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1996MNRAS.282.1211S
Mon. Not. R. Astron. Soc. 282, 1211-1222 (1996)
Infrared colours, distance determination and absolute magnitudes of a
sample of faint cataclysmic variables
L. N. Sproats/ S. B. Howe1l2 * and K. O. Mason!
1Mullard Space Science Laboratory, Holmbury St Mary, Dorking, SU"ey RH5 6NT
2Planetary Science Institute, Astrophysics Group, 620 N. 6th Avenue, Tucson, AZ 85705, USA
Accepted 1996 June 3. Received 1996 May 8; in original form 1995 October 26
ABSTRACT
Key words: novae, cataclysmic variables: - infrared: stars.
1 INTRODUCTION
Cataclysmic variables (CVs) are semidetached binary systems in which a white dwarf primary accretes material from
a late-type secondary star via Roche Lobe overflow. An
accretion disc is expected to form around the primary in
non-magnetic systems. CVs have orbital periods which typically range from ~ 80 min to ~ 10 h (see Warner 1995 for a
recent review).
Our current knowledge of the basic characteristics of
CVs, such as their distance and luminosity, is based largely
on observations of optically bright (V < 16) systems. This
favours nearby and/or luminous members of the class. In an
attempt to redress this bias Howell & Szkody (1990; hereafter HS90) compiled a list of faint (V> 16) CVs located at
high galactic latitude, Ib I > 40°. This list contained 57 stars
categorized as dwarf novae (ONe), 12 nova-like systems
(NLs) , seven AM Her systems and eight classical novae.
The distances of some of these systems implied by adopting
the then best estimates of the mean absolute magnitude of
various CV subclasses (e.g. Warner 1987, hereafter W87)
ranged up to several kpc, placing them in the halo of the
Galaxy.
The HS90 survey also revealed that a number of the faint
*Present address: Department of Physics and Astronomy, University of Wyoming, P.O. Box 3905, University Station, Laramie, WY
82071, USA.
short-period ONe had very large outburst amplitudes, ranging up to ~ 10 mag. HS90 suggested that if these ONe were
at large distances in the Galactic halo, the large outburst
amplitudes could be the result of low metallicity (see Cannizzo, Shafter & Wheeler 1988). On the other hand, if they
are systems in the local neighbourhood Howell, Szkody &
Cannizzo (1995; hereafter HSC95) argue that their large
outburst amplitudes might be due to the viscosity in the
accretion disc being very low during quiescence, so that
virtually all the accretion through the disc occurs during
outburst. This would also imply that the quiescent disc luminosity is low. These large outburst amplitude objects may be
extreme SU UMa stars and have been variously referred to
as TOADs (Tremendous Outburst Amplitude Dwarf
Novae: HSC95) or WZ Sge stars (O'Oonoghue et al. 1991,
and references therein).
In this paper we determine the absolute magnitudes and
distances of 36 faint CVs based on infrared photometry, the
majority taken from the HS90 list of high-latitude CVs. The
aim was to establish whether these systems are truly at large
distances or are instead a population of intrinsically faint
objects that are relatively nearby. We are thus also testing
current ideas about the absolute magnitudes of CVs in quiescence. The stars studied include a number of systems with
large optical outburst amplitudes.
In Section 2 we outline the observations and data analysis
and in Section 3 we describe the methods used in this analysis to determine the distances and M v of our sample. In
© 1996 RAS
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
We have measured J and K colours for a sample of 36 predominantly faint
cataclysmic variables (CVs) at high Galactic latitude. The J-K colours lie blueward
of the main sequence, consistent with moderate (",20 per cent) contamination of
the secondary light by an accretion disc. We use the K magnitudes to estimate the
distances of the stars. We show that most lie within about 400 pc, and only two of the
stars in our sample (RU LMi and DO Leo) are convincingly at Galactic halo
distances. We have determined Mv values based on quiescent V magnitudes. The
long-period dwarf nova systems (P > 3 h) have absolute magnitudes which are
consistent with the previously observed range. However, for the short-period dwarf
novae the range of M v extends to much fainter magnitudes than previously believed
(as faint as Mv= 14). We find that those systems that have large outburst amplitudes
tend to be those with the faintest Mvs in quiescence.
1996MNRAS.282.1211S
1212 L. N Sproats, S. B. Howell and K. 0. Mason
Section 4 we present our results, which are discussed in
Section 5. Our conclusions are given in Section 6.
2 INSTRUMENTATION AND
OBSERVATIONS
J- (1.25 /lm, AA=O.3 Jlffi) andK- (2.2 /lm, AA=O.3 /lm) band
2.1 Data reduction and analysis
As a first step in the reductions, all the data were linearized
to remove the inherent non-linearity present in the
detector. This is routinely carried out when the IR data are
transferred from the telescope to the Joint Astronomy
Centre in Hilo (JACH).
All subsequent data reduction was performed using the
IRCAM software package (Aspin 1991) at JACH. The following procedure was adopted for each object in each band.
Dark frames, obtained nearest in time to the observation,
were subtracted from the five individual frames in the offset-pointing pattern. These five frames were then median
filtered to produce a flat-field frame which was divided into
the individual images from which it was created. This procedure was adopted, rather than obtaining a master flatfield image for the whole of the night, because the sky
spectrum is known to vary on short time-scales (~15 min)
because of varying atmospheric OH emission. The standard
3 DISTANCE AND Mv DETERMINATION
The distance to each CV in our sample was calculated using
Bailey's method (Bailey 1982). Thus
K
SK
logd=- +1- -
5
Rz
+ log5
Ro'
(1)
© 1996 RAS, MNRAS 282,1211-1222
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
images of the halo CV candidates were obtained on 1993
February 16-18 (UT) and September 8-18 (UT) using the
1-5 Jlffi cooled infrared camera, IRCAM2, at the cassegrain
focus of the 3.8-m United Kingdom Infrared Telescope
(UKlRT) at Mauna Kea, Hawaii. Additional observations
were kindly obtained by Colin Aspin (JACH) during a service observing night on UT 1993 September 18. IRCAM2
contains a Santa Barbara Instruments Research Centre
62 x 58 pixel Indium Antimonide (InSb) 2D array which is
operated at a temperature of 35 K (McClean et al. 1986;
McClean 1987). All of our observations were made using an
image scale of 0.62 arcsec pixel-I.
For each of our programme objects we obtained a series
of five individual images in each of the J and K bands,
offsetting the telescope between pointings by 5 arcsec in a
quincunx (to facilitate sky subtraction - see Section 2.1).
The individual images were made up of 200-s (100 s x 2
coadds) and 180-s (30 s x 6 coadds) exposures in J and K
respectively. Dark frames were obtained at regular intervals
throughout the night. Since the dark current is non-linear
with exposure time, the dark-frame exposures were of the
same length as the individual observations.
Observations of commonly used bright (J ~ 6-9, K ~ 6-8)
UKlRT standard stars (obtained from the UKlRT standard
list available at the telescope) were made on each night for
photometric calibration. Exposures for all standard stars
were 1.45 s (0.145 s x 100 coadds) in both the J and K
bands.
TV Crv and A Y Lyr were found to be in outburst when
viewed on the television acquisition screen (in 1993 February and September respectively), and no images were
taken during our scheduled run. Images of A Y Lyr were
subsequently obtained during the service run on UT September 18. However, the star may have been on the tail of a
subsequent outburst which peaked on September 8/9 UT.
star images (2 frames per band) were flat-fielded in a similar
manner, but using a flat-field image from one of the programme objects taken nearby in time. Nightly photometric
zero points (the magnitude of an object which would correspond to 1 DN per second) for the J and K bands were
obtained by performing aperture photometry on the
standard star observations. Our zero points agree with typical UKlRT values to within ±0.1 mag on all nights.
Simple two-dimensional aperture photometry was used
to determine the magnitudes of our programme objects,
most of which were situated in relatively uncrowded fields.
This was done for each of the 5 frames per object allowing
us to determine whether or not any significant intrinsic
variations were present during the total of ~ 1000 s coverage per star. For RZ Leo and GO Com, we obtained two
sets of 5 frames each owing to their faintness.
The photometric precision of our program star measurements was estimated by measuring field stars contained
within a number of programme frames and assumed to be
constant in intrinsic flux. The magnitudes of these 'comparison' stars ranged from J = 12 to 16, similar to our program
stars. The upper limits to variations measured in these constant field stars were 0.05 mag (J, K ~ 12) to 0.09 mag (J,
K ~ 16). In no case, except HU Aqr, an eclipsing AM Her
star, did any of our programme stars show variations during
observation exceeding these 1u limits. The photometric precisions determined for the constant field stars are in good
agreement with those determined by Aspin, Sandell & Russell (1994).
The J and K magnitudes determined for our programme
stars are listed in Table 1, along with other relevant information from the literature.
All data were corrected for airmass using standard stars
observed at a variety of zenith distances on each night. All
our targets were observed at airmass of less than 1.5, and the
corrections are small « 0.1 mag airmass-I at J and K:
Krisciunas et al. 1987). The scatter in the zero points
derived from the standard stars (0.07 mag in J, 0.04 mag in
K) is included in the error estimates for the J and K magnitudes of our programme objects. No corrections were made
for interstellar reddening since at the galactic latitudes of
the bulk of our sample the reddening is small
(E(B - V) < 0.1: Zombeck 1990).
The V magnitudes listed in Table 1 are the best estimates
of the quiescent magnitude, compiled from the literature.
Most, in fact, are based on precise CCD measurements.
While we are aware that there is variability during quiescence of perhaps up to 1 mag (Howell et al. 1990), the
quoted V magnitudes will provide at least an upper limit to
the minimum quiescent absolute magnitude. We note that
V -J colours of main-sequence stars (Bessell & Brett 1988)
range from ~ 1 for mid G stars to ~ 6 for late M stars,
consistent with the spread in values that we see in our
programme stars.
1996MNRAS.282.1211S
Colours, distances and magnitudes for faint CVs
1213
Table 1. Table of observational results.
Object
Alternative
Type
names
(1)
BV416
N Com 1961
HV6292
PG1135+036
N Aqr 1907
CBS-31
AN 87.1911
SVS382
N Leo 1918
PG0917+342
AN 330.1928
AN 143.1935
US943
RE2107-05
PG0244H04
HV1319
HV 8002
Ton 408
PG0134+070
SVS1153
PG1038+155
CBS-119
TAV 0226+39
HV 3899
SVS 52
HV 6276
M LMi 1980
SON 9178
(3)
DN
DN
DN
AM?
DN
DN
DN
DN
DN
DN
DN
NL
DN
DN
DN
DN
AM
NL
DN
DN
DN
DN
NL
DN
NL
DN
DN
DN
DN
DN
DN
DN
DN
DN
DN
DN
80:
80.4:
82,h
84:
85
87
90
91
93 sh
94.8
102
104
105 sh
123 sh
123.8
125
125.02
198
236
236:
252
309
313.2
330
336
355
636
b
V
J
K
J-K
refs
(5)
(6)
(7)
(8)
(9)
(10)
-79.1
42.7
76.5
-39.
60.5
-35.2
64.2
65.1
-4.1
88.7
59.1
44.5
17.6
-23.3
49.3
47.2
-32.6
-42.9
-23.7
-18.8
-71.7
44.5
-53.9
2.2
57.4
53.2
-16.7
-19.1
-12.3
-62.3
-36.3
25.9
-42.3
-52.3
59.9
-12.9
18.0
17.5
20.5
18.5
15.7
17.0
16.0
18.6
17.0
20.0
19.2
14.7
18.0
18.0
19.3
20.7
16.5
15.8
17.0
17.0
13.9
19.0
15.8
17.0
17.0
17.8
16.5
19.0
18.5
18.5
17.2
21.0
17.5
20.8
20.5
21.0
16.95±0.20
16.30±0.19
15.26±0.16
15.96±0.18
15.04±0.15
14.01±0.11
14.42±0.13
15.48±0.16
16.64±0.21
15.90±0.18
17.31±0.22
15.55±0.17
14.63±0.13
13.08±0.08
15.07±0.15
16.1O±0.18
0.65±0.27
1,2,3
13,26
5,12,27
5,31
7
8,9
11,26
4,26,27
13,28
4,27
6,26,29
14,29,32
10
15
16,27,30
4
17
19
20,28,31
18,31
21
4,27
5,30
22,28
23
4,27,33
24
34
25,34
34
26,34
34
26,34
34
34
34
16.79±0.20
15.62±0.16
14.92±0.14
15.24±0.15
16.19±0.18
17.16±0.20
16.81±0.19
17.91±0.23
16.56±0.19
15.45±0.15
13.92±0.1O
16.14±0.18
18.60±0.19
16.10±0.17
14.82±0.14
14.34±0.11
13.95±0.11
12.24±0.08
16.64±0.19
15.37±0.15
14.65±0.12
16.01±0.17
12.94±0.08
18.51±0.25
16.38±0.19
17.1O±0.20
14.98±0.13
16.98±0.21
15.03±0.14
18.48±0.25
17.62±0.23
17.37±0.22
15.31±0.16
14.26±0.12
13.53±0.09
13.11±0.08
11.65±0.05
15.92±0.18
14.48±0.13
13.75±0.10
15.29±0.15
16.01±0.17
12.20±0.05
17.66±0.24
15.26±0.16
16.26±0.19
14.26±0.12
16.23±0.19
14.04±0.11
17.16±0.21
16.68±0.21
16.87±0.21
0.83±0.26
0.58±0.21
0.91±0.18
0.82±0.18
0.71±0.23
0.52±0.29
0.91±0.26
0.60±0.31
1.01±0.25
0.70±0.19
0.84±0.13
1.07±0.23
0.79±0.23
0.56±0.18
0.81±0.14
0.84±0.14
0.59±0.09
0.72±0.26
0.89±0.20
0.90±0.16
0.72±0.22
0.74±0.09
0.85±0.35
1.12±0.24
0.85±0.28
0.72±0.18
0.75±0.28
0.99±0.18
1.32±0.33
0.94±0.31
0.50±0.30
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
WX Cet
DIUMa
ALCorn
AH Eri
T Leo
VY Aqr
SXLMi
BCUMa
UV Per
GO Com
RZ Leo
BK Lyn
AY Lyr
EF Peg
DVUMa
DMDra
HU Aqr
WXAri
ARAnd
UU Aql
WWCet
AR Cnc
AY Psc
AFCam
DO Leo
RULMi
DX And
PQ And
DH Aql
EG Aqr
VZ Aqr
DV Dra
XZ Eri
SZ For
SS LMi
QY Per
(2)
Orbital
period
(mins)
(4)
Notes.
Column (3): DN - dwarf novae; NL - nova-like; AM - AM Her.
Column (4): (sh) superhump period.
Column (6): WW Cet is probably a VY Sci star (Ringwald 1996). A Y Lyr was observed 2 weeks after an outburst
which peaked on 1993 September 8/9. Monitoring observations indicated that it faded to below 15.3 mag by
September 17 indicating that this was a short outburst. Thus our J and K observations are likely to have been made
during quiescence.
Column (10): (1) Mennickent (1994); (2) O'Donoghue et al. (1991); (3) Downes & Margon (1981); (4) Howell et
al. (1990); (5) Szkody et al. (1989); (6) Howell & Szkody (1988); (7) Shafter & Szkody (1984); (8) Della Valle &
Augustein (1990); (9) Augustein (1994); (10) Nogami et al. 1994); (11) Wagner et al. (in preparation); (12)
Nogami & Kato (1995); (13) HSC90; (14) Dobrzycka & Howell (1992); (15) Howell & Liebert (1994); (16) Howell
et al. (1988); (17) Glenn et al. (1994); (18) Shafter (1992); (19) Beuermann et al. (1992); (20) Szkody, Piche &
Feinswog (1990); (21) Thorentson & Freed (1985); (22) Szkody & Howell (1989); (23) Abbott et al. (1990); (24)
Drew et al. (1993); (25) Vogt (1983); (26) Szkody & Howell (1992); (27) Mukai et al. (1990); (28) Szkody (1985);
(29) Howell et al. (1991); (30) Szkody & Howell (1993); (31) Szkody (1987); (32) Skillman & Patterson (1993);
(33) Howell et al. (1994); (34) Downes & Shara (1993).
where K, SK and Rz/Ro are the K-band magnitude, K-band
surface brightness and radius of the secondary star,
respectively.
The radius of the secondary can be estimated from the
orbital period using the empirical relationship of Patterson
(1984):
(2)
From this we can also estimate the spectral type of the
secondary assuming it to be a main-sequence star of radius
R2 (Zombeck 1990). Bailey's method is useful because the
K-band surface brightness, SK, changes relatively slowly with
© 1996 RAS, MNRAS 282, 1211-1222
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
1996MNRAS.282.1211S
1214 L. N. Sproats, S. B. Howell and K. 0. Mason
3.1 Contamination by the accretion disc
Formally, Bailey's method gives a lower limit to the distance
of a CV. This is because the accretion disc also potentially
contributes continuum and emission-line flux in the K band
in addition to that from the secondary.
We can establish a representative value of K%" the percentage of the K-band light contributed by the secondary,
for dwarf novae in quiescence from the results of timeresolved infrared photometric studies that are available in
the literature (Szkody & Mateo 1986, and references
therein; Szkody & Feinswog 1988; Ringwald 1995). In the
majority of cases, K%s has been estimated to be between
50-90 per cent from system to system. The K-band contribution of the secondary in AF Cam, the only one of our
programme stars for which suitable time-resolved infrared
photometry is available, is estimated to be 80 per cent by
Szkody & Mateo (1986). All of the systems contained in the
above samples from the literature have orbital periods
above the period gap. However, both the luminosity of the
secondary and the luminosity of the accretion disc will scale
with orbital period, thus a K%s contribution of 50-90 per
cent is still likely to be a reasonable assumption even below
the gap. This is consistent with the level of disc contamination estimated in Section 4.1. Thus, for DNe we list in Table
2 estimates of distance and absolute magnitude for three
representative values of the fraction of the K-band light
contributed by the secondary, K%s = 100 (the formal limit
provided by Bailey's method), 75 and 50 per cent.
For nova-like systems the situation is less clear, however
because the accretion rate in these systems is probably
higher than in a DN in quiescence, the likelihood is that the
fraction of the K-band light contributed by the secondary is
much less than in the quiescent DNe. For example, only 30
per cent of the light comes from the secondary in IX Vel
(Haug 1988). Because the data on nova-likes is relatively
sparse, we quote only the Bailey limit (corresponding to
K%s=100 per cent) in Table 2.
For the CVs in our sample with no orbital period data, a
lower limit (i.e K%s = 100 per cent) to the distance was
calculated using values of R2 and SK that corresponded to a
minimum period of 80 min and a maximum of 6 h. This is a
reasonable assumption, as the observed periods of the
majority of CVs lie in this range (e.g. Patterson 1984).
3.2 Effects of inclination, bright spot and secondary star
The orbital inclination of the binary, the presence of the
bright spot (the region where the accretion stream impacts
on the outer accretion disc) and the contribution of the
secondary star can all potentially affect m v , and therefore
must in principle be taken into account in determining
Mv·
We have specific estimates of the inclination of seven
stars in our sample - AR Cnc, A Y Psc, DV UMa and
DO Leo, which are high inclination systems with i > 70°,
and T Leo, WW Cet and V503 Cyg which are low-inclination systems (see Table 3 for references). Of the remaining
systems, none are known to be eclipsing, suggesting that
their orbital inclinations are not very high.
W87 gives an expression which corrects Mv for an optically thick disc for the effects of inclination, taking account of
limb darkening (Paczynski & Schwarzenberg-Czerny 1980).
This would suggest that a disc accreting system viewed at
i=O° is ~3.5 mag brighter than one viewed at i=85°. The
results of recent accretion disc modelling by Smak (1994)
agree with those of W87 for high mass transfer rate CVs
(DNe above the period gap and NLs). However, for low
mass transfer rates (~10 17 g S-I: see HSC95) when the disc
is optically thin, Smak finds that M v is much less affected by
the binary inclination and there is no significant reduction in
Mv except when the inclination is near 90°. Since the faint
DNe in our sample were observed at minimum, their discs
are probably optically thin, and we are justified in ignoring
any inclination effects for these systems during quiescence.
The bright spot, seen in photometric light curves as a
brightness 'hump', is generally only prominent in highinclination (i > 70°), short-period systems (e.g. W87). Presumably, however, the bright spot contributes a similar flux
in low-inclination systems spread over all orbital phases. For
maximum uniformity, we have used the mean V magnitude,
averaged over an orbital cycle (but excluding eclipses) for
all systems, irrespective of whether or not they show a prominent orbital hump. These magnitudes should generally be
considered as upper limits to the brightness of the 'disc'
proper. Of course, the bright spot luminosity should
properly be included in any estimate of the total accretion
luminosity of a system.
The secondary star is generally only detected in the Vband in long-period (Porb > 3-4 h) systems, and in these it
© 1996 RAS, MNRAS 282,1211-1222
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
spectral type for late-type stars expected to be found as the
companion to the white dwarf in short-period CVs. We use
the most recent data (Ramseyer 1994) relating Sk and
(V - K) for single field stars. The V - K colour, in turn, is
obtained from the spectral type using the data of Bessell &
Brett (1988).
CVs with periods of between 80 min and ~3 hare
expected to have M7 - M5 dwarf secondaries and we adopt a
value of SK~5.5 ±0.5 for these (based on fig. 2 of Ramseyer
1994); systems with periods from 3-6 h have ~M4-MO
secondaries, and for these we use SK~4.5 ±0.5. Six of the
stars in our sample (BC UMa, DV UMa, HU Aq, VZ Aqr,
AR Cnc and DZ And) have been found to exhibit absorption features in their optical spectrum attributable to the
secondary (Mukai et al. 1990; Szkody & Howell 1992; Drew,
Jones & Woods 1993; Glenn et al. 1994). In these cases the
spectral type can be verified directly. With the exception of
the two longest period systems, AR Cnc and DZ And, this
agrees with the estimate based on Patterson's empirical
relationship. For AR Cnc (Porb =5.15 h) we adopt a spectral
type of M4.5 (Mukai et al. 1990). The longest period system
in our sample (DXAnd, Porb = 10.6 h) has a spectral type of
K1 estimated from its absorption-line spectrum (Drew et al.
1993) and there is evidence that the secondary in this system
is oversized for its spectral type compared with a dwarf star
(see also Section 4.1). We thus combine the measured spectral type with the radius determined from equation (2) to
determine its distance and absolute magnitude.
The uncertainties in the distances were estimated by
propagating the errors on K and SKin quadrature through
equation (1). The error on the individual distances is typically ± 20 per cent.
1996MNRAS.282.1211S
Colours, distances and magnitudes for faint CVs
1215
Table 2. Table of results.
Object
Type
H.t3
Outburst
amplitude
E.W (A)
d
(pc)
(pc)
My
(quiescence)
My
(max)
(5)
(6)
(7)
(8)
185,213,260
81,92,113
11.6,11.3,10.9
12.2,11.8,11.4
3.3
9.3
(1)
(2)
(3)
(4)
WXCet
DIUMa
ON
ON
7.5-8.5
2.5
90
22
188,217,265
107,136,184
z
ALCorn
ON
9
47
187,215,264
182,209,257
14.1,13.7,13.4
4.7
AH Eri
AM?
5
70
113,131,160
71,82,101
13.2,12.9,12.5
7.9
T Leo
ON
5.5
62
76,88,107
62,77,93
11. 7,11.4,11.0
5.9
VY Aqr
ON
8.7
52
97,112,137
56,65,79
12.1,11.8,11.4
3.1
SXLMi
ON
3
64
150,173,211
135,156,190
10.1,9.8,9.4
6.8
BC UMa
ON
7
11
255,294,361
231,267,320
11.6,11.3,10.9
4.3
ON
6
91-98
186,215,263
13,15,19
10.7,10.3,9.9
4.3
ON
7
89
361,417,510
361,417,510
12.2,11.9,11.5
4.9
7.5
22
174,201,246
149,172,211
13.0,12.7,12.2
5.2
114
80
9.2
52,61,74
15,18,23
14.4,14.1,13.8
RZ Leo
ON
BK Lyn
NL
AY Lyr
ON
6
8.1
EF Peg
ON
7.9
172,199,243
68,79,96
11.8,11.5,11.1
3.6
OVUMa
ON
3.2
277,320,391
220,243,296
12.1,11.8,11.3
8.6
OM Ora
ON
5.3
580-700
425-513
12-11.5
6.5
HU Aqr
AM
111,191,231
192,222,272
10.1,9.8,9.3
9.3
WXAri
NL
198
133
UU Aql
ON
5
50-80
225,255,313
70,82,102
10.3,10.0,9.6
ARAnd
ON
6
50
269,310,380
108,125,153
9.9,9.5,9.1
3.5
WW Cet
ON
5
121,140,171
115,133,162
8.5,8.2,7.9
2.8
AR Cnc
ON
4
681,783,959
477,549,672
9.7,9.4,9.0
5.4
AY Psc
NL
565
457
7.0
AFCam
ON
00 Leo
NL
4
25-30
25-32
4.5
425,490,600
16,19,23
8.9,8.5,8.1
878
740
7.3
1273,1469,1799
1019,1176,1440
7.3,7.0,6.6
3
660,761,931
189,219,267
7.5,7.1,6.7
2.6
RULMi
ON
4
OX And
ON
4.5
PQ And
ON
8
352-2828
114-920
11.3-6.7
OH Aql
ON
6
116-937
25-195
13.2-8.6
EG Aqr
ON
4
184-1485
163-1315
12.2-7.6
42
5
VZ Aqr
ON
6
179-251
106-149
12.9-8.3
OV Ora
ON
6
182-1464
79-639
14.7-10.2
XZ Eri
ON
4
66-534
44-359
13.4-8.9
SZ For
ON
2.6
279-2247
221-1178
13.6-9.0
SSLMi
ON
6
355-1801
307-1558
12.7-9.2
QY Per
ON
5.8
244-1966
54-439
14.1-9.5
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
UV Per
GO Com
Notes.
Column (3): obtained from archival data in Oownes & Shara (1993) and Ritter & Kolb (1993).
Column (4): equivalent widths at minimum.
Column (5)-(7): distances and MvS based on values of K,% = 100,75 and 50 per cent respectively. For the novalike stars, only lower limits are given (corresponding to K,% = 100 per cent; see Section 3.1). For the 10 systems
with no orbital period data, distances and M vS are calculated using K, % = 100 per cent for orbital periods of 80 min
and 6 h (see Section 3.1). AY Lyr may have been on the decline from an outburst resulting in an underestimate
of its true instance.
Column (8): Mv at maximum calculated based on quiescent Mv, determined assuming K,%=75 per cent, and
outburst amplitude listed in column (3).
996 RAS, MNRAS 282,1211-1222
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
©1
1996MNRAS.282.1211S
1216 L. N. Sproats, S. B. Howell and K. 0. Mason
Thble 3. Distance and M v range comparisons with previous work.
Object
d(pc)
This Work
z
(pc)
TLeo
BCUMa
DVUMa
HU Aqr
WW Cet
ARCnc
DXAnd
VZ Aqr
76-107
255-361
277-391
111-231
121-171
681-959
660-931
179-251
62-93
231-320
220-296
192-272
115-162
477-672
189-267
106-149
Mv
11.7-11.0
11.6-10.9
12.1-11.3
10.1-9.3
7.9-8.5
9.7-9.0
7.5-6.7
12.9-8.3
Previous Work
Mean z
d(pc)
(pc)
Mv
>110
130-400
275-824
79-380
90-300
330-1330
",630
110-220
>100
240
390
103Q
185
590
",180
100
<11.1
11.0-13.5
9.5-11.9
14.6-11.1
9.4-12.0
7.7-10.7
6.7 b
10.5-12
Reference
2
1
1
4
3
1
5
6
Notes.
(a)z calculated from mean distance of d= 191 pc.
(b) Mv derived assuming d = 630 pc, B = 15.83 and B - V = 0.3.
References: (1) Mukai et al. (1990); (2) Warner (1987); (3) Young & Schneider (1981); (4) Glenn
et al. (1994); (5) Drew et al. (1993); (6) Szkody & Howell (1992).
K1 giant, supporting their suggestion that the secondary in
DX And is evolved.
4.2 Absolute magnitude aud distance
4 RESULTS
4.1
Infrared colours
In Fig. 1 we plot the observed] - K colour of our programme stars (from Table 1) as a function of orbital period.
The predicted J - K colour of the secondary star as a function of orbital period is also shown, where we have adopted
Patterson's empirical period-radius relationship and
assumed that the secondary is a main-sequence star. It can
be seen that there is a spread in the J - K colours of the CV
for a given orbital period, and that these, for the most part,
lie blueward of the predicted secondary] - K colour (a plot
of V -J versus orbital period shows a similar effect).
Although discrepancies between the predicted and
observed secondary star spectral types have been reported
for some systems (see Szkody & Howell 1993, and references therein), these are insufficient to account for the
range in the observed] -K colour.
The cause of the spreads in J - K and in the blue colours
compared with those expected from the secondary star is
almost certainly the residual flux from the accretion disc. A
~ 20 per cent continuum contribution from an accretion
disc in the K band (which has an approximately blackbody
spectral distribution and a] - K colour of ~ - 0.2) would
be enough to perturb the J - K colour by the amount
observed. Furthermore, infrared spectroscopy shows that
the] and K photometric bands include strong emission lines
from the accretion disc (Dhillon & Marsh 1995), which can
also cause the] - K colour to appear bluer than the main
sequence. Infrared photometry alone cannot distinguish
between these two possibilities.
The longest period system in our sample (DX And) lies
significantly redward of the main-sequence-star colour relation (Fig. 1). Drew et al. (1993) deduce a spectral type ofK1
for this star, and its J - K colour is consistent with that of a
The distances and absolute magnitudes derived for the stars
in our sample are listed in Table 2, and are shown graphically in Figs 2 and 3. These are compared with data on a
number of bright CVs (V < 16) as compiled by W87.
In Fig. 2, we plot the distance of each star above the
galactic plane, z, for both our sample and the compilation of
W87, as a function of orbital period. For reference the
horizontal dashed line represents the mean z distance of
old-disc Population II stars (Allen 1973). Most of the stars
in our sample have z < 400 pc with a distribution which is
indistinguishable from that of the stars in W87. Thus, the
majority of stars in our sample do not represent a separate
halo population of CVs.
However, the two stars with the largest z estimates,
DO Leo with z > 740 pc and RU LMi with z > 1019 pc, are
sufficiently far from the Galactic plane so as to constitute
two genuine halo candidates. RU LMi is presently the best
candidate for the most distant DN. For comparison, the
scale height of CVs above the Galactic plane derived by
Patterson (1984) is 190 pc, so the probability of finding a CV
from the Galactic disc population at z > 700 pc is of order
0.1 per cent.
In Fig. 3, Mv(min) is plotted as a function of orbital
period, where we again compare our sample with that of
W87. Also plotted is the linear fit to M v versus orbital period
for DNe at minimum light as calculated by W87 for the stars
in his list. The absolute magnitude expected for disc accretion at rates of 1016 and 1015 g S-1 is also indicated. We have
not plotted the CVs from W87 (his table 2) whose M v were
derived using the Mv(max) -Mv(min) empirical equation
(W87's equation 13).
Among the systems with known orbital periods in our
sample there are 15 DNe with P orb < 2.5 h. The empirical fits
of W87 to his compilation of data on the absolute magnitude of dwarf novae in quiescence would lead us to expect a
© 1996 RAS, MNRAS 282,1211-1222
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
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contributes less than 5 per cent of the total V-band flux (e.g.
AR Cnc: Mukai et al. 1990). The effect of the contribution
to M v from the secondary star is thus likely to be minor.
1996MNRAS.282.1211S
Colours, distances and magnitudes for faint CVs
1217
1.2
M
Main Se'tence Dwarf
Star J -K olours
M2.S
M
DO
u
Nova-Like
I::.
M6
0
1.0
...
:s
-§
Dwarf Nova
0
0
1EJ0
0.8
0
~
0
01::.
0
0.6
00
Typical Error
I::.
Bar
0
G6
0.4
0
200
600
800
Figure 1. The J - K colour of our sample plotted against orbital period. The J - K colour of field dwarf stars predicted to be present in
systems of a given orbital period are also shown.
1
0
· i
.. •
· •1tJ
---~-------.------------------------
CO .El
cR..
~1~
•
'0'
,e
.,
00
S
:a'"
N
~
•
•
I
•
Typical
uncertainty
•
•
•
I
•
200
0
•
0
•
•
1
0
•
• • •
••
g
•
•
o
DN - Our data
• DNfromW87
I::.
NL - Our data
• NLfromW87
400
Orbital Period (min)
600
800
Figure 2. The distance from the Galactic plane, Izl, plotted as a function of orbital period. The distances to a number of bright (V < 16) CVs
from the compilation given by W87 is shown for comparison. The dashed line represents the mean z distance of the old disc Population II
stars (Allen 1973). The z distances of the nova-like stars represent lower limits and the arrows indicate that their true distances are likely to
be greater than that shown.
Among the systems with known orbital periods in our
sample there are 15 ONe withPorb < 2.5 h. The empirical fits
of W87 to his compilation of data on the absolute magnitude of dwarf novae in quiescence would lead us to expect a
mean Mv~9.3 for these short-period systems with a range
(W87) betweenMv=7 and Mv= 10. In contrast we find that
the objects in our sample are distributed relatively uniformly at much fainter absolute magnitudes in the range
M v = 10-14. Taking the W87 data and our sample together,
the overall range of Mv for short-period ONe in quiescence
is from 7 to 14.
At longer orbital periods (Porb > 3 h) we observed six
ONe. In contrast to the short-period systems, the absolute
magnitudes of the long-period CV are within the range seen
by W87, i.e. consistent with his empirical fit.
Our sample also includes four nova-like objects. For
nova-likes, assuming that they are high accretion rate systems, the limits derived from Bailey's method are likely to
be very conservative since the true fractional contribution of
the secondary to the K-band light could be very small. Three
of the nova-like systems in our sample. WX Ari, A Y Psc and
DO Leo, have P Orb > 2.5 h. WX Ari (Porb =3.3 h) has
© 1996 RAS, MNRAS 282,1211-1222
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014
400
Orbital Period (mins)
1996MNRAS.282.1211S
1218 L. N Sproats, S. B. Howell and K. 0. Mason
4
6
o
]
(
8
M= 1.0x 10 16g1s
~ 10
M=1.0x10 15g1s
·a
~
(
'0
:<"'
12
0
~B
I
Typical
uncertainty
DO
14
0
0
o
0
DN - Our data •
6. NL - Our data •
200
DN from W87
NL from W87
400
600
800
Figure 3. The absolute magnitude (at minimum for dwarf novae) plotted as a function of orbital period. Approximate rates of mass transfer
through the disc, absolute magitude of the secondary star [Mv(2)] and the linear fit to the Mv of the dwarf novae given by equation (18) of
W87 are also shown. The absolute magnitudes of the nova-like stars are lower limtis where the arrows indicate that their true M v are likely
to be more luminous than that shown.
Mv<9.3, while AYPsc (Porb =5.2h) has Mv<7.0 and
DO Leo (Porb =5.6 h; see also Section 5.1) has Mv<7.3.
This compares with an expectation for M v of between 4 and
7 for long-period nova-like CVs based on the compilation of
W87. The only short-period nova-like in our sample is
BK Lyn (Porb = 1.7 h) which has Mv< 9.2. The basis of the
classification of BK Lyn as a nova-like is its spectrum which,
at the V magnitude quoted in Table 1, exhibits broad
absorption lines of H, He I and Ca I with variable emission
cores (which are weak except at HIX). Warner does not have
any short-period nova-likes in his sample, but if the true
absolute magnitude of BK Lyn is at the faint end of the
long-period nova-like range (i.e .. Mv=7), the secondary star
would be contributing about 5 per cent of the total K-band
luminosity.
Distance and M v estimates have been made previously for
eight CVs in our sample. In Table 3 we compare our measurements of these stars with those in the literature and find
that they agree well.
5 DISCUSSION
We have obtained J- and K-band colours of a sample of
faint, predominantly high-latitude CVs which were potentially systems residing in the Galactic halo. These include a
number of the TOAD subgroup of DNe that have outburst
amplitudes greater than 6 mag. Our results indicate that the
majority of these faint systems are not in fact in the halo,
rather their distribution above the Galactic plane is indistinguishable from existing data on brighter objects (W87;
Howell & Szkody 1995). A corollary of this result, however,
is that many of these systems have very low absolute luminosities in quiescence, with the absolute magnitude of some
short-period DNe being as faint asMv= 14. The implication
is that the space density of CVs in the local neighbourhood
may be higher than previously thought (cf. Ford & Jacoby
1978; Drissen et at. 1994; HSC95).
5.1
Halo CV
Two stars in our sample, DO Leo and RU LMi, stand out as
being at large distances above the Galactic plane. These
appear to be genuine halo objects. DO Leo is classed as a
nova-like object and RU LMi as a DN. Both have similar
orbital periods of just under 6 h.
DO Leo is an eclipsing system (Abbott et at. 1990). The
eclipse depth is approximately 1.5 mag, and the full width at
half flux level is 5 per cent of the orbital cycle, suggesting a
relatively high inclination. The optical spectrum has He II
4686-A emission as well as Balmer line emission, consistent
with a nova-like classification of the object. Abbott et al.
found no evidence for secondary star features in the spectrum, and did not confirm the earlier report by Green et at.
(1982) of a 'red' continuum component
RU LMi has a rather complex orbital light curve at minimum light with a total range of ~ 1.2 mag (Howell et al.
1990). The latter authors found a mean quiescent magnitude of 17.8, implying an outburst amplitude of 4 mag. The
star is reported as faint as 19.5 (pg) in Downes & Shara
(1993), but this might be consistent with the Howell et al.
result if the measurement was made in the faint part of the
orbital cycle. The optical spectrum of RU LMi is shown in
Mukai et al. (1990) and exhibits a typical CV emission-line
spectrum with no evidence of a secondary star signature.
Our work suggests that the absolute magnitude of
RU LMi is 7, consistent with the upper range found by W87
for long-period DNe in quiescence. DO Leo has an absolute
magnitude which is < 7.3. While this limit is consistent with
the faint end of the values derived by W87 for long-period
nova-likes, the absolute magnitude of this star may also be
© 1996 RAS, MNRAS 282,1211-1222
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
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Orbital Period (min)
1996MNRAS.282.1211S
Colours, distances and magnitudes for faint CVs
suppressed because of its high inclination (cf. W87). In all,
based on the information available to date, neither halo
candidate exhibits properties that would distinguish them
from CVs in the local neighbourhood.
5.2 Faint systems in qniescence
One TOAD system with good outburst coverage is
SW UMa. Howell et al. (1995a) describe three different
outburst types in this star, characterized by different durations and amplitudes. These range from typical superoutburst light curves with large amplitude which last for
16-20 d, to, at the opposite extreme, much lower amplitude
outbursts which last for only ~ 2 d. The larger outbursts are
preceded by a long period of quiescence ( ~ 1500 d). While
coverage for most other TOADs is less complete, Howell et
al. (1995b) give evidence that these characteristics may be
typical for stars in this group.
We also have estimates of both the outburst duration and
recurrence time for WX Cet (20 and 3600 d, respectively:
O'Donoghue et al. 1991), in line with the data on SW UMa.
The outburst amplitUde of WX Cet is 8 mag, which implies
a mean absolute magnitude averaged over the outbursts of
9 (cf. 11.3 in quiescence). The outburst recurrence times of
VY Aqr, AL Com and EF Peg are given as 600, 325 and
200 d, respectively (HSC95; Kholopov & Efremov 1976),
while their outburst amplitudes are 8.7, 9 and 7.9 mag,
respectively. Although we have little information on the
outburst durations of these stars, if we assume that they are
similar to those of WX Cet and SW UMa (i.e. 20 d) we find
mean absolute magnitudes for these systems of 9.7 (11.8),
7.3 (13.7) and 6.8 (11.5), where we list the quiescent
absolute magnitudes in parentheses for comparison. The
absolute magnitude expected for constant accretion
through the disc driven by gravitational radiation in these
short-period systems is about 9.7 (W87). Thus, given the
uncertainties, there is no strong evidence that the mean
accretion rate of these systems falls short of the rate
expected from gravitational radiation. A different rate
might be expected, for instance, if these systems were out of
equilibrium (for example, owing to a nova explosion: Shara
et al. 1986). There is a clear need for improved observational data before this question can be conclusively laid to
rest.
5.3
Large-amplitnde outbursts
A number of the stars in our sample have large-amplitude
optical outbursts. One reason advanced to explain the large
amplitudes is that the quiescent magnitude in these systems
is abnormally low, so that the contrast with the outburst is
greater than for other DNe (HSC95). This would imply that
the TOADs ought to be those stars with the faintest accretion discs in quiescence.
We test this in Fig. 4, where we plot the outburst amplitude of DNe versus their absolute magnitude in quiescence.
In addition to the stars in our sample, we also include data
on three additional systems, WZ Sge, BZ UMa and
AK Cnc, taken from HSC95.
Fig. 4 indicates that the majority of the data would be
consistent with a relationship between outburst amplitude
and the absolute magnitude in quiescence. Four systems,
however, stand out as departing significantly from a linear
relation - A Y Lyr, DV UMa, DI UMa and SX LMi - and it
is important to consider whether these are truly discrepant.
Considering the four systems, in turn, there is doubt as to
whether A Y Lyr was truly in quiescence when we observed
it, as discussed in Section 2. If it was still on the decline from
an outburst, the absolute magnitude we measure for this
© 1996 RAS, MNRAS 282,1211-1222
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
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One of the major results of this work is that a substantial
fraction of faint, short-period CVs in our sample have a low
intrinsic luminosity in quiescence. Using typical accretion
disc brightness profiles, the implied mass transfer rate in
these discs is less than 1.0 x 1015 g S-I. This is near or slightly
below the values used by HSC95 in their accretion disc
models of TOADs. This rate is also much lower than typical
values used for accretion disc models, ~ 1019 _1017 g S-1 (e.g.
Smak 1994). Even at 1017 g S-I, Smak finds that the accretion
discs are optically thin except for the highest of inclinations
and differ markedly from blackbody models. Further theoretical work to understand these 'weak' accretion discs
would be in order.
In Fig. 3 we plot the expected absolute magnitude of the
secondary star [Mv(2)] as a function of orbital period based
on the empirical relation given by P84. We note that our
results are consistent with this relation in the sense that the
absolute magnitUde determined for our short-period programme stars are all brighter than the expected absolute
magnitude of the secondary. However, it is also clear that
the secondary must contribute a large proportion of the
quiescent light in the fainter systems. Similarly, the difference between the measured absolute magnitude and that
expected for the secondary provides an upper limit to the
absolute magnitude of the white dwarf star. Isolated hot
white dwarfs have measured absolute magnitudes as high as
10 (Allen 1973). The fainter systems we are studying clearly
do not contain white dwarfs that are this luminous, implying
instead that the white dwarfs are older and cooler.
We include in Table 2 available data on the equivalent
width of the Hf3 emission-line at or near quiescence. P84
have suggested that there is an o~erall correlation between
the Hf3 equivalent width and absolute magnitude, with
fainter systems showing a higher equivalent width. The
mean equivalent width of the emission lines in our sample is
similar to the mean value of previously studied stars in
quiescence at similar orbital periods. However, the equivalent widths of the lines do not appear to be sensitive to the
absolute magnitude in quiescence, and our data thus show
considerable scatter on P84's plot of equivalent width versus
disc absolute magnitude owing to the large range of quiescent absolute magnitudes we measure.
Another question posed by these intrinsically faint systems, and one which can be addressed by further observations, is to what extent their low luminosity reflects
conditions in the quiescent state alone, or if the overall
accretion rate of the system is also low. To assess this we
need information about the luminosity generated during
outburst, i.e. statistics on outburst amplitudes, durations
and recurrence times. Because these systems are (by definition) very faint during quiescence, however, they have not
been well monitored and information on the outburst duration and recurrence times, in particular, is very incomplete.
1219
1996MNRAS.282.1211S
1220 L. N Sproats, S. B. Howell and K. 0. Mason
10
<> DN
• TOADDN
.
ALCorn
VY Aqr
•
WX Cet •• EF Peg
WZ Sge)l(
BC UMaARAnd
)I(
•
wwCet <>
<>
• GO Com
AKCne
~
AYLyr
UVPer <>
TLeo
UUAql
DXAnd <>
4
RZ Leo
'BZUMa
RULMi<>
<>
AFCarn
<>
ARCne
SXLMi <>
lDvUMa
DIUMa <>
6
8
14
16
Figure 4. The outburst amplitude (A) plotterd as a function of absolute magnitude in quiescence of the dwarf novae in our sample with
known orbital periods. The solid line represents the best linear fit to the data set of outburst amplitude versus absolute magnitude in
quiescence which is given by A =0.73 Mv (min) -1.41. The three asterisk symbols are stars taken from HSC95.
2
<>
e
.§
4
~
.,
:E
<>
.g
B
.~
6
<>
•
TOAD
<> DN
:E'"
<> SXLMi
B
.E!
0
'"
~
8
<> DVUMa
<> DIUMa
10
100
200
400
300
Orbital Period (mins)
500
600
700
Figure 5. The absolute magnitude at maximum of dwarf novae in our sample plotted against orbital period. The linear fit to the M v at
maximum of the dwarf nova given by W87 (his equation 14) is shown and the three outlier stars are labelled (see Section 5.3).
star is likely to be an upper limit on the true value in quiescence (since if the star is not in quiescence the secondary
will be contributing a smaller fraction of the K-band light
than we assume), and thus A Y Lyr would move to the left in
Fig. 4, more in line with the majority of the data points.
DV UMa, on the other hand, is known to be a very highinclination (eclipsing) system, and while its measured
absolute magnitude in quiescence is probably reasonably
representative (see Section 3.2) the high inclination can be
expected to suppress the observed brightness in outburst, so
that the true increase in disc luminosity during outburst will
be underestimated. Thus DV UMa would likely move up in
Fig. 4.
This leaves DI UMa and SX LMi. DI UMa belongs to a
small group of stars referred to variously as the RZ LMi or
ER UMa stars (Kato & Kunjaya 1995; Nogami et al. 1995;
Misselt & Shafter 1995; Robertson et al. 1995). These are
SU UMa stars that exhibit extremely short superoutburst
cycles (20-45 d). Recent evidence (Wagner et ai., private
communication) suggests that SX LMi may have similar
characteristics, although data on this star are less complete.
It is possible therefore that these two systems, along with
the other RZ LMi stars, represent a distinct physical type,
perhaps associated with a high mean mass transfer rate.
Setting these possibly peculiar objects aside, the remaining
data in Fig. 4 do suggest a relationship between outburst
© 1996 RAS, MNRAS 282,1211-1222
© Royal Astronomical Society • Provided by the NASA Astrophysics Data System
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10
12
Absolute Magnitude in Quiescence
1996MNRAS.282.1211S
Colours, distances and magnitudes for faint CVs
ACKNOWLEDGMENTS
We are grateful to our telescope operators, Joel Aycock and
Thor Wold, and to Colin Aspin for his assistance and useful
discussions throughout our program. We also thank Jon
Mittaz for help in reducing the data from the second run.
This work has made use of the SIMBAD data base run by the
CDS at Strabourg and the NASA Astrophysics Data System
(ADS). LNS is grateful for the hospitality of the staff at PSI
during part of the analysis and writing of this paper. We
thank the anonymous referee for useful comments that led
to an improved manuscript.
LNS is supported by a PPARC studentship. SBH acknowledges partial support from NSF grant AST-921971.
UKIRT is operated by the Royal Observatory Edinburgh on
behalf of the PPARC.
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