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Optical observations of close binary systems with a compact component
Augusteijn, T.
Publication date
1994
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Citation for published version (APA):
Augusteijn, T. (1994). Optical observations of close binary systems with a compact
component.
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8 8
AA 59
m
photometric period in the dwarf nova V485
Cen n
T.. Augusteijn, M.H. van Kerkwijk, and J. van Paradijs
AstronomyAstronomy & Astrophysics 267, L55 (1993)
Abstract t
Wee present time resolved V-band CCD photometry of the dwarf nova V485 Cen. Itt is found that the optical brightness during quiescence is modulated with a period off 59m. Brightness variations with the same period may also by present during outburst.. The observed spectrum in quiescence is fairly typical for a dwarf nova. However,, we find that the He I line at 5876 A is much stronger relative to Ha than iss the case in other cataclysmic variables. We discuss the nature of the 59m period inn terms of (i) the rotation period of the white dwarf; (ii) the orbital period of the system. .
8.11 Introduction
Inn this Letter we present spectroscopy and time resolved CCD photometry of the cataclysmic variablee (CV) V485 Cen. This system is classified in the General Catalogue of Variable Stars (Kholopovv et al. 1985) as a U Gem type dwarf nova. Some monitoring of this source has beenn done by members of the Variable Star Section of the Royal Astronomical Society of New Zealandd (e.g. Bateson 1979,1982). About twenty outburst have been observed with a duration off between ~ 1 and ~ 7 days. During outburst the system reaches V~ 14 mag. Due to its relativee faintness many outbursts probably have remained unobserved.
8.22 Observations and Analysis
8.2.11 Photometry
Wee observed V485 Cen during three nights with a GEC CCD attached to the 91cm Dutch telescopee at the European Southern Observatory in Chile. The source was monitored, using a Besselll V filter, for 6.8 hrs on March 31s t starting at 2:13 UT, for 5.7 hrs on April 3r d starting att 1:13 UT, and for 3.0 hrs on April 4t h starting at 1:17 UT. We obtained a total of 187 V filter
exposuress with an integration time of 4m each. Two Bessell B filter exposures of 5m each were takenn on March 31s t and April 3r d starting at 5:11 and 0:56 UT, respectively.
112 2 88 A 59"1 photometric period in the dwarf nova V485 Cen
F i g u r ee 8 . 1 . The differential V magnitudee of V485 Cen with re-spectt to a nearby comparison star ass function of Heliocentric Julian Date.. The magnitude of the com-parisonn star is Vcornp=15.04(2) mag
15.66 15.7 15.8 16.6
Timee (JD0-24487OO.)
D u r i n gg t h e observations a near-real t i m e analysis of the raw C C D frames using a p e r t u r e p h o t o m e t r yy software was performed. All d a t a were differentially reduced with respect to a c o m p a r i s o nn s t a r ( V = 1 5 . 0 4 ( 2 ) , ( B - V ) = 0 . 6 4 ( 4 ) ) located ~ 5 0 " W , a n d ~ 1 0 " S with respect t o t h ee source.
Absolutee calibration was o b t a i n e d from observations of two Landolt C C D fields. E a c h con-t a i n ss a n u m b e r of s con-t a r s w i con-t h a large range in B - V colours. T h e calibracon-tions decon-termined from t h ee t w o fields separately were in excellent a g r e e m e n t , differing by only 0.002 m a g in V a n d 0.020 m a gg in B . T h e s t a n d a r d m a g n i t u d e s of the comparison star were determined differentially from t h ee comparison of several images of the V485 Cen field which were taken close in t i m e a n d in a i r m a s ss t o observations of one of t h e standard fields. All the errors in m a g n i t u d e s a n d colours q u o t e dd in this Letter a r e e s t i m a t e s determined from the r m s values of the observed values, a n d shouldd b e considered as a r o u g h indication of t h e 1-<T errors in these values.
I nn Fig. 8.1 we show t h e differential m a g n i t u d e of t h e source with respect t o the compar-isonn s t a r as function of Heliocentric Julian date. Clearly, t h e source was in o u t b u r s t during t h ee first night of observations. Unfortunately the weather during the following two nights pre-v e n t e dd obserpre-ving. By t h e last two nights the source h a d reached a constant brightness lepre-vel s u p e r p o s e dd on which a r e relatively large variations (Fig. 8.1). Using t h e absolute calibration describedd above we derive for t h e outburst a n d t h e quiescent s t a t e V = 1 5 . 6 1 ( 4 ) , ( B - V ) = 0 . 0 1 ( 6 ) att JD@ =2448712.7300, a n d V=18.04(9), ( B - V ) = 0 . 0 4 ( 1 1 ) at J D0= 2 4 4 8 7 1 5 . 5 4 6 1 respectively
( t h ee colours are r a t h e r u n c e r t a i n as the B a n d V frames are not taking simultaneous a n d [peri-odic]] colour variations could b e present).
Too look for periodic variations we performed a Fourier s p e c t r u m analysis using t h e Lomb-Scarglee m e t h o d (see Press a n d Rybicki 1989 and references therein) of t h e observations from eachh night separately. T h e clear t r e n d seen in the observations of the first night was removed by s u b t r a c t i n gg a low-order p o l y n o m i a l before analysing the d a t a . T h e individual frequency spectra a r ee presented in F i g . 8.2. D u r i n g t h e last t w o nights a strong peak is found at a period of ~ lh
( ~ 2 44 c y / d a y ) . T h e frequency s p e c t r u m for the first night does not show any such strong peaks, b u tt t h e highest peak (at ~ 4 9 c y / d a y ) is consistent with being the second h a r m o n i c of the ~ 2 4 c y / d a yy frequency. T h e second ( a n d the t h i r d ) h a r m o n i c of t h e ~ 2 4 c y / d a y frequency is also p r e s e n tt in t h e frequency s p e c t r u m for the second night.
T h ee frequency s p e c t r u m of t h e second night also shows a peak at a period of ~ 2 . 6 . No clearr p e a k a t this p e r i o d is seen in t h e other nights (however, during t h e last night t h e source
12.6 6 12.7 7 12.8 8 AA _ \ >,, \ ( I ) ) I N N ro o CN N n n
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8.28.2 Observations and Analysis 113 3
F i g u r ee 8.2. Frequency spectra for (fromm top to bottom) the observa-tionss of each night separately, and thee last two nights combined. The observationss of the first night have beenn corrected for the trend present inn the data. The ordinate indicates thee power, normalized on the to-tall variance of the data, as func-tionn of frequency. Indicated are the expectedd positions of peaks of the fundamentall and harmonics for a periodd of 0.041096 days, or 24.333 cy/dayy (see text)
500 100 Frequencyy ( c y / d a y )
wass observed for only 3h) . T h e ~ 2 . 6h period might be due t o slow a t m o s p h e r i c variations in the
differentiall m a g n i t u d e s as a result of t h e colour difference between V485 Cen a n d t h e comparison s t a r .. However, no clear variation is found in t h e differential m a g n i t u d e s of t h e comparison star withh respect t o a second comparison star with V = 1 5 . 0 3 ( 2 ) , ( B - V ) = 1.00(4). Still, we stress t h a t
AA period, should be considered very t e n t a t i v e
att t h e present t h e ~ 2 . 6h period, contrary t o t h e att best.
Alsoo shown in Fig. 8.2 is t h e frequency s p e c t r u m of t h e d a t a from t h e last two nights combined.. T h e peaks present in t h e transforms of the individual nights are now split i n t o several peakss as a result of t h e 1-day spacing between the observations. T h e highest peak corresponds t oo a frequency of 24.307(77) cy/day, b u t frequencies corresponding t o t h e two peaks on either sidee of t h e highest peak can not be excluded. However, in all cases t h a t significant harmonics aree detected, they agree best with t h e above frequency being t h e f u n d a m e n t a l . Especially, the highestt peak in t h e t r a n s f o r m of t h e first night corresponds t o a frequency of 49.00(58) cy/day, whichh excludes t h e f u n d a m e n t a l frequency being either 23.307 or 25.307 c y / d a y at t h e 4.0 and 2.77 <r level, respectively.
Sinusoidall fits with a fixed period of 0.04114 days t o t h e observations in t h e last two nights separatelyy give arrival times of m a x i m u m light of J D ® = 2 4 4 8 7 1 5 . 6 7 4 3 ( 1 4 ) a n d 2448716.6195(17) respectively,, w i t h a corresponding period of 0.041096(96) days ( t h e 1-day aliases being a period off either 0.039383(92) or 0.04296(10) days, see above). In Fig. 8.3 we show t h e observations of thee last t w o nights m o d u l o t h e 0.041096 day period. Two cycles are shown for clarity.
Wee also performed sinusoidal fits t o t h e observations from t h e first a n d second night with halff t h e 59™ period. Unfortunately the errors in t h e arrival times a n d in t h e fundamental period aree too large t o d e t e r m i n e if t h e brightness variations at t h e first h a r m o n i c during t h e first night, whenn t h e source was in o u t b u r s t , were in phase with t h e brightness variations at the same period duringg t h e second night.
1144 8 A 59™ photometric period in t i e dwarf nova V485 Cen
Figuree 8.3. The differential V
magnitudee of V485 Cen obtained duringg the 2n d (circles) and 3r d
(crosses)) nights of observation as functionn of the 0.041096 day period. Thee light curve is shown twice for clarity y
8 . 2 . 22 S p e c t r o s c o p y
H o p i n gg t o clarify t h e n a t u r e of t h e ~ lh p h o t o m e t r i c period, we o b t a i n e d spectra of t h e source on J u l yy 2n d, with t h e New Technology Telescope at t h e E u r o p e a n Southern Observatory in Chile,
u s i n gg t h e ESO M u l t i M o d e I n s t r u m e n t ( E M M I ) .
O n ee s p e c t r u m was o b t a i n e d in t h e red a r m in low dispersion m o d e using grism # 4 a n d a T h o m s o nn C C D , covering 5700-10000A at ~ 5 A p e r pixel. A slit w i d t h of 1" was used, corre-s p o n d i n gg to 2.2 pixel ( l l A ) . T h e corre-slit wacorre-s r o t a t e d corre-such t h a t a corre-s t a r located ~ 3 0 " W , a n d ~ 4 5 " S w i t hh respect t o t h e variable, was also in t h e slit. T h e i n t e g r a t i o n t i m e was 1 5m, which led t o a
signal-to-noisee r a t i o of ~ 2 5 p e r pixel.
T w oo spectra were o b t a i n e d using both a r m s simultaneously in m e d i u m dispersion m o d e (with aa dichroic as b e a m s p l i t t e r ) . In t h e blue a r m grating # 4 was used, which gives a wavelength coveragee of 39005200A at ~ 2 A per pixel. In t h e red a r m grating # 7 was used, covering 7 9 0 0 -9000AA a t l . l A per pixel. T h e slit width was 0.9", which corresponds t o 2.5 pixels (5A) in t h e b l u ee a r m a n d t o 2 pixels (2.2 A) in t h e red a r m . Unfortunately, b o t h spectra are of poor quality, w i t hh signal-to-noise ratios p e r pixel of only 10 a n d 5 in t h e blue and red a r m , respectively.
I nn Fig. 8.4 we show t h e s p e c t r u m of V485 Cen as obtained with t h e grism. T h e s p e c t r u m iss fairly typical for a C V , showing H a , He I 5876, 6678, 7065A, several Paschen lines, and t h e C a nn triplet in emission. T h e equivalent w i d t h s ( E W ' s ) of those H B a l m e r a n d t h e He I lines we weree able t o m e a s u r e a r e : H a = 2 0 . 5 ( 2 . 5 ) A; H/3=9(5) A; H 7 = 2 . 9 ( l . l ) A; H e l 7065=6.0(2.0) A; Hee I 6678=4.5(1.5) A; a n d He I 5876=14.0(3.0) A ( t h e quoted errors are 90% confidence limits). Forr t h e He II 4686 line we o b t a i n e d a 90% confidence upper limit of 2 A. T h e H W Z I of t h e H a
F i g u r ee 8.4. Spectra of V485 Cen andd a second star on the slit ob-tainedd with EMMI at the NTT us-ingg grism # 4 . Also shown is the ra-tioo of these two spectra (see text). Eachh spectrum has been normalized att 8000A and shifted vertically for clarity y
60000 7000 8000 9000 10
4A(A) )
CO O f-> > CM M PI PI .. I ' ' I ' ' 9 9-- »
6 0 0''
\ %
11'
0 0> >
OO O xOo ftft * * x v . . x 00 % »c f i*< !'*»c?* 8> <?"*>? ** a 00 o,, .
, , , , > , ,
1 1 x x « 0 0" \\ *
'' ' 1 . "f," " A ** OO O „ 0 ftft * * * 0 O v x 00 Z> x ex x 00 % ?*M«* SS 0 " * o 0 8 00 0 11 1 1 . 1 1 , .. , 1 0.55 1 1.5 ^0.0410966 day8.38.3 Discussion 115 5
linee is 0 km/s (90% confidence).
Alsoo shown in Fig. 8.4 is the spectrum of the second star on the slit. Its spectral features (alsoo in the blue) and its photometric colour indicate a spectral type in the range of a (slightly reddened)) late-type F to mid-type G main-sequence star. In an attempt to correct for atmo-sphericc effects we divided the two spectra (see Fig. 8.4). The most interesting features of this ratioo spectrum are two 'dips' which appear on either side of the He I emission line at 7065A. Althoughh both 'dips' (partially) coincide with atmospheric features it is hard to understand whyy these two features should remain, whilst the (much stronger) atmospheric feature at 7600A hass nearly completely disappeared. We limit ourselves here only to note that wide, shallow absorptionn features at ~6800 and ~7200A due to molecular absorption, are seen in M-dwarf star.. However, M-dwarfs also show an absorption feature of similar strength at ~6200A which iss not present in the spectrum presented here.
8.33 Discussion
Thee period distribution of CV's shows two striking features: (i) there is a 'period gap' between ~ 2hh and ~ 3h; (ii) the distribution has a cut-off at a rninirnum period of ~ 8 0m. The lower cut-offf in the orbital period distribution can be explained as the result of the secondary becoming degeneratee when its mass becomes less than ~ 0.1 Af© (Paczynski and Sienkiewicz 1981). The valuee of this minimum orbital period depends on the total mass of the system and its chemical compositionn (Paczynski and Sienkiewicz 1981, Rappaport, Joss, and Webbink 1982, Sienkiewicz 1984).. For a system with solar abundance this minimum period is ~ 8 0m.
Theree are presently two groups of CV's known that show periodic brightness variations with periodss of an hour (or less). One group, the AM CVn stars, are characterized by the total absencee of hydrogen in their spectra and extremely short photometric and/or spectroscopic periodss of 17.5-46.5m (Ritter 1990). These systems are probably CV's containing a (helium degenerate)) white dwarf secondary.
Thee other group are the intermediate polars (IP's) in which a strong magnetic field of the whitee dwarf disrupts the inner part of the accretion disk, and funnels matter onto the magnetic poles.. This results in brightness variation in X-rays and the optical at the white dwarf rotation periodd and/or its orbital side bands.
Inn the following we address the question whether the ~ lh photometric period may represent thee white dwarf rotation period, or the orbital period of the system.
8.3.11 A 1 hour rotation period of the white dwarf?
Mostt IP's have spin periods of order ~103 sec, and orbital periods which are substantially longer (typicallyy a factor ~10 Ritter 1990; Barrett et al. 1988).
Veryy long orbital periods (e.g. conforming to the ratio ~10 between orbital and spin periods seenn in many IP's, see Barrett et al. 1988) can probably be excluded for V485 Cen, since at ann orbital period of ~10h we would have detected the spectral signatures of a giant companion star.. The quiescent ( B - V ) indicates that the orbital period is ;$6h (see Echevarria and Jones 1984).. The only indication we found for a periodic signal other than the lh period is the very
tentativee ~2.6h period in our second night of observations (see Sect. 8.2.1).
IP'ss rarely show dwarf nova outbursts (small outbursts have been observed for V1223Sgr: Vann Amerongen and Van Paradijs 1990; TV Col: Szkody and Mateo 1984). The outburst of V4855 Cen reported here was not noted by members of the Variable Star Section of the RASNZ, butt an outburst was detected only two months later indicating that these outbursts are not very rare. .
116 6 88 A 59™ photometric period in the dwarf nova V485 Cen
F i g u r ee 8.5. The equivalent width off Ha vs. He I 5876A for several cataclysmicc variables as observed byy Williams (1983). The symbol givess the position of V485 Cen as determinedd from the observations presentedd here. Indicated are the 90%% confidence limits
9 7m) ,, which shows regular (recurrence time ~ 6 0 0 days, see R i t t e r 1990) dwarf novae o u t b u r s t
off c o m p a r a b l e a m p l i t u d e t o those of V485 Cen. A p a r t from regular eclipses, no strong orbital p h o t o m e t r i cc variations are found in EX Hya. This might be t h e case for V485 Cen as well.
8 . 3 . 22 A 1 hour orbital p e r i o d ?
T h ee t o t a l range in b r i g h t n e s s , a n d t h e changes with the average brightness in t h e s t r e n g t h of t h e v a r i a t i o n ss at t h e f u n d a m e n t a l frequency and its harmonics (see Fig. 8.1 and 8.2) are reminiscent off similar variations observed in t h e AM C V n t y p e variables P G 1 3 4 6 + 0 8 2 (Wood et al. 1987) a n dd V 8 0 3 Cen ( O ' D o n o g h u e a n d Kilkenny 1989).
T h ee detection of h y d r o g e n emission lines in t h e s p e c t r u m of V485 Cen (see Fig. 8.4) excludes t h ee possibility t h a t it is a n A M C V n system. However, the secondary might be a degenerate c o n t a i n i n gg (some) hydrogen. A n o t h e r possibility is t h a t the companion in V485 Cen is not d e g e n e r a t e ,, b u t a hydrogen-deficient main-sequence star. For a n orbital period of ~ lh, this wouldd imply a hydrogen fraction of t h e secondary of less t h a n 30% (Sienkiewicz 1984).
Iff V485 Cen were either such system, one might - p e r h a p s naively, since line s t r e n g t h does n o tt d e p e n d on a b u n d a n c e in a simple way - expect its helium lines t o be relatively strong. To i n v e s t i g a t ee this possibility we have taken t h e E W ' s of the strongest H a n d He I lines in our s p e c t r u mm of V485 Cen (i.e. H a a n d He I 5876A) a n d compared these t o 153 values for 69 CV's byy Williams (1983).
F r o mm Fig. 8.5 we see t h a t t h e strength of the He I 5876A (relative t o H a ) is stronger in V485 C e nn t h a n in any of t h e systems studied by Williams (1983). A possible explanation for this resultt might be t h a t there is a strong absorption component in H a , as often seen in dwarf nova d u r i n gg o u t b u r s t s . F r o m t h e few R images w e have at the time of our spectroscopic observations wee o b t a i n R = 1 7 . 4 ( 4 ) m a g . Observed ( V - R ) colours of CV's are in the range 0.1-0.8 (e.g. E c h e v a r r i aa 1984). If we a s s u m e t h a t V485 C e n has a similar ( V - R ) colour, we derive t h a t given t h ee V m a g n i t u d e in quiescence (see Sect. 8.2.1) t h e source was at most slightly brighter t h a n d u r i n gg quiescence.
Forr all phases during a dwarf nova outburst it is found t h a t the emission is weaker t h a n the a b s o r p t i o nn t h e higher t h e q u a n t u m number of t h e Balmer lines (La Dous 1990). If we assume t h a tt n o He I 5876 A a b s o r p t i o n is present, we require a lower limit of ~ 3 0 A absorption in H a forr V 4 8 5 Cen t o have similar relative strength in these lines compared to other C V ' s . F r o m t h ee E W ' s of t h e B a l m e r lines we measured, and t h e shape of the s p e c t r u m itself, we find no i n d i c a t i o nn for t h e presence of such strong B a l m e r absorption features.
Itt could also he t h a t H a a b s o r p t i o n originates from t h e white dwarf primary, independent of t h ee o u t b u r s t s t a t e . However, we do not find any indication of t h e very wide and strong hydrogen a b s o r p t i o nn features expected in t h a t case.
Inn s u m m a r y , V485 Cen is a very special t y p e of CV. To determine the orbital period conclu-sivelyy we are planning t o perform time-resolved spectroscopy. Further p h o t o m e t r i c observations
References s 117 7 aree also planned.
AA cknowledgements
Wee thank V.S. Dhillon for supplying the 'PERIOD' analysis package, which we used for part off our data analysis. TA acknowledges support by the Netherlands Foundation for Research in Astronomyy (NFRA) with financial aid from the Netherlands Organisation for Scientific Research (NWO). .
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