Astronomy & Astrophysics manuscript no. tovmassian˙englisheditor c© ESO 2018
November 2, 2018
On the orbital period of the magnetic Cataclysmic Variable
HS 0922+1333
G. H. Tovmassian ⋆ and S.V. Zharikov
Observatorio Astrónomico Nacional SPM, Instituto de Astronomı́a, Universidad Nacional Autónoma de México, Ensenada, BC,
México⋆⋆;
e-mail: gag,zhar@astrosen.unam.mx
ABSTRACT
Context. The object HS 0922+1333 was visited briefly in 2002 in a mini survey of low accretion rate polars (LARPs) in order to test
if they undergo high luminosity states similar to ordinary polars. On the basis of that short observation the suspicion arose that the
object might be an asynchronous polar (Tovmassian et al. 2004). The disparity between the presumed orbital and spin period appeared
to be quite unusual.
Aims. We performed follow-up observations of the object to resolve the problem.
Methods. New simultaneous spectroscopic and photometric observations spanning several years allowed measurements of radial
velocities of emission and absorption lines from the secondary star and brightness variations due to synchrotron emission from the
primary.
Results. New observations show that the object is actually synchronous and its orbital and spin period are equal to 4.04 hours.
Conclusions. We identify the source of confusion of previous observations to be a high velocity component of emission line arousing
from the stream of matter leaving L1 point.
Key words. stars: - cataclysmic variables - magnetic, individual: - stars: HS 0922+1333
1. Introduction
Magnetic cataclysmic variables (CV) are accreting binary sys-
tems in which material transfers from a dwarf secondary star
onto a magnetic (∼5 < B < ∼250 MG) white dwarf (WD)
through Roche lobe overflow. Polars or AM Her systems with
magnetic fields larger than ∼ 10 MG stand out among magnetic
CVs because the spin period of the primary WD is synchronized
with the orbital period of the system. Unlike non-magnetic or
low-magnetic accreting binaries, they have neither a disk nor
the capacity to accumulate the transferred matter, so the bulk of
flux of these systems comes from the accretion flow, particularly
around magnetic poles. Therefore, their luminosity is sensitive
to the mass transfer rate Ṁ. Polars are known to have highs and
lows in their luminosity state, which is directly dependent on Ṁ.
In recent years a number of polars were identified with ex-
tremely low accretion rates. They are commonly called LARPs,
a name coined by Schwope et al. (2002). Their mass accretion
rate is estimated to be about a few 10−13 M⊙/yr, two orders
of magnitude below the average for CVs and they are distin-
guished for their prominent cyclotron emission lines on top of
otherwise featureless blue continua. The first two LARPs, in-
cluding the subject of this study, were discovered in the course
of the Hamburg QSO survey, thanks to a broad variable fea-
ture in the spectra subsequently identified with cyclotron lines
(Reimers et al. 1999; Reimers & Hagen 2000, hereafter RH).
Later, another newly identified magnetic CV from the list of
Send offprint requests to: G. Tovmassian
⋆ Visiting research fellow at Center for Astrophysics and Space
Sciences, University of California, San Diego, 9500 Gilman Drive, La
Jolla, CA 92093-0424, USA
⋆⋆ PO Box 439027, San Diego, CA, 92143-9024, USA
ROSAT sources (RX J1554.2+2721) was spotted in the low
state with a spectrum identical to LARPs (Tovmassian et al.
2001, 2004). Intrigued by that discovery, we conducted a blitz
campaign to check if canonical LARPs, namely HS 1023+3900
and HS 0922+1333, might be caught in a high state as well.
Since both objects had only recently been discovered and had
very limited observational coverage, we obtained one full binary
orbital period of spectral observations. Our instrumental setup
provided higher spectral resolution than the original discovery
observation by RH. We observed emission from the Hα line ap-
parently arising from the irradiated surface of the secondary star
facing the hot accreting spot on the WD and Na i infrared dou-
blet from the cooler parts of the secondary star (Tovmassian et al.
2004). The derived radial velocity (RV) curve from that obser-
vation did not fold well with the period estimated in the dis-
covery paper. However, the limited time coverage undermined
our ability to measure the period properly. We could only state
that the period might be exceeding what was reported by RH by
at least 1.14 times, corresponding to Pspin/Porb=0.88. It should
be noted that RH determined their period from the cyclotron
hump cycles and thus, they measured the WD spin period rather
than the binary orbital period. It would be quite usual to find
some degree of de-synchronization between the spin period of
the WD and orbital period. Nevertheless, the difference in peri-
ods was too large for an asynchronous polar and too small for an
intermediate polar. The latter mostly follow the empirical ratio
Pspin/Porb ∼ 0.1. In rare cases, Pspin/Porb ∼ 0.25 (see e.g. Norton
et al. 2004). There are also theoretical restrictions on a kind of
ratio that was indicated by our observation as evident from the
Norton et al. (2004) paper. Therefore, we conducted a new se-
ries of observations in order to establish the orbital period of the
compact binary and to classify it properly. This brief paper anal-
http://arxiv.org/abs/0704.0961v1
2 Tovmassian & Zharikov: Orbital period of the HS 0922+1333
yses a combined set of observations and discusses the reasons
that led us to an erroneous conclusion in 2004.
In Sect.2 we describe our observations and the data reduc-
tion. The data analysis and the results are presented in Sect.3,
and conclusions are drawn in Sect.4.
2. Observations and reduction
Sets of observations were collected over a four-year period and
analyzed. All observations of HS 0922+1333 reported here were
obtained at the Observatorio Astrónomico Nacional San Pedro
Martir, México. The B&Ch spectrograph installed at the 2.1 me-
ter telescope was used for the extensive spectroscopy, while a
1.5 m telescope was used to obtain simultaneous photometry
during the 2003 March run. In the first observations, upon which
Tovmassian et al. (2004) depended, we used a 600 l/mm grating
centered in the optical IR range (6200 – 8340 Å) to achieve a
spectral resolution of 4.2 Å FWHM in a sequence of 900 sec ex-
posures covering one orbital period. The controversy over the
periods led us to re-observe the object during three nights in
March 2003. This time we utilized the highest available grat-
ing of 1200 l/mm. The spectral resolution reached 2.2 Å FWHM
covering the 6100 – 7200 Å range. Later we collected more ob-
servations with lower resolution to refine the orbital period and
properly classify the secondary star.
In all observations an SITe 1024 × 1024 24 µm pixel CCD
was used to acquire the data. The slit width was usually set to 2.′′0
and oriented in the E–W direction. He-Ar arc lamp exposures
were taken at the beginning and end of each run for wavelength
calibration.
In 2003 March observations we conducted simultaneously
with differential photometry. Exposure times were 40–60 sec
with an overall time resolution of about 80–100 sec using the
Johnson-Cousins Rc filter.
The reduction of data was done in a fairly standard manner.
The bulk of reduction was performed using IRAF1 procedures,
except for removing of cosmic rays by a corresponding program
in MIDAS2, as this is an easier and more reliable tool. The bi-
ases were taken at the beginning and end of the night and were
subtracted after being combined using the CCD overscan area
for control of possible temperature-related variations during the
night. We did not do flat field correction for spectral observations
and used blank sky images taken at twilight for direct images.
The flux calibration was done by observing a spectrophotomet-
ric standard star. Feige 34 was observed during a 2002 run and
G191-B2B during the rest of the observations.
The wavelength calibration is routinely done by observing
a He–Ar arc lamp at the beginning and end of a sequence on
the object or every 2 hours if the sequence is too long. Then the
wavelength solutions calculated for each arc-lamp exposure and
an average of preceding and succeeding images are applied to
the object observed in between. The wavelength solutions are
usually good to a few 1/10 of an Angstrom, while deviations due
to the telescope position and flexations of the spectrograph can
exceed that by an order of magnitude. Usually, that does not pose
a problem since we work with moderate resolutions and the am-
plitude of radial velocity variation is on the order of hundreds
of km/sec. The sensible way of checking and correcting wave-
length calibration is to measure the night sky lines. We mea-
1 http://iraf.noao.edu
2 ESO-MIDAS is the acronym for the European Southern
Observatory Munich Image Data Analysis System which is developed
and maintained by the European Southern Observatory
Fig. 1. The CLEANed power spectrum of the RV variation is
presented by a solid line. The dashed line is the power spectrum
of photometric data. Vertical axes on the right side correspond
to the photometric power scale.
sured several lines by selecting unblended ones located close to
Hα and Na I. The measurements of sky lines show a clear trend
and indicate the scale of errors that one can incur depending on
the telescope inclination. Although the trend is unusually steep,
reaching 30 km/sec over 4 hours of observation, the scatter of
points around a linear fit is relatively small, which defines the
error of the measurements (rms) and is ≤8 km/sec. Nevertheless,
the error bars in the corresponding plots reflect the entire range
of deviation just to demonstrate the scale of corrections applied
to the data. The deviations of the linear fit to the measured night
sky lines (with an average of 2 night sky lines around each mea-
sured line) from the rest value were used to correct the wave-
length calibration by the corresponding amount.
3. Orbital period and system parameters
We measured the Hα line in the 2002 spectra with single
Gaussian fits. The resulting RV curve was reasonably smooth
and sinusoidal, but the ends of the curve would not overlap when
folded with the period reported by RH (see Fig. 5 in 2004). We
speculated that the actual orbital period is longer than the one de-
rived from the photometry. However, the measurements of new
spectra obtained in 2003 do not show such a large discrepancy,
and the period analysis of the combined dataset easily reveals
that the true period is indeed 4.0395 hours and coincides with the
photometric period derived from the synchrotron lines variabil-
ity within errors of measurement. The combination of data taken
years apart and several nights in a row each year allowed us to
determine the period very precisely. We applied the CLEAN pro-
cedure (Roberts et al. 1987) to sort out the alias periods resulting
from the uneven data sampling and daily gaps and obtained a
strong peak in the power spectrum at the 5.94131 ± 0.00065 cy-
cles/day, corresponding to a 4.0395 ± 0.0001 hour period (see
Fig. 1). Simultaneous with spectroscopy, we obtained photome-
Tovmassian & Zharikov: Orbital period of the HS 0922+1333 3
Table 1. Log of observations of HS 0922+1333
Date HJD+ Telescope Instrument/Grating Range/Band Exp.Time Duration
Spectroscopy 24530000 Num. of Integrations
2002 02 04 2309 2.1m B&Ch1 600l/mm 6200-8340Å 900s×19 4.5h
2003-03-25 2723 2.1m B&Ch 1200l/mm 6100-7200Å 900s×13 3.3h
2003-03-26 2724 2.1m B&Ch 1200l/mm 6100-7200Å 900s×6 1.3h
2003-03-27 2725 2.1m B&Ch 1200l/mm 6100-7200Å 900s×15 2.6h
2005-10-27 3670 2.1m B&Ch 400l/mm 6100-9200Å 900s×15 2.2h
2005-10-29 3672 2.1m B&Ch 400l/mm 6100-9200Å 900s×11 1.7h
2005-10-30 3673 2.1m B&Ch 400l/mm 6100-9200Å 900s×8 1.2h
2006-01-18 3673 2.1m B&Ch 400l/mm 5800-8900Å 1800s×8 3.7h
Photometry
2003-03-26 2724 1.5m RUCA2 R 120s×107 3.8h
2003-03-27 2725 1.5m RUCA R 120s×99 3.4h
2003-03-28 2726 1.5m RUCA R 120s×99 3.4h
1 B&Ch - Boller & Chivens spectrograph (http://haro.astrospp.unam.mx/Instruments/bchivens/bchivens.htm)
2 RUCA - CCD photometer (http://haro.astrospp.unam.mx/Instruments/laruca/laruca intro.htm)
Fig. 2. The light curve of HS 0922+1333 obtained in filter Rc
and folded with the orbital period. The phasing is according to
spectroscopic data.
try in Rc band that partially includes the strongest cyclotron line.
It is a dominant contributor to the light curve (Fig.2), so we can
use it to determine the spin period of the WD. The power spec-
trum calculated from photometry gives exactly the same result,
but the peak is broader, because the data lacks a longer time base.
In Fig.1 the power spectra of spectral and photometric data are
presented together.
It is clear that this system is a synchronous magnetic cata-
clysmic variable. The spin period of its white dwarf primary is
locked with the orbital and is not shorter, as suspected earlier.
We explored the cause of confusion. First of all, we corrected all
measured radial velocities using the night sky lines to remove
the trends. This decreased the gap a little between points in the
2002 data in phases 0.0 through 0.2 where they were not overlap-
ping. But even taking errors related to the wavelength calibration
into account, they still do not fold properly (see the open (blue)
square symbols in Fig.3). What is more interesting, however, is
that the amplitude of the radial velocity variation has a much
higher value in the 2002 data than in 2003 data, as measured
with single Gaussians. The careful examination of the 2003 data,
with twice the spectral resolution than in 2002, reveals that at the
bottom of the Hα emission line there is a weak and broad bump
present in most phases. We de-blended the Hα line from 2003
observations using two Gaussian components in the IRAF splot
procedure. The result is shown on the right side of the left panel
of Fig.3 only (positive phases). The strong, narrow component
basically coincides with the single Gaussian measurements. But
the weak broad component appears to show a much larger am-
plitude and reveals itself mainly between phases 0.3 through 0.9.
This component is clearly identified as the heated matter leaving
the Lagrangian L1 point, the nozzle where the accretion stream
forms. Outflowing matter has intrinsic velocity, so at phase 0.75
when the secondary star reaches maximum velocity toward the
observer, it tilts the weight of the emission line toward larger
velocity. Its phasing appears to be similar to a high-velocity
component (HVC) detected routinely in polars (Schwope et al.
1997, Tovmassian et al. 1999) that originates in the ballistic part
of the stream. In the lower-resolution spectra this component
could not be separated, therefore the radial velocity curve be-
came stretched and deformed. That and the short time coverage
limited to just a little over one orbital period led to the misinter-
pretation of the 2002 spectral data. It is very interesting that we
were able to distinguish the accretion flow onset. So far these ob-
jects have been known to show only the synchrotron humps as an
evidence of accretion processes taking place in them (Schwope
et al. 2002).
The RV curve derived from sodium lines (see Fig.3) gives the
measure of the rotation of the center of mass of the secondary in
the orbital plane, while the narrow component of the Hα line
originates from the front side of the elliptically distorted sec-
ondary. The ephemerides of HS 0922+1333 from the RV mea-
surements can be described as
T0 = HJD 2452308.336+ 0.
d168313[200]× E,
where T0 corresponds to the −/+ crossing of the RV curve as fol-
lows from the fitting of sinusoid to the RV measurements of Hα
and sodium lines separately according to the following equation:
V(t) = γ + K × sin(2π(t − t0)/Porb)
http://haro.astrospp.unam.mx/Instruments/bchivens/bchivens.htm
http://haro.astrospp.unam.mx/Instruments/laruca/laruca_intro.htm
4 Tovmassian & Zharikov: Orbital period of the HS 0922+1333
Fig. 3. The radial velocity curve of Hα line (left panel) and of Na i doublet (right). The open squares (blue) in the left panel represent
2002 data obtained with lower spectral resolution. The filled squares (red) are measurements of 2003 observations with single
Gaussian fitting. The error bars on the left side of the plots reflect the amplitude of wavelength corrections. The points are placed at
the correct positions after trend removal. The right side of the plot presents measurements of the 2003 data but with double Gaussian
de-blending of the line. The filled (green) circles are from the stronger line component originating at the irradiated secondary, the
open circles correspond to a much weaker component coming from the stream. In the right panel, measurements of the Na I lines
are presented from 2002 observations. The filled square symbols denote RV of λ 8197 Å measured with Gaussian deblending, after
velocity correction with sky lines. The open squares and triangles are measurements of the same doublet with single gaussians
(squares λ 8185Å and triangles λ 8197 Å). The diamonds are measurements of the λ 8185 Å line from 2006 observations. The curve
is a result of sin fit to the combined data. Note that scales of y-axes of panels are different.
Table 2. Radial velocity parameters of HS 0922+1333.
Line γ K Residuals
km/sec km/sec km/sec
Hα 36.6±7 132±12 25.1
Na I 8185Å -81±11 162±17 29.5
Na I 8197Å -65±13 139±20 32.7
Corresponding numbers derived from the fitting are pre-
sented in the Table 2. Unfortunately, due to the large errors there
is no marked difference between the semi-amplitude of radial
velocities between Hα and Na I lines. Otherwise, knowing the
spectral type of the secondary, we could deduce the basic pa-
rameters of the binary since that difference reflects the size of
the Roche lobe of the secondary.
The spectrum of the secondary in the absence of an accretion
disk is clearly seen, and in the phases when the magnetic ac-
creting spot that is radiating strong synchrotron emission is self-
eclipsed, one can see undisturbed secondary spectrum in the near
infrared range. In the Fig. 4 the flux calibrated spectrum of the
object obtained at phase 0.5 is presented. Overplotted are stan-
dard spectra of M3 to M5 main sequence stars (Pickles 1998)
normalized to the object. The WD’s contribution has not been
removed. However, at wavelengths above 6500 Å, its contribu-
tion is apparently insignificant and a good accordance emerges
between the object and M4 V standard star. This is also consis-
tent with what is expected from the Porb - spectral type II rela-
tion (Beuermann 2000), although the secondary is a M3.5 star
according to RH. The masses of secondaries in systems with pe-
Fig. 4. The spectrum of HS 0922+1333is presented by the solid
line. For comparison the standard spectra of M3-M5 stars are
plotted from the Pickles (1998)
Tovmassian & Zharikov: Orbital period of the HS 0922+1333 5
V (km/s)
V (km/s)
V (km/s)
V (km/s)
V (km/s)
Fig. 5. The Doppler maps of HS 0922+1333. On the top the to-
mograms of Hα emission line are presented in two panels with
different contrast levels to emphasize the concentration of the
emitting region on the facing side of the secondary on the left
and possibly some trace of mass transfer stream on the right.
The curved lines in the top panels correspond to the stream tra-
jectory, with numbers in the top right panel indicating stream
azimuth. The tomogram corresponding to the Na I line is placed
below in the right corner. The circle-shaped emission around the
center of mass is caused by the presence of the component of the
doublet line. In the bottom left corner, the observed and recon-
structed trailed spectra of Hα line (above) and Na I (below) are
presented.
riods similar to HS 0922+1333 range from 0.35 to 0.42 M⊙, in
those cases where the mass could be estimated precisely. Such
a secondary would follow the empirical mass-period and radius-
period relations from Smith & Dhillon (1998)
M2/M⊙ = 0.126(11) P(h) − 0.11(4)
R2/R⊙ = 0.117(4) P(h) − 0.041(18), (1)
Observations with higher resolution in the near IR will per-
mit investigators to precisely measure the difference between the
RV of Hα originating at the facing side of the secondary and
sodium absorption lines reflecting the motion of the center of
mass. Subsequently, it should allow for estimating the observed
radius of the star to check the possibility that it fills the Roche
lobe. For now, we can only assume that the mass transfer pro-
ceeds in a way similar to other polars, based on the detection of
a high velocity component in the emission line. Its presence can
also be illustrated by constructing Doppler tomograms.
Doppler tomography (Marsh & Horne 1988; Marsh 2001) is
a powerful tool in cases like this, where the origin of line profiles
is bound to the orbital plane and the system has relatively high
inclination. We constructed Doppler maps, or tomograms, using
both the Hα emission line and the Na I λ8197Å absorption line
to prove the accuracy of our estimate of the binary parameters.
The tomograms in Fig.5 show that the Hα line is mostly confined
to the front side of the secondary, while the sodium absorption
fills the entire body of the secondary. However, the difference is
not very obvious. The reason for that appears to be the the lower
spectral resolution and fewer spectra employed.
4. Conclusions
1. We have determined the 4.0395 hours spectroscopic period
of the LARP HS 0922+1333 based on the radial velocity
measurements of Hα emission line originating at the irra-
diated secondary star. The derived value coincides within
measurement errors with the spin period of the system, thus
proving that the object is a synchronized polar.
2. The profiles of the Hα emission line in higher-spectral res-
olution observations turned out to be complex. They are
formed basically on the irradiated surface of the secondary
star, but they also show a small contribution from the matter
in close proximity to the L1 point. The matter escaping the
secondary shows RVs with higher velocity and a different
phase.
3. The Doppler tomograms tend to confirm detection of a
stream of transfer matter.
The parameters of the system that we have obtained are in-
teresting in the context of the model proposed by Webbink and
Wickramasinghe (2005). According to it, the LARPs are rela-
tively young and are still approaching their first Roche lobe over-
flow. The accretion is due to the capture of the wind material
from the secondary by the strong magnetic field of the primary.
We think that we see evidence of a faint stream common to the
polars that transfer material through the L1 point which is usually
due to the Roche lobe overflow. However, the wind will proba-
bly also cause a flow of matter through the same trajectory, so
it is difficult to say if the observation runs against the model.
The precise measurement of the secondary star size may help to
clarify this.
Acknowledgements. This study was supported partially by grant 25454 from
CONACyT. GT acknowledges the UC-MEXUS fellowship program enabling
him to visit CASS UCSD. The authors are grateful to the anonymous referee for
careful reading of the manuscript and valuable comments. We thank L.Valencic
for help in language related issues.
References
Beuermann, K. 2000, New Astronomy Review, 44, 93
Demircan, O., & Kahraman, G. 1991, Ap&SS, 181, 313
Marsh, T. R., Horne, K. 1988, MNRAS, 235, 269
Marsh, T. R. 2001, ”Doppler Tomography”, Astrotomography, Indirect Imaging
Methods in Observational Astronomy, Edited by H.M.J. Boffin, D. Steeghs
and J. Cuypers, Lecture Notes in Physics, vol. 573, p.1
Norton, A. J., Wynn, G. A., & Somerscales, R. V. 2004, ApJ, 614, 349
Pickles, A. J. 1998, PASP, 110, 863
Roberts, D. H., Lehar, J., Dreher, J. W. 1987, AJ, 93, 968
Reimers, D., & Hagen, H.-J. 2000, A&A, 358, L45
Reimers, D., Hagen, H.-J., & Hopp, U. 1999, A&A, 343, 157
Schwope, A. D., Mantel, K.-H., & Horne, K. 1997, A&A, 319, 894
Schwope, A. D., Brunner, H., Hambaryan, V., & Schwarz, R. 2002, ASP
Conf. Ser. 261: The Physics of Cataclysmic Variables and Related Objects,
261, 102
Smith, D. A., Dhillon, V. S. 1998, MNRAS, 301, 767
Tovmassian, G. H., et al. 1999, ASP Conf. Ser. 157: Annapolis Workshop on
Magnetic Cataclysmic Variables, 157, 133
Tovmassian, G. H., Greiner, J., Zharikov, S. V., Echevarrı́a, J., & Kniazev, A.
2001, A&A, 380, 504
Tovmassian, G., Zharikov, S., Mennickent, R., & Greiner, J. 2004, ASP
Conf. Ser. 315: IAU Colloq. 190: Magnetic Cataclysmic Variables, 315, 15
6 Tovmassian & Zharikov: Orbital period of the HS 0922+1333
Warner, B. 1995, Cataclysmic variable stars, Cambridge Astrophysics Series,
Cambridge, New York: Cambridge University Press, 1995
Webbink, R. F., & Wickramasinghe, D. T. 2005, Astronomical Society of the
Pacific Conference Series, 330, 137
	Introduction
	Observations and reduction
	Orbital period and system parameters
	Conclusions

  Context: The object HS 0922+1333 was visited briefly in 2002 in a mini survey
of low accretion rate polars (LARPs) in order to test if they undergo high
luminosity states similar to ordinary polars. On the basis of that short
observation the suspicion arose that the object might be an asynchronous polar
(Tovmassian et al. 2004). The disparity between the presumed orbital and spin
period appeared to be quite unusual. Aims: We performed follow-up observations
of the object to resolve the problem. Methods: New simultaneous spectroscopic
and photometric observations spanning several years allowed measurements of
radial velocities of emission and absorption lines from the secondary star and
brightness variations due to synchrotron emission from the primary. Results:
New observations show that the object is actually synchronous and its orbital
and spin period are equal to 4.04 hours. Conclusions: We identify the source of
confusion of previous observations to be a high velocity component of emission
line arousing from the stream of matter leaving L1 point.

Introduction
Magnetic cataclysmic variables (CV) are accreting binary sys-
tems in which material transfers from a dwarf secondary star
onto a magnetic (∼5 < B < ∼250 MG) white dwarf (WD)
through Roche lobe overflow. Polars or AM Her systems with
magnetic fields larger than ∼ 10 MG stand out among magnetic
CVs because the spin period of the primary WD is synchronized
with the orbital period of the system. Unlike non-magnetic or
low-magnetic accreting binaries, they have neither a disk nor
the capacity to accumulate the transferred matter, so the bulk of
flux of these systems comes from the accretion flow, particularly
around magnetic poles. Therefore, their luminosity is sensitive
to the mass transfer rate Ṁ. Polars are known to have highs and
lows in their luminosity state, which is directly dependent on Ṁ.
In recent years a number of polars were identified with ex-
tremely low accretion rates. They are commonly called LARPs,
a name coined by Schwope et al. (2002). Their mass accretion
rate is estimated to be about a few 10−13 M⊙/yr, two orders
of magnitude below the average for CVs and they are distin-
guished for their prominent cyclotron emission lines on top of
otherwise featureless blue continua. The first two LARPs, in-
cluding the subject of this study, were discovered in the course
of the Hamburg QSO survey, thanks to a broad variable fea-
ture in the spectra subsequently identified with cyclotron lines
(Reimers et al. 1999; Reimers & Hagen 2000, hereafter RH).
Later, another newly identified magnetic CV from the list of
Send offprint requests to: G. Tovmassian
⋆ Visiting research fellow at Center for Astrophysics and Space
Sciences, University of California, San Diego, 9500 Gilman Drive, La
Jolla, CA 92093-0424, USA
⋆⋆ PO Box 439027, San Diego, CA, 92143-9024, USA
ROSAT sources (RX J1554.2+2721) was spotted in the low
state with a spectrum identical to LARPs (Tovmassian et al.
2001, 2004). Intrigued by that discovery, we conducted a blitz
campaign to check if canonical LARPs, namely HS 1023+3900
and HS 0922+1333, might be caught in a high state as well.
Since both objects had only recently been discovered and had
very limited observational coverage, we obtained one full binary
orbital period of spectral observations. Our instrumental setup
provided higher spectral resolution than the original discovery
observation by RH. We observed emission from the Hα line ap-
parently arising from the irradiated surface of the secondary star
facing the hot accreting spot on the WD and Na i infrared dou-
blet from the cooler parts of the secondary star (Tovmassian et al.
2004). The derived radial velocity (RV) curve from that obser-
vation did not fold well with the period estimated in the dis-
covery paper. However, the limited time coverage undermined
our ability to measure the period properly. We could only state
that the period might be exceeding what was reported by RH by
at least 1.14 times, corresponding to Pspin/Porb=0.88. It should
be noted that RH determined their period from the cyclotron
hump cycles and thus, they measured the WD spin period rather
than the binary orbital period. It would be quite usual to find
some degree of de-synchronization between the spin period of
the WD and orbital period. Nevertheless, the difference in peri-
ods was too large for an asynchronous polar and too small for an
intermediate polar. The latter mostly follow the empirical ratio
Pspin/Porb ∼ 0.1. In rare cases, Pspin/Porb ∼ 0.25 (see e.g. Norton
et al. 2004). There are also theoretical restrictions on a kind of
ratio that was indicated by our observation as evident from the
Norton et al. (2004) paper. Therefore, we conducted a new se-
ries of observations in order to establish the orbital period of the
compact binary and to classify it properly. This brief paper anal-
http://arxiv.org/abs/0704.0961v1
2 Tovmassian & Zharikov: Orbital period of the HS 0922+1333
yses a combined set of observations and discusses the reasons
that led us to an erroneous conclusion in 2004.
In Sect.2 we describe our observations and the data reduc-
tion. The data analysis and the results are presented in Sect.3,
and conclusions are drawn in Sect.4.
2. Observations and reduction
Sets of observations were collected over a four-year period and
analyzed. All observations of HS 0922+1333 reported here were
obtained at the Observatorio Astrónomico Nacional San Pedro
Martir, México. The B&Ch spectrograph installed at the 2.1 me-
ter telescope was used for the extensive spectroscopy, while a
1.5 m telescope was used to obtain simultaneous photometry
during the 2003 March run. In the first observations, upon which
Tovmassian et al. (2004) depended, we used a 600 l/mm grating
centered in the optical IR range (6200 – 8340 Å) to achieve a
spectral resolution of 4.2 Å FWHM in a sequence of 900 sec ex-
posures covering one orbital period. The controversy over the
periods led us to re-observe the object during three nights in
March 2003. This time we utilized the highest available grat-
ing of 1200 l/mm. The spectral resolution reached 2.2 Å FWHM
covering the 6100 – 7200 Å range. Later we collected more ob-
servations with lower resolution to refine the orbital period and
properly classify the secondary star.
In all observations an SITe 1024 × 1024 24 µm pixel CCD
was used to acquire the data. The slit width was usually set to 2.′′0
and oriented in the E–W direction. He-Ar arc lamp exposures
were taken at the beginning and end of each run for wavelength
calibration.
In 2003 March observations we conducted simultaneously
with differential photometry. Exposure times were 40–60 sec
with an overall time resolution of about 80–100 sec using the
Johnson-Cousins Rc filter.
The reduction of data was done in a fairly standard manner.
The bulk of reduction was performed using IRAF1 procedures,
except for removing of cosmic rays by a corresponding program
in MIDAS2, as this is an easier and more reliable tool. The bi-
ases were taken at the beginning and end of the night and were
subtracted after being combined using the CCD overscan area
for control of possible temperature-related variations during the
night. We did not do flat field correction for spectral observations
and used blank sky images taken at twilight for direct images.
The flux calibration was done by observing a spectrophotomet-
ric standard star. Feige 34 was observed during a 2002 run and
G191-B2B during the rest of the observations.
The wavelength calibration is routinely done by observing
a He–Ar arc lamp at the beginning and end of a sequence on
the object or every 2 hours if the sequence is too long. Then the
wavelength solutions calculated for each arc-lamp exposure and
an average of preceding and succeeding images are applied to
the object observed in between. The wavelength solutions are
usually good to a few 1/10 of an Angstrom, while deviations due
to the telescope position and flexations of the spectrograph can
exceed that by an order of magnitude. Usually, that does not pose
a problem since we work with moderate resolutions and the am-
plitude of radial velocity variation is on the order of hundreds
of km/sec. The sensible way of checking and correcting wave-
length calibration is to measure the night sky lines. We mea-
1 http://iraf.noao.edu
2 ESO-MIDAS is the acronym for the European Southern
Observatory Munich Image Data Analysis System which is developed
and maintained by the European Southern Observatory
Fig. 1. The CLEANed power spectrum of the RV variation is
presented by a solid line. The dashed line is the power spectrum
of photometric data. Vertical axes on the right side correspond
to the photometric power scale.
sured several lines by selecting unblended ones located close to
Hα and Na I. The measurements of sky lines show a clear trend
and indicate the scale of errors that one can incur depending on
the telescope inclination. Although the trend is unusually steep,
reaching 30 km/sec over 4 hours of observation, the scatter of
points around a linear fit is relatively small, which defines the
error of the measurements (rms) and is ≤8 km/sec. Nevertheless,
the error bars in the corresponding plots reflect the entire range
of deviation just to demonstrate the scale of corrections applied
to the data. The deviations of the linear fit to the measured night
sky lines (with an average of 2 night sky lines around each mea-
sured line) from the rest value were used to correct the wave-
length calibration by the corresponding amount.
3. Orbital period and system parameters
We measured the Hα line in the 2002 spectra with single
Gaussian fits. The resulting RV curve was reasonably smooth
and sinusoidal, but the ends of the curve would not overlap when
folded with the period reported by RH (see Fig. 5 in 2004). We
speculated that the actual orbital period is longer than the one de-
rived from the photometry. However, the measurements of new
spectra obtained in 2003 do not show such a large discrepancy,
and the period analysis of the combined dataset easily reveals
that the true period is indeed 4.0395 hours and coincides with the
photometric period derived from the synchrotron lines variabil-
ity within errors of measurement. The combination of data taken
years apart and several nights in a row each year allowed us to
determine the period very precisely. We applied the CLEAN pro-
cedure (Roberts et al. 1987) to sort out the alias periods resulting
from the uneven data sampling and daily gaps and obtained a
strong peak in the power spectrum at the 5.94131 ± 0.00065 cy-
cles/day, corresponding to a 4.0395 ± 0.0001 hour period (see
Fig. 1). Simultaneous with spectroscopy, we obtained photome-
Tovmassian & Zharikov: Orbital period of the HS 0922+1333 3
Table 1. Log of observations of HS 0922+1333
Date HJD+ Telescope Instrument/Grating Range/Band Exp.Time Duration
Spectroscopy 24530000 Num. of Integrations
2002 02 04 2309 2.1m B&Ch1 600l/mm 6200-8340Å 900s×19 4.5h
2003-03-25 2723 2.1m B&Ch 1200l/mm 6100-7200Å 900s×13 3.3h
2003-03-26 2724 2.1m B&Ch 1200l/mm 6100-7200Å 900s×6 1.3h
2003-03-27 2725 2.1m B&Ch 1200l/mm 6100-7200Å 900s×15 2.6h
2005-10-27 3670 2.1m B&Ch 400l/mm 6100-9200Å 900s×15 2.2h
2005-10-29 3672 2.1m B&Ch 400l/mm 6100-9200Å 900s×11 1.7h
2005-10-30 3673 2.1m B&Ch 400l/mm 6100-9200Å 900s×8 1.2h
2006-01-18 3673 2.1m B&Ch 400l/mm 5800-8900Å 1800s×8 3.7h
Photometry
2003-03-26 2724 1.5m RUCA2 R 120s×107 3.8h
2003-03-27 2725 1.5m RUCA R 120s×99 3.4h
2003-03-28 2726 1.5m RUCA R 120s×99 3.4h
1 B&Ch - Boller & Chivens spectrograph (http://haro.astrospp.unam.mx/Instruments/bchivens/bchivens.htm)
2 RUCA - CCD photometer (http://haro.astrospp.unam.mx/Instruments/laruca/laruca intro.htm)
Fig. 2. The light curve of HS 0922+1333 obtained in filter Rc
and folded with the orbital period. The phasing is according to
spectroscopic data.
try in Rc band that partially includes the strongest cyclotron line.
It is a dominant contributor to the light curve (Fig.2), so we can
use it to determine the spin period of the WD. The power spec-
trum calculated from photometry gives exactly the same result,
but the peak is broader, because the data lacks a longer time base.
In Fig.1 the power spectra of spectral and photometric data are
presented together.
It is clear that this system is a synchronous magnetic cata-
clysmic variable. The spin period of its white dwarf primary is
locked with the orbital and is not shorter, as suspected earlier.
We explored the cause of confusion. First of all, we corrected all
measured radial velocities using the night sky lines to remove
the trends. This decreased the gap a little between points in the
2002 data in phases 0.0 through 0.2 where they were not overlap-
ping. But even taking errors related to the wavelength calibration
into account, they still do not fold properly (see the open (blue)
square symbols in Fig.3). What is more interesting, however, is
that the amplitude of the radial velocity variation has a much
higher value in the 2002 data than in 2003 data, as measured
with single Gaussians. The careful examination of the 2003 data,
with twice the spectral resolution than in 2002, reveals that at the
bottom of the Hα emission line there is a weak and broad bump
present in most phases. We de-blended the Hα line from 2003
observations using two Gaussian components in the IRAF splot
procedure. The result is shown on the right side of the left panel
of Fig.3 only (positive phases). The strong, narrow component
basically coincides with the single Gaussian measurements. But
the weak broad component appears to show a much larger am-
plitude and reveals itself mainly between phases 0.3 through 0.9.
This component is clearly identified as the heated matter leaving
the Lagrangian L1 point, the nozzle where the accretion stream
forms. Outflowing matter has intrinsic velocity, so at phase 0.75
when the secondary star reaches maximum velocity toward the
observer, it tilts the weight of the emission line toward larger
velocity. Its phasing appears to be similar to a high-velocity
component (HVC) detected routinely in polars (Schwope et al.
1997, Tovmassian et al. 1999) that originates in the ballistic part
of the stream. In the lower-resolution spectra this component
could not be separated, therefore the radial velocity curve be-
came stretched and deformed. That and the short time coverage
limited to just a little over one orbital period led to the misinter-
pretation of the 2002 spectral data. It is very interesting that we
were able to distinguish the accretion flow onset. So far these ob-
jects have been known to show only the synchrotron humps as an
evidence of accretion processes taking place in them (Schwope
et al. 2002).
The RV curve derived from sodium lines (see Fig.3) gives the
measure of the rotation of the center of mass of the secondary in
the orbital plane, while the narrow component of the Hα line
originates from the front side of the elliptically distorted sec-
ondary. The ephemerides of HS 0922+1333 from the RV mea-
surements can be described as
T0 = HJD 2452308.336+ 0.
d168313[200]× E,
where T0 corresponds to the −/+ crossing of the RV curve as fol-
lows from the fitting of sinusoid to the RV measurements of Hα
and sodium lines separately according to the following equation:
V(t) = γ + K × sin(2π(t − t0)/Porb)
http://haro.astrospp.unam.mx/Instruments/bchivens/bchivens.htm
http://haro.astrospp.unam.mx/Instruments/laruca/laruca_intro.htm
4 Tovmassian & Zharikov: Orbital period of the HS 0922+1333
Fig. 3. The radial velocity curve of Hα line (left panel) and of Na i doublet (right). The open squares (blue) in the left panel represent
2002 data obtained with lower spectral resolution. The filled squares (red) are measurements of 2003 observations with single
Gaussian fitting. The error bars on the left side of the plots reflect the amplitude of wavelength corrections. The points are placed at
the correct positions after trend removal. The right side of the plot presents measurements of the 2003 data but with double Gaussian
de-blending of the line. The filled (green) circles are from the stronger line component originating at the irradiated secondary, the
open circles correspond to a much weaker component coming from the stream. In the right panel, measurements of the Na I lines
are presented from 2002 observations. The filled square symbols denote RV of λ 8197 Å measured with Gaussian deblending, after
velocity correction with sky lines. The open squares and triangles are measurements of the same doublet with single gaussians
(squares λ 8185Å and triangles λ 8197 Å). The diamonds are measurements of the λ 8185 Å line from 2006 observations. The curve
is a result of sin fit to the combined data. Note that scales of y-axes of panels are different.
Table 2. Radial velocity parameters of HS 0922+1333.
Line γ K Residuals
km/sec km/sec km/sec
Hα 36.6±7 132±12 25.1
Na I 8185Å -81±11 162±17 29.5
Na I 8197Å -65±13 139±20 32.7
Corresponding numbers derived from the fitting are pre-
sented in the Table 2. Unfortunately, due to the large errors there
is no marked difference between the semi-amplitude of radial
velocities between Hα and Na I lines. Otherwise, knowing the
spectral type of the secondary, we could deduce the basic pa-
rameters of the binary since that difference reflects the size of
the Roche lobe of the secondary.
The spectrum of the secondary in the absence of an accretion
disk is clearly seen, and in the phases when the magnetic ac-
creting spot that is radiating strong synchrotron emission is self-
eclipsed, one can see undisturbed secondary spectrum in the near
infrared range. In the Fig. 4 the flux calibrated spectrum of the
object obtained at phase 0.5 is presented. Overplotted are stan-
dard spectra of M3 to M5 main sequence stars (Pickles 1998)
normalized to the object. The WD’s contribution has not been
removed. However, at wavelengths above 6500 Å, its contribu-
tion is apparently insignificant and a good accordance emerges
between the object and M4 V standard star. This is also consis-
tent with what is expected from the Porb - spectral type II rela-
tion (Beuermann 2000), although the secondary is a M3.5 star
according to RH. The masses of secondaries in systems with pe-
Fig. 4. The spectrum of HS 0922+1333is presented by the solid
line. For comparison the standard spectra of M3-M5 stars are
plotted from the Pickles (1998)
Tovmassian & Zharikov: Orbital period of the HS 0922+1333 5
V (km/s)
V (km/s)
V (km/s)
V (km/s)
V (km/s)
Fig. 5. The Doppler maps of HS 0922+1333. On the top the to-
mograms of Hα emission line are presented in two panels with
different contrast levels to emphasize the concentration of the
emitting region on the facing side of the secondary on the left
and possibly some trace of mass transfer stream on the right.
The curved lines in the top panels correspond to the stream tra-
jectory, with numbers in the top right panel indicating stream
azimuth. The tomogram corresponding to the Na I line is placed
below in the right corner. The circle-shaped emission around the
center of mass is caused by the presence of the component of the
doublet line. In the bottom left corner, the observed and recon-
structed trailed spectra of Hα line (above) and Na I (below) are
presented.
riods similar to HS 0922+1333 range from 0.35 to 0.42 M⊙, in
those cases where the mass could be estimated precisely. Such
a secondary would follow the empirical mass-period and radius-
period relations from Smith & Dhillon (1998)
M2/M⊙ = 0.126(11) P(h) − 0.11(4)
R2/R⊙ = 0.117(4) P(h) − 0.041(18), (1)
Observations with higher resolution in the near IR will per-
mit investigators to precisely measure the difference between the
RV of Hα originating at the facing side of the secondary and
sodium absorption lines reflecting the motion of the center of
mass. Subsequently, it should allow for estimating the observed
radius of the star to check the possibility that it fills the Roche
lobe. For now, we can only assume that the mass transfer pro-
ceeds in a way similar to other polars, based on the detection of
a high velocity component in the emission line. Its presence can
also be illustrated by constructing Doppler tomograms.
Doppler tomography (Marsh & Horne 1988; Marsh 2001) is
a powerful tool in cases like this, where the origin of line profiles
is bound to the orbital plane and the system has relatively high
inclination. We constructed Doppler maps, or tomograms, using
both the Hα emission line and the Na I λ8197Å absorption line
to prove the accuracy of our estimate of the binary parameters.
The tomograms in Fig.5 show that the Hα line is mostly confined
to the front side of the secondary, while the sodium absorption
fills the entire body of the secondary. However, the difference is
not very obvious. The reason for that appears to be the the lower
spectral resolution and fewer spectra employed.
4. Conclusions
1. We have determined the 4.0395 hours spectroscopic period
of the LARP HS 0922+1333 based on the radial velocity
measurements of Hα emission line originating at the irra-
diated secondary star. The derived value coincides within
measurement errors with the spin period of the system, thus
proving that the object is a synchronized polar.
2. The profiles of the Hα emission line in higher-spectral res-
olution observations turned out to be complex. They are
formed basically on the irradiated surface of the secondary
star, but they also show a small contribution from the matter
in close proximity to the L1 point. The matter escaping the
secondary shows RVs with higher velocity and a different
phase.
3. The Doppler tomograms tend to confirm detection of a
stream of transfer matter.
The parameters of the system that we have obtained are in-
teresting in the context of the model proposed by Webbink and
Wickramasinghe (2005). According to it, the LARPs are rela-
tively young and are still approaching their first Roche lobe over-
flow. The accretion is due to the capture of the wind material
from the secondary by the strong magnetic field of the primary.
We think that we see evidence of a faint stream common to the
polars that transfer material through the L1 point which is usually
due to the Roche lobe overflow. However, the wind will proba-
bly also cause a flow of matter through the same trajectory, so
it is difficult to say if the observation runs against the model.
The precise measurement of the secondary star size may help to
clarify this.
Acknowledgements. This study was supported partially by grant 25454 from
CONACyT. GT acknowledges the UC-MEXUS fellowship program enabling
him to visit CASS UCSD. The authors are grateful to the anonymous referee for
careful reading of the manuscript and valuable comments. We thank L.Valencic
for help in language related issues.
References
Beuermann, K. 2000, New Astronomy Review, 44, 93
Demircan, O., & Kahraman, G. 1991, Ap&SS, 181, 313
Marsh, T. R., Horne, K. 1988, MNRAS, 235, 269
Marsh, T. R. 2001, ”Doppler Tomography”, Astrotomography, Indirect Imaging
Methods in Observational Astronomy, Edited by H.M.J. Boffin, D. Steeghs
and J. Cuypers, Lecture Notes in Physics, vol. 573, p.1
Norton, A. J., Wynn, G. A., & Somerscales, R. V. 2004, ApJ, 614, 349
Pickles, A. J. 1998, PASP, 110, 863
Roberts, D. H., Lehar, J., Dreher, J. W. 1987, AJ, 93, 968
Reimers, D., & Hagen, H.-J. 2000, A&A, 358, L45
Reimers, D., Hagen, H.-J., & Hopp, U. 1999, A&A, 343, 157
Schwope, A. D., Mantel, K.-H., & Horne, K. 1997, A&A, 319, 894
Schwope, A. D., Brunner, H., Hambaryan, V., & Schwarz, R. 2002, ASP
Conf. Ser. 261: The Physics of Cataclysmic Variables and Related Objects,
261, 102
Smith, D. A., Dhillon, V. S. 1998, MNRAS, 301, 767
Tovmassian, G. H., et al. 1999, ASP Conf. Ser. 157: Annapolis Workshop on
Magnetic Cataclysmic Variables, 157, 133
Tovmassian, G. H., Greiner, J., Zharikov, S. V., Echevarrı́a, J., & Kniazev, A.
2001, A&A, 380, 504
Tovmassian, G., Zharikov, S., Mennickent, R., & Greiner, J. 2004, ASP
Conf. Ser. 315: IAU Colloq. 190: Magnetic Cataclysmic Variables, 315, 15
6 Tovmassian & Zharikov: Orbital period of the HS 0922+1333
Warner, B. 1995, Cataclysmic variable stars, Cambridge Astrophysics Series,
Cambridge, New York: Cambridge University Press, 1995
Webbink, R. F., & Wickramasinghe, D. T. 2005, Astronomical Society of the
Pacific Conference Series, 330, 137
	Introduction
	Observations and reduction
	Orbital period and system parameters
	Conclusions
