Mon. Not. R. Astron. Soc. 332, 34–36 (2002) Radial velocity study of the pulsating subdwarf B star and possible Type Ia supernova progenitor KPD 193012752 Vincent M. Woolf,1P C. Simon Jeffery1 and Donald L. Pollacco2 1Armagh Observatory, College Hill, Armagh BT61 9DG 2Queens University Belfast, Belfast BT7 1NN Accepted 2001 December 11. Received 2001 December 10; in original form 2001 November 30 Downloaded from https://academic.oup.com/mnras/article/332/1/34/974706 by guest on 05 October 2021 ABSTRACT We report the results of a high-time-resolution radial velocity study of the subdwarf B star and possible Type Ia supernova progenitor KPD 1930þ2752. There were no significant peaks in the power spectrum of the velocity curve above our detection limit, about 4 km s21, at the frequencies where peaks arising from pulsation were present in the photometric data of previous researchers. We report an orbital velocity amplitude, 348:5 ^ 1kms21, in agreement with that reported by previous investigators. We find an orbital period of P ¼ 0:095 093 08 ^ 0:000 000 15 d based on our data and the ephemeris of Maxted et al. Key words: binaries: spectroscopic – stars: individual: KPD 193012752 – subdwarfs – stars: variables: other. measured velocity and brightness variations will be useful in 1 INTRODUCTION identifying the pulsation modes so that asteroseismology can be Subdwarf B (sdB) stars are a class of evolved hot stars with used to make accurate determinations of the stellar parameters. The physical properties that place them between the main sequence photometric amplitudes of the KPD 1930þ2752 pulsations are and the white dwarf track in a colour–magnitude diagram. Since similar to those seen in other sdBV stars for which radial velocity the discovery that the sdB star EC 1402622647 pulsates (Kilkenny variations were measured, so it seemed likely that similar velocity et al. 1997), the number of known pulsating sdB stars (EC 14026 or measurements would be possible for this star. sdBV stars) with published data has increased to 19 (O’Donoghue et al. 1999; Piccioni et al. 2000; Bille`res et al. 2000; Silvotti et al. 2 OBSERVATIONS AND REDUCTION 2000; Østensen et al. 2001). KPD 1930þ2752 is a subdwarf B star (Downes 1986). It was We obtained spectra of KPD 1930 þ 2752 at a high time resolution identified as a pulsating sdB (EC 14026 or sdBV) star which varies (12.1-s repetition) using the 4.2-m William Herschel Telescope at multiple frequencies by Bille`res et al. (2000). They also found (WHT) on the nights of 2000 June 20 and 21. We used the blue arm that the largest amplitude photometric frequency is best explained of the ISIS spectrograph with the R1200B grating. The charge- as ellipsoidal variations caused by a close binary orbit coupled device (CCD) was read out in drift mode (Rutten et al. ðP ¼ 8217:8sÞ. This was confirmed by Maxted, Marsh & North 1997). Observations were made over a period of 7.2 h starting : (2000), who found a spectroscopic orbit with the same period as the at HJD ¼ 2451716 43394 and a period of 5.2 h starting at : ellipsoidal variation. The large orbital velocity and a typical sdB HJD ¼ 2451717 42469, with short breaks for CuAr arc calibration ¼ mass imply a minimum total binary mass of 1:47 ^ 0:01 M(.An spectra. Spectral resolution was l/Dl 5000 and the spectral orbital inclination smaller than 908 would imply a larger mass. range was 4020 , l , 4430 A. The exposure timings come from Maxted et al. suggest that the characteristics of the binary make it a the clock of the CCD controller, which is synchronized daily with good candidate to become a Type Ia supernova. In a recent report, the GPS-calibrated time service of the observatory. The clock Ergma, Fedorova & Yungelson (2001) say that their evolution normally drifts by up to 1 s during the night. models indicate that the most likely result of the eventual merger The data were reduced using scripted standard IRAF routines to will be a high-mass white dwarf, but with a mass below the make the bias, flat-field and sky corrections, extract the one- Chandrasekhar mass. dimensional spectra, and apply the wavelength calibration from In previous work it has been possible to measure the velocity CuAr arc spectra. variations of sdBV stars resulting from pulsation (Jeffery & Pollacco 2000; O’Toole et al. 2000; Woolf et al. 2002). The 3 ANALYSIS AND RESULTS Our first attempt to determine radial velocities from the spectra and PE-mail: [email protected] search for periods in the velocity variations followed the method q 2002 RAS Radial velocities in KPD 1930þ2752 35 outlined by Jeffery & Pollacco (2000) using IDL programs. The to a 108-s phasing difference over the 6 d covered by the data of large velocity variations caused by orbital motion introduced Bille`res et al. The difference is probably the result of the improved problems with the normally straightforward determination of precision possible with the longer time baseline, more than 64 d the velocity shifts from cross-correlation peaks. For the small- (.680 orbits), rather than a change in orbital period. amplitude velocity variations present with other sdBV stars The orbit of the stars should be decaying as a result of we defined a fitting window of about 10 pixels (about 280 km s21). gravitational radiation at a rate predicted by The velocity variations of KPD 1930þ2752 move the cross- correlation peak outside this window. Simply using a much larger 96 PG 3 M 2m P_ ¼ 2 ¼ 21:7 £ 10210 dyr21; window does not solve the problem because fits to the peak become 5 c 5 a 4 much less precise when they are affected by points far from the peak centre. We found best results by shifting a 35-pixel where M is the total mass of the binary, m is the reduced mass, and window to follow the approximate velocity shifts and ran several a is the separation of the centres of mass of the stars (i.e. the sum of iterations until the velocity amplitude found for the orbital motions their distances from the centre of mass of the system) (Shapiro & matched the amplitude introduced in the fitting window. Using Teukolsky 1983), and we use M1 ¼ 0:5M(, M2 ¼ 0:96 M( and Downloaded from https://academic.oup.com/mnras/article/332/1/34/974706 by guest on 05 October 2021 5 this method, we found a projected orbital velocity amplitude of a ¼ a1 þ a2 ¼ 6:92 £ 10 km. It should be possible to measure 359 ^ 8kms21. the effect of the period change on cycle timings after a few Because of worries about the automatic fitting routine, we used decades: a 1-min cycle arrival timing difference owing to orbital the IRAF routine FXCOR to do the cross-correlation ‘manually’. We decay should be present after about 46 yr. The period uncertainty applied the velocity corrections found with the automatic routine to will need to be reduced by a factor of 100 for such a measurement the individual spectra and co-added them to use as the cross- to be possible. Other effects, e.g. magnetic fields, will also correlation template. Using FXCOR allowed us to eliminate bad probably affect the orbit. pixels more reliably and discard spectra where the cross- The Fourier transform power spectrum of the velocity data with correlation function peak looked especially bad. The radial the orbital motion removed (Fig. 2) is shown in Fig. 3. We did not velocities thus found are shown in Fig. 1. A sine wave was fitted unambiguously detect peaks in the power spectrum above our to the velocity curve thus produced using the downhill simplex detection limit, about 4 km s21, with frequencies at which Bille`res 2 program AMOEBA (Press et al. 1992) to minimize the x difference. et al. (2000) reported photometric variations caused by stellar Using this method we found a projected orbital velocity amplitude pulsation. KPD 1930þ2752 is fainter ðB ¼ 13:75Þ than the sdBV of 348 ^ 1kms21. We believe that this amplitude is more reliable stars for which high-speed spectroscopy using 4-m telescopes was than the one found using our automatic fitting routine. The velocity able to detect velocity variations arising from pulsation: PG curve with the orbital velocity fit subtracted is shown in Fig. 2. 1605þ072 ðB ¼ 13:01Þ, KPD 2109þ4401 ðB ¼ 13:17Þ and Because of the noise in our data, we cannot improve upon the PB 8783 ðB ¼ 12:5Þ (Jeffery & Pollacco 2000; Woolf et al. orbital period determined by Bille`res et al. (2000) using our data 2002). The non-detection is probably due to higher observational alone. However, by including the T0 found by Maxted et al. (2000) noise, not smaller pulsation amplitudes: comparing the photo- an improvement in accuracy is possible. We fitted a sine wave metric amplitudes arising from pulsation in KPD 1930þ2752 with to our (FXCOR-derived) data assuming the period from Bille`res et al. the photometric and velocity amplitudes of KPD 2109þ4401 and We then determined the period using the phasing we found and PB 8783, stars with comparable T and log g, indicates that 21 T0 ¼ 2451651:6466 from Maxted et al. ðP ¼ Dtn , where n is we should expect maximum pulsational velocity amplitudes in the number of orbits during the time Dt ). We then fitted a sine KPD 1930þ2752 of between 2 and 3 km s21. However, since wave to our data assuming the period found in the first iteration. period and amplitude changes are often observed for pulsations of After four iterations the input and output periods matched.