X-ray spectral hardness and the lack of UV continuum imply that
the compact object is almost certainly not a white dwarf
\cite[]{garcia83}. A black hole companion is also unlikely since the thermal
X-ray component in black hole candidates is typically an order of
magnitude lower in temperature than that measured in \srcx. 
However, no rapid  quasi-periodic or coherent X-ray  
variations have been confirmed \cite[]{masetti01}, which is
rather unusual for an accreting neutron star.

Based on {\it Einstein}\/ position measurements, \cite{garcia83} proposed
the association of \srcx\ with the 8th magnitude M2 III star \srco.  In a
recent re-analysis of {\it ROSAT}\/ HRI data, however, \cite{mg01} found
only a 10\% chance that the two sources are associated, but did not
propose an alternative candidate.  Previous optical spectroscopy ruled out
radial velocity variations with amplitudes greater than $5\ \kms$ on time
scales of 40 min to $\sim 1$~yr \cite[]{garcia83}, strongly suggesting
that either the orbital period is significantly longer than one year, or
that the system is viewed very close to face-on (inclination angle
$i\approx0$).

\srco\ may be only the second symbiotic-like binary known to
contain a neutron star, after the well-studied M-giant/X-ray pulsar
system V2116~Oph/GX~1+4.  However, these systems are quite different in
X-rays (Table \ref{comptable}).  Furthermore, the optical spectrum of V2116~Oph
exhibits extremely strong, variable emission lines, probably powered by UV
photons originating from an accretion disk \cite[]{chak97:opt}.  In
contrast, \srco\ has much weaker line emission.  In fact, optically
the spectrum is very close to that of an isolated M giant
\cite[e.g.][]{garcia83}.

Here we describe optical and X-ray observations which go some way towards
resolving the uncertainties surrounding this system.  We compare the
properties of \srco/\srcx\ with those of V2116~Oph/GX~1+4, and discuss
what may be inferred about the former in light of our results.

%%%%
\centerline{\epsfxsize=8.5cm\epsfbox{f1.eps}}
  \figcaption{Radial velocity measurements of \srco\ and
the corresponding $404\pm3$~d orbital solution. The top panel shows the
radial velocity measurements versus time and the estimated $1\sigma$
uncertainties.  The bottom panel shows the measurements folded on the
orbital period.  In both panels the eccentric orbital model is overplotted
as a solid line (see Table \ref{orbtbl}).
 \label{orbit} }
%%%%
\bigskip

\section{Observations}

\subsection{Radial velocity variations} \label{sec:radvel}

Echelle spectra of \srco\ were obtained using intensified Reticon
detectors on the 1.5-m Tillinghast Reflector at the Whipple Observatory on
Mt. Hopkins, Arizona, and also the 1.5-m Wyeth reflector at Oak Ridge
Observatory, Massachusetts. The observations were made between 1982 April
12 and 1996 October 26, with the 18 observations before 1992 October 7
taken at Mt. Hopkins, and those since taken at Oak Ridge.  The instruments
are described in detail by \cite{latham92}. The spectra cover a 44~\AA\
bandpass centered around 5200~\AA, with a resolution of approximately $8.5\
\kms$. Each observation consisted of a 10--20 minute integration of \srco,
with 90~s Th-Ar comparison scans before and after.  The positions of
around 30 strong spectral lines in each calibration scan were fit with a
fifth-order polynomial, and the pre- and post-observation scans were then
combined to calibrate the wavelength scale for each spectrum. Typical
residuals to the polynomial fit were 0.5~pixel, corresponding to 0.01~\AA,
or $0.5\ \kms$. Approximately 25 lines were common to the entire set of
wavelength solutions.  The radial velocity was measured by
cross-correlating each observed spectrum with an M-star template spectrum
\cite[e.g.][]{td79} and measuring the mean offset.  Quadratic
least-squares fits to the peak in the cross-correlation function give a
radial velocity accurate to better than $1\ \kms$.

% asmfold,'symbiot.orb',/plot,pmax=1000,/binerr,mixboot=33000
% Peak  10.54 at P= 408.50 +/-   7.86 d
% Running bootstrap simulations with   33000 trials, please wait...
% Est. significance from bootstrap simulations is 5.45e-04 ( 3.27 sigma)
The radial velocity measurements show significant variations of
approximately $1\ \kms$ about the mean systemic velocity of $-48.7\pm0.1\
\kms$ (Figure \ref{orbit}). We searched the data for periodic signals in
the range 50--1000~d by calculating the Lomb-Scargle periodogram
\cite[e.g.][]{nr}, and found a marginally significant peak near a period
of 410~d (Lomb-normalized power 10.5; estimated significance from
bootstrap simulations, $3.3\sigma$).  Fitting an elliptical orbit to the
data gave a solution with orbital period $404\pm3$~d and radial velocity
amplitude $K=0.75\pm0.12\ \kms$ (Table \ref{orbtbl}; solid lines in Figure
\ref{orbit}).  The reduced-$\chi^2$ ($\chi^2_\nu$) for the fit was 1.26,
indicating a marginally acceptable fit. While the best-fit eccentricity
was 0.26, the deviations from circularity were significant at less than
the $2\sigma$ level.
% According to the Lucy-Sweeney test, a non-zero $e$ is significant at the
% 90\% confidence level.]
% eccentricity $e=0.4\pm0.2$

There was also evidence for variations in the radial velocity on
longer time scales.  We calculated $\chi^2_\nu$ for an eccentric orbit
%is plotted 
as a function of trial period for periods up to 4000~d.
%in Figure \ref{orbit}, bottom panel. Local minima where the curve
%approaches $\chi^2_r = 1$ are 
%indicative of significant power at the relevant period; 
While the 404~d signal gave rise to a $\chi^2_\nu$ value close to 1, lower
values were achieved at substantially longer periods around 1930~d
($\chi^2_\nu=1.04$), 2370~d (1.02) and 3700~d (0.91).  
%These orbital solutions are
%also plotted in Figure \ref{orbit}, upper panel (dotted/dashed lines).
The irregular spacing of the measurements did not allow us to distinguish
clearly between these periodicities; eccentric orbital solutions with any
one of the three periods can account for the long-term radial velocity
variations.  The observations do not cover more than one full cycle for
these longer-period signals, so more observations are needed to
determine whether they are true periodic variations or aperiodic trends
related to irregular red-giant variability.

\subsection{RXTE/ASM observations} \label{sec:xte}

%first by xxx, and then more recently <- couldn't find any other ref. 
\cite{masetti01} noted tentative evidence for an approximately
$400$~d periodicity in \xte\/ All-Sky Monitor (ASM) data.  To
investigate the 400-d periodicity, we obtained 1-d averaged ASM
measurements of \srcx\ between 1996 January 5 and
\asmend\ (MJD 50087--\asmendmjd) from the ASM WWW site at
{\tt http://xte.mit.edu}.  The mean quiescent ASM count rate was $0.2\
\cts$.  During the flare observed between 1997 September--December
(MJD 50700--50800), the 1-d averaged rate reached $2.5\ \cts$ (Figure
\ref{asm}, top panel).

% asmfold results for data as of 2002 April 18
% IDL> asmfold,'x1704+240_current.1day',/plot,hifac=.1,ofac=32,perr=1000,
%        /binerr,include=800
%   Peak  14.05 at P= 382.64 +/-   3.06 d, sig. 1.09e-04
%   Amplitude =    0.11674 +/-   0.29090 count/s (32.29% rms)
%   Epoch of maximum = (51139.59839 + n * 382.64) +/- ***** d
%   Mean period = 384.17 +/- 26.56 d
To avoid a biased period measurement due to the presence of the 1997--98
outburst, we initially calculated a Lomb-normalized periodogram over the
interval from the end of the outburst (around 1998 March 28) to \asmend.
The result was a peak with Lomb-normalized power 14.05 (equivalent to
3.9$\sigma$), at $383\pm30$~d.  We estimated the period
uncertainty by folding the ASM data on the detected period, and
calculating 1000 sets of statistically equivalent data sets with the same
time sampling, average profile and noise properties per bin as measured in
the original data set. The variance in the detection periods of the
simulated data sets is then a measure of the uncertainty in the actual
detection.

% adding errthresh=0.5 gives
%   Peak  24.04 at P= 392.04 +/-   6.43 d, sig. 4.21e-08
%   Amplitude =    0.09852 +/-   0.01191 count/s (28.43% rms)
%   Epoch of maximum = (51116.46778 + n * 392.04) +/-  7.68 d
%   Mean period = 394.96 +/- 14.10 d
% Folding on the best (ASM) period of 400d, and excluding post-outburst,
% gives (see inset.pro)
% Sine amplitude =    0.10754 +/-   0.01186 count/s (30.98% rms)
% Total RMS = 32.42%
To obtain a more precise estimate of the quiescent X-ray periodicity, we
excluded those measurements with errors of $>0.5\ \cts$.  Around 13\% of
the 1-d averaged ASM measurements had errors above this threshold,
possibly due to source variability on time scales less than a day, but
more likely arising from low-significance detections due to solar
contamination or other sub-optimal observing conditions.  These
low-significance measurements appeared to introduce a bias to the period
measurements, and also tended to contribute disproportionately to the
estimated uncertainty in the period.  A Lomb-normalized periodogram
calculated on this subset of data following the outburst resulted in a
peak power of 24.04 (estimated significance $4.2\times10^{-8}$, equivalent to
$5.5\sigma$) at $392\pm14$~d.
% The standard deviation of the period measured from 1000 simulated
% datasets was 14~d.
To confirm the above results,
we used the phase-dispersion minimization technique of \cite{stell78}
and folded the ASM light curve on a grid of periods between 200 to 1000 d.
We found that the X-ray flux was most likely modulated with a period of
$404 \pm 13$ d.  Combining the results
from the latter two period-search techniques,
% our two most reliable measurements (the
% periodogram with high-significance data, and the PDM period determination,
% which takes into account the flux error bars, in both cases using the
% post-outburst data only),
we determine the period of the X-ray oscillation
to be $400 \pm 20$~d with the epoch of maximum at MJD $51140 \pm 10$~d.
The mean profile of the post-outburst ASM measurements folded on the 400~d
period is also shown in Figure \ref{asm} (inset, bottom panel). The
profile was approximately sinusoidal with an amplitude of
$0.108 \pm 0.012$~ASM~$\cts$, or 31\% rms. With the periastron time from the
orbital fit used as a reference phase for the fold, we find that on
average the X-ray maximum occurs 0.1 in phase ($\sim40$~d) after
periastron.

For comparison, we also calculated an oversampled Lomb-normalized
periodogram of the entire ASM data set (including the 1997--98 outburst),
which revealed a significant periodicity at \pall~d (identical to that
measured following the outburst).  
% The peak power was \pwrall, with estimated formal significance \sigall.
% Naturally, it would be preferable to use the entire ASM data set in order
% to measure the time variability of the X-ray emission from the source.
% However,
The 1997--98 outburst was unusual; no similar events were
observed by the ASM in $\approx6.5$~yr of monitoring. Thus, it may have
occurred due to a mechanism distinct to some extent from the periodic
behavior of the quiescent emission. Furthermore, there is evidence that
the behaviour of the persistent emission prior to the outburst was
different than afterwards. Although only a little more than one 400-d
interval was covered before the outburst, the persistent flux was weakly
anticorrelated compared to the mean profile measured following the
outburst.  Applying the same error threshold ($\sigma<0.5\
\cts$) to the full data set resulted in an increase in both the peak
power and period, to 51.1 (estimated significance $1.2\times10^{-20}$ or
$>8\sigma$) at 385~d. 
% The Monte-Carlo method of error determination is problematic for the
% full ASM dataset, since in the simulated data the enhanced count rate
% from the single outburst peak is spread across all the times
% corresponding to phase maximum (the method assumes a profile which is
% constant in time), thus exaggerating the underlying periodic variation of
% the persistent emission.
%If the orbit is slightly eccentric, 
The shorter measured periods using the full ASM data set may be attributed
to aliasing resulting from a delay between the nearest phase of X-ray
maximum for the persistent emission and the peak of the 1997--98 outburst.
If \srcx\ is similar to wind-accreting neutron stars in high-mass
binaries, we may expect large outbursts to correspond only approximately
with periastron phase \cite[e.g.][]{swr86}.  In fact, the peak of the 1997
outburst falls around MJD 50780, approximately $40$~d after the nearest
epoch of maximum flux according to the 400-d periodicity detected in
the post-1997 outburst ASM measurements, and approximately $70$ days after
periastron passage according to the orbital solution discussed in
\S\ref{sec:radvel} (Figure \ref{asm}).
% JENO: might add this bit to the discussion??
% In
% addition, the X-ray modulation appears to only be present after the
% outburst.  Such a phenomenon could occur if, for example, X-ray
% illumination of the red giant during the outburst triggered an episode
% of enhanced mass loss from the red giant.
% [*ref. King or some other papers
% %on illumination-induced mass loss?], 
% The accretion onto the neutron star would then be modulated at the
% orbital period 
% %with a larger amplitude 
% after the outburst due to the increased contrast between the nebular
% density at apastron vs. periastron. 
%[*But material leaving the red
%giant with a velocity of roughly 40 km/s should be lost from the binary
%system on a time scale of months.  So why does the X-ray modulation
%persist for over 4 years?  Could the red giant just have moved into a
%phase of persistently higher mass loss beginning with a burst in
%1997?]
% (assumingan eccentric orbit).

%%One of the
%%calculated times of phase maximum was synchronous with the 1997
%%outburst, while at several of the other times less significant
%%increases in the ASM count rate was observed.  
%Consequently, the resulting standard
%deviation of \errall~d of the period detections of the simulated
%datasets must be considered a lower limit to the actual uncertainty on
%the period for the full ASM dataset.  


%The outburst could have occured at any
%orbital phase, and not necessarily near periastron.
%As we have argued, the 1997--98 outburst may arise from a somewhat
%different mechanism than the quiescent flux modulation.  
%Qualitatively it is understandable how 
%The presence of the outburst in the flux history may significantly
%bias the period measurement, particularly if the phasing of the
%outburst is slightly different from the phase of maximum of the
%persistent modulation.  

%While the periodicity was consistent with that obtained from the full
%dataset, the estimated uncertainty was much larger at
%30~d. 

%We note that applying the same error threshold on the full
%dataset has a similar effect, in that the period was longer and the
%estimated error was smaller.

% For the full dataset, then, we adopted a second ``bootstrap'' method in
% which randomly-selected subsets of the data with half the number of
% measurement points were analysed in the same manner as the full dataset.
% The standard deviation of the period measurements from 1000 of these
% subsets was 6~d; if the uncertainty scales according to $\sqrt N$, where
% $N$ is the number of points (as might be expected), then the estimated
% uncertainty for the full dataset is $\approx4$~d.



%% old text
%% Times of periastron:
%%        49105.302       49508.618       49911.933       50315.249
%%        50718.565       51121.880       51525.196       51928.512
%We selected those measurements with estimated error $<0.5\ \cts$ and
%computed the Lomb-normalized periodogram.  A highly significant peak at a
%period of 384.6~d was found, with power 46.03. The mean profile was
%approximately sinusoidal with an amplitude of $0.11\ \cts$ (corresponding
%to 30\% RMS). We repeated the analysis with the same data averaged over
%intervals of up to 5~d. Although the period of the peak did not vary, the
%peak power decreased with increasing window length, to 21.3 at 5~d.
%Bootstrap simulations suggest that this result is significant at
%$>4\sigma$ significance.  We note that the 1997 outburst occurred just
%after one of the predicted times of periastron passage according to the
%403~d model (50720 MJD).

\section{Discussion}

%\subsection{Correlated Variations}

We have found evidence for significant variation in the radial
velocity measurements from \srco, which we fit with an eccentric orbit
with period $404\pm3$~d.  We have also found evidence for a
significant variation in the X-ray flux from \srcx.  Our best
estimate of the X-ray period is
%$392\pm14$~d  
$400 \pm 20$ d.  Hence, 
%we conclude that 
the periods measured from the radial velocity and ASM X-ray flux
variations are consistent.
%between 380--390~d.
%We adopt the value of $392\pm14$~d as the best estimate of
%the periodicity at which the quiescent ASM flux varies,   
It is compelling to suggest that these two periodicities arise from a
common source.
% While the most precise period measured from the full ASM
% dataset appears to be inconsistent with the period from radial velocity
% measurements to a high degree of significance, simple sensitivity tests
% provide evidence for the dependence of the results on the presence of the
% 1997 outburst.
%Indeed, the 1-d ASM averages with error $<0.5\ \cts$ since the end of the
%outburst suggest a significant modulation of the quiescent flux with
%period $392\pm14$~d, which (thanks to the relatively large
%uncertainty) is consistent with the period obtained from radial velocity
%measurements. 
%However, the full ASM dataset gives a shorter periodicity at
%\pall~d which (given the estimated lower limit on the error) may be
%inconsistent with the period obtained from the radial velocity
%measurements at a significance level of at most $3\sigma$.
% Since outbursts like that observed in 1997--98 do not repeat on
% $\sim400$~d time scale (recall that only one has been observed over the
% 7-year lifetime of {\it RXTE}) it seems likely that such events
Naturally, it is also possible that the agreement between the two
periodicities is merely coincidence, and that \srco\ and \srcx\ are 
unrelated.  Estimating the probability of such a coincidence depends
upon the origin of the periodicities in each star, as well as the
intrinsic distributions of these phenomena.
% (which are presumably at least potentially measurable given a large
%enough sample of similar sources). 
In practice, it is not possible to make such an estimate with much
confidence.
%Since a variety of reasonable assumptions about the origin of the two
%periodicities (see below) tend to lead to the measured values
%representing extreme cases 
However, 
%as mentioned above, 400 d is long for a red giant pulsation period.
%Hence, it is already unlikely to find a pulsating red giant with such
%a period, much less so close to an X-ray source oscillating with the
%same frequency.  The other option, that HD 154791 is a $P=400$-d
%binary, spatially coincident with an X-ray source that also has a
%400-d period, also seems highly unlikely.  
it is clear that the likelihood of an unrelated red giant and X-ray
source being at nearly the same measured position at high galactic latitude,
and also both having a periodicity 
%%%%
\centerline{\epsfxsize=8.5cm\epsfbox{f2.eps}}
 \figcaption{\xte/ASM measurements of \srcx. The top
panel shows the 1~d averaged ASM measurements between 1996 January 5 and
\asmend\ (MJD 50087--\asmendmjd).  The horizontal lines (-- --) bracket
the estimated times of phase maximum for the best period estimate (400~d;
see \S\ref{sec:xte}).  The open triangles indicate the predicted times of
periastron passage for the eccentric orbital fit ($P_{\rm orb}=404$~d;
Table \ref{orbtbl}).  The bottom panel shows the periodogram calculated
from a subset of the ASM data after the end of the outburst (1998 March
28), where the measurement error was $<0.5\ \cts$. The dashed vertical
line shows the candidate 400~d period.  The inset shows the folded ASM
profile using the time of periastron from the orbital fit as the reference
phase.
 \label{asm} }
%%%%
\bigskip
\noindent
that would be unusual for either type
of source separately, is extremely small.  The detection of a 400-d
modulation from both stars therefore supports the original suggestion that
HD 154791 is the optical counterpart of \srcx.

% old text
% While these two periodicities appear to disagree at around the $4\sigma$
% level, we note that 1) the \xte/ASM flux history only spans around 5
% cycles of the 385/403~d period; 2) only one large outburst has been
% observed over this time. We expect that a much more accurate estimate of
% the period would result if another outburst was observed by the \xte/ASM.
% Furthermore, the small uncertainty on the period from the radial velocity
% measurements is due to the long baseline, but the measurements between
% MJD~45040--45760 do not give good phase coverage and are not particularly
% well correlated with the orbital model (Figure \ref{orbit}, top panel).
% The uncertainty on the period estimate using just the data after MJD~48400
% is four times higher at 8~d, which would reduce the discrepancy between
% the radial velocity and ASM periodicities to $2.3\sigma$.

\subsection{Cause of the Correlation} \label{sec:causeofcor}

%In that case, the association of \srcx\ and \srco\ would appear to be
%confirmed. 
There are several plausible mechanisms for a correspondence between
the radial velocity and the persistent X-ray flux.
%, depending primarily on the origin of the radial velocity
%variations.  
We consider two possibilities for the origin of the 404-d period in
the \srco\, radial velocity curve: orbital motion, and pulsational
variation of the M-giant.  Precise optical photometric monitoring
could distinguish between these two possibilities, since the maximum
brightness occurring at any phase besides that of maximum recessional
velocity would favour orbital motion.

If the radial velocity variations are due to orbital motion, the 400-d
X-ray flux modulation may arise from changes in the accretion rate
onto the neutron star as it follows an elliptical orbit around the
companion and moves through different nebular densities. Such
modulation is common in high-mass wind-accreting X-ray binaries
\cite[e.g.][]{swr86}.  Moreover, the times of periastron passage in
our best-fit orbital model are very well correlated with the times of
increased X-ray flux.  For a 1.4~$M_\odot$ neutron star and red giant
mass $\la 2.5\ M_\odot$, the 404~d period and low amplitude of the
radial velocity variations imply that the inclination is
$i\la2^\circ$.  Although the {\it a priori} probability of observing
such a system is $\la7\times10^{-4}$, such a low inclination could
explain the lack of X-ray pulsations or QPOs from the source (see 
\S\ref{gxcomp}).
%On the other hand, this does not explain 
If the 400-d variation is in fact due to orbital motion, the
longer-term variations in the radial velocities may be related to
irregular variations in the red giant \cite[]{pwh01}.

Alternatively, red-giant pulsations with $P=400$~d could 
%If, on the other hand, the radial velocity variations are due to
%red-giant pulsation, it may be possible that these pulsations 
modulate the stellar wind, and hence the accretion rate onto the neutron
star.  The IRAS fluxes of HD~154791 ($f_{12}=2.77$, $f_{25}=0.736$,
$f_{60}<0.4$, $f_{100}<1.0$, where $f_{12}$ is the flux density $f_\nu$ at
12 microns in Jy) indicate that is it not a Mira variable
\cite[]{kenyon88}.  However, non-Mira M-giants also pulsate.  Although
typical fundamental pulsation periods lie in the range 20--200~d,
variations with a period up to an order of magnitude larger are also seen
\cite[]{pwh01}.  If the 400-d variations are due to pulsation, one of the
slower trends observed in the radial velocity variations could be related
to orbital motion.  In this case, the possible range of orbital
inclinations becomes much more likely.  An orbital period of 2370~d gives
an inclination of 4--$6^\circ$ while for a period of 3700~d the likely
range is 5--$7^\circ$.  Also, the low accretion rate onto the neutron star
would be more understandable if the orbital period is longer than 400 d.
However, a longer orbital period alone cannot account for the discrepancy
in luminosity between 4U~1700+24 and GX~1+4.

%, even at the closest possible distance for the latter source.  
Finally, the shape of the 400-d radial velocity variation argues
against a red-giant pulsation.  Pulsation light curves are typically
non-sinusoidal (in the sense that more time is spent with brightness
below the median value than above it), and similar behavior would be
expected for the radial velocity variations.
%[*what about the expected amplitude of velocity
%variations from pulsation?]  

\subsection{Comparison with GX~1+4}
\label{gxcomp}

If the 400~d periodicity in the \srco\, radial velocities is due to
orbital motion, then the orbital periods of \srco/\srcx\ and GX~1+4 may be
quite similar (see Table \ref{comptable}). This similarity makes the
contrast between their other observed properties puzzling. The X-ray
luminosity of GX~1+4 appears to be two to three orders of magnitude
greater than that of \srcx.  In GX~1+4, there is clear evidence for an
accretion disk \cite[]{chak97:opt}, whereas in \srcx\, there is no
measurable UV continuum \cite[]{garcia83}.  GX~1+4 has a rich optical
emission-line spectrum, and \srco\, shows very little in the way of
emission lines.  Finally, GX~1+4 is a 2-minute X-ray pulsar, whereas no
coherent or quasi-period oscillations have been convincingly detected from
\srcx.
%In particular, why is the X-ray luminosity of \srcx\ so much
%lower?

One explanation for the differing X-ray luminosities is a larger
mass-loss rate from the late-type giant in GX~1+4.  \cite{chak97:opt}
find that the M6 giant in GX~1+4 is probably near the tip
of the first-ascent red-giant branch.  \srco\, has a spectral type M2,
%[*most of the time?, but see Dal Fiume et al. 1990 (Il Nuovo Cimento,
%13, 481), who find M5-7 in observations between 1984-89],
and so would be expected to be losing mass at a lower rate. 
The relative lack of optical/UV emission lines suggests that either
the ionized nebula in \srco/\srcx\ is substantially smaller and/or
less dense than in GX~1+4, or there are not enough UV photons to power
nebular line emission in \srcx/\srco.  Either of these conditions can
plausibly arise if the mass loss from the companion is much smaller.
%[*However, the X-ray luminosity of \srcx\, implies just $10^{-12}
%\msyr$ onto a 1.4$\msun$ NS, which is too low for any reasonable red
%giant wind at a distance of ~an AU.  Masetti et
%al. claim this rate is reasonable for an efficiency factor of $\eta
%\approx 10^{-4}$, but $\eta$ should not be so low for accretion from a
%slow wind.  A low $\eta$ is typical for accretion from a Be-star wind
%(check), but not the red-giant wind in a symbiotic.  Is it crazy to
%consider a black hole, where 
%some of the energy can be advected into the hole?]  
% Another factor contributing to a higher mass-transfer rate in GX~1+4 is
Moreover,
if the binary separation is smaller in
GX~1+4 than in \srcx/\srco\ a larger fraction of the red giant wind could
be accreted.  The higher X-ray luminosity in GX~1+4
could also illuminate the red giant and increase the mass loss from
the red giant further, in a sense producing a feedback effect.  The
observed variability of the optical spectrum of V2116~Oph is evidence for
the significant role played by X-ray heating. On the other hand, in
\srcx/\srco\ the X-ray luminosity is perhaps low enough that very little
illumination of the red giant or enhanced mass loss is expected.
%[*then why the large change in spectral type from M2 to M5-7 to M2?
%Do we not believe this?].  
Even during the 1997 X-ray outburst, \cite{tomasella97} note that
the optical spectrum was unchanged.
%[* The accretion rate could also be low because of some
%sort of barrier (Masetti?), i.e. if the accretion rate has dropped and the NS is
%spinning down (the co-rotation radius is less than the magnetospheric
%radius where the pressures balance).
% Finally, if the magnetic field is substantially lower in
% \srco/\srcx\, the magnetosphere may be smaller than in GX~1+4 and thus
% the compact object could present a smaller target for the
% expelled material.
% [*For this to be a factor in explaining the
% luninosity difference, the size of the GX~1+4 magnetosphere must be
% larger than the Bondi accretion radius.  Is it?]

%%[*but the UV line emission at least seems variable, since Garcia et
%%al. 1983 found high-ionization emission lines, and later obserations
%%have not] 
%Radio observations in both sources can provide only upper limits on
%the nebular emission
%\cite[]{wend95,fen00}, [*although Schaefer (1986; PASP, 98,556) finds
%IR excess in \srco\, and Gaudenzi \& Polcaro (1999; A\&A 347,473) find
%anomalous reddening which they ascribe to intrinsic nebulosity].  As
%discussed above, \srco\, could have a lower mass-loss rate than the
%red giant in GX~1+4, and so a smaller/ less dense nebula would be
%expected.  There is also no continuum UV emission detected from \srco[*ref],
%%and  hard X-rays do not produce much line emission in GX~1+4
%%(Chakrabarty \& Roche 1997 - go over this), 
%so both factors could play a role in the low emission line strengths. 

%The optical lines in GX~1+4 are powered by UV emission from an
%accretion disk \cite[]{chak97:opt}, whereas \srco\, shows a derth of
%UV emission.
%%[*except that Garcia et al. 1983 report high-ionization-state emission
%%lines, which later disappear (Dal Fiume et al. )].  
%This discrepancy suggests that \srcx\ has either a much smaller (or
%completely absent) disk, a disk that is severely truncated by a strong
%magnetic field, or a disk that is faint because of the low accretion
%rate.  A neutron star (with a weak magnetic field) presents a very
%small target, so it is difficult for accreted wind material to reach
%the NS without orbiting it first and forming a disk.
%%[*Using the calculation of Livio (ref), the wind velocity relative to
%%the NS would have to be greater than approximately 200 km/s (check
%%and include scalings) to inhibit disk formation.  Since typical wind
%%velocities from a red giant are on the order of tens of km/s, and a NS
%%in a 400-d orbit around a 1.5 $\msun$ red giant will be moving at xx
%%km/s, one would expect a disk to form.  ]
% For a disk to extend all the way to the NS surface with a low
% accretion rate onto the NS such as we infer in \srcx, the magnetic
% field must be extremely [*unreasonably?] low.  On the other hand, if
% the magnetic field is more typical for a NS, then the accretion flow
% becomes chanelled along magnetic field lines before it reaches the NS
% surface.  This effect makes it potentially harder to form a disk,
% since the magnetic radius could be larger than the accretion radius.
% If a disk does form, the inner parts will be truncated, and the disk
% will therefore appear cooler [*possibly explaining the lack of UV
% emission].
%One may also ask if there is any reason for the
%disk to be especially small or faint.  If the disk is small, there may
%be significant differences in the way circumstellar material is
%captured by the neutron star.  
% Evidence from HEXTE for magnetic fields?

% Consideration of the magnetic field in \srcx\, brings us finally to the
% issue of the lack of rapid pulations.
Finally, the lack of pulsations or QPOs in \srcx\ remains puzzling.
\cite{garcia83} and \cite{mg01}
proposed a 900-s QPO, but this QPO was not confirmed in high-quality
\xte\/ data \cite[]{masetti01}.  High-precision optical photometry has also
revealed a lack of pulsations (down to a limit of 0.5\% fractional
variation), or rapid variations of any kind \cite[]{sbh01}.  This behavior
is again in contrast to GX~1+4, where \cite{jab97} find an oscillation in
the optical emission with an amplitude of from 0.4 to 4.5\%, plus
stochastic variations.  Given that the optical light from GX~1+4 has a
significant contribution from an accretion disk, the lack of rapid optical
variations from \srco\, provides evidence that any contribution from a
disk in \srcx\, is negligible in the optical regime as well as the UV
regime.  The lack of X-ray pulsations is expected if the 400-d radial
velocity variations are due to orbital motion.  As discussed in 
\S\ref{sec:causeofcor}, the low amplitude of the radial velocity
variations requires $i\sim 0$, in which case any magnetic hot spots on the
neutron star will be continuously in view (assuming the NS spin axis is
aligned with angular momentum axis of the binary).  For higher and {\it a
priori} more probable values of $i$, the non-detection of pulsations
may be explained if the
% requires either that the neutron star spin is very slow, or that the
accretion flow is not appreciably disrupted by the magnetic field of the
neutron star above the surface \cite[e.g.][]{gl79a,gl79b}, and spreads the
accreting material more or less evenly over the neutron star surface.
Given the extremely low luminosity (and hence accretion rate) of this
source, standard calculations of ram and $B$-field pressure balance would
require that the surface magnetic field strength of the neutron star be
$\la10^6$~G, far below the canonical value for low-mass X-ray binaries or
even old neutron stars in recycled ms pulsars. A more likely scenario
consistent with larger $i$ is that the neutron star spin is very slow.

% Even if the system is not
% face-on, the NS spin is probably slow -- possibly slow enough to be
% difficult to detect with RXTE.   The low accretion rate
% and short accretion history mean that the NS has probably not
% experienced enough torque to increase the spin significantly.
%The accretion rate onto the neutron star is so low that the NS could
%be spinning very slowly.  In equilibrium, the neutron star spins at
%the Kepler period at the radius where the ram pressure of the incoming
%material balances the magnetic pressure of the dipole field.  This
%equilibrium period goes like $\mdot^{-3/7}$ and $B^{6/7}$.  If we
%scale from the 2-min spin period of GX~1+4, we find that the spin
%period of the neutron star in \srcx\, could be approximately 20 times
%longer for the same strength magnetic field, and several hundred times
%longer for a magnetic field just an order of magnitude larger.

In conclusion, all scenarios for \srcx/\srco\, have some difficulties.
We have presented some options, but more observational constraints are
needed to select beween them.  Additional observations to determine
whether the 400-d radial velocity variations are due to orbital motion
or red-giant pulsation would also enable one to better interpret the
differences between \srcx/\srco\, and GX~1+4.
%We slightly favor the long orbit NS
%model [*do we?], which raises the question of the evolutionary path of
%this sytem.

% Radio observations (from Wendker 1995) provide upper limits of $f_{1418} <
% 0.8$~mJy and $f_{5000} < 1.1$~mJy, where $f_{5000}$ is the flux density at
% 5000~MHz.  I had hoped that radio fluxes might give us some indication of
% the mass-loss rate from the red giant, but since a radio flux measurement
% only gives a lower limit to the mass-loss rate, radio upper limits do not
% significantly constrain the mass-loss rate.  If, however, we wanted to
% blindly plug the 5~Ghz flux into the standard relationship for symbiotic
% stars (which I don't even think is applicable given the atypically hard
% spectrum from the compact object in 4U1700), we would get Mdot (from the
% red giant) = 9e-9 msun/yr (ref Seaquist \& Taylor 1990, ApJ 349, 313).

% P.S. 
% I forgot to mention that the (highly suspect) mass-loss estimate from
% the red giant scales as the relative wind velocity, and I assumed 10 km/s.

\acknowledgments

We are grateful to Dave Latham and Robert Stefanik for their assistance
with the data from the Wyeth telescope, and to Phil Uttley and Deepto
Chakrabarty for useful discussions.  This work was funded in part by NSF
grant INT-9902665 to J.L.S., and also the NASA Long Term Space
Astrophysics program under grant NAG 5-9184 (PI: Chakrabarty).  This
research has made use of data obtained through the High Energy
Astrophysics Science Archive Research Center Online Service, provided by
the NASA/Goddard Space Flight Center.

%% Appendix material should be preceded with a single \appendix command.
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% Use these commands to generate the bibliography using bibtex
% Then for submission, comment them out and read in the .bbl instead

% \bibliography{all}
% \bibliographystyle{apj}
\begin{thebibliography}{24}
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\bibitem[{{Press} {et~al.}(1996){Press}, {Teukolsky}, {Vetterling}, \&
  {Flannery}}]{nr}
{Press}, W.~H., {Teukolsky}, S.~A., {Vetterling}, W.~T., \& {Flannery}, B.~P.
  1996, Numerical Recipes in \protect{Fortran 77}: The Art of Scientific
  Computing, 2nd edn. (Cambridge, New York, Melbourne: Cambridge University
  Press), 1447

\bibitem[{{Sokoloski} {et~al.}(2001){Sokoloski}, {Bildsten}, \& {Ho}}]{sbh01}
{Sokoloski}, J.~L., {Bildsten}, L., \& {Ho}, W.~C.~G. 2001, \mnras, 326, 553

\bibitem[{{Stella} {et~al.}(1986){Stella}, {White}, \& {Rosner}}]{swr86}
{Stella}, L., {White}, N.~E., \& {Rosner}, R. 1986, \apj, 308, 669

\bibitem[{{Stellingwerf}(1978)}]{stell78}
{Stellingwerf}, R.~F. 1978, \apj, 224, 953

\bibitem[{{Tomasella} {et~al.}(1997){Tomasella}, {Munari}, {Tomov}, {Kolev},
  {Mikolajewski}, \& {Rejkuba}}]{tomasella97}
{Tomasella}, L., {Munari}, U., {Tomov}, T., {Kolev}, D., {Mikolajewski}, M., \&
  {Rejkuba}, M. 1997, IBVS, 4537, 1

\bibitem[{{Tonry} \& {Davis}(1979)}]{td79}
{Tonry}, J. \& {Davis}, M. 1979, \aj, 84, 1511

\end{thebibliography}

\clearpage

%% Tables should be submitted one per page, so put a \clearpage before
%% each one.

%% Two options are available to the author for producing tables:  the
%% deluxetable environment provided by the AASTeX package or the LaTeX
%% table environment.  Use of deluxetable is preferred.
%%

%% Three table samples follow, two marked up in the deluxetable environment,
%% one marked up as a LaTeX table.


\begin{deluxetable}{lccl}
% \renewcommand{\arraystretch}{1.5}
% \begin{center}
\tablewidth{0pt}
   \tablecaption{Comparison of the properties of the two known symbiotic
neutron-star binaries.  \label{comptable} }
%   \begin{tabular}{lccl}
  \tablehead{ \colhead{Parameter} & \colhead{4U~1700+24} & \colhead{GX 1+4} &
    \colhead{References} }
%\hline
\startdata
  Distance (kpc) & $0.42\pm0.2$ & 3--6/12--15 \tablenotemark{a} & [1,2] \\
  $P_{\rm spin}$ (s) & ?   & $\approx100$--137.7 & [3,4] \\
  $P_{\rm orb}$ (d) & $404\pm3$? & 304? & [5,6,7] \\
  $L_{\rm X}$ (2--10~keV, $10^{36}\ \eps$) & (0.2--$10)\times10^{-3}$
    & $\sim10$/100\tablenotemark{b} & [1,2] \\
  $\dot{M}$ ($10^{16}\ {\rm g\,s^{-1}}$)\tablenotemark{c}
    & (0.1--$5)\times10^{-3}$ & 5/50\tablenotemark{b}\\
  Companion spectra & M2 III & M3--6 III & [1,2] \\
  $L_{\rm opt}/L_{\rm X}$ & 200 & 0.25 & [1] \\
\enddata
%\hline
%  \end{tabular}
% \end{center}
\tablenotetext{a}{Depending upon the evolutionary status of the mass donor, first
ascent red giant/beginning inital ascent of the asymptotic giant branch.}

\tablenotetext{b}{Depending upon the distance}

\tablenotetext{c}{Assuming a neutron star with $R_*=10$~km and $M=1.4M_\sun$.}

\tablerefs{
 1. \cite{masetti01}; 2. \cite{chak97:opt}; 3. \cite{damle88,leahy89}; 4.
\cite{jab02}; 5. this paper; 6.
\cite{cut86}; 7. \cite{per99} } 

\end{deluxetable}

% Latest orbital fit parameters using eorbit
% Times of periastron:
%      45044.0      45448.3      45852.6      46256.8      46661.1 47065.4
%      47469.6      47873.9      48278.2      48682.4      49086.7 49491.0
%      49895.2      50299.5      50703.8      51108.1      51512.3 51916.6
%      52320.9      52725.1
% Running simulation for     1000 fake timeseries
% vel0 =  -48.69961 +/-   0.094295 (sys   0.091996) km/s
% K =   0.75 +/-  0.12km/s (sys  0.12)
% Porb =  404.26685 +/-   2.28534 (sys   2.95384) d
% a_x sin i = 4.17e+06 +/- 6.88e+05 (sys 6.80e+05) km
% Mass function = 1.77e-05 +/- 8.76e-06 (sys 8.66e-06) solar masses
% T_0 =    45044.04688 +/- 42.83641 (sys 82.91969) d
% eccentricity =    0.25569 +/-  0.14931 (sys  0.17676)
% longitude of periastron =  -95.86815 +/-  40.59164 (sys  +/-   1.27651) deg
% Chi-sq =  95.96 (76 dof)

\begin{table}
 \begin{center}
% \tablewidth{0pt}
\caption{Orbital parameters for \srco\ \label{orbtbl} }
 \begin{tabular}{lc}
 \tableline
Radial velocity amplitude, $K$ (km$\,{\rm s}^{-1}$) & $0.75 \pm 0.12$ \\
Projected semimajor axis, $a_{\rm X} \sin i$ ($10^6$~km) & $4.2\pm0.7$ \\
Orbital period, $P_{\rm orb}$ (d) & $404\pm3$ \\
Epoch of periastron, $T_0$ (MJD) & $49090 \pm 80$ \\
Eccentricity, $e$ & $0.26\pm0.15$ \\
Longitude of periastron, $\omega$ (degrees) & $260 \pm 40$ \\
Mass function, $f_{\rm O}$ ($10^{-5}\ M_\odot$) & $1.8\pm0.9$ \\
Model fit $\chi^2$ & 95.96 (76 dof) \\
\tableline
 \end{tabular}
\end{center}
\end{table}

%% The following command ends your manuscript. LaTeX will ignore any text
%% that appears after it.

\end{document}

%%
%% End of file `ms.tex'.

astro-ph0207594.extracted_bib
\begin{thebibliography} is followed by an empty set of
%% curly braces.  If you forget this, LaTeX will generate the error
%% "Perhaps a missing \item?".
%%
%% thebibliography produces citations in the text using \bibitem-\cite
%% cross-referencing. Each reference is preceded by a
%% \bibitem command that defines in curly braces the KEY that corresponds
%% to the KEY in the \cite commands (see the first section above).
%% Make sure that you provide a unique KEY for every \bibitem or else the
%% paper will not LaTeX. The square brackets should contain
%% the citation text that LaTeX will insert in
%% place of the \cite commands.

%% We have used macros to produce journal name abbreviations.
%% AASTeX provides a number of these for the more frequently-cited journals.
%% See the Author Guide for a list of them.

%% Note that the style of the \bibitem labels (in []) is slightly
%% different from previous examples.  The natbib system solves a host
%% of citation expression problems, but it is necessary to clearly
%% delimit the year from the author name used in the citation.
%% See the natbib documentation for more details and options.

% Use these commands to generate the bibliography using bibtex
% Then for submission, comment them out and read in the .bbl instead

% \bibliography{all}
% \bibliographystyle{apj}
\begin{thebibliography}{24}
\expandafter\ifx\csname natexlab\endcsname\relax\def\natexlab#1{#1}\fi

\bibitem[{{Chakrabarty} \& {Roche}(1997)}]{chak97:opt}
{Chakrabarty}, D. \& {Roche}, P. 1997, \apj, 489, 254

\bibitem[{{Cooke} {et~al.}(1978){Cooke}, {Ricketts}, {Maccacaro}, {Pye},
  {Elvis}, {Watson}, {Griffiths}, {Pounds}, {McHardy}, {Maccagni}, {Seward},
  {Page}, \& {Turner}}]{2acat}
{Cooke}, B.~A., {Ricketts}, M.~J., {Maccacaro}, T., {Pye}, J.~P., {Elvis}, M.,
  {Watson}, M.~G., {Griffiths}, R.~E., {Pounds}, K.~A., {McHardy}, I.,
  {Maccagni}, D., {Seward}, F.~D., {Page}, C.~G., \& {Turner}, M.~J.~L. 1978,
  \mnras, 182, 489

\bibitem[{{Cutler} {et~al.}(1986){Cutler}, {Dennis}, \& {Dolan}}]{cut86}
{Cutler}, E.~P., {Dennis}, B.~R., \& {Dolan}, J.~F. 1986, \apj, 300, 551

\bibitem[{{Damle} {et~al.}(1988){Damle}, {Kunte}, {Naranan}, {Sreekantan}, \&
  {Leahy}}]{damle88}
{Damle}, S.~V., {Kunte}, P.~K., {Naranan}, S., {Sreekantan}, B.~V., \& {Leahy},
  D.~A. 1988, Advances in Space Research, 8, 415

\bibitem[{{Forman} {et~al.}(1978){Forman}, {Jones}, {Cominsky}, {Julien},
  {Murray}, {Peters}, {Tananbaum}, \& {Giacconi}}]{4ucat}
{Forman}, W., {Jones}, C., {Cominsky}, L., {Julien}, P., {Murray}, S.,
  {Peters}, G., {Tananbaum}, H., \& {Giacconi}, R. 1978, \apjs, 38, 357

\bibitem[{{Garcia} {et~al.}(1983){Garcia}, {Baliunas}, {Elvis}, {Fabbiano},
  {Patterson}, {Schwartz}, {Doxsey}, {Koenigsberger}, {Swank}, \&
  {Watson}}]{garcia83}
{Garcia}, M., {Baliunas}, S.~L., {Elvis}, M., {Fabbiano}, G., {Patterson}, J.,
  {Schwartz}, D., {Doxsey}, R., {Koenigsberger}, G., {Swank}, J., \& {Watson},
  M.~G. 1983, \apj, 267, 291

\bibitem[{{Ghosh} \& {Lamb}(1979{\natexlab{a}})}]{gl79a}
{Ghosh}, P. \& {Lamb}, F.~K. 1979{\natexlab{a}}, \apj, 232, 259

\bibitem[{{Ghosh} \& {Lamb}(1979{\natexlab{b}})}]{gl79b}
---. 1979{\natexlab{b}}, \apj, 234, 296

\bibitem[{{Jablonski} {et~al.}(2002){Jablonski}, {Cieslinski}, {Braga}, \&
  {Pereira}}]{jab02}
{Jablonski}, F., {Cieslinski}, D., {Braga}, J., \& {Pereira}, M.~G. 2002,
  \iaucirc, 7895

\bibitem[{{Jablonski} {et~al.}(1997){Jablonski}, {Pereira}, {Braga}, \&
  {Gneiding}}]{jab97}
{Jablonski}, F.~J., {Pereira}, M.~G., {Braga}, J., \& {Gneiding}, C.~D. 1997,
  \apjl, 482, L171

\bibitem[{{Jahoda} {et~al.}(1996){Jahoda}, {Swank}, {Giles}, {Stark},
  {Strohmayer}, {Zhang}, \& {Morgan}}]{xte96}
{Jahoda}, K., {Swank}, J.~H., {Giles}, A.~B., {Stark}, M.~J., {Strohmayer}, T.,
  {Zhang}, W., \& {Morgan}, E.~H. 1996, \procspie, 2808, 59

\bibitem[{{Kenyon} {et~al.}(1988){Kenyon}, {Fernandez-Castro}, \&
  {Stencel}}]{kenyon88}
{Kenyon}, S.~J., {Fernandez-Castro}, T., \& {Stencel}, R.~E. 1988, \aj, 95,
  1817

\bibitem[{{Latham}(1992)}]{latham92}
{Latham}, D.~W. 1992, in ASP Conf. Ser. 32: IAU Colloq. 135: Complementary
  Approaches to Double and Multiple Star Research, 110

\bibitem[{{Leahy}(1989)}]{leahy89}
{Leahy}, D.~A. 1989, \baas, 21, 749

\bibitem[{{Masetti} {et~al.}(2002){Masetti}, {Dal Fiume}, {Cusumano}, {Amati},
  {Bartolini}, {Del Sordo}, {Frontera}, {Guarnieri}, {Orlandini}, {Palazzi},
  {Parmar}, {Piccioni}, \& {Santangelo}}]{masetti01}
{Masetti}, N., {Dal Fiume}, D., {Cusumano}, G., {Amati}, L., {Bartolini}, C.,
  {Del Sordo}, S., {Frontera}, F., {Guarnieri}, A., {Orlandini}, M., {Palazzi},
  E., {Parmar}, A.~N., {Piccioni}, A., \& {Santangelo}, A. 2002, \aap, 382, 104

\bibitem[{{Morgan} \& {Garcia}(2001)}]{mg01}
{Morgan}, W.~A. \& {Garcia}, M.~R. 2001, \pasp, 113, 1386

\bibitem[{{Percy} {et~al.}(2001){Percy}, {Wilson}, \& {Henry}}]{pwh01}
{Percy}, J.~R., {Wilson}, J.~B., \& {Henry}, G.~W. 2001, \pasp, 113, 983

\bibitem[{{Pereira} {et~al.}(1999){Pereira}, {Braga}, \& {Jablonski}}]{per99}
{Pereira}, M.~G., {Braga}, J., \& {Jablonski}, F. 1999, \apjl, 526, L105

\bibitem[{{Press} {et~al.}(1996){Press}, {Teukolsky}, {Vetterling}, \&
  {Flannery}}]{nr}
{Press}, W.~H., {Teukolsky}, S.~A., {Vetterling}, W.~T., \& {Flannery}, B.~P.
  1996, Numerical Recipes in \protect{Fortran 77}: The Art of Scientific
  Computing, 2nd edn. (Cambridge, New York, Melbourne: Cambridge University
  Press), 1447

\bibitem[{{Sokoloski} {et~al.}(2001){Sokoloski}, {Bildsten}, \& {Ho}}]{sbh01}
{Sokoloski}, J.~L., {Bildsten}, L., \& {Ho}, W.~C.~G. 2001, \mnras, 326, 553

\bibitem[{{Stella} {et~al.}(1986){Stella}, {White}, \& {Rosner}}]{swr86}
{Stella}, L., {White}, N.~E., \& {Rosner}, R. 1986, \apj, 308, 669

\bibitem[{{Stellingwerf}(1978)}]{stell78}
{Stellingwerf}, R.~F. 1978, \apj, 224, 953

\bibitem[{{Tomasella} {et~al.}(1997){Tomasella}, {Munari}, {Tomov}, {Kolev},
  {Mikolajewski}, \& {Rejkuba}}]{tomasella97}
{Tomasella}, L., {Munari}, U., {Tomov}, T., {Kolev}, D., {Mikolajewski}, M., \&
  {Rejkuba}, M. 1997, IBVS, 4537, 1

\bibitem[{{Tonry} \& {Davis}(1979)}]{td79}
{Tonry}, J. \& {Davis}, M. 1979, \aj, 84, 1511

\end{thebibliography}