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 Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014 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 Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014 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 Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014 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 Downloaded from http://mnras.oxfordjournals.org/ by guest on October 6, 2014 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. 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Stars 4155 O'Donoghue D., Chen A, Winkler H., Marang F., Mittaz J. P. D., 1991, MNRAS, 250, 363 © 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 amplitude and quiescent absolute magnitude. In particular, the explanation of the large outburst amplitude (TOAD) systems as arising in systems with faint quiescent absolute magnitudes is borne out by the data. Accretion disc models developed to explain the large, infrequent outbursts in the TOADs (HSC95) suggest that the viscosity of the disc material during the minimum (cold) state must be very low, even at low Mvalues such as 10- 11 Mo yr-l. Furthermore, the results presented by Smak (1994) indicate that these types of discs, while accumulating material, mostly remain optically thin. Thus, it is reasonable to expect that the accretion disc contribution to the total M v of these systems is correspondingly low, leading to the observational results seen in Fig. 4. We can check the consistency of our interpretation by examining the implied absolute magnitude at maximum of DNe in our sample, derived by combining the absolute magnitude in quiescence, which we measure using Bailey's method, with the outburst amplitude. In Fig. 5 we plot the values so obtained as a function of orbital period, and compare them with the empirical relationship determined by W87 (his equation 13). The agreement is good, apart from the three stars noted above as peculiar (the RZ LMi stars SX LMi and DI UMa, and the high-inclination system DV UMa). This demonstrates that a consistent picture emerges for the majority of DNe in which their peak luminosities are similar, but there is a much wider range of brightness in quiesence. 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