Spectroscopy of Nine Cataclysmic Variable Stars 1
Holly A. Sheets, John R. Thorstensen, Christopher J. Peters, and Ann B. Kapusta
Department of Physics and Astronomy
6127 Wilder Laboratory, Dartmouth College
Hanover, NH 03755-3528;
h.a.sheets@dartmouth.edu
Cynthia J. Taylor
The Lawrenceville School
P.O. Box 6008, Lawrenceville, NJ 08648
ABSTRACT
We present optical spectroscopy of nine cataclysmic binary stars, mostly dwarf no-
vae, obtained primarily to determine orbital periods Porb. The stars and their periods
are LX And, 0.1509743(5) d; CZ Aql, 0.2005(6) d; LU Cam, 0.1499686(4) d; GZ Cnc,
0.0881(4) d; V632 Cyg, 0.06377(8) d; V1006 Cyg, 0.09903(9) d; BF Eri, 0.2708804(4)
d; BI Ori, 0.1915(5) d; and FO Per, for which Porb is either 0.1467(4) or 0.1719(5) d.
Several of the stars proved to be especially interesting. In BF Eri, we detect the
absorption spectrum of a secondary star of spectral type K3 ±1 subclass, which leads
to a distance estimate of ∼ 1 kpc. However, BF Eri has a large proper motion (∼ 100
mas yr−1), and we have a preliminary parallax measurement that confirms the large
proper motion and yields only an upper limit for the parallax. BF Eri’s space velocity
is evidently large, and it appears to belong to the halo population. In CZ Aql, the
emission lines have strong wings that move with large velocity amplitude, suggesting
a magnetically-channeled accretion flow. The orbital period of V1006 Cyg places it
squarely within the 2- to 3-hour ‘gap’ in the distribution of cataclysmic binary orbital
periods.
Subject headings: novae, cataclysmic variables — stars: individual (LX And, CZ Aql,
LU Cam, GZ Cnc, V632 Cyg, V1006 Cyg, BF Eri, BI Ori, FO Per) — stars: distances
— binaries: close — binaries: spectroscopic
1Based on observations obtained at the MDM Observatory, operated by Dartmouth College, Columbia University,
Ohio State University, Ohio University, and the University of Michigan.
http://arxiv.org/abs/0704.0948v1
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1. Introduction
Cataclysmic variables (CVs) are binary star systems in which the secondary, usually a late-
type main sequence star, fills its Roche lobe and loses mass to the white dwarf primary (Warner
1995). CVs are long-lived systems that are stable against mass transfer, so the mass transfer must
be driven by gradual changes in the orbit, or in the secondary star, or both. It is commonly
believed that the evolution of most CVs is driven by the slow loss of angular momentum from the
orbit, most likely through magnetic braking of the co-rotating secondary star, at least at longer
orbital periods Porb where gravitational radiation is ineffective (Andronov & Pinsonneault 2004
give a recent discussion). The loss of angular momentum constricts the Roche critical lobe around
the secondary and causes the system to transfer mass as it evolves toward shorter Porb. In this
scenario, Porb serves as a proxy measurement for the system’s evolutionary state. Correct and
complete orbital period measurements are fundamental to any accurate theory of CV evolution.
Given the usefulness of Porb, it is fortunate that it can usually be measured accurately and precisely.
This paper presents optical spectroscopy of the nine CVs listed in Table 1. We took these
observations mostly for the purpose of finding orbital periods using radial velocities (none of these
systems are known to eclipse). The long cumulative exposures also allowed us to look for any
unusual features. The Catalog and Atlas of Cataclysmic Variables Archival Edition (Downes et al.
2001) 2 lists seven of the stars as dwarf novae, one as either a dwarf nova or a DQ Her star, and one
simply as a cataclysmic, possibly a dwarf nova similar to U Gem or SS Cygni (type UGSS). Except
for CZ Aql, for which we confirm a 4.8-hour candidate period suggested by Cieslinski et al. (1998),
all of these objects lacked published orbital periods when we began working on them. Subsequently
Tappert & Bianchini (2003) found Porb = 0.0883 d for GZ Cancri; we had communicated our
advance findings to these authors so they could disambiguate their period determination.
2. Observations, Reductions, and Analysis
2.1. Observations
All our spectra were taken at the MDM Observatory on Kitt Peak, Arizona, using either the
1.3m McGraw-Hill telescope or the 2.4m Hiltner telescope. The earliest observations we report
here are from 1995, and the latest were obtained 2007 January. Table 2 gives a journal of the
observations.
At the 1.3m we used the Mark III spectrograph and a SITe 1024 × 1024 CCD detector. The
spectral resolution is 5.0 Å, covering a range of either 4480 to 6760 Å with 2.2 Å pixel−1 for the
2001 December BF Eri data, or 4646 to 6970 Å with 2.3 Å pixel−1 for the remaining data. The
2Available at http://archive.stsci.edu/prepds/cvcat/index.html; this had been called the Living Edition until its
author retired and ceased updates.
http://archive.stsci.edu/prepds/cvcat/index.html
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2.4m spectra, except for those of FO Per, were obtained with the modular spectrograph and a
SITe 20482 CCD detector, with 2.0 Å pixel−1, over a range of 4210 to 7500 Å and with a spectral
resolution of 3.5 Å. The relatively small number of 2.4 m spectra of FO Per were taken with a
LORAL 20482-pixel detector, and cover from 4285 to 6870 Å at 1.25 Å pixel−1.
2.2. Reductions
For the most part we reduced the spectra using standard IRAF3 procedures. The wavelength
calibration was based on exposures of Hg, Ne, and Xe lamps. Prior to 2003 we took lamp exposures
through the night and whenever the telescope was moved. For the 2.4m data from 2003 to the
present, we used lamp exposures taken in twilight to find the shape of the pixel-to-wavelength
relation, and set the zero point individually for each nighttime exposure using the OI λ5577 night-
sky feature. The apparent velocity of the telluric OH emission bands at the far red end of the
spectrum, found with a cross-correlation routine, provided a check; although these are far from
the feature used to set the zero point, their apparent velocity typically remain within 10 km s−1
of zero. Because of the increased efficiency of this technique, we attempted to use it at the 1.3m
telescope also, during the 2004 June/July observing run. For unknown reasons the results were
unsatisfactory. To salvage the Hα emission velocities from that run, we determined a correction by
cross-correlating the night-sky emission features in the 6200-6625 Å range with a well-calibrated
night-sky spectrum obtained with a similar instrument. The correction was calculated for each
individual spectrum and then applied to each measured velocity, and it did reduce the scatter
somewhat, evidently because the wavelength range used includes the Hα emission line for which
we measured velocities.
On all our runs, we observed flux standards during twilight when the sky was clear, and applied
the resulting calibration to the data. The reproducibility of these observations suggests that our
fluxes are typically accurate to ±20 per cent. We also took short exposures of bright O and B
stars in twilight to map the telluric absorption features and divide them out approximately from
our program object spectra. Before flux calibration, we divided our program star spectra by a
mean hot-star continuum, in order to remove the bulk of the response variation. Table 1 lists
V magnitudes synthesized from our mean spectra, using the IRAF sbands task and the passband
tabulated by Bessell (1990); clouds, losses at the slit, and calibration errors make these uncertain
by a few tenths of a magnitude, but they do give a rough indication of the brightness of each system
at the time of our observation.
3IRAF is distributed by the National Optical Astronomy Observatories.
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2.3. Analysis
Except for a few spectra taken in outburst (which show weak emission or absorption on a
strong continuum), all of the stars show the prominent emission lines. Figs. 1 and 2 show averaged
spectra, and Table 3 gives the equivalent width and FWHM of each line measured for each star
from its averaged spectrum.
Two stars, BF Eri and BI Ori, showed the spectral features of a late-type star. To quantify the
secondary contribution in these objects, we began by preparing averaged flux-calibrated spectra
(in BF Eri’s case the secondary’s radial velocity curve was measurable, so we shifted the individual
spectra to the secondary’s rest frame before averaging). Over time we have used the 2.4 m and
modular spectrograph to collect spectra of K and M stars classified by Keenan & McNeil (1989) or
Boeshaar (1976). The wavelength coverage and spectral resolution of these data are similar to the
1.3m data. We applied a range of scaling factors to the library spectra, subtracted them from the
averaged spectra, and examined the results by eye to estimate a range of spectral types and scaling
factors giving acceptable cancellation of the late-type features.
We use the spectral type and secondary flux to estimate the distance in the following manner.
We begin by finding the surface brightness of the secondary star in V , on the assumption that the
surface brightness is similar to that of main-sequence stars of the same spectral type; the Barnes-
Evans relation for late-type stars is discussed by Beuermann (2006). Combining the known Porb
with the assumption that the secondary fills its Roche critical lobe yields the secondary’s radius
R2 as a function of its mass, M2. In the relevant range of mass ratio, R2 ∝ M
2 , approximately,
and the dependence on M1 is weak enough to ignore. We generally do not know M2, so we guess
at a generous allowable range for this parameter using evolutionary simulations by Baraffe & Kolb
(2000) as a guideline; the weakness of the dependence of R2 on M2 means that this (rather ques-
tionable) step does not dominate the error budget. Combining the surface brightness with R2
yields the absolute magnitude MV . Subtracting this from the apparent magnitude measured for
the secondary star gives a distance modulus. The reddening maps of Schlegel et al. (1998) then
can be used to estimate the extinction. Note carefully that we do not assume that the secondary
is a ‘normal’ main-sequence star; we assume only that the secondary’s surface brightness is similar
to field stars of the same spectral type. The normalization of the secondary’s contribution also
depends on the assumption that the spectral features used to judge the subtraction are similar in
strength to those of a normal star.
As noted earlier, the immediate aim of our observations was to find orbital periods from radial
velocity time series data. The Hα emission line is usually the strongest feature, and it generally
gives good results in dwarf novae. All the emission-line velocities reported here are of Hα.
We measured radial velocities of Hα emission using convolution methods described by Schneider & Young
(1980) and Shafter (1983). In this technique one convolves an antisymmetric function with the line
profile, and takes the zero of the convolution (where the two sides of the line contribute equally)
as the line center. For the antisymmetric function with which the spectrum is convolved, we used
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either the derivative of a Gaussian with adjustable width, or positive and negative Gaussians of
adjustable width offset from each other by an adjustable separation. Uncertainties in the convolu-
tion velocities are estimated by propagating forward the counting-statistics errors in the individual
data channels; in practice, these are lower limits to the true uncertainties, since the line profile
can vary in ways unrelated to the orbital modulation. The choice of convolution parameters is
dictated by the shape and width of the line, and in practice the parameters are adjusted to give
the best detection of the orbit. The physical interpretation of CV emission lines is complicated
and controversial (see, e.g., Shafter 1983, Marsh 1988, Robinson 1992), but in almost all cases the
emission-line periodicity accurately reflects Porb (though Araujo-Betancor et al. 2005 describe a
noteworthy exception to this rule). A sample of the radial velocities for each object are listed in
Table 4, while the full tables can be found online.
One of our systems, BF Eri, has a K-type absorption component in its spectrum. We measured
velocities of this using the cross-correlation radial velocity package described by Kurtz & Mink
(1998), using the region from 5000 to 6500 Å, and excluding the region containing the He I λ5876
emission line and and the NaD absorption complex. For a cross-correlation template spectrum, we
used the a velocity-compensated sum of many observations of IAU velocity standards taken with
the same instrument, as described in Thorstensen et al. (2004).
We searched for periods in all the velocity time series using the “residualgram” method
(Thorstensen et al. 1996); the resulting periodograms are given in Figs. 3 and 4. At the best
candidate periods we fitted least-squares sinusoids of the form v(t) = γ + K sin[2π(t − T0)/P ].
Fig. 5 shows the velocities folded on the best-fitting periods, and Table 5 gives the parameters of
these fits. Because of limitations of the sampling (e.g., the need to observe only at night from a
single site), a single periodicity generally manifests as a number of alias frequencies. To assess the
confidence with which we could assert that the strongest alias is the true period, we used a Monte
Carlo test described by Thorstensen & Freed (1985).
The alias problem can be particularly irksome over longer timescales; in this case the uncertain
number of cycles elapsed between observing runs causes fine-scale “ringing” in the periodogram.
The individual periods have tiny error bars, because of the large time span covered, but the am-
biguity in period means that a realistic error bar – one that covers the range of possibilities – is
much larger. In those cases, the period uncertainties given in Table 5 are estimated by analyzing
data from the individual observing runs separately. When only two observing runs are available,
the allowable fine-scale frequencies are well-described by a fitting formula
Porb = (t2 − t1)/n.
Here t1 and t2 are the epochs of blue-to-red velocity crossing observed on the two runs, and n
is the integer number of cycles that have passed between t1 and t2. The allowed range of n is
determined from the weighted average of the periods derived from separate fits to the two runs’
data. When more than two observing runs are available, the situation becomes more complex. In
some happy cases there are enough overlapping constraints that only a single, very precise period
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remains tenable. We were able to find such precise periods for LX And, LU Cam, and BF Eri.
3. Notes on Individual Objects
We discuss the stars in alphabetical order by constellation.
3.1. LX Andromedae
LX And was first identified as a variable star (RR V-3) in the Lick RR Lyrae search (Kinman et al.
1982). It was classified incorrectly as an RV Tauri star, and its dwarf nova nature was unrecognized
until the photometric study by Uemura et al. (2000). Morales-Rueda & Marsh (2002) obtained
spectra of LX And as part of their study of dwarf novae in outburst and determined the equivalent
widths and FWHMs of the Balmer and He II lines. Our mean spectrum appears typical for a dwarf
nova at minimum light.
Because of the large hour-angle span, the radial velocity time series leaves no doubt about
the daily cycle count, which is near 6.6 cycle d−1. The several observing runs constrain the fine-
scale period in a more complicated way, but the Monte Carlo test indicates that a precise period
of 0.1509743(5) d is preferred with about 98 per cent confidence. Two other candidate periods
separated from this by 1 cycle per 53.2 d in frequency are much less likely.
3.2. CZ Aquilae
Very little has been published on CZ Aql, which is listed in the Archival Edition as a U-
Gem dwarf nova. Cieslinski et al. (1998) included the star in their spectroscopic study of irregular
variables, and noted a probable 4.8 hour period and emission lines typical of dwarf novae. Our
velocities confirm the suggested 4.8-hour period, but we cannot determine a unique cycle count
between our observing runs.
While the spectrum superficially resembles that of a dwarf nova, a closer look reveals interesting
behavior. Fig. 6, constructed using methods described by Taylor et al. (1999), presents our spectra
as a phase-averaged greyscale image. There is a striking broad component in the stronger Balmer
and HeI lines that shows a large velocity excursion, with the red wing of Hα reaching to +3100 km
s−1 at phase 0.3 (where phase 0 corresponds to the blue-to-red crossing of the line core). The broad
components around Hβ and λ6678 move in phase with those of Hα and range from 900 to 2600
and −2500 to −900 km s−1 and 700 to 2100 and −1000 to −600 km s−1, respectively. The wings
of λ5876 also move in phase with the others, but the red edge is difficult to follow at its minimum
because of interference from the NaD absorption lines, which are stationary and hence interstellar.
The maximum of the red edge is 3200 km s−1, while the blue edge ranges from −1500 to −600 km
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s−1. The blueward wing of all these lines is noticeably weaker than the redward wing.
Other emission lines present include HeII λ4686, HeI λ4713 and λ4921, and, very weakly, FeII
λ5169. We also detect unidentified emission lines at λ6344, as is also seen in LS Peg (Taylor et al.
1999), and at λ5046. The strength of the λ5780 diffuse interstellar band (Jenniskens & Desert
1994) and the NaD lines suggest that a good deal of interstellar material lies along the line of sight,
and that the luminosity is relatively high.
High-velocity wings reminiscent of the ones seen here have been seen in V795 Her (Casares et al.
1996; Dickinson et al. 1997), LS Peg (Taylor et al. 1999), V533 Her (Thorstensen & Taylor 2000),
and RX J1643+34 (Patterson et al. 2002), all of which are SW Sex stars. We do not, how-
ever, detect another SW Sex characteristic, namely phase-dependent absorption in the HeI lines
(Thorstensen et al. 1991). The orbital periods of most SW Sex stars are shorter than 4 hours, so
CZ Aql’s 4.8-h period would be unusually long for an SW Sex star.
3.3. LU Camelopardalis
Jiang et al. (2000) obtained the first spectrum of this dwarf nova in a follow-up study of CV
candidates from the ROSAT All Sky Survey. We found no other published spectroscopic studies.
Our velocities constrain the period to a unique value, 0.1499685(7) d. The averaged spectrum
shows a rather strong, blue continuum, which may indicate a state somewhat above true minimum.
3.4. GZ Cancri
Jiang et al. (2000) confirmed the cataclysmic nature of GZ Cnc by obtaining the first spectrum
of the object. Kato et al. (2002) suggested that this star, originally labeled as a dwarf nova, could
possibly be an intermediate polar (DQ Her star), based on similarities in its long-term photometric
behavior to that of other intermediate polars. Tappert & Bianchini (2003) conducted a photometric
and spectroscopic study of the system. Using advance results from the present study to help decide
the daily cycle count, they found Porb = 0.08825(28) d, or 2.118(07) h, placing the system near
the lower edge of the so-called gap in the CV period distribution – a dearth of systems in the
period range from roughly 2 to 3 hr. Tappert & Bianchini (2003) also saw characteristics that
could indicate an intermediate polar classification, but did not claim their evidence was definitive
on this point.
Almost all our observations come from two observing runs a year apart. The full set of
velocities strongly indicates an orbital frequency near 11.4 cycle d−1, with the Monte Carlo test
giving a discriminatory power greater than 0.99 for the choice of daily cycle count. However, the
number of cycles between the two observing runs is not determined. Precise periods that fit the
combined data set are given by P = [349.785(3) d]/n, where n is the integer number of cycle counts;
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n = 3972± 8 corresponds to roughly 1 standard deviation. While our period agrees well with that
of Tappert & Bianchini (2003), our data neither support nor disprove the claim that GZ Cnc may
be an intermediate polar.
3.5. V632 Cygni
Liu et al. (1999) offer the only published spectrum of this dwarf nova. They measured the
equivalent widths and integrated line fluxes of the Balmer, HeI, and HeII emission lines and sug-
gested that the orbital period is likely short based on the very strong Balmer emission. Our spec-
trum appears similar to theirs, and our measured flux level is also nearly the same. The periodigram
in Fig. 3 clearly favors an orbital frequency near 15.7 cycles d−1, with a discriminatory power of 95
per cent and a correctness likelihood near unity. This confirms the suggestion of Liu et al. (1999)
that the period is rather short and suggests that it is an SU UMa-type dwarf nova.
3.6. V1006 Cygni
Bruch & Schimpke (1992) present the only published spectrum we know of, and characterized
it as a “textbook example” of a dwarf nova spectrum. They noted a slightly blue continuum with
strong Balmer and He I emission, as well as clear He II λ4686 and Fe II emission. Our spectrum
(Fig. 1) is similar to theirs both in appearance and normalization, and our line measurements
(Table 3) are also comparable.
The periodogram (Fig. 4) indicates a frequency near 10.1 cycles d−1, and the Monte Carlo test
confirms that the daily cycle count is securely determined. Most of our data are from 2004 June,
but we returned in 2005 June/July to confirm the unusual period indicated in the earlier data.
The periods found by analyzing the two runs separately are consistent within their uncertainties.
As with GZ Cnc, there are multiple choices for the cycle count between the two observing runs;
the best-fitting periods are given by P = [369.006(4) d]/n, where n = 3726 ± 4 corresponds to
1 standard deviation. Including a few velocities from other observing runs suggests that n is
slightly larger, perhaps 3728. In any case, the period amounts to 2.38 h, which places V1006 Cyg
firmly in the period gap (Warner 1995), where there is apparently a true scarcity of dwarf novae
(Hellier & Naylor 1998).
3.7. BF Eri
The first evidence that BF Eridani was a cataclysmic variable came when an Einstein X-
ray source, 1ES0437-046, was matched to the variable (Elvis et al. 1992). Schachter et al. (1996)
confirmed this match and presented an optical spectrum. Kato (1999) and the Variable Star
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Observers’ League in Japan (VSOLJ) found photometric variability characteristic of a dwarf nova.
The spectrum of BF Eri (Fig. 2) shows a significant contribution from a K star along with the
usual dwarf-nova emission lines. Normally, this suggests that Porb > 6 h. Nearly all our spectra
yielded good cross-correlation radial velocity measurements as well as emission-line velocities. The
absorption- and emission-line velocities independently give a period near 6.50 h (Table 5), in ac-
cordance with expectation based on the spectrum. There is no ambiguity in cycle count over the
5-year span of the observations, so the period is precise to a few parts per million. Fig. 6 shows a
phase-resolved average of the BF Eri spectra, with the absorption spectrum shifting in antiphase
to the emission lines.
If the emission-line velocities faithfully trace the primary’s center-of-mass motion, and the
absorption-line velocities also trace the secondary’s motion, then the two velocity curves should be
exactly one-half cycle out of phase. In BF Eri, we find a shift of 0.515 ± 0.007 cycles between the
two curves, consistent with 0.5 cycles, so we feel emboldened to explore the system dynamics.
Masses can only be derived when the orbital inclination is known, as in eclipsing systems. To
see if BF Eri might eclipse, we derived differential magnitudes from images that were taken for
astrometry (discussed below) and plotted them as a function of orbital phase. Some images were
taken at the phase at which an eclipse would appear, but no evidence for an eclipse was found.
Limits on the depth and duration of the eclipse are difficult to quantify because the data were taken
in short bursts in the presence of strong intrinsic variability, so a weak eclipse cannot be ruled out,
but the photometry does suggest that the inclination is not close to edge-on.
Because the system apparently does not eclipse, we cannot derive masses; rather, we find broad
constraints on the inclination by assuming astrophysically reasonable masses for the components.
Taken at face value, the velocity amplitudes K imply a mass ratio q = M2/M1 = 0.60 ± 0.03. If
we arbitrarily choose a white dwarf mass M1 = 0.9 M⊙ (so that M2 = 0.53 M⊙), the observed K
velocities imply i = 50 degrees. To find a rough lower limit on the inclination, we consider a massive
white dwarf (M1 = 1.2 M⊙) and, ignoring the constraint on q for the moment, take M2 = 0.4 M⊙;
this yields i = 40 degrees. For a rough upper limit, we assume M1 = 0.6 M⊙ and M2 = 0.4 M⊙,
which gives i = 67 degrees.
The decomposition procedure described earlier yielded a spectral type of K3 ±1 subclass; the
result of the subtraction is shown in Fig. 2. Using the V passband tabulated by Bessell (1990) and
the IRAF sbands task, we find a synthetic V = 16.9± 0.3 for the K star’s contribution. Taking the
range of plausible secondary star masses to be 0.4 to 0.8 M⊙ yields R2 = 0.7± 0.1 R⊙ at this Porb.
Combining this with the surface brightness expected at this spectral type yields MV = 6.8 ± 0.4
for the secondary. If there is no significant interstellar extinction, we have m − M = 10.1 ± 0.5,
or a distance of approximately 1100 ± 300 pc. The dust maps of Schlegel et al. (1998) give a total
E(B − V ) = 0.062 in this direction. Assuming that BF Eri is beyond the Galactic dust and taking
AV /E(B − V ) = 3.3 gives an extinction-corrected (m−M)0 = 9.9, and a distance estimate of 950
(+250,−200) pc.
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We can also estimate a distance using the relation found by Warner (1987) between Porb, i,
and the absolute magnitude at maximum light MV (max). Using our inclination constraints, the
Warner relation predicts MV max = 3.9±0.7 at this orbital period. The General Catalog of Variable
Stars (Kholopov et al. 1999) lists mp = 13.2 at maximum light; taking this to be similar to Vmax
yields m−M = 9.3, or 9.1 corrected for extinction, which corresponds to 660 pc.
Given these distance estimates, it is surprising that BF Eri has a very substantial proper
motion. The Lick proper motion survey (Hanson et al. 2004) gives [µX , µY ] = [+34,−97] mas
yr−1. We have begun a series of parallax observations with the Hiltner 2.4m telescope using the
protocols described by Thorstensen (2003); so far we have five epochs from 2005 November and 2007
January. The proper motion relative to the background stars is [µX , µY ] = [32,−111] mas yr
and the parallax is not detected, with a nominal value of 1 ± 2 mas. The parallax determination
is very preliminary, but given the data so far we estimate the lower limit on the distance based on
the astrometry alone to be ∼ 200 pc.
At the nominal 950 pc distance derived from the secondary star, a 100 mas yr−1 proper
motion corresponds to a transverse velocity vT = 451 km s
−1. This is implausibly large, so we
are left wondering how we might have overestimated the distance. One effect might be as follows.
Our distance is based on the secondary’s apparent brightness, and we estimate the secondary’s
contribution to the total light by searching for the best cancellation of its features. If the secondary’s
absorption lines are weaker than those in the spectral-type standards, we would underestimate the
secondary’s contribution. In our best decomposition, the secondary is about 2.2 magnitudes fainter
than the total light in V . Assuming (unrealistically) that all the light is from the secondary would
therefore decrease the distance modulus by 2.2 magnitudes, to a distance of 340 pc.
We do not yet have enough information to resolve the conundrum posed by BF Eri’s unmis-
takably large proper motion and its apparently large distance, but a reasonable compromise might
be to put it at something like 400-500 pc, with an underluminous, low-metallicity secondary. The
cross-correlation velocities of the secondary have a zero point determined to ±5 km s−1, more or
less, and give a substantial systemic velocity of −72±3 km s−1, or −86 km s−1 in the local standard
of rest. If the star is at 450 pc, its space velocity with respect to the local standard of rest is ∼250
km s−1, with Galactic components [U, V,W ] = [−180,−180,−3] km s−1, that is, the velocity is
mostly parallel to the Galactic plane and lags far behind the rotation of the Galactic disk. This
would put BF Eri on a highly eccentric orbit; these are halo-population kinematics (even though
the star remains close to the plane). These kinematics would be qualitatively consistent with the
weak-line conjecture used earlier.
3.8. BI Ori
Szkody (1987) published the first quiescent spectrum of BI Orionis, which showed emission
lines typical of dwarf novae. Morales-Rueda & Marsh (2002) show a spectrum in outburst and
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note the possible presence of weak HeII λ4686 emission.
Only the 2006 January velocities are extensive enough for period finding; they give Porb =
4.6 hr, with no significant ambiguity in the daily cycle count. The average spectrum shows the
usual emission lines; M-dwarf absorption features are also visible, though the signal-to-noise of the
individual spectra was not adequate for finding absorption-line velocities. Using the procedures
described earlier, we estimate the secondary’s spectral type to be M2.5 ± 1.5, with the secondary
alone having V = 20.0 ± 0.4. Assuming that the secondary’s mass lies in the broad range from
0.2 to 0.6 M⊙, its radius at this Porb would be 0.35 to 0.6 R⊙. Combining this with the surface
brightness derived from the spectral types gives an absolute magnitude MV = 10.2 ± 1.0. The
distance modulus, uncorrected for extinction, is therefore m − M = +9.8 ± 1.1, corresponding
to 910(+600,−360) pc. Schlegel et al. (1998) estimate a total reddening E(B − V ) = 0.11 in this
direction; assuming that BI Ori lies beyond all the dust, and taking AV /E(B−V ) = 3.3 reduces the
distance to ∼ 770 pc. At maximum light, BI Ori has mp = 13.2 (Kholopov et al. 1999). Assuming
the color is neutral, we find MV = 3.4 ± 1.1 at maximum. At BI Ori’s period, the Warner (1987)
relation predicts MV > 3.6 (with the brightest value corresponding to i = 0). This agrees broadly
with our nominal value based on the secondary star’s distance, but is a little fainter, suggesting
that BI Ori is not too far from face-on, or a little closer than our nominal distance, or both.
3.9. FO Per
FO Persei was apparently discovered by Morgenroth (1939), but its cataclysmic nature was
not immediately recognized. Bruch (1989) obtained spectra and gave equivalent widths for the
Balmer lines for two different nights of observations, between which the continuum changed from
relatively flat to inclined toward the red.
The emission lines in FO Per are rather narrow (Fig. 1, Table 3). This is often taken to indicate
a low orbital inclination. The velocity amplitude K is small, so that K/σ ≈ 1.6 for the best fits
(Table 5). Because of this, the daily cycle count remains ambiguous; the orbital frequency is either
5.8 or 6.8 cycle d−1, corresponding to Porb of 3.52 or 4.13 hr. CVs with periods in the 3-4 hour
range tend to be novalike variables (Shafter 1992), whereas FO Per is a dwarf nova; thus the 4.13
hr period is more likely a priori.
4. Summary
We have determined the orbital periods of eight CVs without significant daily cycle count
ambiguity; for FO Per, the period is narrowed to two choices. For three of the systems we find
high-precision periods by establishing secure cycle counts over long baselines.
While most of these objects are similar to others already known, three stand out as especially
– 12 –
interesting. CZ Aql shows asymmetric, high-velocity wings around the Balmer and HeI λ5876 and
λ6678 lines, possibly indicating a magnetic system. BF Eri’s proper motion of ∼ 100 mas yr−1 is
surprising in view of the large distance indicated by its secondary spectrum and by the Warner
relation; even if it is somewhat nearer than these indicators suggest, its kinematics are not typical of
disk stars. Finally, the orbital period of V1006 Cyg places it squarely in the middle of the so-called
period gap between 2 and 3 hours.
Acknowledgments. We are most grateful for support from the National Science Foundation
through grants AST-9987334 and AST-0307413. Bill Fenton took most of the spectra of GZ Cnc,
and J. Cameron Brueckner assisted with the BF Eri spectroscopy. Some of the astrometric images of
BF Eri were obtained by Sébastien Lépine and Michael Shara of the American Museum of Natural
History. We would like to thank the MDM Observatory staff for their skillful and conscientious
support. Finally, we are grateful to the Tohono O’odham for leasing us their mountain for a while,
so that we may study the glorious universe in which we all live.
– 13 –
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This preprint was prepared with the AAS LATEX macros v5.2.
– 16 –
Fig. 1.— Plots of the average flux-calibrated spectra for eight of the stars studied here. The weak
features seen near λ5577 are artifacts caused by imperfect subtraction of the strong [OI] night-sky
emission.
– 17 –
Fig. 2.— Plot of the averaged spectrum of BF Eri (top) and the spectrum after scaled late-type
(K3V) star has been subtracted (bottom). The spectra have been shifted into a rest frame before
averaging and do not include the 2006 March or 2007 January data.
– 18 –
Fig. 3.— Periodigrams for most of the stars studied here. The vertical axis in each case is the
inverse of chi-square for the least-squares best fitting sinusoid at each trial frequency. When data
from more that one observing run are combined, the periodigram can require hundreds of thousands
of points to resolve the fine-scale ringing; in those cases, the curve shown is formed by connecting
local maxima of the periodogram with straight lines. In those cases the right-hand panel gives a
close-up view of the region around the highest peak, revealing the alias structure resulting from
different choices of cycle count between the observing runs. The periodogram of BF Eri is for the
absorption-line velocities.
– 19 –
Fig. 4.— Periodigrams for the remainder of the stars, plotted in the same manner as the previous
figure. Because the choice of daily cycle count for FO Per remains ambiguous, we have not chosen
to enlarge either peak region.
– 20 –
Fig. 5.— Radial velocities plotted as a function of phase using the adopted orbital periods. For
CZ Aql, GZ Cnc, and V1006 Cyg, the number of cycle counts between observing runs is unknown,
and the exact period chosen to fold the velocities is one of a number of possibilities. The two
plots for FO Per are for different choices of the daily cycle count, and each of these in turn is also
an arbitrary choice among many finely-spaced periods. In BF Eri, both emission and absorption
velocities are plotted; the absorption velocities are shown with error bars.
– 21 –
Fig. 6.— Phase-averaged spectra of CZ Aql (top two panels) and BF Eri (bottom two panels),
presented as a greyscale. The scale is inverted, so that emission is represented by darker shades.
The two CZ Aql spectra are scaled differently to show the line cores (top) and the extent of the the
line wings. Note the NaD lines in CZ Aql remain stationary, indicating an interstellar origin. The
feature at λ6280 is telluric. BF Eri’s spectrum is plotted in two overlapping sections; the K-star’s
orbital motion is plainly visible.
– 22 –
Table 1. Stars Observed
Star α2000
a δ2000 Epoch
b Vobs
c maxd min
[hh:mm:ss] [◦:′:”] [mag.] [mag.] [mag.]
LX And 2:19:44.08 +40:27:22.3 2006.7 16.3 13.5p 16.4p
CZ Aql 19:19:58.21 −07:10:55.2 2003.4 15.4 13.p 15.p
LU Cam 5:58:17.86 +67:53:46.2 2002.0 16.3 14.v (16.v
GZ Cnc 9:15:51.68 +09:00:49.6 2000.3 15.4 13.1v 15.4v
V632 Cyg 21:36:04.22 +40:26:19.4 2000.5 17.9 12.6p 17.5p
V1006 Cyg 19:48:47.20 +57:09:22.8 2000.5 17.8 15.4p 17.0p
BF Eri 4:39:29.96 −04:35:59.5 2006.2 14.8 13.5p 15.5p
BI Ori 5:23:51.77 +01:00:30.6 2002.8 17.1 13.2p 16.7p
FO Per 4:08:34.98 +51:14:48.5 2004.0 17.1 11.8v 16.v
aPositions measured from images taken at the 2.4m Hiltner telescope, using
astrometric solutions from fits to USNO A2.0 (Monet et al. 1996) or UCAC 2
(Zacharias et al. 2004) stars. Uncertainties are of order 0.1 arcsec.
bThe date of the image used in the position measurement. The coordinate
system (equator and equinox) is J2000 in all cases.
cSynthesized from spectra, as described in text.
dTaken from the GCVS (Kholopov et al. 1999). Photgraphic magnitudes
flagged with ‘p’, visual with ‘v’.
– 23 –
Table 2. Journal of Observations
data N HA start HA end telescope
(UT) [hh:mm] [hh:mm]
LX And
2004 Jan 13 1 +2:35 +2:35 2.4m
2004 Mar 02 3 +3:45 +4:00 2.4m
2004 Nov 18 1 +0:29 +0:29 1.3m
2004 Nov 19 26 −3:36 +4:17 1.3m
2004 Nov 19 2 +2:31 +2:35 2.4m
2004 Nov 20 13 −3:17 +3:15 1.3m
2004 Nov 20 1 −1:23 −1:23 2.4m
2006 Jan 19 14 +1:31 +4:10 1.3m
2006 Jan 22 8 +3:31 +4:46 1.3m
2007 Jan 26 3 +1:36 +1:58 1.3m
2007 Jan 27 15 +0:58 +3:39 1.3m
CZ Aql
2005 Jul 02 3 +1:40 +2:04 1.3m
2005 Jul 04 48 −3:28 +3:17 1.3m
2005 Jul 05 13 −1:55 −0:06 1.3m
2005 Jul 06 12 −3:57 +2:28 1.3m
2005 Sep 03 2 +1:23 +1:32 1.3m
2005 Sep 07 2 −0:03 +0:05 1.3m
2005 Jun 28 2 −0:01 +0:04 2.4m
2006 Jun 18 2 +2:05 +2:10 2.4m
2006 Jun 19 2 +0:36 +0:40 2.4m
2006 Jun 23 3 −1:24 −1:12 2.4m
– 24 –
Table 2—Continued
data N HA start HA end telescope
(UT) [hh:mm] [hh:mm]
LU Cam
2002 Jan 22 2 −1:40 −1:19 2.4m
2002 Jan 23 8 −1:33 +2:31 2.4m
2002 Jan 24 25 −3:15 +5:47 2.4m
2002 Feb 18 2 +2:20 +2:28 2.4m
2002 Feb 19 2 +2:41 +2:50 2.4m
2002 Feb 20 4 −0:02 +3:43 2.4m
2002 Feb 22 1 +2:18 +2:18 2.4m
2004 Jan 16 2 +2:49 +2:53 2.4m
2004 Jan 17 1 +0:40 +0:40 2.4m
2004 Jan 19 1 −0:32 −0:32 2.4m
2004 Mar 02 1 +0:53 +0:53 2.4m
2004 Mar 07 6 +2:29 +3:14 2.4m
2004 Nov 19 4 +2:48 +3:27 2.4m
2005 Mar 21 1 +1:12 +1:12 2.4m
2005 Mar 22 2 +1:21 +1:30 2.4m
2005 Sep 09 1 −2:00 −2:00 2.4m
2005 Sep 12 1 −1:49 −1:49 2.4m
2006 Jan 09 2 +1:49 +1:55 2.4m
GZ Cnc
2000 Apr 07 2 +4:26 +4:32 2.4m
2000 Apr 08 1 +0:27 +0:27 2.4m
2000 Apr 10 5 −0:31 +4:32 2.4m
2000 Apr 11 15 +2:19 +4:22 2.4m
2001 Mar 24 2 +4:39 +4:49 2.4m
2001 Mar 25 21 −1:10 +3:10 2.4m
2001 Mar 26 20 −0:50 +4:27 2.4m
2001 Mar 27 2 −0:08 +0:01 2.4m
2001 Mar 28 3 +0:14 +0:25 2.4m
V632 Cyg
2005 Jul 07 2 +1:00 +1:16 1.3m
2005 Jul 08 10 −5:12 +0:33 1.3m
2005 Jul 09 18 −5:03 +1:09 1.3m
2005 Jul 10 18 −5:09 +1:07 1.3m
2005 Jul 11 3 +0:52 +1:19 1.3m
– 25 –
Table 2—Continued
data N HA start HA end telescope
(UT) [hh:mm] [hh:mm]
V1006 Cyg
2003 Jun 22 1 +1:06 +1:06 2.4m
2004 Jun 24 5 +0:52 +1:56 1.3m
2004 Jun 25 5 −1:43 +0:34 1.3m
2004 Jun 25 1 +4:06 +4:06 2.4m
2004 Jun 26 10 −4:25 +1:56 1.3m
2004 Jun 27 3 −3:00 −2:00 1.3m
2004 Jun 28 4 −2:25 +1:02 1.3m
2004 Jun 28 1 +0:28 +0:28 2.4m
2004 Jun 29 5 +0:55 +1:59 1.3m
2004 Jun 29 1 +4:07 +4:07 2.4m
2004 Jun 30 18 −4:39 +2:26 1.3m
2004 Jul 01 10 −3:58 +1:53 2.4m
2005 Jul 05 13 +0:58 +3:11 1.3m
2004 Jun 30 4 −3:41 +3:52 2.4m
2004 Jul 01 12 −4:07 +2:17 1.3m
2005 Jul 03 27 −3:57 +1:20 1.3m
2005 Jul 05 13 +0:58 +3:11 1.3m
BF Eri
2001 Dec 18 3 +2:35 +2:56 1.3m
2001 Dec 19 10 −3:06 +4:04 1.3m
2001 Dec 20 12 −3:49 +2:00 1.3m
2001 Dec 21 2 +3:04 +3:14 1.3m
2001 Dec 22 5 −3:13 −2:31 1.3m
2001 Dec 23 12 −3:18 +4:35 1.3m
2001 Dec 24 13 −3:17 +3:08 1.3m
2001 Dec 25 14 −2:20 +1:07 1.3m
2001 Dec 26 8 −2:10 +3:05 1.3m
2001 Dec 27 18 −2:39 +4:14 1.3m
2002 Jan 19 1 +1:27 +1:27 2.4m
2002 Jan 20 2 −1:33 +2:13 2.4m
2002 Jan 22 1 −2:02 −2:02 2.4m
2002 Feb 21 2 +1:26 +1:35 2.4m
2002 Feb 22 2 +0:40 +0:49 2.4m
2002 Oct 26 2 −0:11 +0:05 2.4m
2002 Oct 31 1 +3:21 +3:21 2.4m
2003 Feb 02 1 +0:04 +0:04 2.4m
2005 Sep 11 2 −0:57 −0:40 1.3m
2006 Mar 16 1 +1:56 +1:56 1.3m
2006 Mar 17 5 +2:15 +2:57 1.3m
– 26 –
Table 2—Continued
data N HA start HA end telescope
(UT) [hh:mm] [hh:mm]
2007 Jan 28 9 −1:46 −0:13 1.3m
BI Ori
2006 Jan 20 32 −2:54 +4:05 1.3m
2006 Jan 21 22 −2:47 +2:21 1.3m
2006 Jan 23 6 +1:03 +2:08 1.3m
FO Per
1995 Oct 09 9 −4:40 −3:26 2.4m
1995 Oct 10 5 −5:22 +1:26 2.4m
1996 Dec 19 14 +1:38 +4:03 1.3m
1996 Dec 20 5 +2:08 +3:43 1.3m
2001 Dec 18 9 −2:35 +4:36 1.3m
2004 Nov 18 1 −0:11 −0:11 1.3m
2004 Nov 19 9 +2:52 +5:08 1.3m
2004 Nov 19 2 +3:24 +2:48 2.4m
2004 Nov 20 20 −0:56 +5:10 1.3m
2006 Jan 10 12 −0:35 +2:35 1.3m
2006 Jan 10 1 +4:21 +4:21 2.4m
2006 Jan 11 18 −1:29 +2:41 1.3m
2006 Jan 11 1 −3:12 −3:12 2.4m
2006 Jan 12 4 −1:16 −0:36 1.3m
2006 Jan 13 10 +3:44 +5:37 1.3m
2006 Jan 16 11 −1:36 +1:00 1.3m
– 27 –
Table 3. Spectral Features in Quiescence
E.W.a Flux FWHM b
Feature (Å) (10−16 erg cm−2 s1) (Å)
LX And
Hβ 45 690 18
HeI λ4921 4 60 25
HeI λ5015 3 50 20
Fe λ5169 2 20 14
HeI λ5876 11 120 19
Hα 54 560 17
HeI λ6678 4 40 19
CZ Aql
Hβ 21 670 18
HeI λ4921 1 40 12
HeI λ5015 2 50 13
HeI λ5876 7 160 15
NaD −1 −16 · · ·
Hα 61 1250 27
HeI λ6678 5 90 16
LU Cam
Hγ 10 170 12
HeI λ4471 2 30 10
Hβ 14 190 12
HeI λ4921 1 20 12
HeI λ5015 2 20 13
Fe λ5169 1 10 12
HeI λ5876 4 50 12
Hα 27 240 13
HeI λ6678 3 20 15
HeI λ7067 3 20 · · ·
GZ Cnc
Hγ 26 940 26
HeI λ4471 8 260 28
HeII λ4686 5 140 46
Hβ 36 1040 25
HeI λ4921 5 140 26
HeI λ5015 4 100 28
Fe λ5169 2 60 26
HeI λ5876 9 200 27
– 28 –
Table 3—Continued
E.W.a Flux FWHM b
Feature (Å) (10−16 erg cm−2 s1) (Å)
Hα 38 790 25
HeI λ6678 4 90 31
HeI λ7067 3 60 32
V632 Cyg
Hβ 80 260 24
HeI λ4921 6 20 27
HeI λ5015 8 20 26
Fe λ5169 5 10 26
HeI λ5876 28 70 27
Hα 113 260 27
HeI λ6678 15 30 32
V1006 Cyg
Hβ 74 250 27
HeI λ4921 8 30 30
HeI λ5015 8 20 28
Fe λ5169 8 30 28
HeI λ5876 26 70 34
Hα 108 250 31
HeI λ6678 11 30 38
BF Eri
HeII λ4686 6 240 46
Hβ 23 1060 22
HeI λ5015 2 110 29
HeI λ5876 5 240 21
Hα 27 1200 22
HeI λ6678 3 120 27
BI Ori
Hβ 34 180 32
HeI λ4921 6 30 36
HeI λ5015 6 30 43
Fe λ5169 5 30 39
HeI λ5876 6 30 31
Hα 36 190 31
FO Per
– 29 –
Table 3—Continued
E.W.a Flux FWHM b
Feature (Å) (10−16 erg cm−2 s1) (Å)
Hβ 24 160 9
HeI λ4921 2 20 7
HeI λ5015 3 20 7
Fe λ5169 2 10 7
HeI λ5876 5 30 7
Hα 29 170 9
HeI λ6678 2 10 9
aEmission equivalent widths are counted as positive.
bFrom Gaussian fits.
– 30 –
Table 4. Radial Velocities
Star time a vabs σvabs vemn σvemn
(km s−1) (km s−1) (km s−1) (km s−1)
LX And 53017.7061 · · · · · · −44 −11
LX And 53066.6160 · · · · · · −57 −8
LX And 53066.6208 · · · · · · −37 −8
LX And 53066.6263 · · · · · · −77 −7
aHeliocentric Julian date of mid-integration, minus 2400000.
Note. — All emission-line velocities are of Hα. Emission-line velocity
uncertainties are derived from counting statistics and should be regarded as
lower limits. Table 4 is published in its entirety in the electronic version of
the Publications of the Astronomical Society of the Pacific. A short portion
is shown here for guidance regarding its form and content.
– 31 –
Table 5. Fits to Radial Velocities
Star Algorithma T0
b P K γ N σc
(d) (km s−1) (km s−1) (km s−1)
LX And G2,21,7 53754.6861(12) 0.1509743(5) 81(4) −48(3) 87 15
CZ Aql G2,18,8 53557.880(2) 0.2005(6)d 193(15) 5(10) 89 61
LU Cam D,15 52327.7421(14) 0.1499686(4) 57(4) 44(3) 66 14
GZ Cnc G2,15,9 51992.8928(13) 0.0881(4)d 79(7) 22(5) 71 26
V632 Cyg D,28 53560.9746(13) 0.06377(8) 62(8) −49(5) 51 28
V1006 Cyg G2,20,9 53187.9091(16) 0.09904(9)d 89(8) −11(6) 120 44
BF Eri emission G2,21,7 52574.0027(18) 0.2708801(6) 109(5) −91(3) 126 24
BF Eri absorption · · · 52573.8632(9) 0.2708805(4) 182(4) −72(3) 117 20
BF Eri mean: · · · · · · 0.2708804(4) · · · · · · · · ·
BI Ori 53756.541(3) G2,35,9 0.1915(5) 131(13) 24(9) 60 44
FO Per (shorter) D,11 52261.872(3) 0.1467(4)d 27(3) −49(2) 131 17
FO Per (longer) D,11 52261.893(3) 0.1719(5)d 27(3) −45(2) 131 17
Note. — Parameters of sinusoidal least-squares fits to the velocity timeseries, of the form v(t) = γ +K sin(2π(t −
T0)/P ). The quoted parameter uncertainties are based on the assumption that the scatter of the data around the best
fit is a realistic estimate of the velocity uncertainty (Cash 1979). In practice this is more conservative than assuming
that counting statistics uncertainties are realistic.
aCode for the convolution function used to derive emission line velocities; D = derivative of a Gaussian, G2 =
double-Gaussian function (see text). For the D algorithm the number that follows gives the line full-width at half-
maximum, in Å, for which the function is optimized; for the G2 algorithm the two numbers are respectively the
separation of the two Gaussians and their individual FWHMa, again in Å.
bHeliocentric Julian Date minus 2400000. The epoch is chosen to be near the center of the time interval covered
by the data, and within one cycle of an actual observation.
cRoot-mean-square residual of the fit.
dThe period determination in this case is complicated by unknown numbers of cycles between observing runs; the
uncertainty given here is an estimate based on fits to individual runs. Only certain values within the period range
given here are allowed; see text for details.
This figure "f6.png" is available in "png"
 format from:
http://arxiv.org/ps/0704.0948v1
http://arxiv.org/ps/0704.0948v1
	Introduction
	Observations, Reductions, and Analysis
	Observations
	Reductions
	Analysis
	Notes on Individual Objects
	LX Andromedae
	CZ Aquilae
	LU Camelopardalis
	GZ Cancri
	V632 Cygni
	V1006 Cygni
	BF Eri
	BI Ori
	FO Per
	Summary

  We present optical spectroscopy of nine cataclysmic binary stars, mostly
dwarf novae, obtained primarily to determine orbital periods Porb. The stars
and their periods are LX And, 0.1509743(5) d; CZ Aql, 0.2005(6) d; LU Cam,
0.1499686(4) d; GZ Cnc, 0.0881(4) d; V632 Cyg, 0.06377(8) d; V1006 Cyg,
0.09903(9) d; BF Eri, 0.2708804(4) d; BI Ori, 0.1915(5) d; and FO Per, for
which Porb is either 0.1467(4) or 0.1719(5) d.
  Several of the stars proved to be especially interesting. In BF Eri, we
detect the absorption spectrum of a secondary star of spectral type K3 +- 1
subclass, which leads to a distance estimate of approximately 1 kpc. However,
BF Eri has a large proper motion (100 mas/yr), and we have a preliminary
parallax measurement that confirms the large proper motion and yields only an
upper limit for the parallax. BF Eri's space velocity is evidently large, and
it appears to belong to the halo population. In CZ Aql, the emission lines have
strong wings that move with large velocity amplitude, suggesting a
magnetically-channeled accretion flow. The orbital period of V1006 Cyg places
it squarely within the 2- to 3-hour "gap" in the distribution of cataclysmic
binary orbital periods.

Introduction
Cataclysmic variables (CVs) are binary star systems in which the secondary, usually a late-
type main sequence star, fills its Roche lobe and loses mass to the white dwarf primary (Warner
1995). CVs are long-lived systems that are stable against mass transfer, so the mass transfer must
be driven by gradual changes in the orbit, or in the secondary star, or both. It is commonly
believed that the evolution of most CVs is driven by the slow loss of angular momentum from the
orbit, most likely through magnetic braking of the co-rotating secondary star, at least at longer
orbital periods Porb where gravitational radiation is ineffective (Andronov & Pinsonneault 2004
give a recent discussion). The loss of angular momentum constricts the Roche critical lobe around
the secondary and causes the system to transfer mass as it evolves toward shorter Porb. In this
scenario, Porb serves as a proxy measurement for the system’s evolutionary state. Correct and
complete orbital period measurements are fundamental to any accurate theory of CV evolution.
Given the usefulness of Porb, it is fortunate that it can usually be measured accurately and precisely.
This paper presents optical spectroscopy of the nine CVs listed in Table 1. We took these
observations mostly for the purpose of finding orbital periods using radial velocities (none of these
systems are known to eclipse). The long cumulative exposures also allowed us to look for any
unusual features. The Catalog and Atlas of Cataclysmic Variables Archival Edition (Downes et al.
2001) 2 lists seven of the stars as dwarf novae, one as either a dwarf nova or a DQ Her star, and one
simply as a cataclysmic, possibly a dwarf nova similar to U Gem or SS Cygni (type UGSS). Except
for CZ Aql, for which we confirm a 4.8-hour candidate period suggested by Cieslinski et al. (1998),
all of these objects lacked published orbital periods when we began working on them. Subsequently
Tappert & Bianchini (2003) found Porb = 0.0883 d for GZ Cancri; we had communicated our
advance findings to these authors so they could disambiguate their period determination.
2. Observations, Reductions, and Analysis
2.1. Observations
All our spectra were taken at the MDM Observatory on Kitt Peak, Arizona, using either the
1.3m McGraw-Hill telescope or the 2.4m Hiltner telescope. The earliest observations we report
here are from 1995, and the latest were obtained 2007 January. Table 2 gives a journal of the
observations.
At the 1.3m we used the Mark III spectrograph and a SITe 1024 × 1024 CCD detector. The
spectral resolution is 5.0 Å, covering a range of either 4480 to 6760 Å with 2.2 Å pixel−1 for the
2001 December BF Eri data, or 4646 to 6970 Å with 2.3 Å pixel−1 for the remaining data. The
2Available at http://archive.stsci.edu/prepds/cvcat/index.html; this had been called the Living Edition until its
author retired and ceased updates.
http://archive.stsci.edu/prepds/cvcat/index.html
– 3 –
2.4m spectra, except for those of FO Per, were obtained with the modular spectrograph and a
SITe 20482 CCD detector, with 2.0 Å pixel−1, over a range of 4210 to 7500 Å and with a spectral
resolution of 3.5 Å. The relatively small number of 2.4 m spectra of FO Per were taken with a
LORAL 20482-pixel detector, and cover from 4285 to 6870 Å at 1.25 Å pixel−1.
2.2. Reductions
For the most part we reduced the spectra using standard IRAF3 procedures. The wavelength
calibration was based on exposures of Hg, Ne, and Xe lamps. Prior to 2003 we took lamp exposures
through the night and whenever the telescope was moved. For the 2.4m data from 2003 to the
present, we used lamp exposures taken in twilight to find the shape of the pixel-to-wavelength
relation, and set the zero point individually for each nighttime exposure using the OI λ5577 night-
sky feature. The apparent velocity of the telluric OH emission bands at the far red end of the
spectrum, found with a cross-correlation routine, provided a check; although these are far from
the feature used to set the zero point, their apparent velocity typically remain within 10 km s−1
of zero. Because of the increased efficiency of this technique, we attempted to use it at the 1.3m
telescope also, during the 2004 June/July observing run. For unknown reasons the results were
unsatisfactory. To salvage the Hα emission velocities from that run, we determined a correction by
cross-correlating the night-sky emission features in the 6200-6625 Å range with a well-calibrated
night-sky spectrum obtained with a similar instrument. The correction was calculated for each
individual spectrum and then applied to each measured velocity, and it did reduce the scatter
somewhat, evidently because the wavelength range used includes the Hα emission line for which
we measured velocities.
On all our runs, we observed flux standards during twilight when the sky was clear, and applied
the resulting calibration to the data. The reproducibility of these observations suggests that our
fluxes are typically accurate to ±20 per cent. We also took short exposures of bright O and B
stars in twilight to map the telluric absorption features and divide them out approximately from
our program object spectra. Before flux calibration, we divided our program star spectra by a
mean hot-star continuum, in order to remove the bulk of the response variation. Table 1 lists
V magnitudes synthesized from our mean spectra, using the IRAF sbands task and the passband
tabulated by Bessell (1990); clouds, losses at the slit, and calibration errors make these uncertain
by a few tenths of a magnitude, but they do give a rough indication of the brightness of each system
at the time of our observation.
3IRAF is distributed by the National Optical Astronomy Observatories.
– 4 –
2.3. Analysis
Except for a few spectra taken in outburst (which show weak emission or absorption on a
strong continuum), all of the stars show the prominent emission lines. Figs. 1 and 2 show averaged
spectra, and Table 3 gives the equivalent width and FWHM of each line measured for each star
from its averaged spectrum.
Two stars, BF Eri and BI Ori, showed the spectral features of a late-type star. To quantify the
secondary contribution in these objects, we began by preparing averaged flux-calibrated spectra
(in BF Eri’s case the secondary’s radial velocity curve was measurable, so we shifted the individual
spectra to the secondary’s rest frame before averaging). Over time we have used the 2.4 m and
modular spectrograph to collect spectra of K and M stars classified by Keenan & McNeil (1989) or
Boeshaar (1976). The wavelength coverage and spectral resolution of these data are similar to the
1.3m data. We applied a range of scaling factors to the library spectra, subtracted them from the
averaged spectra, and examined the results by eye to estimate a range of spectral types and scaling
factors giving acceptable cancellation of the late-type features.
We use the spectral type and secondary flux to estimate the distance in the following manner.
We begin by finding the surface brightness of the secondary star in V , on the assumption that the
surface brightness is similar to that of main-sequence stars of the same spectral type; the Barnes-
Evans relation for late-type stars is discussed by Beuermann (2006). Combining the known Porb
with the assumption that the secondary fills its Roche critical lobe yields the secondary’s radius
R2 as a function of its mass, M2. In the relevant range of mass ratio, R2 ∝ M
2 , approximately,
and the dependence on M1 is weak enough to ignore. We generally do not know M2, so we guess
at a generous allowable range for this parameter using evolutionary simulations by Baraffe & Kolb
(2000) as a guideline; the weakness of the dependence of R2 on M2 means that this (rather ques-
tionable) step does not dominate the error budget. Combining the surface brightness with R2
yields the absolute magnitude MV . Subtracting this from the apparent magnitude measured for
the secondary star gives a distance modulus. The reddening maps of Schlegel et al. (1998) then
can be used to estimate the extinction. Note carefully that we do not assume that the secondary
is a ‘normal’ main-sequence star; we assume only that the secondary’s surface brightness is similar
to field stars of the same spectral type. The normalization of the secondary’s contribution also
depends on the assumption that the spectral features used to judge the subtraction are similar in
strength to those of a normal star.
As noted earlier, the immediate aim of our observations was to find orbital periods from radial
velocity time series data. The Hα emission line is usually the strongest feature, and it generally
gives good results in dwarf novae. All the emission-line velocities reported here are of Hα.
We measured radial velocities of Hα emission using convolution methods described by Schneider & Young
(1980) and Shafter (1983). In this technique one convolves an antisymmetric function with the line
profile, and takes the zero of the convolution (where the two sides of the line contribute equally)
as the line center. For the antisymmetric function with which the spectrum is convolved, we used
– 5 –
either the derivative of a Gaussian with adjustable width, or positive and negative Gaussians of
adjustable width offset from each other by an adjustable separation. Uncertainties in the convolu-
tion velocities are estimated by propagating forward the counting-statistics errors in the individual
data channels; in practice, these are lower limits to the true uncertainties, since the line profile
can vary in ways unrelated to the orbital modulation. The choice of convolution parameters is
dictated by the shape and width of the line, and in practice the parameters are adjusted to give
the best detection of the orbit. The physical interpretation of CV emission lines is complicated
and controversial (see, e.g., Shafter 1983, Marsh 1988, Robinson 1992), but in almost all cases the
emission-line periodicity accurately reflects Porb (though Araujo-Betancor et al. 2005 describe a
noteworthy exception to this rule). A sample of the radial velocities for each object are listed in
Table 4, while the full tables can be found online.
One of our systems, BF Eri, has a K-type absorption component in its spectrum. We measured
velocities of this using the cross-correlation radial velocity package described by Kurtz & Mink
(1998), using the region from 5000 to 6500 Å, and excluding the region containing the He I λ5876
emission line and and the NaD absorption complex. For a cross-correlation template spectrum, we
used the a velocity-compensated sum of many observations of IAU velocity standards taken with
the same instrument, as described in Thorstensen et al. (2004).
We searched for periods in all the velocity time series using the “residualgram” method
(Thorstensen et al. 1996); the resulting periodograms are given in Figs. 3 and 4. At the best
candidate periods we fitted least-squares sinusoids of the form v(t) = γ + K sin[2π(t − T0)/P ].
Fig. 5 shows the velocities folded on the best-fitting periods, and Table 5 gives the parameters of
these fits. Because of limitations of the sampling (e.g., the need to observe only at night from a
single site), a single periodicity generally manifests as a number of alias frequencies. To assess the
confidence with which we could assert that the strongest alias is the true period, we used a Monte
Carlo test described by Thorstensen & Freed (1985).
The alias problem can be particularly irksome over longer timescales; in this case the uncertain
number of cycles elapsed between observing runs causes fine-scale “ringing” in the periodogram.
The individual periods have tiny error bars, because of the large time span covered, but the am-
biguity in period means that a realistic error bar – one that covers the range of possibilities – is
much larger. In those cases, the period uncertainties given in Table 5 are estimated by analyzing
data from the individual observing runs separately. When only two observing runs are available,
the allowable fine-scale frequencies are well-described by a fitting formula
Porb = (t2 − t1)/n.
Here t1 and t2 are the epochs of blue-to-red velocity crossing observed on the two runs, and n
is the integer number of cycles that have passed between t1 and t2. The allowed range of n is
determined from the weighted average of the periods derived from separate fits to the two runs’
data. When more than two observing runs are available, the situation becomes more complex. In
some happy cases there are enough overlapping constraints that only a single, very precise period
– 6 –
remains tenable. We were able to find such precise periods for LX And, LU Cam, and BF Eri.
3. Notes on Individual Objects
We discuss the stars in alphabetical order by constellation.
3.1. LX Andromedae
LX And was first identified as a variable star (RR V-3) in the Lick RR Lyrae search (Kinman et al.
1982). It was classified incorrectly as an RV Tauri star, and its dwarf nova nature was unrecognized
until the photometric study by Uemura et al. (2000). Morales-Rueda & Marsh (2002) obtained
spectra of LX And as part of their study of dwarf novae in outburst and determined the equivalent
widths and FWHMs of the Balmer and He II lines. Our mean spectrum appears typical for a dwarf
nova at minimum light.
Because of the large hour-angle span, the radial velocity time series leaves no doubt about
the daily cycle count, which is near 6.6 cycle d−1. The several observing runs constrain the fine-
scale period in a more complicated way, but the Monte Carlo test indicates that a precise period
of 0.1509743(5) d is preferred with about 98 per cent confidence. Two other candidate periods
separated from this by 1 cycle per 53.2 d in frequency are much less likely.
3.2. CZ Aquilae
Very little has been published on CZ Aql, which is listed in the Archival Edition as a U-
Gem dwarf nova. Cieslinski et al. (1998) included the star in their spectroscopic study of irregular
variables, and noted a probable 4.8 hour period and emission lines typical of dwarf novae. Our
velocities confirm the suggested 4.8-hour period, but we cannot determine a unique cycle count
between our observing runs.
While the spectrum superficially resembles that of a dwarf nova, a closer look reveals interesting
behavior. Fig. 6, constructed using methods described by Taylor et al. (1999), presents our spectra
as a phase-averaged greyscale image. There is a striking broad component in the stronger Balmer
and HeI lines that shows a large velocity excursion, with the red wing of Hα reaching to +3100 km
s−1 at phase 0.3 (where phase 0 corresponds to the blue-to-red crossing of the line core). The broad
components around Hβ and λ6678 move in phase with those of Hα and range from 900 to 2600
and −2500 to −900 km s−1 and 700 to 2100 and −1000 to −600 km s−1, respectively. The wings
of λ5876 also move in phase with the others, but the red edge is difficult to follow at its minimum
because of interference from the NaD absorption lines, which are stationary and hence interstellar.
The maximum of the red edge is 3200 km s−1, while the blue edge ranges from −1500 to −600 km
– 7 –
s−1. The blueward wing of all these lines is noticeably weaker than the redward wing.
Other emission lines present include HeII λ4686, HeI λ4713 and λ4921, and, very weakly, FeII
λ5169. We also detect unidentified emission lines at λ6344, as is also seen in LS Peg (Taylor et al.
1999), and at λ5046. The strength of the λ5780 diffuse interstellar band (Jenniskens & Desert
1994) and the NaD lines suggest that a good deal of interstellar material lies along the line of sight,
and that the luminosity is relatively high.
High-velocity wings reminiscent of the ones seen here have been seen in V795 Her (Casares et al.
1996; Dickinson et al. 1997), LS Peg (Taylor et al. 1999), V533 Her (Thorstensen & Taylor 2000),
and RX J1643+34 (Patterson et al. 2002), all of which are SW Sex stars. We do not, how-
ever, detect another SW Sex characteristic, namely phase-dependent absorption in the HeI lines
(Thorstensen et al. 1991). The orbital periods of most SW Sex stars are shorter than 4 hours, so
CZ Aql’s 4.8-h period would be unusually long for an SW Sex star.
3.3. LU Camelopardalis
Jiang et al. (2000) obtained the first spectrum of this dwarf nova in a follow-up study of CV
candidates from the ROSAT All Sky Survey. We found no other published spectroscopic studies.
Our velocities constrain the period to a unique value, 0.1499685(7) d. The averaged spectrum
shows a rather strong, blue continuum, which may indicate a state somewhat above true minimum.
3.4. GZ Cancri
Jiang et al. (2000) confirmed the cataclysmic nature of GZ Cnc by obtaining the first spectrum
of the object. Kato et al. (2002) suggested that this star, originally labeled as a dwarf nova, could
possibly be an intermediate polar (DQ Her star), based on similarities in its long-term photometric
behavior to that of other intermediate polars. Tappert & Bianchini (2003) conducted a photometric
and spectroscopic study of the system. Using advance results from the present study to help decide
the daily cycle count, they found Porb = 0.08825(28) d, or 2.118(07) h, placing the system near
the lower edge of the so-called gap in the CV period distribution – a dearth of systems in the
period range from roughly 2 to 3 hr. Tappert & Bianchini (2003) also saw characteristics that
could indicate an intermediate polar classification, but did not claim their evidence was definitive
on this point.
Almost all our observations come from two observing runs a year apart. The full set of
velocities strongly indicates an orbital frequency near 11.4 cycle d−1, with the Monte Carlo test
giving a discriminatory power greater than 0.99 for the choice of daily cycle count. However, the
number of cycles between the two observing runs is not determined. Precise periods that fit the
combined data set are given by P = [349.785(3) d]/n, where n is the integer number of cycle counts;
– 8 –
n = 3972± 8 corresponds to roughly 1 standard deviation. While our period agrees well with that
of Tappert & Bianchini (2003), our data neither support nor disprove the claim that GZ Cnc may
be an intermediate polar.
3.5. V632 Cygni
Liu et al. (1999) offer the only published spectrum of this dwarf nova. They measured the
equivalent widths and integrated line fluxes of the Balmer, HeI, and HeII emission lines and sug-
gested that the orbital period is likely short based on the very strong Balmer emission. Our spec-
trum appears similar to theirs, and our measured flux level is also nearly the same. The periodigram
in Fig. 3 clearly favors an orbital frequency near 15.7 cycles d−1, with a discriminatory power of 95
per cent and a correctness likelihood near unity. This confirms the suggestion of Liu et al. (1999)
that the period is rather short and suggests that it is an SU UMa-type dwarf nova.
3.6. V1006 Cygni
Bruch & Schimpke (1992) present the only published spectrum we know of, and characterized
it as a “textbook example” of a dwarf nova spectrum. They noted a slightly blue continuum with
strong Balmer and He I emission, as well as clear He II λ4686 and Fe II emission. Our spectrum
(Fig. 1) is similar to theirs both in appearance and normalization, and our line measurements
(Table 3) are also comparable.
The periodogram (Fig. 4) indicates a frequency near 10.1 cycles d−1, and the Monte Carlo test
confirms that the daily cycle count is securely determined. Most of our data are from 2004 June,
but we returned in 2005 June/July to confirm the unusual period indicated in the earlier data.
The periods found by analyzing the two runs separately are consistent within their uncertainties.
As with GZ Cnc, there are multiple choices for the cycle count between the two observing runs;
the best-fitting periods are given by P = [369.006(4) d]/n, where n = 3726 ± 4 corresponds to
1 standard deviation. Including a few velocities from other observing runs suggests that n is
slightly larger, perhaps 3728. In any case, the period amounts to 2.38 h, which places V1006 Cyg
firmly in the period gap (Warner 1995), where there is apparently a true scarcity of dwarf novae
(Hellier & Naylor 1998).
3.7. BF Eri
The first evidence that BF Eridani was a cataclysmic variable came when an Einstein X-
ray source, 1ES0437-046, was matched to the variable (Elvis et al. 1992). Schachter et al. (1996)
confirmed this match and presented an optical spectrum. Kato (1999) and the Variable Star
– 9 –
Observers’ League in Japan (VSOLJ) found photometric variability characteristic of a dwarf nova.
The spectrum of BF Eri (Fig. 2) shows a significant contribution from a K star along with the
usual dwarf-nova emission lines. Normally, this suggests that Porb > 6 h. Nearly all our spectra
yielded good cross-correlation radial velocity measurements as well as emission-line velocities. The
absorption- and emission-line velocities independently give a period near 6.50 h (Table 5), in ac-
cordance with expectation based on the spectrum. There is no ambiguity in cycle count over the
5-year span of the observations, so the period is precise to a few parts per million. Fig. 6 shows a
phase-resolved average of the BF Eri spectra, with the absorption spectrum shifting in antiphase
to the emission lines.
If the emission-line velocities faithfully trace the primary’s center-of-mass motion, and the
absorption-line velocities also trace the secondary’s motion, then the two velocity curves should be
exactly one-half cycle out of phase. In BF Eri, we find a shift of 0.515 ± 0.007 cycles between the
two curves, consistent with 0.5 cycles, so we feel emboldened to explore the system dynamics.
Masses can only be derived when the orbital inclination is known, as in eclipsing systems. To
see if BF Eri might eclipse, we derived differential magnitudes from images that were taken for
astrometry (discussed below) and plotted them as a function of orbital phase. Some images were
taken at the phase at which an eclipse would appear, but no evidence for an eclipse was found.
Limits on the depth and duration of the eclipse are difficult to quantify because the data were taken
in short bursts in the presence of strong intrinsic variability, so a weak eclipse cannot be ruled out,
but the photometry does suggest that the inclination is not close to edge-on.
Because the system apparently does not eclipse, we cannot derive masses; rather, we find broad
constraints on the inclination by assuming astrophysically reasonable masses for the components.
Taken at face value, the velocity amplitudes K imply a mass ratio q = M2/M1 = 0.60 ± 0.03. If
we arbitrarily choose a white dwarf mass M1 = 0.9 M⊙ (so that M2 = 0.53 M⊙), the observed K
velocities imply i = 50 degrees. To find a rough lower limit on the inclination, we consider a massive
white dwarf (M1 = 1.2 M⊙) and, ignoring the constraint on q for the moment, take M2 = 0.4 M⊙;
this yields i = 40 degrees. For a rough upper limit, we assume M1 = 0.6 M⊙ and M2 = 0.4 M⊙,
which gives i = 67 degrees.
The decomposition procedure described earlier yielded a spectral type of K3 ±1 subclass; the
result of the subtraction is shown in Fig. 2. Using the V passband tabulated by Bessell (1990) and
the IRAF sbands task, we find a synthetic V = 16.9± 0.3 for the K star’s contribution. Taking the
range of plausible secondary star masses to be 0.4 to 0.8 M⊙ yields R2 = 0.7± 0.1 R⊙ at this Porb.
Combining this with the surface brightness expected at this spectral type yields MV = 6.8 ± 0.4
for the secondary. If there is no significant interstellar extinction, we have m − M = 10.1 ± 0.5,
or a distance of approximately 1100 ± 300 pc. The dust maps of Schlegel et al. (1998) give a total
E(B − V ) = 0.062 in this direction. Assuming that BF Eri is beyond the Galactic dust and taking
AV /E(B − V ) = 3.3 gives an extinction-corrected (m−M)0 = 9.9, and a distance estimate of 950
(+250,−200) pc.
– 10 –
We can also estimate a distance using the relation found by Warner (1987) between Porb, i,
and the absolute magnitude at maximum light MV (max). Using our inclination constraints, the
Warner relation predicts MV max = 3.9±0.7 at this orbital period. The General Catalog of Variable
Stars (Kholopov et al. 1999) lists mp = 13.2 at maximum light; taking this to be similar to Vmax
yields m−M = 9.3, or 9.1 corrected for extinction, which corresponds to 660 pc.
Given these distance estimates, it is surprising that BF Eri has a very substantial proper
motion. The Lick proper motion survey (Hanson et al. 2004) gives [µX , µY ] = [+34,−97] mas
yr−1. We have begun a series of parallax observations with the Hiltner 2.4m telescope using the
protocols described by Thorstensen (2003); so far we have five epochs from 2005 November and 2007
January. The proper motion relative to the background stars is [µX , µY ] = [32,−111] mas yr
and the parallax is not detected, with a nominal value of 1 ± 2 mas. The parallax determination
is very preliminary, but given the data so far we estimate the lower limit on the distance based on
the astrometry alone to be ∼ 200 pc.
At the nominal 950 pc distance derived from the secondary star, a 100 mas yr−1 proper
motion corresponds to a transverse velocity vT = 451 km s
−1. This is implausibly large, so we
are left wondering how we might have overestimated the distance. One effect might be as follows.
Our distance is based on the secondary’s apparent brightness, and we estimate the secondary’s
contribution to the total light by searching for the best cancellation of its features. If the secondary’s
absorption lines are weaker than those in the spectral-type standards, we would underestimate the
secondary’s contribution. In our best decomposition, the secondary is about 2.2 magnitudes fainter
than the total light in V . Assuming (unrealistically) that all the light is from the secondary would
therefore decrease the distance modulus by 2.2 magnitudes, to a distance of 340 pc.
We do not yet have enough information to resolve the conundrum posed by BF Eri’s unmis-
takably large proper motion and its apparently large distance, but a reasonable compromise might
be to put it at something like 400-500 pc, with an underluminous, low-metallicity secondary. The
cross-correlation velocities of the secondary have a zero point determined to ±5 km s−1, more or
less, and give a substantial systemic velocity of −72±3 km s−1, or −86 km s−1 in the local standard
of rest. If the star is at 450 pc, its space velocity with respect to the local standard of rest is ∼250
km s−1, with Galactic components [U, V,W ] = [−180,−180,−3] km s−1, that is, the velocity is
mostly parallel to the Galactic plane and lags far behind the rotation of the Galactic disk. This
would put BF Eri on a highly eccentric orbit; these are halo-population kinematics (even though
the star remains close to the plane). These kinematics would be qualitatively consistent with the
weak-line conjecture used earlier.
3.8. BI Ori
Szkody (1987) published the first quiescent spectrum of BI Orionis, which showed emission
lines typical of dwarf novae. Morales-Rueda & Marsh (2002) show a spectrum in outburst and
– 11 –
note the possible presence of weak HeII λ4686 emission.
Only the 2006 January velocities are extensive enough for period finding; they give Porb =
4.6 hr, with no significant ambiguity in the daily cycle count. The average spectrum shows the
usual emission lines; M-dwarf absorption features are also visible, though the signal-to-noise of the
individual spectra was not adequate for finding absorption-line velocities. Using the procedures
described earlier, we estimate the secondary’s spectral type to be M2.5 ± 1.5, with the secondary
alone having V = 20.0 ± 0.4. Assuming that the secondary’s mass lies in the broad range from
0.2 to 0.6 M⊙, its radius at this Porb would be 0.35 to 0.6 R⊙. Combining this with the surface
brightness derived from the spectral types gives an absolute magnitude MV = 10.2 ± 1.0. The
distance modulus, uncorrected for extinction, is therefore m − M = +9.8 ± 1.1, corresponding
to 910(+600,−360) pc. Schlegel et al. (1998) estimate a total reddening E(B − V ) = 0.11 in this
direction; assuming that BI Ori lies beyond all the dust, and taking AV /E(B−V ) = 3.3 reduces the
distance to ∼ 770 pc. At maximum light, BI Ori has mp = 13.2 (Kholopov et al. 1999). Assuming
the color is neutral, we find MV = 3.4 ± 1.1 at maximum. At BI Ori’s period, the Warner (1987)
relation predicts MV > 3.6 (with the brightest value corresponding to i = 0). This agrees broadly
with our nominal value based on the secondary star’s distance, but is a little fainter, suggesting
that BI Ori is not too far from face-on, or a little closer than our nominal distance, or both.
3.9. FO Per
FO Persei was apparently discovered by Morgenroth (1939), but its cataclysmic nature was
not immediately recognized. Bruch (1989) obtained spectra and gave equivalent widths for the
Balmer lines for two different nights of observations, between which the continuum changed from
relatively flat to inclined toward the red.
The emission lines in FO Per are rather narrow (Fig. 1, Table 3). This is often taken to indicate
a low orbital inclination. The velocity amplitude K is small, so that K/σ ≈ 1.6 for the best fits
(Table 5). Because of this, the daily cycle count remains ambiguous; the orbital frequency is either
5.8 or 6.8 cycle d−1, corresponding to Porb of 3.52 or 4.13 hr. CVs with periods in the 3-4 hour
range tend to be novalike variables (Shafter 1992), whereas FO Per is a dwarf nova; thus the 4.13
hr period is more likely a priori.
4. Summary
We have determined the orbital periods of eight CVs without significant daily cycle count
ambiguity; for FO Per, the period is narrowed to two choices. For three of the systems we find
high-precision periods by establishing secure cycle counts over long baselines.
While most of these objects are similar to others already known, three stand out as especially
– 12 –
interesting. CZ Aql shows asymmetric, high-velocity wings around the Balmer and HeI λ5876 and
λ6678 lines, possibly indicating a magnetic system. BF Eri’s proper motion of ∼ 100 mas yr−1 is
surprising in view of the large distance indicated by its secondary spectrum and by the Warner
relation; even if it is somewhat nearer than these indicators suggest, its kinematics are not typical of
disk stars. Finally, the orbital period of V1006 Cyg places it squarely in the middle of the so-called
period gap between 2 and 3 hours.
Acknowledgments. We are most grateful for support from the National Science Foundation
through grants AST-9987334 and AST-0307413. Bill Fenton took most of the spectra of GZ Cnc,
and J. Cameron Brueckner assisted with the BF Eri spectroscopy. Some of the astrometric images of
BF Eri were obtained by Sébastien Lépine and Michael Shara of the American Museum of Natural
History. We would like to thank the MDM Observatory staff for their skillful and conscientious
support. Finally, we are grateful to the Tohono O’odham for leasing us their mountain for a while,
so that we may study the glorious universe in which we all live.
– 13 –
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This preprint was prepared with the AAS LATEX macros v5.2.
– 16 –
Fig. 1.— Plots of the average flux-calibrated spectra for eight of the stars studied here. The weak
features seen near λ5577 are artifacts caused by imperfect subtraction of the strong [OI] night-sky
emission.
– 17 –
Fig. 2.— Plot of the averaged spectrum of BF Eri (top) and the spectrum after scaled late-type
(K3V) star has been subtracted (bottom). The spectra have been shifted into a rest frame before
averaging and do not include the 2006 March or 2007 January data.
– 18 –
Fig. 3.— Periodigrams for most of the stars studied here. The vertical axis in each case is the
inverse of chi-square for the least-squares best fitting sinusoid at each trial frequency. When data
from more that one observing run are combined, the periodigram can require hundreds of thousands
of points to resolve the fine-scale ringing; in those cases, the curve shown is formed by connecting
local maxima of the periodogram with straight lines. In those cases the right-hand panel gives a
close-up view of the region around the highest peak, revealing the alias structure resulting from
different choices of cycle count between the observing runs. The periodogram of BF Eri is for the
absorption-line velocities.
– 19 –
Fig. 4.— Periodigrams for the remainder of the stars, plotted in the same manner as the previous
figure. Because the choice of daily cycle count for FO Per remains ambiguous, we have not chosen
to enlarge either peak region.
– 20 –
Fig. 5.— Radial velocities plotted as a function of phase using the adopted orbital periods. For
CZ Aql, GZ Cnc, and V1006 Cyg, the number of cycle counts between observing runs is unknown,
and the exact period chosen to fold the velocities is one of a number of possibilities. The two
plots for FO Per are for different choices of the daily cycle count, and each of these in turn is also
an arbitrary choice among many finely-spaced periods. In BF Eri, both emission and absorption
velocities are plotted; the absorption velocities are shown with error bars.
– 21 –
Fig. 6.— Phase-averaged spectra of CZ Aql (top two panels) and BF Eri (bottom two panels),
presented as a greyscale. The scale is inverted, so that emission is represented by darker shades.
The two CZ Aql spectra are scaled differently to show the line cores (top) and the extent of the the
line wings. Note the NaD lines in CZ Aql remain stationary, indicating an interstellar origin. The
feature at λ6280 is telluric. BF Eri’s spectrum is plotted in two overlapping sections; the K-star’s
orbital motion is plainly visible.
– 22 –
Table 1. Stars Observed
Star α2000
a δ2000 Epoch
b Vobs
c maxd min
[hh:mm:ss] [◦:′:”] [mag.] [mag.] [mag.]
LX And 2:19:44.08 +40:27:22.3 2006.7 16.3 13.5p 16.4p
CZ Aql 19:19:58.21 −07:10:55.2 2003.4 15.4 13.p 15.p
LU Cam 5:58:17.86 +67:53:46.2 2002.0 16.3 14.v (16.v
GZ Cnc 9:15:51.68 +09:00:49.6 2000.3 15.4 13.1v 15.4v
V632 Cyg 21:36:04.22 +40:26:19.4 2000.5 17.9 12.6p 17.5p
V1006 Cyg 19:48:47.20 +57:09:22.8 2000.5 17.8 15.4p 17.0p
BF Eri 4:39:29.96 −04:35:59.5 2006.2 14.8 13.5p 15.5p
BI Ori 5:23:51.77 +01:00:30.6 2002.8 17.1 13.2p 16.7p
FO Per 4:08:34.98 +51:14:48.5 2004.0 17.1 11.8v 16.v
aPositions measured from images taken at the 2.4m Hiltner telescope, using
astrometric solutions from fits to USNO A2.0 (Monet et al. 1996) or UCAC 2
(Zacharias et al. 2004) stars. Uncertainties are of order 0.1 arcsec.
bThe date of the image used in the position measurement. The coordinate
system (equator and equinox) is J2000 in all cases.
cSynthesized from spectra, as described in text.
dTaken from the GCVS (Kholopov et al. 1999). Photgraphic magnitudes
flagged with ‘p’, visual with ‘v’.
– 23 –
Table 2. Journal of Observations
data N HA start HA end telescope
(UT) [hh:mm] [hh:mm]
LX And
2004 Jan 13 1 +2:35 +2:35 2.4m
2004 Mar 02 3 +3:45 +4:00 2.4m
2004 Nov 18 1 +0:29 +0:29 1.3m
2004 Nov 19 26 −3:36 +4:17 1.3m
2004 Nov 19 2 +2:31 +2:35 2.4m
2004 Nov 20 13 −3:17 +3:15 1.3m
2004 Nov 20 1 −1:23 −1:23 2.4m
2006 Jan 19 14 +1:31 +4:10 1.3m
2006 Jan 22 8 +3:31 +4:46 1.3m
2007 Jan 26 3 +1:36 +1:58 1.3m
2007 Jan 27 15 +0:58 +3:39 1.3m
CZ Aql
2005 Jul 02 3 +1:40 +2:04 1.3m
2005 Jul 04 48 −3:28 +3:17 1.3m
2005 Jul 05 13 −1:55 −0:06 1.3m
2005 Jul 06 12 −3:57 +2:28 1.3m
2005 Sep 03 2 +1:23 +1:32 1.3m
2005 Sep 07 2 −0:03 +0:05 1.3m
2005 Jun 28 2 −0:01 +0:04 2.4m
2006 Jun 18 2 +2:05 +2:10 2.4m
2006 Jun 19 2 +0:36 +0:40 2.4m
2006 Jun 23 3 −1:24 −1:12 2.4m
– 24 –
Table 2—Continued
data N HA start HA end telescope
(UT) [hh:mm] [hh:mm]
LU Cam
2002 Jan 22 2 −1:40 −1:19 2.4m
2002 Jan 23 8 −1:33 +2:31 2.4m
2002 Jan 24 25 −3:15 +5:47 2.4m
2002 Feb 18 2 +2:20 +2:28 2.4m
2002 Feb 19 2 +2:41 +2:50 2.4m
2002 Feb 20 4 −0:02 +3:43 2.4m
2002 Feb 22 1 +2:18 +2:18 2.4m
2004 Jan 16 2 +2:49 +2:53 2.4m
2004 Jan 17 1 +0:40 +0:40 2.4m
2004 Jan 19 1 −0:32 −0:32 2.4m
2004 Mar 02 1 +0:53 +0:53 2.4m
2004 Mar 07 6 +2:29 +3:14 2.4m
2004 Nov 19 4 +2:48 +3:27 2.4m
2005 Mar 21 1 +1:12 +1:12 2.4m
2005 Mar 22 2 +1:21 +1:30 2.4m
2005 Sep 09 1 −2:00 −2:00 2.4m
2005 Sep 12 1 −1:49 −1:49 2.4m
2006 Jan 09 2 +1:49 +1:55 2.4m
GZ Cnc
2000 Apr 07 2 +4:26 +4:32 2.4m
2000 Apr 08 1 +0:27 +0:27 2.4m
2000 Apr 10 5 −0:31 +4:32 2.4m
2000 Apr 11 15 +2:19 +4:22 2.4m
2001 Mar 24 2 +4:39 +4:49 2.4m
2001 Mar 25 21 −1:10 +3:10 2.4m
2001 Mar 26 20 −0:50 +4:27 2.4m
2001 Mar 27 2 −0:08 +0:01 2.4m
2001 Mar 28 3 +0:14 +0:25 2.4m
V632 Cyg
2005 Jul 07 2 +1:00 +1:16 1.3m
2005 Jul 08 10 −5:12 +0:33 1.3m
2005 Jul 09 18 −5:03 +1:09 1.3m
2005 Jul 10 18 −5:09 +1:07 1.3m
2005 Jul 11 3 +0:52 +1:19 1.3m
– 25 –
Table 2—Continued
data N HA start HA end telescope
(UT) [hh:mm] [hh:mm]
V1006 Cyg
2003 Jun 22 1 +1:06 +1:06 2.4m
2004 Jun 24 5 +0:52 +1:56 1.3m
2004 Jun 25 5 −1:43 +0:34 1.3m
2004 Jun 25 1 +4:06 +4:06 2.4m
2004 Jun 26 10 −4:25 +1:56 1.3m
2004 Jun 27 3 −3:00 −2:00 1.3m
2004 Jun 28 4 −2:25 +1:02 1.3m
2004 Jun 28 1 +0:28 +0:28 2.4m
2004 Jun 29 5 +0:55 +1:59 1.3m
2004 Jun 29 1 +4:07 +4:07 2.4m
2004 Jun 30 18 −4:39 +2:26 1.3m
2004 Jul 01 10 −3:58 +1:53 2.4m
2005 Jul 05 13 +0:58 +3:11 1.3m
2004 Jun 30 4 −3:41 +3:52 2.4m
2004 Jul 01 12 −4:07 +2:17 1.3m
2005 Jul 03 27 −3:57 +1:20 1.3m
2005 Jul 05 13 +0:58 +3:11 1.3m
BF Eri
2001 Dec 18 3 +2:35 +2:56 1.3m
2001 Dec 19 10 −3:06 +4:04 1.3m
2001 Dec 20 12 −3:49 +2:00 1.3m
2001 Dec 21 2 +3:04 +3:14 1.3m
2001 Dec 22 5 −3:13 −2:31 1.3m
2001 Dec 23 12 −3:18 +4:35 1.3m
2001 Dec 24 13 −3:17 +3:08 1.3m
2001 Dec 25 14 −2:20 +1:07 1.3m
2001 Dec 26 8 −2:10 +3:05 1.3m
2001 Dec 27 18 −2:39 +4:14 1.3m
2002 Jan 19 1 +1:27 +1:27 2.4m
2002 Jan 20 2 −1:33 +2:13 2.4m
2002 Jan 22 1 −2:02 −2:02 2.4m
2002 Feb 21 2 +1:26 +1:35 2.4m
2002 Feb 22 2 +0:40 +0:49 2.4m
2002 Oct 26 2 −0:11 +0:05 2.4m
2002 Oct 31 1 +3:21 +3:21 2.4m
2003 Feb 02 1 +0:04 +0:04 2.4m
2005 Sep 11 2 −0:57 −0:40 1.3m
2006 Mar 16 1 +1:56 +1:56 1.3m
2006 Mar 17 5 +2:15 +2:57 1.3m
– 26 –
Table 2—Continued
data N HA start HA end telescope
(UT) [hh:mm] [hh:mm]
2007 Jan 28 9 −1:46 −0:13 1.3m
BI Ori
2006 Jan 20 32 −2:54 +4:05 1.3m
2006 Jan 21 22 −2:47 +2:21 1.3m
2006 Jan 23 6 +1:03 +2:08 1.3m
FO Per
1995 Oct 09 9 −4:40 −3:26 2.4m
1995 Oct 10 5 −5:22 +1:26 2.4m
1996 Dec 19 14 +1:38 +4:03 1.3m
1996 Dec 20 5 +2:08 +3:43 1.3m
2001 Dec 18 9 −2:35 +4:36 1.3m
2004 Nov 18 1 −0:11 −0:11 1.3m
2004 Nov 19 9 +2:52 +5:08 1.3m
2004 Nov 19 2 +3:24 +2:48 2.4m
2004 Nov 20 20 −0:56 +5:10 1.3m
2006 Jan 10 12 −0:35 +2:35 1.3m
2006 Jan 10 1 +4:21 +4:21 2.4m
2006 Jan 11 18 −1:29 +2:41 1.3m
2006 Jan 11 1 −3:12 −3:12 2.4m
2006 Jan 12 4 −1:16 −0:36 1.3m
2006 Jan 13 10 +3:44 +5:37 1.3m
2006 Jan 16 11 −1:36 +1:00 1.3m
– 27 –
Table 3. Spectral Features in Quiescence
E.W.a Flux FWHM b
Feature (Å) (10−16 erg cm−2 s1) (Å)
LX And
Hβ 45 690 18
HeI λ4921 4 60 25
HeI λ5015 3 50 20
Fe λ5169 2 20 14
HeI λ5876 11 120 19
Hα 54 560 17
HeI λ6678 4 40 19
CZ Aql
Hβ 21 670 18
HeI λ4921 1 40 12
HeI λ5015 2 50 13
HeI λ5876 7 160 15
NaD −1 −16 · · ·
Hα 61 1250 27
HeI λ6678 5 90 16
LU Cam
Hγ 10 170 12
HeI λ4471 2 30 10
Hβ 14 190 12
HeI λ4921 1 20 12
HeI λ5015 2 20 13
Fe λ5169 1 10 12
HeI λ5876 4 50 12
Hα 27 240 13
HeI λ6678 3 20 15
HeI λ7067 3 20 · · ·
GZ Cnc
Hγ 26 940 26
HeI λ4471 8 260 28
HeII λ4686 5 140 46
Hβ 36 1040 25
HeI λ4921 5 140 26
HeI λ5015 4 100 28
Fe λ5169 2 60 26
HeI λ5876 9 200 27
– 28 –
Table 3—Continued
E.W.a Flux FWHM b
Feature (Å) (10−16 erg cm−2 s1) (Å)
Hα 38 790 25
HeI λ6678 4 90 31
HeI λ7067 3 60 32
V632 Cyg
Hβ 80 260 24
HeI λ4921 6 20 27
HeI λ5015 8 20 26
Fe λ5169 5 10 26
HeI λ5876 28 70 27
Hα 113 260 27
HeI λ6678 15 30 32
V1006 Cyg
Hβ 74 250 27
HeI λ4921 8 30 30
HeI λ5015 8 20 28
Fe λ5169 8 30 28
HeI λ5876 26 70 34
Hα 108 250 31
HeI λ6678 11 30 38
BF Eri
HeII λ4686 6 240 46
Hβ 23 1060 22
HeI λ5015 2 110 29
HeI λ5876 5 240 21
Hα 27 1200 22
HeI λ6678 3 120 27
BI Ori
Hβ 34 180 32
HeI λ4921 6 30 36
HeI λ5015 6 30 43
Fe λ5169 5 30 39
HeI λ5876 6 30 31
Hα 36 190 31
FO Per
– 29 –
Table 3—Continued
E.W.a Flux FWHM b
Feature (Å) (10−16 erg cm−2 s1) (Å)
Hβ 24 160 9
HeI λ4921 2 20 7
HeI λ5015 3 20 7
Fe λ5169 2 10 7
HeI λ5876 5 30 7
Hα 29 170 9
HeI λ6678 2 10 9
aEmission equivalent widths are counted as positive.
bFrom Gaussian fits.
– 30 –
Table 4. Radial Velocities
Star time a vabs σvabs vemn σvemn
(km s−1) (km s−1) (km s−1) (km s−1)
LX And 53017.7061 · · · · · · −44 −11
LX And 53066.6160 · · · · · · −57 −8
LX And 53066.6208 · · · · · · −37 −8
LX And 53066.6263 · · · · · · −77 −7
aHeliocentric Julian date of mid-integration, minus 2400000.
Note. — All emission-line velocities are of Hα. Emission-line velocity
uncertainties are derived from counting statistics and should be regarded as
lower limits. Table 4 is published in its entirety in the electronic version of
the Publications of the Astronomical Society of the Pacific. A short portion
is shown here for guidance regarding its form and content.
– 31 –
Table 5. Fits to Radial Velocities
Star Algorithma T0
b P K γ N σc
(d) (km s−1) (km s−1) (km s−1)
LX And G2,21,7 53754.6861(12) 0.1509743(5) 81(4) −48(3) 87 15
CZ Aql G2,18,8 53557.880(2) 0.2005(6)d 193(15) 5(10) 89 61
LU Cam D,15 52327.7421(14) 0.1499686(4) 57(4) 44(3) 66 14
GZ Cnc G2,15,9 51992.8928(13) 0.0881(4)d 79(7) 22(5) 71 26
V632 Cyg D,28 53560.9746(13) 0.06377(8) 62(8) −49(5) 51 28
V1006 Cyg G2,20,9 53187.9091(16) 0.09904(9)d 89(8) −11(6) 120 44
BF Eri emission G2,21,7 52574.0027(18) 0.2708801(6) 109(5) −91(3) 126 24
BF Eri absorption · · · 52573.8632(9) 0.2708805(4) 182(4) −72(3) 117 20
BF Eri mean: · · · · · · 0.2708804(4) · · · · · · · · ·
BI Ori 53756.541(3) G2,35,9 0.1915(5) 131(13) 24(9) 60 44
FO Per (shorter) D,11 52261.872(3) 0.1467(4)d 27(3) −49(2) 131 17
FO Per (longer) D,11 52261.893(3) 0.1719(5)d 27(3) −45(2) 131 17
Note. — Parameters of sinusoidal least-squares fits to the velocity timeseries, of the form v(t) = γ +K sin(2π(t −
T0)/P ). The quoted parameter uncertainties are based on the assumption that the scatter of the data around the best
fit is a realistic estimate of the velocity uncertainty (Cash 1979). In practice this is more conservative than assuming
that counting statistics uncertainties are realistic.
aCode for the convolution function used to derive emission line velocities; D = derivative of a Gaussian, G2 =
double-Gaussian function (see text). For the D algorithm the number that follows gives the line full-width at half-
maximum, in Å, for which the function is optimized; for the G2 algorithm the two numbers are respectively the
separation of the two Gaussians and their individual FWHMa, again in Å.
bHeliocentric Julian Date minus 2400000. The epoch is chosen to be near the center of the time interval covered
by the data, and within one cycle of an actual observation.
cRoot-mean-square residual of the fit.
dThe period determination in this case is complicated by unknown numbers of cycles between observing runs; the
uncertainty given here is an estimate based on fits to individual runs. Only certain values within the period range
given here are allowed; see text for details.
This figure "f6.png" is available in "png"
 format from:
http://arxiv.org/ps/0704.0948v1
http://arxiv.org/ps/0704.0948v1
	Introduction
	Observations, Reductions, and Analysis
	Observations
	Reductions
	Analysis
	Notes on Individual Objects
	LX Andromedae
	CZ Aquilae
	LU Camelopardalis
	GZ Cancri
	V632 Cygni
	V1006 Cygni
	BF Eri
	BI Ori
	FO Per
	Summary
