PDF hosted at the Radboud Repository of the Radboud University Nijmegen


 Todd Williams
 6 months ago
 Views:
Transcription
1 PDF hosted at the Radboud Repository of the Radboud University Nijmegen The following full text is a preprint version which may differ from the publisher's version. For additional information about this publication click this link. Please be advised that this information was generated on and may be subject to change.
2 Mon. Not. R. A stron. Soc. 000, 111 (2010) Printed 15 O ctober 2010 (MN M ex style file v2.2) Kepler observations o f th e beam ing bin ary K P D arxiv: vl [astroph.sr] 13 Oct 2010 S. Bloemen1*, T. R, Marsh2, R, H. 0stensen1, S. Charpinet3, G. Fontaine4, P. Degroote1, U. Heber5, S. D. Ka.waler6, C. Aerts1,7, E. M. Green8, J. Telting9, P. Brassard4, B. T. Gänsicke2, G. Handler10, D. W. K urtz11, R, Silvotti12, V. Van Grootel3, J. E. Lindberg8,13, T. Pursim o8, P. A. Wilson8,14, R, L. Gilliland15, H. Kjeldsen16, J. ChristensenDalsgaard16, W. J. Borucki17, D. Koch17, J. M. Jenkins18, T. C. Klaus19 1 Instituut voor Sterrenkunde, Katholieke U niversiteit Leuven, Celestijnenlaan 200D, B3001 Leuven, Belgium 2 D epartm ent of Physics, University of Warwick, Coventry CV4 7AL, UK 3Laboratoire d Astrophysique de ToulouseTarbes, Université de Toulouse, CNRS, 14 Av. E. Belin, Toulouse, France 4 Départem ent de Physique, Université de Montréal, C.P. 6128, Succ. CentreVille, M ontréal, Québec H 3C 3J7, Canada 5 Dr. R em eissternw arte & E C A P Astronom isches Institut, Univ. Erlang ennürnberg, Sternw artstr. 7, Bamberg, Germ any 6 D epartm ent of Physics & Astronom y, Iowa State University, Am es, IA 50011, USA 7D epartm ent of Astrophysics, IM A PP, Radboud University Nijmegen, PO Box 9010, N L6500 GL Nijmegen, the Netherlands 8 Steward Observatory, University of Arizona, 933 N orth Cherry Avenue, Tucson, A Z 85721, USA 9 Nordic Optical Telescope, Santa Cruz de La Palma, Spain 10 In stitu t für Astronom ie, Universität Wien, Türkenschanzstrasse 17, 1180 Wien, A ustria 11 Jeremiah Horrocks Institute of Astrophysics, University of Central Lancashire, PR1 2HE, UK 12 IN AFO sservatorio Astronom ico di Torino, Strada dell Osservatorio 20, Pino Torinese, Italy 13 Centre fo r Star and P lanet Formation, Natural H istory Museum of Denmark, U niversity of Copenhagen, 0 s te r Voldgade 57, DK1350 Copenhagen, Denmark 14Institute of Theoretical Astrophysics, University of Oslo, P.O. Box 1029 B lindem, N0315 Oslo, Norway 15 Space Telescope Science Institute, 3700 San M artin Drive, Baltimore, M D 21218, USA 16D epartm ent of Physics and Astronom y, Aarhus University, D K Aarhus C, Denm ark 17N A SA A m es Research Center, M S , M offett Field, CA 94035, USA 18SE T I In stitu te/n A S A A m es Research Center, M S , M offett Field, CA 94035, USA 19 Orbital Sciences Corp., N A SA A m es Research Center, M S , M offett Field, CA 94035, USA Accepted 2010 August 18. Received 2010 August 17; in original form 2010 June RAS
3 2 S. Bloemen et al. A B ST R A C T The Kepler Mission has acquired 33.5 d of continuous oneminute photometry of KPD , a shortperiod binary system that consists of a subdwarf B star (sdb) and a white dwarf. In the light curve, eclipses are clearly seen, with the deepest occurring when the compact white dwarf crosses the disc of the sdb (0.4 per cent) and the more shallow ones (0.1 per cent) when the sdb eclipses the white dwarf. As expected, the sdb is deformed by the gravitational field of the white dwarf, which produces an ellipsoidal modulation of the light curve. Spectacularly, a very strong Doppler beaming (also known as Doppler boosting) effect is also clearly evident at the 0.1 per cent level. This originates from the sdb s orbital velocity, which we measure to be ± 1.9km s_1 from supporting spectroscopy. We present light curve models that account for all these effects, as well as gravitational lensing, which decreases the apparent radius of the white dwarf by about 6 per cent when it eclipses the sdb. We derive system parameters and uncertainties from the light curve using Markov Chain Monte Carlo simulations. Adopting a theoretical white dwarf massradius relation, the mass of the subdwarf is found to be 0.47 ± 0.03 Mq and the mass of the white dwarf 0.59 ± 0.02 M q. The effective tem perature of the white dwarf is ± 300 K. W ith a spectroscopic effective tem perature of Teg = ± 250 K and a surface gravity of log# = 5.43 ± 0.04, the subdwarf has most likely exhausted its core helium, and is in a shell He burning stage. The detection of Doppler beaming in Kepler light curves potentially allows one to measure radial velocities without the need of spectroscopic data. For the first time, a photom etrically observed Doppler beaming amplitude is compared to a spectroscopically established value. The sdb s radial velocity amplitude derived from the photom etry (1 6 8 ± 4 k m s~ 1) is in perfect agreement with the spectroscopic value. After subtracting our best model for the orbital effects, we searched the residuals for stellar oscillations but did not find any significant pulsation frequencies. K ey words: binaries: close  binaries: eclipsing  stars: individual (KPD )  subdwarfs  Kepler. 1 IN T R O D U C T IO N Subdwarf B stars are mostly assumed to be extrem e horizonta l branch stars, i.e., core helium burning stars w ith a thin inert hydrogen envelope (Heber 1986; Saffer et al. 1994). In order to reach such high tem peratures and surface gravities, the progenitor m ust have lost almost its entire hydrogen envelope. The m ajority of sdbs is expected to have lost its envelope via binary interaction channels, as elaborated by H an et al. (2002, 2003). Our target, KPD (KIC ), is a. subdw arf B star (sdb) w ith a white dwarf (WD) companion in a (8) day orbit (MoralesRueda et al. 2003), which identifies th e theoretical form ation channel for this system as the second commonenvelope ejection channel of Han et al. (2002, 2003). In this scenario the white dwarf is engulfed by the sdb progenitor as it ascends the first giant branch. The white dwarf will deposit its angular m om entum in the atm osphere of the giant and spin up the envelope until it is ejected. There are two subchannels to this scenario, depending on the initial mass of the progenitor. If sufficiently massive, it will ignite helium nondegeneratively, and the resulting extended horizontal branch (EHB) star will have a mass of ~ 0.35 M,t>. The more common scenario, starting w ith a roughly solarmass giant, produces an EHB star with a mass th a t must be very close to the helium flash mass of 0.47 M,t>. A third possibility occurs when the white * steven.bloem encsster.kuleuven.be dwarf companion ejects the envelope before the core has attained sufficient mass to ignite helium. In this case the remaining core will evolve directly to the white dwarf cooling track. On its way it crosses the domain of the EHB stars, b u t w ithout helium ignition the period in which it appears as an sdb star is brief, making this channel a very small contributor to the sdb population. For a recent extensive review on hot subdw arf stars, their evolution and observed properties, see Heber (2009). The exact physical details involved in commonenvelope ejection are not well understood. This uncertainty is commonly em bodied in the efficiency param eter a, which denotes the am ount of orbital energy used to eject the envelope (see e.g. de Kool 1990; Hu et al. 2007). Eclipsing subdwarf binaries could help constrain the perm itted values of a, but studies have hitherto been ham pered by the fact th a t both sdb +W D and sdb +M dw arf binaries have virtually invisible companions and are therefore single lined, leaving the masses indeterm inate. Firm ly establishing the param eters of both com ponents of a postce system therefore has substantial implications not just for confirming th a t our formation scenarios are correct, b u t also in order to tune future binary population synthesis studies by confining th e a param eter. The target studied here, KPD , is an sdb star discovered by the K itt P eak D ow n es survey (Downes 1986). KPD has a Fband m agnitude of (Allard et al. 1994), a ^m agnitude of ± (Wesemael et al. 1992) and a K p (K e p le r) m agnitude of
4 Kepler observations of the beaming binary K PD T he sta r was included in th e radial velocity survey of M oralesrueda et al. (2003), who found th e targ et to be a spectroscopic binary w ith a period of (8) d and a velocity am plitude K i = 167 ± 2 k in s  1. T hey also con cluded th a t th e sdb prim ary should be in a poste H B stage of evolution, due to its relatively low surface gravity placing it above th e canonical EHB in th e HR diagram. T his implies th a t th e sdb exhausted all available helium in its core and is now in a shell helium burning stage. Assuming th e sdb mass to be 0.5 M,t>, they found a m inim um mass of M,t> for th e com panion. In this pap er we present th e first light curve of K PD obtained from space. T he targ et was ob served for 33.5 d by th e Kepler Mission, and th e light curve reveals sufficient low level features to perm it purely p hoto m etric m easurem ents of velocities, radii and masses of b o th com ponents. A review of th e Kepler Mission and its first results is given in Koch et al. (2010). We com bine the Kepler photom etry w ith new and old spectroscopic m easurem ents, use light curve modelling to estim ate th e system param eters and M arkov C hain M onte Carlo (MCMC) sim ulations to establish th e uncertainties. T he relativistic D oppler beam ing effect is clearly detected in th e light curve, and can be used to determ ine th e orbital velocity of th e prim ary. T his effect, which is also known as D oppler boosting, was recently noted in a Kepler light curve of KOI74 by van K erkw ijk et al. (2010). We present th e first com parison of a radial velocity am plitude as derived from the am plitude of the D oppler beam ing to th e spectroscopi cally determ ined value. We also use th e spectroscopic d a ta to provide a revised ephemeris, as well as to determ ine the effective tem perature, surface gravity and helium fraction of the atm osphere. A fter detrending th e Kepler light curve w ith our best model for th e orbital effects, we search the residuals for stellar oscillations. 3 analysis is shown on Fig. 1. T he tim e span of th e d ataset is B M JD (T D B ) to T he high signaltonoise spectra from G reen et al. (2008) were used to derive T eg and log g\ see Section 4 for details. To refine th e orbital period d eterm ination of Mora.lesRueda. et al. (2003), sp ectra were collected w ith the 2.56m Nordic O ptical Telescope (NOT) on 2009 December 5, 9 and 10. Using th e A LFO SC spectrograph and a 0.5 arcsec slit, we obtained eleven spectra w ith exposure tim es of 300 s and a of resolution R ~ 2000, covering a wavelength region of Á. We m easured th e radial velocities using m ultig aussian fits (M oralesrueda et al. 2003); the values are listed in Table 1. We first fitted th e d a ta of Mora.lesRueda. et al. (2003) and th e new d a ta separately to determ ine th eir RM S scatter. From these fits we found th a t it was necessary to add 3.6 and 4.1 k m s  1 in q u ad ra tu re to th e uncertainties of th e two d atasets to deliver a u n it \ 2 Per degree of freedom; these values probably reflect system atic errors due to incom plete filling of th e slit. We th en carried out a leastsquares sinusoidal fit to th e combined d ataset finding a best fit of \ 2 = 21.6 (25 points). T he next best alias had \ 2 = 39, and so we consider our best alias to be th e correct one. This gave th e following spectroscopic ephemeris: B M JD (T D B ) = (62) (16) ;, m arking th e tim es when th e sdb is closest to E arth. T he corresponding radial velocity am plitude was A'i = ± 1.9 k in s  1. T he W D s spectrum will slightly reduce th e ob served velocity am plitude of th e sdb. We estim ate th a t this effect is less th a n l k m s  1. T he radial velocity m easure m ents and th e fit are shown in Fig. 2. From fitting light curve models to th e Kepler photom etry (see Section 3.4) we derived th e following photom etric ephemeris: B M JD (T D B ) = (25) (96) ;. 2 O BSER V A TIO N S, R A D IA L V E L O C ITIES A N D U P D A T E D E P H E M E R IS We used 33.5 d of Q1 short cadence Kepler d a ta w ith a tim e resolution of 59 s. A review of th e characteristics of the first short cadence d atasets is presented by Gilliland et al. (2010b). T he d a ta were delivered to us through th e K ASOC (Kepler Asteroseismic Science O perations C enter) w ebsite1. T he level of contam ination of th e fluxes by other stars is poorly known. We used th e raw fluxes and assum ed zero con tam ination, which is justified by th e absence of o ther signif icantly bright sources w ithin te n arcsec of K PD We applied a barycentric correction to th e Kepler tim ings and converted them from U T C to barycentric dynam ical tim e (TD B). T he raw d a ta show a. ~ 2 per cent downward tren d over the 33.5 d which we assume is instrum en tal in ori gin. We removed this variation by fitting and dividing out a spline function. O ut of th e original dataset of 49170, we rejected 54 points because of a bad quality flag. T hen, after initial light curve m odel fits to be described below, we re jected another 87 points because they differed by more th a n 3.5 a from our model. T he full light curve we used for our 1 h ttp ://k a s o c.p h y s.a u.d k /k a s o c / 2010 RAS, M N R A S 0 0 0, 111 T he two independent periods agree to w ithin th eir uncer tainties. T he 9 yr baseline of th e spectroscopic ephemeris gives a more precise value and henceforth we fix it at this value in our lightcurve models. T he cycle count betw een the spectroscopic and photom etric ephem erides (using th e spec troscopic period) is , an integer to w ithin th e uncertainties. T he Keplerbñsed zeropoint is th e more precise one and is therefore retained as a free p aram eter in our models. 3 LIG H T C U R V E A N A L Y SIS T he Kepler light curve we analyse in th is p aper (Fig. 1) re veals th a t K PD is an eclipsing binary. We g raph ically determ ined th e eclipse d ep th s and durations. T he eclipses of th e W D by th e sdb are per cent deep and th e eclipses of th e sdb by th e w hite dw arf per cent. T h e d u ratio n of th e eclipses at half m axim um d ep th is in orbital phase units. T here is a clear asym m etric ellipsoidal m odulation p a tte rn in which th e flux m ax im um after th e deeper eclipses is larger th a n th e m axim um 2 B M JD (T D B ) refers to B arycent.riccorrected M odified Ju lian D a te on th e B arycentric D ynam ical T im escale.
5 4 S. Bloemen et al. «3 QJ CO "cs d Ph Fh o z ; _ r r T i T  f ' j ' 1 ' ; r, i ; 1 " i " ' ! > 1 1 '  i' ', ' 1 i >. i 1 r \ i! '] i i i. 1 j... 1 : W 'V '7' ^ ; i... i... 1 'j :/ '1 '' 1' ' } ',, 1 j i V  },, 1 : ' l jr A.! î ' ' l r» f \ ' Ï i i i,* I. 1. T T.. i L [ I 1 1_ i i i i i 1 i 1. 1 i i i i i j i i i! i i i i. 1 1 i 1 i i i i I I _ r ' \ M " \ ' \ T! \ r Î! 1 ' I 1 b * L. i i i _ 1 1 * 5? j i i i i j i 1 1, I 1 i Time (d) F ig u re 1. Kepler light curve of K PD after detrending and rem oving outliers. I T able 1. R adial velocities of th e sdb in K PD determ ined from N O T spectroscopy. B M JD (TDB) RV (km s  1) ± ± ± ± ± ± ± ± ± ± ± 6.5 after the shallower eclipses. We attribute this to Doppler beaming, see Section 3.2. To determ ine the system properties, we modelled the light curve w ith the LCURVE code w ritten by TRM (for a description of the code, see Copperw heat et al. 2010, A p pendix A). This code uses grids of points to model the two stars, taking into account limb darkening, gravity darkening, Doppler beaming and gravitational lensing when the white dw arf eclipses its companion. It assumes the effipsoidalfy deformed star to be in corotation with the binary orbit, which is usually a good assum ption because of the large tidal interactions between the two binary components. Reprocessing of light from the sdb by the WD is included in the light curve models as well ( reflection effect ). To speed up the com putation of the models used in this paper we implemented a new option whereby a finer grid can be used along the track of the white dwarf as it ecfipses the sdb. This reduces the overall num ber of points needed to model the fight curve to the dem anding precision required to modef the K e p le r data. In addition, we only used the finefyspaced grid during the ecfipse phases, taking care to make the modef values continuous when changing between Orbital phase F ig u re 2. R adial velocity curve of th e sdb in K PD B oth our new radial velocity m easurem ents and th e ones from M oralesr ueda et al. (2003) are shown, folded on th e orbital period. T he error bars of th e datapoint.s show th e uncertainties after adding 3.6 k m s ~ 1 and 1.1 km s in q u adratu re to th e values of M oralesr ueda et al. (2003) and our new d a ta respectively. We find a radial velocity am plitude of A'i = ± 1.9 km s 1. grids by applying norm alisation factors (very close to unity) to the coarse grid fluxes. We used ~ (~ ) grid points for the fine (coarse) grids for the sdb and 3000 for the white dwarf. To modef the finite exposures more accurately during the ecfipses, where smearing occurs due to the 1 m integration time, we calculated 7 points for each exposure (i.e., one point for every ~ 1 0 s) and took their trapezoidal average. 3.1 G ravity darkening and limb darkening coefficients We used model spectra to com pute the gravity darkening coefficient (GDC) of th e sdb and the fimb darkening coefficients for both the sdb and the white dwarf, which are all important param eters for the modeling of a close binary s light curve. The GDC is needed to model the effects of the white dw arf s gravity on the sdb, which slightly distorts the sdb s shape. The bolometric flux from a stellar surface depends on the focal gravity as T 4 oc g l3b in which /3f, is the bolo
6 Kepler observations of the beaming binary K PD m etric GDC. For radiative stars, 3b = 1 (von Zeipel 1924). We observe th e bandlim ited stellar flux, not th e bolom etric flux, and hence we require a different coefficient defined by I (X g^k. T he GDC for th e K e p le r bandpass, 3k was com puted from K d^ log ^I d lo g g d log I d log g d log I d log T 9 log T d lo g g (1 ) in which I is th e photonw eighted bandpassintegrated spe cific intensity at fj, = 1 and = We used a grid of sdb atm osphere models calculated from th e LTE m odel atm osphere grid of H eber et al. (2000) using th e Linfor program (Lemke 1997) and assum ed T eg = K, l o g g = 5.5, l o g (?? H e / n H ) = 1.5 and l o g ( Z / Z,i ) = 2. To estim ate th e interstellar reddening, we com pared th e ob served B V colour of 0.20±0.01 mag (Allard et al. 1994) w ith th e colours expected from a m odel atm osphere. We found an intrinsic colour of B V = 0.26 m ag and conse quently adopted a reddening of E ( B V ) = To account for this interstellar reddening th e m odel spectra were red dened following Cardelli et al. (1989). T he gravity darkening coefficient was found to be 3k = Using a model for the same set of param eters, we com p uted limb darkening coefficients for th e sdb. We adopted th e 4param eter limb darkening relation of C laret (2004, equation 5) and determ ined a i = 0.818, an = 0.908, a.3 = and a.4 = For th e w hite dwarf, angledependent model spectra were calculated using th e code of Gansicke et al. (1995) for Xefr = K (estim ated from a com parison of model sur face brightnesses given initial light curve fits) and log g = 7.8. We adopted th e same limb darkening law as for the sdb and found a.\ = 0.832, an, = 0.681, a.3 = and a 4 = D oppler b eam ing factor T he asym m etry in K PD s ellipsoidal m odulation p a tte rn is th e result of D oppler beam ing. D oppler beam ing is caused by th e sta rs radial velocity shifting th e spectrum, m odulating th e photon emission rate and beam ing th e pho tons som ew hat in th e direction of m otion. T he effect was, as far as we are aware, first discussed in Hills & Dale (1974) for ro tatio n of w hite dwarfs and by S hakura & Postnov (1987) for o rbital m otion in binaries. It was first observed by M axted et al. (2000). Its expected detection in K e p le r light curves was suggested and discussed by Loeb & G audi (2003) and Zucker et al. (2007). Van K erkw ijk et al. (2010) report th e detection of D oppler beam ing in th e long ca dence K e p le r light curve of th e binary KOI74. For th e first tim e, they m easured th e radial velocity of a binary com po nent from th e photom etrically detected beam ing effect. T he m easured radial velocity am plitude, however, did n o t m atch th e am plitude as expected from th e mass ratio derived from th e ellipsoidal m odulation in th e light curve. T he derived velocity of th e prim ary of KOI74 is yet to be confirmed spectroscopically. For K PD , radial velocities are available which allows th e first spectroscopic check of a pho tom etrically determ ined radial velocity. For radial velocities th a t are much smaller th a n the speed of light, the observed flux F \ is related to th e em itted flux F 0 a a s 2010 RAS, M N R A S 0 0 0, 111 Fx = F0,x 1  B c 5 (2) w ith B th e beam ing factor B = 5 + d in F a / d i n A (Loeb & G audi 2003). T he beam ing factor th u s depends on th e spectrum of th e sta r and th e w avelength of th e obser vations. For th e b roadband K e p le r photom etry, we use a p h o to n weighted bandpassintegrated beam ing factor f e \ X F \ B dx (B) = ƒ exxfx dx (3) in which e \ is th e response function of th e K e p le r bandpass. We determ ined th e beam ing factor from a series of fully m etal lineblanketed LTE models (Heber et al. 2000, see also Section 3.1) w ith m etallicities ranging from log (Z /Z,~,) = 2 to + 1, as well as from N LTE models w ith zero m etals and w ith B lanchette m etal com position (see Section 4 of th is p a per for more inform ation ab o u t th e NLTE models). W ith o u t taking reddening into account, th e beam ing factor is found to be ( B ) = 1.30 ± T he u n certainty incorporates the dependence of th e beam ing factor on th e model grids and th e u n certainty on th e sd B s effective tem p eratu re, gravity and, m ost im portantly, metallicity. T he m etal com position of th e m odel atm ospheres is a poorly known factor th a t can only be constrained w ith highresolution spectroscopy. T his tim e, th e effect of reddening has to be accounted for by changing th e spectral response accordingly instead of reddening th e model atm osphere spectrum. Using a red dened spectrum would in th is case erroneously imply th a t th e reddening is caused by m aterial th a t is comoving w ith th e sdb star. W ith reddening, th e beam ing factor is d eter m ined to be ~ lower. R eddening th u s only m arginally affects th e beam ing of K PD b u t should certainly be taken into account in case of higher reddening values. T here are three contributions to th e beam ing factor. T he enhanced p hoton arrival rate of an approaching source contributes + 1 to th e beam ing factor. A berratio n also in creases th e num ber of photons th a t is observed from an approaching source, adding + 2 to th e beam ing factor be cause of th e squared relation betw een norm al angle and solid angle. Finally, w hen th e sdb comes tow ards us, an observed w avelength A0 corresponds to an em itted wave length Ae = A0 ( l + wr /c). Since sdbs are blue, looking at a longer w avelength reduces th e observed flux which coun teracts th e oth er beam ing factor com ponents. In case of an infinite tem p eratu re R ayleighjeans spectrum th is D oppler shift contribution to th e beam ing factor would be 2. For th e prim ary of K PD , we find a contribution of ~ 1.70 which brings th e to ta l beam ing factor to ~ T he contribution of th e D oppler shift does not always have to be negative; a red spectrum could actually increase the effect of beaming. 3.3 Light curve m odel A typical fit to th e d a ta is shown in Fig. 3, w ith th e different contributions sw itched on, onebyone. From th e residuals (b o tto m panel) it is clear th a t th e model reproduces the variations at th e orbital period very well. W hen we fit the light curve outside th e eclipses w ith sine curves to represent th e reflection effect, ellipsoidal m odulation and th e beaming, th e phase of th e ellipsoidal m odulation is found to be off
7 6 S. Bloemen et al. by ± in orbital phase units. We do not know the origin of this offset, which also gives rise to the shallow structure th a t is left in the residuals. O ur best fits have X2 = for d ata points. The significance of the Doppfer beaming is obvious, and even the more subtfe gravitational lensing effect is very significant, although it cannot be independently deduced from the d ata since it is highly degenerate w ith changes in the white dw arf s radius and tem perature. One p art of the gravitational lensing is caused by light from the sdb th a t is bent around the white dwarf, effectivefy making the white dwarf appear smaller. The second lensing contribution is a magnification effect which is caused by the altered area of the sdb th a t is visible, given th a t surface brightness is conserved by the lensing effect. In the case of KPD , the first p art of the fensing is the most im portant. Lensing effects in compact binaries were discussed in e.g. Maeder (1973), Goufd (1995), M arsh (2001) and Agof (2002, 2003). Sahu & Gilliland (2003) explored the expected influence of microiensing effects on fight curves of com pact binaries and planetary systems observed by K epler. They found th a t the fensing effect of a typical white dwarf at 1 AU of a main sequence star will swamp the eclipse signal. A transit of a planet, which is of similar size but a lot less massive, can therefore easily be distinguished from an eclipse by a white dwarf. In the case of K PD , the separation of the two com ponents is a lot less. An eclipse is still seen, but with reduced depth. For the most likely system param eters, gravitational lensing reduces the eclipse depth by ~ 1 2 per cent which is equivalent to a ~ 6 per cent reduction of the apparent white dwarf radius. The effect of gravitational lensing is implemented in our light curve modelling code following M arsh (2001). 3.4 M arkov Chain M onte Carlo sim ulation The param eters which determ ine models can be fixed by minimisation of \ 2 If the signaltonoise is high, a quadratic approxim ation around the point of minimum \ 2 can lead to the uncertainties of, and correlations between, the bestfit param eters. The K e p le r d ata have superb signaltonoise, but owing to the very shallow depths of the eclipses the quadratic approxim ation does not work well. Strong correlations between several param eters play a significant roie in this probiem. The duration of the ecfipses essentially fix the scaled radius of the sdb star (which we take to be the prim ary) ri = R i/a, where a is the binary separation. The scaled radius is a function of orbital inclination i, r i = The depth of the eclipse of the sdb by the white dwarf fixes the ratio of radii rn r \ = R n /R 1, so rn is also a function of orbital inclination. The duration of the ingress and egress features provides an independent constraint on rn as a function of i, which can break the degeneracy. In this case, however, one is lim ited by a com bination of signaltonoise and the minutelong cadence which is not sufficient to resofve the ingress/egress features. Under these circumstances, a Markov Chain Monte Carlo (MCMC) m ethod can be very valuable. The MCMC m ethod allows one to build up a sequence of models in which the fitting param eters, which we denote by the vector a, have a probabifity distribution matching the Bayesian posterior probabifity of th e param eters given the data, P(a d). Orbital phase F ig u re 3. Phasefolded light curve of K PD (green, datapoint.s grouped by 30) and our best fitting m odel (black). In th e to p panel, only th e eclipses and reflection effects are m odelled. In th e second panel, ellipsoidal m odulation is added. In th e th ird panel, gravitational lensing is taken into account as well, which affects th e depth of th e eclipse at orbital phase 0.5. T he b ottom panels show th e full m odel taking into account D oppler beam ing and th e residuals (grouped by 30 in green and grouped by 600 in black). From long chains of models one can then calculate variances and plot confidence regions. The MCMC m ethod also has the side benefit of hefping w ith the minimisation which can become difficuft when param eters are highfy correlated. For d ata in the form of independent Gaussian random variables, this probability can be w ritten as P (a d ) (X P (a )e ~ x //2, (4) i.e., the product of one s prior knowledge of the model param eters and a factor depending upon the goodness of fit as expressed in \ 2 We impfemented the MCMC m ethod foffowing procedures along the line of Collier Cameron et al. (2007). We incorporated prior inform ation in two ways. In all cases we used our constraint K i = 164 ± 2 k m s~ 1. Using our own spectroscopic analysis and the results of MoralesRueda et al. (2003), we decided to put also the fol
8 where K i and Ti are the values in the current MCMC model under test. The period of the binary orbit was kept fixed at the spectroscopically determ ined value (see Section 2). The param eters th a t were kept free during the modelling are the scaled stellar radii R i/a. and R n /a, the mass ratio q, the inclination i, the effective tem perature of the WD Tn, the beaming factor and the zero point of the ephemeris. The radial velocity scale (which leads to the masses M i and M n ) and the effective tem perature of the prim ary Ti were included in the fits as well, but w ith the spectroscopically allowed range as a prior constraint. Note th a t the beaming factor (B ) is kept as a free param eter, which allows the code to fit the Doppler beaming am plitude while we constrain the allowed range of K i. By comparing the MCMC results for (B ) w ith the theoretical beaming factor, we can check if the beam ing am plitude is consistent w ith our expectations. As explained above, this led to param eter distributions w ith strong correlations between R i, R n, M i, M n, q, etc. The massradius relations for the two stars are shown in Fig. 4. The favoured param eters for the secondary (left panel) clearly show th a t it is a white dwarf and the allowed distribution nicely crosses the expected track of m assradius (solid green line) which we calculated from the zero tem perature relation of Eggleton (quoted in Verbunt & R appaport 1988), inflated by a factor of We estim ated this factor from the cooling models of Holberg & Bergeron (2006) for white dwarfs w ith a mass between 0.5 and 0.7 M,t> (white dwarfs w ith masses outside this range are ruled out by the massradius relation) using our preferred tem perature for the white dwarf of around K and assuming an envelope th a t consists of M h = 1 0 ~ 4 Mwd and Mg_e = 1 0 ~ 2 Mwd Given th a t the secondary is a white dwarf and given th a t the secondary s theoretical m assradius relation intersects the massradius distribution very nearly within the 1 a region, we also undertook MCMC runs where we added the prior constraint th a t the secondary had to m atch the white dwarf MR relation to w ithin an RMS of 5 per cent. This was added in exactly the same m anner as the K i and Ti constraints. The system param eters we derive from these MCMC runs are listed in Table 2. The m assradius relation after applying the constraint is shown in Fig. 5. Especially after applying the white dwarf massradius relation constraint, the sdb s m assradius relation fits perfectly with the one defined by the surface gravity derived from spectroscopy in Section 4. The correlation coefficients between the different param eters are given in Table 3. The binary s inclination, its mass ratio and the stellar radii are highly correlated. 3.5 V ariability in residuals One of the goals of the K e p le r M issio n is to allow detailed asteroseismic studies of pulsating stars. The asteroseismology programme is discussed in Gilliland et al. (2010a). For more inform ation about the search for pulsations in compact objects w ith K epler, see 0sten sen et al. (2010a). Kepler observations of the beaming binary K PD lowing constraint on the effective tem perature of the sdb: T able 2. Properties of K PD T he orbital period and Ti = ± 400 K. These two constraints were applied by th e effective tem p erature of th e sdb were derived from spectroscopy. T he other param eters are obtained by m odelling th e com puting the following modified version of \ 2 Kepler light curve. T he uncertainties on these values are deter 2 ( K i 164 \ 2 /T i \ 2 2 In(P(a d)) = x + ( 5 ) + ( 4C 0 ) (=) m ined by M CM C analysis, using th e prior constraint th a t th e w hite dw arf m assradius relation has to m atch th e Eggleton relation to w ithin 5 per cent RMS. P rim ary (sdb) Secondary (WD) Porb (d) (16) <? i (deg) R (R 0 ) M (M 0 ) T e S (K) ± ± 300 T able 3. C orrelation coefficients of th e different param eters th a t were varied in th e MCM C sim ulations, after applying th e Eggleton m assradius relation constraint. R n i T i T n <?  R i R n i T i T n The Fourier transform of the original light curve of an eclipsing close binary like K PD is highly contam inated by frequencies and their harmonics due to the binary orbit. Subtraction of a good model of the binary signatures of th e light curve allows one to get rid of this contam ination. Since a num ber of sdbs have been found to be multiperiodically pulsating (for a review on asteroseismology of EHB stars, see 0stensen 2009), we checked the residuals of the light curve for signs of pulsations. Using the analysis m ethod and significance criteria outlined in Degroote et al. (2009), 8 significant frequencies were found, which are listed in Table 4. ƒ2 is a known artefact frequency caused by an eclips T able 4. Significant, variability frequencies in th e residuals of the light curve of K PD T he signal to noise (S/N ) value was determ ined by dividing th e am plitude of th e peak by th e uncertainty on th e am plitude. Frequency (d ) A m plitude (jttmag) S /N h = h/2 h instrum ental ƒ = h! 4 h = 2/2 h = 3/2 ƒ instrum ental fl = 2 /orf, h
9 8 S. Bloemen et al. Mwd (Mo ) MSdB (M0) F ig u re 4. Massradius relations of the W D (left) and sdb (right). One tenth of our MCMC models are shown (black dots). The contour plots show the regions in which 68, 95 and 99 per cent of the models reside. The contours that are shown are somewhat artificially broadened by the binning process. The Eggleton massradius relation, inflated by a factor 1.08 to allow for the finite W D temperature (see text for details), is shown as a solid green track on the left plot. The massradius relation intersects the Eggleton relation very nearly within its l<r region. On the right panel the solid green line gives the massradius relation for log<? = 5.45, the dashed lines for lo g Q = 5.40 (left) and log<? = 5.50 (right) ^ Q K Mwd (Mo ) Mscib (M0 ) F ig u re 5. Figure equivalent to Fig. 4 but for an MCMC' run with a prior constraint that the white dwarf massradius relation has to match the Eggleton relation to within 5 per cent RMS. ing binary th a t was used as one of the fineguidance stars during Q1 (see Haas et al. 2010; Jenkins et al. 2010). Four other frequencies ( /i, f z, ƒ4 and /b) are related to fn. The highest frequency, f e, is related to the processing of the long cadence d ata (see Gilliland et al. 2010b). ƒ 7 is the first harmonic of the orbital frequency of KPD , which indicates th a t there is still a weak orbital com ponent left after subtracting our light curve model, fg is not related to any of the other frequencies and corresponds to a period th at is too long to arise from stellar pulsations of the WD or the sdb. If it is real, the signal might result from the rotation of the WD, or it might be due to a background star. The best candidate peak for p mode pulsations of the sdb is at /u.hz with an am plitude of 37 mag, but further d ata is needed to confirm th a t th e sdb is pulsating. From ground based data, 0stensen et al. (2010b) did not detect pulsations, w ith a lim it of 0.68 mma. This is consistent w ith the Kepler photometry. 4 S P E C T R O S C O P IC A N A LY SIS Low resolution high S /N spectra for K PD were taken w ith the B&C spectrograph at Steward O bservatory s 2.3m Bok telescope on K itt Peak, as part of a long term homogeneous survey of hot subdw arf stars (Green et al. 2008), in Septem ber and O ctober The spectrograph param eters, observational procedures, and reduction techniques were kept the same for all the observing nights. The 400/m m grating, blazed at 4889 A, gives a resolu
10 Kepler observations of the beaming binary K PD tion of R = 560 over the w avelength region A, when used w ith the 2.5 a.rcsec slit. T he spectra were taken during clear or m ostly clear conditions w ith integration tim es betw een 1050 and 1200 s. A pproxim ately 1000 bias and flatfield images were obtained for each run. T he d a ta reduc tions were perform ed using sta n d ard IR A F tasks, and each night was fluxcalibrated separately. T he individual spectra were crosscorrelated against a supertem plate to determ ine th e relative velocity shifts, and th en shifted and com bined into a single spectrum. A lthough th e resolution is ra th e r low, th e S /N is quite high: 221/pixel or 795/resolution element for the com bined spectrum. T he continuum fit to th e com bined fluxcalibrated spectrum were done w ith great care to select regions devoid of any weak lines, including expected unresolved lines of heavier elements. T he final K PD spectrum was fitted using two separate grids of NLTE models designed for sdb stars, in order to derive th e effective tem perature, surface grav ity and H e/h ratio. T he first set of models assum ed zero m etals, while th e second included an adopted d istribution of m etals based on th e analysis of FUSE spectra of five sdb stars by B lanchette et al. (2008), see also Van G rootel et al. (2010). From th e set of models w ithout m etals, we derive log g = 5.45 ± 0.04, Teff = ± 220 K and log(h e/h ) = A ssum ing th e B lanchette com position, we find log g = 5.43 ± 0.04, T es = ± 250 K and log(h e/h ) = 1.36 ± These results are in good agreem ent w ith lo g g = , Teff = K, log(h e/h ) = determ ined by M oralesrueda et al. (2003) and log g = , Tes = K determ ined by Geier et al. (2010) using different spectra and model grids. T he fit definitely improves when going from th e zerom etal solution (Fig. 6 ) to th e B lanchette com position (Fig. 7), although there still rem ains a slight B alm er problem, especially noticeable in th e core of H/3. T here are definitely m etals in th e spectrum of K PD : the strongest features are 1) an unresolved C III + N II complex around 4649 A (com pare th e two figures for th a t featu re), and 2) another weaker com plex (C III + O II) in th e blue wing of H<5 th a t th e B lanchette m odel reproduces quite well. All of th e m ajor discrepancies betw een th e spectra and th e models are due to strong interstellar absorption: th e K line of C a II in th e blue wing of He, the Ca II H line in th e core of He, and th e N a l doublet strongly affecting th e red wing of He I It is reassuring th a t th e derived atm ospheric p a ram eters are not too strongly dependent on th e presence of m etals, as m ight be expected for such a hot star, particularly one in which downwards diffusion of m etals is im portant. 5 D IS C U S S IO N T he beam ing factor we derived for K PD using M CM C runs is ( B ) = , which is in perfect agree m ent w ith th e theoretically expected value calculated in Sec tion 3.2. T he uncertainty on th e beam ing factor is a direct reflection of th e uncertainty on th e spectroscopic radial ve locity am plitude of th e sdb. If, contrary to our assum ption, th e K e p le r fluxes would be severely contam inated by light from other (constant) stars, th e observed beam ing factor would be lower. T he d istribution of beam ing factors from our M CMC com putations is shown in Fig. 8. If th e radial ve 2010 RAS, M N RA S 0 0 0, Beaming factor F ig u r e 8. D istrib u tio n of th e sd B s beam ing fa cto r for an MCMC' ru n (using th e M R c o n strain t for th e W D ; see te x t for details). T h e beam ing factor is found to be (B) = 1.33 ± 0.02, w hich is in agreem ent w ith th e th eo re tic ally expected ( B) = 1.30 ± locity would be m easured from th e D oppler beam ing am pli tude, using th e theoretical beam ing factor, th e result would be k m s _ 1 com pared to k m s  1 derived from spectroscopy. T he u n certainty on th e photom etric ra dial velocity is dom inated by th e u n certainty on th e th e o re t ical beam ing factor, prim arily due to its dependence on the poorly known m etallicity of th e sdb, and to a lesser extent due to th e uncertainties on th e sd B s effective tem p eratu re and surface gravity. U nder th e assum ption of corotation, we find a projected ro tatio n al velocity of th e sdb of i>sin(i) = k m s  1. From spectroscopy and using LTE models w ith te n tim es Solar metallicity, Geier et al. (2010) found i>sin(i) = l.o k m s 1, which is in agreem ent w ith our photom etric re sult. We conclude th a t th e assum ption of corotation is likely to be correct. T he spectroscopically determ ined surface gravity of the sdb (log g = and using atm osphere models w ith and w ithout m etals respectively) agrees per fectly w ith th e surface gravity of we derived from th e m assradius d istrib u tio n of our light curve models. As concluded earlier by M oralesrueda et al. (2003), th e sdb is probably in a posteh B phase. T his is illustrated in Fig. 9, which shows th e zero age extended horizontal branch (ZAEHB) and th e term inal age extended horizon ta l branch (TAEHB) for an sdb w ith a typical core mass of 0.47 M,7,, together w ith evolutionary tracks for different hydrogen envelope masses (10~4, 10~3, 2 x 10~3, 3 x 10~ 3 and 4 x 1 0 ~ 3 M (7,) from Kawa.ler & H ostler (2005). Because of its low surface gravity, th e sdb com ponent of KPD1946 falls in a region of th e Teg log g plane relatively far from th e center of th e instability strip. However, at least one pulsator exists in th is region of th e T eg log g plane, V338 Ser, th a t should be in a poste H B phase (see 0 ste n sen 2009, Fig. 3). Moreover, tran sien t pu lsato rs w ith varying pulsation am plitudes th a t can go down to undetectable val ues in a p articu lar epoch m ight exist (see th e case of K IC in 0 ste n sen et al. 2010a). For these reasons, and because we found at least one candidate p m ode pulsation frequency, it is w orth continuing a photom etric m onitoring by Kepler. T he w hite dw arf mass implies th a t it is a CO w hite dwarf. T he progenitor of th e sdb m ust have been th e less massive sta r in th e original binary and by th e tim e it reached th e ZAEHB, th e w hite dw arf was already cooling. T he ac
11 10 S. Bloemen et al. <D > cd 0) KPD Teff = K [ 221 K] log g = [0.042].lo g N(He)/N(H) = [0.047], l l l l NLTE: H, He AA (A) 50 F ig u re 6. Fit (bold lines) to the spectral lines of KPD using NLTE sdb models assuming zero metals. X plh CD > CD 0^ KPD Teff = K [ 248 K] log g = [0.040].lo g N(He)/N(H) = [0.043], l l l l NLTE: H, He, C /10, N, 0 /1 0, S i/1 0, S, Fe AA (A) 50 F ig u re 7. F it (bold lines) to th e spectral lines of K PD using NLTE sdb models assum ing a B lanchette m etal composition. cretion of m aterial by the white dwarf does not change the W D s internal energy content significantly (see e.g. related work on cataclysmic variables by Townsley & Bildsten 2002). The cooling tim e of the white dwarf therefore sets an upper lim it to the tim e since th e sdb was on the ZA EHB. For our best estim ates of the tem perature and mass of the white dwarf, the cooling tracks of Holberg & Bergeron (2006) indicate th a t it has been cooling for about 155 to 170 M yr (depending on the unknown envelope composition). The sdb s evolution from the Zero Age Extended Horizontal Branch to its current postehbphase took 125 to 145 Myr (Kawaler & Hostler 2005, depending on the exact current evolutionary stage), which means th a t the sdb must have formed very shortly after the white dwarf. 6 S U M M A R Y We have analysed a 33.5d short cadence K e p le r light curve of K PD , as well as low resolution spectroscopy. In the light curve, prim ary and seconday eclipses, ellipsoidal m odulation and Doppler beaming are detected. We model the binary light curve, taking into account the Doppler beam ing and gravitational lensing effects. System param eters and uncertainties are determ ined using Markov Chain M onte Carlo simulations. The binary is found to consist of a 0.59 ±0.02 M,t>white dwarf and a 0.47 ± 0.03 M,t>postEHB sdb star. The surface gravity and corotation rotational velocity of the sdb as derived from the light curve models are found to be consistent w ith spectroscopic values. The observed Doppler beaming
12 Kepler observations of the beaming binary K PD F ig u re 9. T he sdb of K PD in th e Teg lo g g plane. T he theoretical zero age and term inal age extended horizontal branches for a 0.17 \l. are shown, together w ith evolutionary tracks for different envelope thicknesses (10~4, 10 \ 2 X 10 \ 3 x 10~3 and 4 X 10~3 T he sdb is found to be in th e post EHB phase. am plitude is in perfect agreement w ith the am plitude expected from spectroscopic radial velocity measurements. It would thus have been possible to derive the radial velocity am plitude of the sdb from the Kepler light curve directly. Subtracting a good light curve model allowed us to search for stellar oscillations. No significant stellar variability of the sdb or white dwarf could be detected yet. At least one candidate p mode pulsation frequency was found, however, and the sdb can also possibly be a transient pulsator. K PD continues to be observed by K epler. A cknowledgm ents We thank the referee M artin van Kerkwijk for his helpfuf suggestions. The authors gratefully acknowledge everybody who has contributed to make the K e p le r M issio n possible. Funding for the K e p le r M issio n is provided by NASA s Science Mission Directorate. P art of the d ata presented here have been taken using ALFOSC, which is owned by the Instituto de Astrofísica de Andalucía (IAA) and operated at the Nordic O ptical Telescope (Observatorio del Roque de los Muchachos, La Palm a) under agreem ent between IAA and the NBIfAFG of the Astronom i cal Observatory of Copenhagen. This research also made use of d ata taken w ith the Bok telescope (Steward Observatory, K itt Peak). The research leading to these results has received funding from the European Research Council under th e European C om m unity s Seventh Fram e work Program m e (FP7/ )/ER C grant agreement n (PR O SPERITY ), as well as from the Research Councif of K.U.Leuven grant agreement GOA/2008/04. During this research TRM was supported under grants from the U K s Science and Technology Facilities Council (STFC, ST/F002599/1 and PP/D /1). R EFER ENCES A gol E., 2002, A pj, 579, 430, 2003, A pj, 594, 449 Allard F., Wesemael F., Fontaine G., Bergeron P., Lamontagne R., 1994, A J, 107, 1565 (K) Blanchette J., Chayer P., Wesemael F., Fontaine G., Fontaine M., Dupuis J., K ruk J. W., Green E. M., 2008, A pj, 678, 1329 Cardelli J. A., Clayton G. C., M athis J. S., 1989, A pj, 345, 245 Claret A., 2004, A&A, 428, 1001 Collier Cameron A., Wilson D. M., West R. G., et al., 2007, MNRAS, 380, 1230 Copperwheat C. M., M arsh T. R., Dhillon V. S., Littlefair S. P., Hickman R., Gänsicke B. T., Southworth J., 2010, MNRAS, 402, 1824 de Kool M., 1990, A pj, 358, 189 Degroote P., A erts C., Ollivier M., et al., 2009, A&A, 506, 471 Downes R. A., 1986, A pjs, 61, 569 Gänsicke B. T., Beuerm ann K., de M artino D., 1995, A&A, 303, 127 Geier S., Heber U., Podsiadlowski P., Edelm ann H., Napiwotzki R., Kupfer T., Mueller S., 2010, ArXiv: Gilliland R. L., Brown T. M., ChristensenDalsgaard J., et al., 2010a, PASP, 122, 131 Gilliland R. L., Jenkins J. M., Borucki W. J., et al., 2010b, ApJ, 713, L160 Gould A., 1995, A pj, 446, 541 Green E. M., Fontaine G., Hyde E. A., For B., Chayer P., 2008, ASPC, 392, 75 Haas M. R., B atalha N. M., Bryson S. T., et al., 2010, ApJ, 713, L115 Han Z., Podsiadlowski P., M axted P. F. L., M arsh T. R., 2003, MNRAS, 341, 669 Han Z., Podsiadlowski P., M axted P. F. L., M arsh T. R., Ivanova N., 2002, MNRAS, 336, 449 Heber U., 1986, A&A, 155, 33, 2009, ARA&A, 47, 211 Heber U., Reid I. N., W erner K., 2000, A&A, 363, 198 Hills J. G., Dale T. M., 1974, A&A, 30, 135 Holberg J. B., Bergeron P., 2006, AJ, 132, 1221 Hu H., Nelemans G., 0stensen R., A erts C., Vuckovic M., Groot P. J., 2007, A&A, 473, 569 Jenkins J. M., Caldwell D. A., Chandrasekaran H., et al., 2010, A pj, 713, L120 Kawaler S. D., Hostler S. R., 2005, A pj, 621, 432 Koch D. G., Borucki W. J., Basri G., et al., 2010, ApJ, 713, L79 Lemke M., 1997, A&AS, 122, 285 Loeb A., Gaudi B. S., 2003, A pj, 588, L117 Maeder A., 1973, A&A, 26, 215 M arsh T. R., 2001, MNRAS, 324, 547 M axted P. F. L., M arsh T. R., N orth R. C., 2000, MNRAS, 317, L41 M oralesrueda L., M axted P. F. L., M arsh T. R., N orth R. C., Heber U., 2003, MNRAS, 338, 752 0sten sen R., Sifvotti R., C harpinet S., et al., 2010a, MN RAS, in press 0stensen R. H., 2009, Communications in Asteroseismology, 159, 75 0stensen R. H., Oreiro R., Solheim J., Heber U., Silvotti R., GonzálezPérez J. M., Ulla A., Pérez Hernández F., RodríguezLópez C., Telting J. H., 2010b, A&A, 513, A6 Saffer R. A., Bergeron P., Koester D., Liebert J., 1994, ApJ, 432, 351 Sahu K. C., Gilliland R. L., 2003, A pj, 584, RAS, M NRAS 0 00, 111