Communications in Asteroseismology

Volume 141 January, 2002

Editor: Michel Breger, Turk¨ enschanzstraße 17, A - 1180 Wien, Austria Layout and Production: Wolfgang Zima and Renate Zechner Editorial Board: Gerald Handler, Don Kurtz, Jaymie Matthews, Ennio Poretti http://www.deltascuti.net

COVER ILLUSTRATION: Representation of the COROT Satellite, which is primarily devoted to under- standing the physical processes that determine the internal structure of stars, and to building an empirically tested and calibrated theory of stellar evolution. Another scientific goal is to observe extrasolar planets.

British Library Cataloguing in Publication data. A Catalogue record for this book is available from the British Library. All rights reserved ISBN 3-7001-3002-3 ISSN 1021-2043 Copyright °c 2002 by Austrian Academy of Sciences Vienna Contents

Multiple frequencies of θ2 Tau: Comparison of ground-based and space measurements by M. Breger 4 PMS stars as COROT additional targets by M. Marconi, F. Palla and V. Ripepi 13 The study of pulsating stars from the COROT exoplanet field data by C. Aerts 20 10 Aql, a new target for COROT by L. Bigot and W. W. Weiss 26 Status of the COROT ground-based photometric activities by R. Garrido, P. Amado, A. Moya, V. Costa, A. Rolland, I. Olivares and M. J. Goupil 42 Mode identification using the exoplanetary camera by R. Garrido, A. Moya, M. J.Goupil, C. Barban, C. van’t Veer-Menneret, F. Kupka and U. Heiter 48 COROT and the late stages of stellar evolution by T. Lebzelter, H. Pikall and F. Kerschbaum 51 Pulsations of Luminous Blue Variables by E. A. Dorfi and A. Gautschy 57 By-product of the V1162 Ori multisite campaign: the new δ Scuti star GSC 4778 00324 by P. Lampens, P. van Cauteren, P. Niarchos, K. Gazeas, V. Manimanis, T. Arentoft, P. Wils, A. Bruch, R. Garrido and R. Shobbrook 65 Strømgren photometry of SX Phe = HD 223065 by A. Stankov, D. Sinachopoulos, E. Elst and M. Breger 72 Comm. in Asteroseismology Vol. 141, 2002

Multiple frequencies of θ2 Tau: Comparison of ground-based and space measurements

M. Breger

Institut fur¨ Astronomie, Univ. Wien, Turk¨ enschanzstr. 17, A–1180 Wien, Austria

Abstract

The satellite photometry of the δ Scuti star θ2 Tau obtained with the Wide- Field Infrared Explorer (WIRE) led to the detection of 12 frequencies of pulsation (Poretti et al. 2002). We examine the available ground-based photometry in the literature to check whether these frequencies are also present. The re-analysis of the 1986 Delta Scuti Network data leads to 10 statistically significant frequen- cies (amplitude signal/noise ratio greater than 4.0). 9 of these 10 frequencies are in excellent agreement with those detected from space, while the 10th fre- quency is seen as a peak in the WIRE residuals, though not at a statistically significant level. Previous reports on amplitude variability are confirmed.

Introduction

A number of lengthy observing campaigns covering individual δ Scuti variables have shown that the majority of these pulsating variables on and near the main sequence pulsate with a large number of simultaneously excited nonradial p modes. Furthermore, long-term variability of the pulsation amplitudes of these nonradial modes with a time scale of years has been discovered for many, but not all of these nonradial pulsators. The Delta Scuti Network specializes in multisite observations of these stars. The network is a collaboration of astronomers located at observatories spaced around the globe in order to avoid regular, daily observing gaps. So far, 21 campaigns have been carried out. The most recent campaign covered BI CMi, for which 29 frequencies of pulsation derived from 1024 hours (177 nights) of photometry were discovered (Breger et al. 2002). The variability of the star θ2 Tau was discovered by Horan (1977, 1979) and confirmed by Duerbeck (1978). In order to study the multiple frequencies M. Breger 5 in this star, the Delta Scuti Network undertook two campaigns (Breger et al. 1987, 1989). The data could be supplemented by additional measurements by Kovacs & Paparo (1989). Altogether, five frequencies of pulsation were detected and interpreted to be due to mainly nonradial pulsation. The size and distribution of the frequencies suggested p modes with values of ` = 0 to 2, since higher ` values would not be seen photometrically due to cancellation effects across the stellar surface. A mode with a higher ` value was also detected spectroscopically by Kennelly & Walker (1996). Furthermore, the Delta Scuti Network reobserved θ2 Tau during 1994, but the data are at present unpublished because of unresolved instrumental difficulties at a few of the participating sites (in the former Soviet republics). Furthermore, satellite data have also now become available (see below). θ2 Tau is a 140.728 d binary system with known orbital elements (Ebbig- hausen 1959, Torres, Stefanik & Latham 1997). The two components have similar temperatures. The primary component is evolved (A7III), while the secondary is fainter by 1.10 mag (Peterson, Stefanik & Latham 1993) and still on the main sequence. Both stars are situated inside the instability strip. It was shown by Breger et al. (1989) that the dominant pulsation modes originate in the primary component because the predicted orbital light-time effects for the primary match the observed shifts in the light curves of up to several minutes. This conclusion agrees with the expectations from the values of the observed frequencies (12 – 15 cd−1). These are compatible with those expected from an evolved star (i.e., the primary) and incompatible with those expected for main-sequence δ Scuti stars (i.e., the secondary star in the binary system). We note here that for frequency analyses, the light-time corrections in the binary system need to be applied to all data covering more than a few weeks.

The 2000 satellite data

During 2000 August, θ2 Tau was monitored extensively with the star camera on the Wide-Field Infrared Explorer satellite (WIRE). This remarkable, pioneering study from space led to the discovery of 12 independent frequencies of pulsation down to the 0.5 mmag level (Poretti et al. 2002). The satellite data are free of the 1 cd−1 aliasing often present in ground-based data, but include aliasing at the orbital frequency of 15 cd−1. The 12 detected frequencies include the 5 frequencies previously found from the ground, but with different amplitudes. The difference in the effective wavelengths between the space and ground- based data can be expected to lead to different amplitudes between the studies. However, the relative strengths between the modes also changed severely, which confirms the existence of amplitude variability in θ2 Tau. Poretti et al. also report a frequency of 26.19 cd−1 and regard this frequency 6 Multiple frequencies of θ2 Tau as a spurious term caused by the length of the duty cycle. This explanation may be too cautious: we find that some of the terrestrial data also show this variability with an identical frequency value. This may be due to pulsation of the fainter secondary. The pulsation of the secondary will be examined in detail at a future time and is consequently ignored in the rest of this paper (which concentrates on the main frequency region of variability.)

The 1982 – 1986 ground-based photometric data

The question arises whether these frequencies determined from the satellite data are also present in the ground-based data. A re-examination of the old data is in order for an additional reason: The published multifrequency analyses were performed at a time at which experimentally determined limits for the extraction of frequencies from multisite photometric data were unavailable and the statistical criteria for the acceptance or rejection of additional modes were less developed. Inspection of Figs. 3 and 4 in Breger et al. (1989) confirms that the published five-frequency solution does remove almost all the power in the power spectra. Could there exist additional statistically significant peaks? One of the most important questions in the examination of multiperiodicity concerns the decision as to which of the detected peaks in the power spectrum can be regarded as variability intrinsic to the star. Due to the presence of nonrandom errors in photometric observations and because of observing gaps, the predictions of standard statistical false-alarm tests give answers which are considered by us to be overly optimistic. In a previous paper (Breger et al. 1993) we have argued that a ratio of amplitude signal/noise = 4.0 provides a useful criterion for judging the reality of a peak. This corresponds to a power signal/noise ratio of 12.6. Subsequent analyses comparing independent data sets have confirmed that this criterion is an excellent predictor of intrinsic vs. possible noise peaks, as long as it is not applied to very small data sets or at low frequencies, where the errors of measurement are far from random. In the present study, the noise was calculated by averaging the amplitudes (oversam- pled by a factor of 20) over 5 cd−1 regions centered around the frequency under consideration. The following data are available: 1982: 10 nights, 1983: 27 nights (some multisite), 1985: 3 nights, 1986: 17 nights (many multisite). New pulsation frequency analyses of the available data in various annual and multiyear combinations were performed with a package of computer programs with single-frequency and multiple-frequency techniques (programs PERIOD, Breger 1990, PERIOD98, Sperl 1998). These programs utilize Fourier as well as multiple-least-squares algorithms. The latter technique fits a number of simultaneous sinusoidal variations in the magnitude domain and does not rely M. Breger 7 on prewhitening. For the purposes of presentation and initial searches, however, prewhitening is required if the low-amplitude modes are to be seen. Therefore, in the presentation of the results (see below), the various power spectra are presented as a series of panels, each with additional frequencies removed relative to the panel above. The results of the multiperiodicity analyses of the 1982–1986 data can be summarized as follows: (i) Little additional information beyond the published five-frequency solution can be extracted out of the 1982/1983 data. (ii) Considerable more information is available in the 1986 November data due to the multisite nature and the extremely high accuracy of the data. (iii) Adding the three nights from 1985 (Kovacs & Paparo 1989) to the 1986 data does not lower the noise level in the power spectrum of the combined data nor improve the significance level of the detected frequencies. (iv) Due to the different quality of the annual data sets and possible ampli- tude variability, the combined 1982–1986 solution shows higher noise than the 1986 data alone. We therefore select the 1986 data (only) for a more detailed re-analysis.

Multiple frequencies present in the 1986 photometric data

During 1986 November, coordinated photoelectric measurements of the star θ2 Tau were obtained at four observatories on three continents. The cam- paign represented the 3rd campaign of the Delta Scuti Network. The following telescopes were used: (i) the 0.6 meter telescope of Sierra Nevada Observatory, Spain (observer R. Garrido), (ii) the 0.6 meter reflector of Xinglong Station of Beijing Observatory, China (observers Huang Lin, Jiang Shi-yang and Guo Zi-he), (iii) the 0.9 meter telescope of McDonald Observatory, Texas, USA (ob- server M. Frueh) (iv) the 0.5 meter reflector at Piszk´estet˝o, the mountain station of Konkoly Observatory, Hungary (observer M. Paparo). Photomultiplier detectors were used together with V or Str¨omgren y filters. The spectral window of the data is very clean, because the measurements were obtained on three continents. In particular, the 1 c/d aliasing, common in ground-based measurements, is quite small (see Fig. 1). The power spectrum of the 1986 θ2 Tau data was computed from frequency values of zero to the Nyquist frequency. Here we have already applied the orbital light-time corrections for the primary, although this correction is not critical for the 21 d duration of the observations. The highest power levels of the stellar 8 Multiple frequencies of θ2 Tau

-3 -2 -1 0 1 2 3

1 Spectral Window

0

Ground-based f1 data of θ2 Tau (1986) 40 f2 2

f3 g

a 20 m i l l i m

n i

r 0 e

w f4 o

P Data - 3f

4 (3 frequencies prewhitened)

2

0 10 11 12 13 14 15 16 Frequency in cycles per day

Figure 1: Power spectra of the 1986 measurements of θ2 Tau. The top panel shows the spectral window, while the bottom panel presents the results after prewhitening the three dominant frequencies found in the middle panel. M. Breger 9

Data - 4f f5 f6 1.0 f7

0.5

0.0 Data - 7f 0.6 f8 f9

0.4 2

g 0.2 a m i l l

i 0.0 m

n 0.4

i f10

r Data - 9f e w o

P 0.2

0.0

Data - 10f Significance level

0.2

0.0 10 11 12 13 14 15 16 Frequency in cycles per day

Figure 2: Power spectrum of the 1986 measurements of θ2 Tau after successive prewhitening of the previously detected frequencies. The four panels show the results after prewhitening 4, 7, 9 and 10 frequencies, f1 to f10, determined in previous panels. The curve in the bottom panel represents the statistical significance limit. Peaks below the limit are rejected. 10 Multiple frequencies of θ2 Tau

Table 1: The frequency spectrum of θ2 Tau Ground-based (1986) WIRE (2000) Frequency Amplitude Frequency Amplitude cd−1 millimag S/N cd−1 millimag

f1 13.23 6.5 44 13.23 2.5 f2 13.69 5.0 33 13.70 6.5 f3 13.48 2.5 17 13.49 1.3 f4 14.32 2.3 15 14.32 3.2 f5 12.17 1.2 8.2 12.13 0.6 f6 14.61 1.1 7.5 14.61 0.7 f7 12.84 1.0 6.9 12.83 1.5 f8 11.74 0.8 5.7 11.73 0.6 - - 11.77 1.5 f9 12.38 0.8 5.4 12.40 0.8 f10 13.81 0.7 4.5 - - - - 10.86 1.1 - - 12.70 0.7

data are found in the 10 to 16 cd−1 region. Since this is also the frequency range in which the frequencies from the satellite photometry were detected, we concentrate on this region. Fig. 1 also shows that the four dominant frequencies are found at 13.23, 13.69, 13.48 and 14.32 cd−1. The four modes are old friends and were already found in previous analyses.

Fig. 2 shows the power spectra after the dominant four frequencies were prewhitened through a simultaneous four-frequency least-squares optimization to the data. These solutions optimize the frequency values, amplitudes, phases of ten sinusoids and determine the overall zero-point. The analysis was repeated after prewhitening multifrequency solutions containing more frequencies. Alto- gether, 10 frequencies exceeding the significance limits could be detected. For the five frequencies known to high precision from the previous Delta Scuti Net- work campaigns, we have used the previous frequency values.

The adopted frequencies and amplitudes are shown in Table 1. The new ten-frequency solution fits the observed data well and leads to residuals of only §2.0 mmag per single observation. M. Breger 11

Comparison between the ground-based and satellite results

The re-analysis of the 1986 ground-based photometry has increased the number of frequencies from 5 to 10. Table 1 shows the excellent agreement between the results of the ground-based and satellite data sets: 9 frequencies are identical. Two frequencies found in both data sets deserve additional comments: −1 (i) The values determined for f5 are 12.17 and 12.13 cd , respectively. The difference of 0.04 cd−1 is not significant, since both data sets are ∼ 20 d long (i.e., 1/T ∼ 0.05 cd−1). (ii) The WIRE satellite photometry led to the detection of a close frequency pair at 11.73 and 11.77 cd−1. The ground-based data only reveals only a single −1 mode, f8, at 11.74 cd . Without additional data, we are unable to resolve the problem. Although the frequency pair is near the limit of frequency resolution, one cannot automatically reject the result without additional tests (e.g., see the discussion on close frequencies in δ Scuti stars by Breger & Bischof 2002). It is conceivable that due to an accidentally favorable phasing of the two modes during one of the two data sets, the pair was detected in only that data set. Amplitude variability of one of the two modes in the pair is also a possible explanation. There exist three other frequencies, which are found at statistically signifi- cant levels in only of the two data sets: (i) The 1986 data reveals an additional mode at 13.81 cd−1, which has not been seen before. The power spectrum of the WIRE residuals shows a peak at 13.83 cd−1 (see Fig. 6 of Poretti et al. 2002) and is therefore probably also present in that data set. (ii) The mode at 12.70 cd−1 found by WIRE may also be present in the 1986 ground-based data. A peak is seen at that value, but below the level of significance. We have tested the possibility that the small amount of 1 cd−1 aliasing with the peak at 11.74 cd−1 might affect the amplitude. Multifrequency solutions with both or even three frequencies (11.73, 11.77 and 12.70 cd−1) fail to raise the amplitude significantly. (iii) The mode at 10.86 cd−1 found by WIRE is not present in the 1986 data, although it is near a small cluster of power in the residuals (see bottom panel of Fig. 2).

Conclusion

The re-analysis of the previously published photometry obtained from observa- tories situated on the ground leads to excellent agreement with the results found by the WIRE satellite: 9 frequencies are detected independently in both data sets at a high level of statistical significance. The 9 modes show considerable 12 Multiple frequencies of θ2 Tau amplitude variability between 1986 and 2000.Two additional modes are present with statistically significant amplitudes in one data set and definite peaks in the power spectrum of the other data set. The WIRE mode at 10.86 cd−1 was not seen in 1986. Amplitude variability is the most probable explanation.

Acknowledgments. This investigation has been supported by the Austrian Fonds zur F¨orderung der wissenschaftlichen Forschung under project number P14546-PHY. It is a pleasure to acknowledge helpful discussions with Ennio Poretti.

References Breger, M. 1990, Communication in Asteroseismology (Vienna), 20, 1 Breger, M. 1993, in ’Stellar Photometry - Current Techniques and Future Developments’, ed. Butler, C. J., Elliott, I., Cambridge University Press, 106 Breger, M. 2000, ASP Conf. Ser., 210, 1 Breger, M., Bischof, K. 2002, A&A, in press Breger, M., Lin, H., Jiang, S.-Y., et al. 1987, A&A 175, 117 Breger, M., Garrido, R., Lin, H., et al. 1989, A&A 214, 209 Breger, M., Stich, J., Garrido, R., et al. 1993, A&A 271, 482 Breger, M., Handler, G., Garrido, R., et al. 1999, A&A 349, 225 Breger, M., Garrido, R., Handler, G., et al. 2002, MNRAS 329, 531 Duerbeck, H. W. 1978, IBVS, 1412, 1 Ebbighausen, E. G. 1959, Pub. Dom. Astrophys. Obs. 11, 235 Horan, S. 1977, IBVS 1232, 1 Horan, S. 1979, AJ 84, 1770 Kennelly, E. J., Walker, G. A. H. 1996, PASP 108, 327 Kovacs, G., Paparo, M. 1989, MNRAS 237, 201 Peterson, D. M., Stefanik, R. P., Latham, D. W. 1993, AJ 105, 2260 Poretti, E., Buzasi, D., Laher, R., et al. 2002, A&A 382, 157 Sperl, M. 1998, Communications in Asteroseismology (Vienna) 111, 1 Torres, G., Stefanik, R. P., Latham, D. W. 1997, ApJ 485, 167 Comm. in Asteroseismology Vol. 141, 2002

PMS stars as COROT additional targets

M. Marconi1, F. Palla2, V. Ripepi1

1Osservatorio Astronomico di Capodimonte, Via Moiariello 16, 80131 Napoli, Italy 2Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy

Abstract

The importance of pre-Main Sequence (PMS) δ Scuti studies is discussed and recent developments both from the theoretical and observational point of view are reviewed. The open problems in the current status of knowledge of these young pulsators are mainly connected to the limitations of short time baseline ground observations. COROT could definitely help to solve these problems by allowing accurate and continuous observations. A list of possible PMS δ Scuti candidates for a COROT additional program has been identified and exploratory ground-based observations are planned.

Introduction

PMS stars with masses M ≥ 1.5M¯, called Herbig Ae/Be stars, cross the pulsation instability strip during their contraction toward the main sequence. The estimated crossing times range from ∼ 0.05 τKH (1.5M¯) to ∼ 0.1 τKH (4M¯), where τKH is the Kelvin-Helmholtz timescale. In spite of the relatively short crossing time, a number of Herbig stars have the appropriate values of colors to fall within the instability strip boundaries of the well known evolved variable stars. The existence of pulsating stars among intermediate Herbig stars was origi- nally suggested by Breger (1972) who discovered two PMS pulsator candidates in the young open cluster NGC 2264. This initial finding was confirmed by sub- sequent observations of a δ Scuti-like pulsation in the Herbig Ae stars HR5999 (Kurtz & Marang 1995) and HD104237 (Donati et al. 1997). This evidence stimulated the first theoretical investigation of the PMS in- stability strip, based on nonlinear convective hydrodynamical models (Marconi 14 PMS stars as COROT additional targets

& Palla 1998). As a result, the topology of the PMS instability strip for the first three radial modes was identified and a list of possible PMS pulsating candidates was provided on the basis of spectral types. New observational studies were then devoted to the search for δ Scuti-type photometric variations with periods of minutes to several hours and amplitudes less then few tenth of magnitudes among Herbig stars inside or close to the theoretical instability boundaries. Kurtz & Muller (1999) reobserved HD104237 and confirmed its photomet- ric variability; Marconi et al. (2000) identified two new PMS pulsating stars, namely HD35929 and V351 Ori; Pigulski et al. (2000) discovered the pulsating nature of two members (BL50 and HP57) of NGC 6823; Kurtz & Catala (2001) confirmed the δ Scuti pulsation and periodicity of HR5999; Kurtz & Muller (2001) identified the PMS δ Scuti star HD142666; Marconi et al. (2001) pro- vided new observations and frequency information on V351Ori; finally, Pinheiro et al. (2001) observed the pulsational variation of the luminosity of V346 Ori. In order to properly understand the nature and intrinsic properties of PMS δ Scuti stars, an extensive comparison of the observed pulsation properties with theoretical predictions is required (e.g. Marconi et al. 2000, 2001). In particular, the comparison between empirical and theoretical periods and period- ratios is expected to provide important constraints on stellar mass, whereas the comparison of the position in the color-magnitude diagram suggested by observations and/or the pulsational analysis with PMS and post-MS tracks may allow to constrain the evolutionary state.

Position of PMS δ Scuti stars in the HR diagram

The main characteristics, including the observed frequencies, of the ten known or suspected δ Scuti stars are reported in Table 1. For most of these objects, the reported periodicities are affected by the uncertainty due to the 1 day alias, so that they may not represent the actual periods. Moreover, for some objects the observations are too scarce and the errors too large to allow accurate frequency determinations. However, let’s assume that the published frequencies are indeed correct and compare them with the predictions of linear non adiabatic pulsation models (see Marconi et al. 2000, 2001 and references therein). Such a comparison allows us to constrain the position of observed stars in the HR diagram, as shown in Fig. 1. In fact, at fixed solar composition, the mass, luminosity, and effective temperature are varied in order to simultaneously satisfy the observed periodicities and the PMS evolutionary constraint. It is worth noting that the solution of such a fitting procedure is unique M. Marconi, F. Palla and V. Ripepi 15

Table 1: Known or suspected PMS δ Scuti stars. Star F1 F2 V (mag) Sp. Type V351Ori 15.49§0.23 11.89§0.24 8.9 A7 V346Ori 34.2§4.9 21.2§3.8 10.1 A5 HD104237 33.0§0.2 8.1 A5 HR5999 4.81§0.01 7.0 A7 HP57 12.7256§0.0002 15.5244§0.0003 14.6 BL50 13.9175§0.0005 9.8878§0.0009 14.5 V588 Mon 9.09§0.82 9.75 A7 V599 Mon 8.06§0.39 10.3 F2 HD142666 21.43§3.00 8.8 A8 HD35929 5.10§0.13 8.1 A5

Figure 1: Position in the HR diagram of observed PMS δ Scuti stars, as predicted on the basis of linear nonadiabatic computations. The PMS evolutionary tracks by Palla & Stahler 1993 and the nonlinear instability strip by Marconi & Palla (1998), for the three lowest radial modes, are shown for comparison. only for candidates for which several periodicities are available from observa- tions (V351 Ori, V346Ori, HP57, BL50). In fact, two relations are needed to 16 PMS stars as COROT additional targets

Table 2: The position in the HR diagram and the pulsation mode(s) of PMS δ Scuti stars. The modes are indicated with F (Fundamental), FO (First Overtone), SO (Second Overtone), TO (Third Overtone), FoO (Fourth Overtone), FiO (Fifth Overtone). Star M/M¯ log L/L¯ log Te Source mode(s) V351Ori 1.8 1.142 3.866 Marconi et al. 2001 F,FO V346Ori 1.5 0.74 3.863 Pinheiro et al. 2001 F,SO HD104237 2.2 1.5 3.93 van den Ancker 1998 FiO HR5999 4.0 2.12 3.845 Marconi & Palla 1998 SO HP57 2.0 1.25 3.857 this paper FO/SO BL50 2.5 1.6 3.86 this paper SO/FoO V588 Mon 3.5 2.05 3.913 this paper TO V599 Mon 2.5 1.51 3.84 this paper FO HD142666 1.7 1.03 3.88 Natta et al. 1997 FO HD35929 3.4 1.92 3.857 Marconi et al. 2000 FO HD35929 3.8 2.06 3.851 Marconi et al. 2000 SO properly disentangle the luminosity and effective temperature contributions to periods, the mass being determined via comparison with the evolutionary tracks. In case of a single observed periodicity, the comparison with the nonlinear in- stability strip (e.g. for HR5999, see Marconi & Palla 1998) and/or independent information on the intrinsic parameters (e.g. for HD104237 and HD142666) are taken into account and used to remove degeneracies. The predicted pulsa- tion modes for all the selected candidates are reported in Table 2 together with references for the luminosity and effective temperature values. For HD35929, the solutions found by Marconi et al. (2000) are reported (see also Fig. 1). We also notice that a certain ambiguity on the evolutionary state holds for stars close to the MS (e.g. V351 Ori). In this region of the HR diagram the PMS and post-MS tracks are not well separated at fixed stellar masses. In such cases a nonradial pulsation analysis could help solve the problem, as nicely discussed by Suran et al. (2001). As a final comment on Fig. 1, it is interesting to note that only two stars (V588 Mon and HD104237) bluer than the second overtone blue edge of the theoretical instability strip are actually predicted to pulsate in overtones higher than the second one.

How COROT could help

The 1 day alias problem affecting most of the observed PMS δ Scuti stars could be easily removed with continuous observations. If these observations are made from space, the additional uncertainty due to weather conditions is M. Marconi, F. Palla and V. Ripepi 17

Table 3: PMS δ Scuti candidates for COROT observations. Star RA(J2000) Dec(J2000) V Sp. Type RNO63 06 07 52.41 -05 16 04.4 13.3 F6e LKH338 06 10 47.05 -06 12 47.9 15.1 F2e LKH339 06 10 54.43 -06 14 39.4 13.6 F2e V588Mon 06 39 05.90 09 41 03.4 9.75 A7 V589Mon 06 39 28.45 06 42 04.8 10.3 F2 VSB2 06 39 41.55 09 34 40.2 13.33 F7 W121 06 40 56.50 09 54 10.4 10.8 F8e V360Mon 06 41 06.18 09 36 22.9 13.29 F8e VV Ser 18 28 47.96 00 08 39.7 11.87 B,Ae Par21 19 29 00.81 09 38 38.7 14.16 A5 also removed. In the case of COROT, a high photometric precision is also foreseen, such that accurate period measurements and the identification of low amplitude pulsations will be possible.

PMS δ Scuti candidates for COROT observations

According to the adopted Scenario 4 of the COROT (Center=102.5 deg, i.e. RA=6h 50min, see Minutes of the 6th Scientific Committee) mission, the two possible fields in the Galactic anticenter and center directions are those reported in Figs. 2 and 3. On the basis of these fields, we have identified a list of PMS candidates (open circles in Figs. 2 and 3) by searching in the Herbig catalogue (Th´e, Winter & Perez 1994). The coordinates, V-magnitudes, and spectral types of the selected candidates are listed in Table 3. Unfortunately, none of the PMS candidates is located close to the main targets of the asteroseismological project (filled circles in Figs. 2 and 3). For this reason, we propose these objects as the targets of a COROT additional program. The first three objects of the list (RNO63, LKH338, LKH339) are so close to each other that they can be observed within a single COROT field. The fourth and the fifth stars are the two PMS δ Scuti stars already discovered by Breger (1972) in NGC2264 (also shown in the figure). We notice that these two stars are the brightest ones in the list. The three following objects (VSB2, W121, V360 Mon) are again observable within a single field. Most of these candidates are too faint (V>9 mag) to be observed in the Astroseismology program configuration and too bright (V<16.5 mag) for the Exoplanet program setup. However, fainter magnitudes than V=9 mag are probably measurable in the Astroseismological configuration even by loosing a bit in precision (1 part over 10,000 is enough for our goals). 18 PMS stars as COROT additional targets

Figure 2: Position of PMS δ Scuti candidates (open circles) in the anticenter direction COROT field. The main targets of the asteroseismological project (filled circles) are shown for comparison.

Figure 3: As in Fig. 2 but in the Galactic center direction

Up to now, only V588 Mon and V589 Mon are known to be δ Scuti-like variable stars, but the search for variability in the other candidates by means of ground based observations is planned. M. Marconi, F. Palla and V. Ripepi 19

Conclusions

With the recent development of both observational and theoretical studies, PMS δ Scuti stars have become an important tool for understanding the structure and the evolutionary properties of young intermediate-mass stars. The number of known or suspected pulsating PMS objects has significantly increased during the last few years, but for most of them the detected periodicities are quite uncertain due to alias problems and/or poor data number. Continuous and accurate observations from space, as those expected with the advent of the COROT mission, would definitely overcome these problems. In this paper, we have proposed PMS δ Scuti stars as the targets of a COROT additional program and provided a list of suitable candidates.

Acknowledgments. We kindly thank the COROT working group on PMS δ Scuti stars: C. Catala, J.M. Alcal´a, E. Antonello, S. Bernabei, M. Breger, E. Covino, D.F.M. Folha, M.J. Goupil, L. Mantegazza, E. Paunzen, F.J.G. Pinheiro, E. Poretti, W. Weiss, F. Zerbi, K. Zwintz.

References Breger, M. 1972, ApJ 171, 539 Donati, J.-F., Semel, M., Carter, B. D., Rees, D. E., Cameron, A. C. 1997, MNRAS 291, 658 Kurtz, D. W., Catala, C. 2001, A&A 369, 981 Kurtz, D. W., Marang, F. 1995, MNRAS 276 191 Kurtz, D. W., Muller,¨ M. 1999, MNRAS 310, 1071 Kurtz, D. W., Muller,¨ M. 2001, MNRAS 325, 1341 Marconi, M., Palla, F. 1998, ApJ 507, L141 Pigulski, A., KoÃlaczkowski, Z., Kopacki, G. 2000a, AcA 50, 113 Marconi, M., Ripepi, V., Alcal`a, J. M., Covino, E., Palla, F., Terranegra, L. 2000, A&A 355, L35 Marconi, M., Ripepi, V., Bernabei, S., Palla, F., Alcal`a, J. M., Covino, E., Terranegra, L. 2001, A&A 372, L21 Natta, A., Grinin, V. P., Mannings, V., Ungerechts, H. 1997, ApJ 491, 885 Palla, F., Stahler S. W. 1993, ApJ 418, 414 Pinheiro, F. J. G., Marconi, M., Ripepi, V., Folha, D., Palla, F. 2001, IAU Coll. 185, in press Suran, M., Goupil, M., Baglin, A., Lebreton, Y., Catala, C. 2001, A&A 372, 233 Th´e, P. S., de Winter, D., Perez, M. R. 1994, A&AS 104, 315 van den Ancker, M. E., Th´e, P. S., de Winter, D. 1996, A&A 309, 809 Comm. in Asteroseismology Vol. 141, 2002

The study of pulsating stars from the COROT exoplanet field data

C. Aerts

Instituut voor Sterrenkunde, Catholic University of Leuven, Belgium

Abstract

In this paper we present our proposal to search for periodically variable stars in the exoplanet field data. In order to do so one first has to perform a quasi- automated frequency search in the data. Once the periodically variable stars are found, they have to be classified, e.g. by means of a statistical scheme analoguous to the one we developed for the classification of new B-type variables discovered from the HIPPARCOS data. After the classification of the variable stars, we plan to study the β Cep stars and the slowly pulsating B stars (SPBs) with the goal to derive constraints on their internal structure.

Goal of the project

In this project we propose to make full use of the COROT exoplanet field data for asteroseismology. The observation strategy adopted by the planet hunters is ideally suited for asteroseismology of some classes of pulsators, i.e. this project is a free by-product of the search for planets within the COROT framework. The individual integration times of 32 seconds for the CCD frames of the exoplanet field will in any case result in a time sampling of the order of minutes during 5 consecutive months. This is largely sufficient to derive very accurately the pulsation frequencies for several types of κ driven modes in main-sequence stars, such as the β Cep stars, the slowly pulsating B stars (hereafter termed SPBs), the γ Doradus stars and most δ Scuti stars. The expected photometric accuracy will be at least an order of magnitude higher than the one of current ground- based datasets of such types of stars. COROT is by far better suited to perform our project than MOST and MONS. MOST will observe only a few very bright stars, while the MONS C. Aerts 21

Field Monitor will gather data for a few thousand stars, which will be observed during one month. COROT is therefore the only mission resulting in data with a total time base significantly longer than a month. Particularly, main-sequence gravity-mode pulsators, i.e. γ Doradus stars and SPBs, cannot be studies by MOST and MONS, as they have beat periods of several months or even years. The concrete goal of our project is twofold: 1) the search for, and classifi- cation of, new variable stars; 2) the derivation of detailed frequency spectra of the new variables and their physical interpretation.

Search for new variables

The discovery of new variables in the COROT exoplanet fields will be quite straightforward in view of the excellent time sampling. Such high temporal resolution is beyond comparison with e.g. the one of the HIPPARCOS mis- sion. The long time base of the data of the latter mission already allowed the discovery of thousands of new periodically variable stars, despite the relatively poor sampling (Eyer, 1998). This led to the publication of the HIPPARCOS catalogue of periodic variable stars (van Leeuwen et al. 1997, Grenon et al. 1997). A considerable amount of new variables will be found among the 30 000 stars which will all be monitored during 5 months by COROT. The applica- tion of classical period search algorithms will result in the detection of many new periodically variable stars of different kinds. In view of the high-temporal resolution and the high duty cycle, Fourier analysis will suffice to perform the frequency search.

Classification of the new variables: discriminant analysis

The newly discovered periodic variables need to be classified, i.e. the origin of their variability has to be decided upon. Waelkens et al. (1998) have developed a classification scheme for the new HIPPARCOS B-type variables. Their method is based upon multivariate discriminant analysis. With this method, one makes use of different known parameters of the stars. One of these parameters is of course the frequency. Besides this quantity, Waelkens et al. (1998) used the Geneva parameters X, Y, Z. These parameters are reddening-free quantities derived from the Geneva colours. Since the COROT exoplanet data will contain three colours, the method developed by Waelkens et al. is also applicable to the COROT data after minor modification. It was also already successfully applied by Aerts et al. (1998) in their search for new γ Doradus candidates from the HIPPARCOS data. 22 The study of pulsating stars from the COROT exoplanet field data

The essential ingredients of discriminant analysis are the following: In gen- eral, multivariate classification schemes use the knowlegde of stellar parameters in an n-dimensional space. The classification of new variables is best done in two steps: the assignment of a new variable to a known class according to a distance measure and subsequently a check through bivariate plots. Discrim- inant analysis is only possible when prototypes of classes are available, such that a “definition” of a known class can be formulated. The unknown ob- jects will then be assigned to one of these so-called calibration classes. This assignment is obtained by the determination of discriminating axes in the mul- tidimensional space, in such a way that optimal separation of the predefined classes is attained. For each star the Mahalanobis or generalised distance to the discriminating axes is used for assignment to one of the definition classes. However, stars with a relatively large Mahalanobis distance still might have a different origin of variability than those of the class to which they were as- signed. A definite classification is then made by means of bivariate plots within the n-dimensional parameter space. The stars with a large Mahalanobis dis- tance having also another position in bivariate plots compared to the calibration stars are probably of a different nature and cannot be assigned to one of the calibration classes. A typical example could be a Be star, a CP star or a su- pergiant. For these stars it is often not possible to provide a clear definition of the class as there are e.g. very different timescales in the variability of Be stars, there are supergiants with and without microvariations, etc. It is better not to use them as calibration classes. All new variables of such kind will then at first be assigned to a wrong class but their different nature will be derived from the large Mahalanobis distance and/or a deviating position in the bivariate plots. Before applying the discriminant analysis one needs to choose the parame- ters to be used for the definition of the classes. In the framework of our proposal these are the frequency and the three COROT colours. This implies that we have to determine the COROT colours for the known variables of the different types of pulsators, in order to find the definition of each class.

Interpretation of the variability

A next step in our project consists of the detailed interpretation of the frequen- cies in terms of asteroseismology. In order for this to be possible, we are obliged to limit ourselves to the multiperiodic variables with correct mode identification. The latter remains a serious problem for κ driven modes outside the asymptotic regime. As we will be dealing with faint stars having long beat periods, spec- troscopic mode identification methods will not be of any help. The best way to identify modes from frequency spectra is in that circumstance the detection of rotationally- or magnetically-induced frequency multiplets. C. Aerts 23

Figure 1: Frequency analysis of the simulated signal for the typical β Cep star with a time span of 20 days (top) and of 5 months (bottom). The dashed lines indicate the input frequencies.

We have performed some preliminary simulations in order to derive the ca- pabilities of COROT with respect to the detection of multiple periods. In view of our expertise we have limited the simulation study to pulsating B stars, i.e. β Cep stars and SPBs. We have simulated a periodic signal typical for a β Cep star and an SPB respectively. For none of these stars the mode selection mech- anism is understood and we have no idea how many modes are excited. The WIRE data of β Crucis have clearly indicated that more modes are excited than thought up to now from ground-based photometry in this β Cep star (Cuypers et al. 2002). This is also the case for the δ Scuti star θ2 Tauri (Poretti et al. 2002). Taking these recent results into account we have selected 10 frequencies out of the pool of observed frequencies for the two classes of stars for our simu- lation. For the prototypical β Cep star it concerns frequencies between 3 and 6 c d−1 and for the SPB between 0.3 and 0.8 c d−1. We have included observed doublet and triplet frequencies. The amplitudes were taken between 50µmag and 20mmag, which is the maximum amplitude observed for such stars. No noise was added at this stage. 24 The study of pulsating stars from the COROT exoplanet field data

Figure 2: Frequency analysis of the simulated signal for the typical SPB with a time span of 20 days (top) and of 5 months (bottom). The dashed lines indicate the input frequencies.

We generated light curves with the following two different time samplings : 1. every 3 minutes during 20 days, 2. every 20 minutes during 5 months. Such sampling is quite conservative as COROT should do better. The Fourier spectra for the simulated signals of the β Cep star, respectively SPB, are shown in Figs. 1 and 2. We are particularly interested to see whether COROT would be able to result in the detection of the multiplets. This will not be the case for the data obtained during the exploratory phases, as can be derived from the upper panels of Figs. 1 and 2. The runs of 5 months do result easily in the derivation of the multiplet frequencies of the β Cep star, which is very important. This will imply that the COROT data will in principle allow us to derive the internal rotation law in stars that will end their lives as type II supernovae. Most of the close frequencies of the SPB are also recovered, albeit not all of them. This corresponds to the well-known fact that the beat periods of such stars can be of the order of a year. At least 7 of the 10 frequencies are easily recovered, though, in our exercise. C. Aerts 25

Of course, our very basic simulations have to be redone, by convolving the generated signal with a realistic noise-budget valid for COROT. This will be done in a forthcoming study. However, the simple exercise allows us to derive the following conclusions : • the time span of the COROT exploratory programme is too short to disentangle the frequencies of pulsating B stars; • rotational splittings will easily be found for β Cep stars in the SISMO core programme, as the sampling is much better than the one we have taken in this study; the core programme will also allow the detection of a very large fraction of the excited gravity modes of SPBs; • the exoplanet data are ideally suited for detecting the larger-amplitude modes of β Cep stars and of SPBs. All the conclusions for the SPB-type pulsations are also valid for the gravity modes in γ Doradus stars. We did not include such stars in the simulation as the observational inventory available for them does not reach the same level of completeness as the ones of the β Cep stars and SPBs.

Future work

We have already stressed the potential of the COROT exoplanet data for as- teroseismology of massive stars at previous COROT meetings. With our basic study provided here, we have shown that our goal is indeed feasible. In the near future we will perform a more extensive simulation study with a realistic noise budget. We will particularly concentrate on the derivation of the detection limit of modes for the exoplanet data. One of the other main goals of this future simulation study is to derive the optimal rebinning for the analysis of gravity modes in SPBs and γ Doradus stars. References Aerts, C., Eyer, L., Kestens, E. 1998, A&A 337, 790 Cuypers, J., Aerts., C., Buzasi, D., Catanzarite, J., Conrow, T., Laher, R. 2002, in ASP Conf. Ser., Vol. 259, Radial and Nonradial Pulsations as Probes of Stellar Physics, eds. C. Aerts, T.R. Bedding and J. Christensen-Dalsgaard, (San Francisco ASP), in press Eyer, L. 1998, PhD Thesis, Geneva Observatory Grenon, M. et al. 1997, ESA SP-1200, Volume 12 Poretti, E., Buzasi, D., Laher, R., Catanzarite, J., Conrow, T. 2002, A&A, in press (astro-ph/0111361) van Leeuwen, F. et al. 1997, ESA SP-1200, Volume 11 Waelkens, C., Aerts, C., Kestens, E., Grenon, M., Eyer, L. 1998, A&A 330, 215 Comm. in Asteroseismology Vol. 141, 2002

10 Aql, a new target for COROT

L. Bigot1, W. W. Weiss2

1Observatoire de la Cˆote d’Azur, BP 4229, Nice, France. 2Institut fur¨ Astronomie, Univ. Wien, Turk¨ enschanzstr. 17, 1180 Wien, Austria.

Abstract

The satellite COROT (e.g., Baglin & Foing 1999) will be devoted to the seismic investigation of stars. It will provide new insights into our knowledge of stellar structure. Among the variable stars, the group of roAp stars is unique. Indeed, the oscillations of these stars are dominated by the effects of magnetic fields, which are typically three orders of a magnitude higher than in the Sun. The study of the oscillations in roAp stars provides a chance to test and to improve our knowledge of the structure of these extremely peculiar stars. The spirit of this article is to argue in favor of 10 Aql (HD 176232, HR 7167) as an interesting target for COROT. We then present basic aspects of roAp stars and in particular recent theoretical developments concerning the magnetic field and rotation effects on the mode properties. We also propose an interpretation of the frequencies already observed for this star.

The Ap and roAp stars

The group of Ap stars is one of the most studied. They are Main Sequence stars of about 2 M¯ but with spectra showing abnormal lines of some elements. These peculiarities were first discovered more than a century ago and were the first application of spectroscopy to stellar astrophysics. During the twentieth century these stars indeed revealed strongly abnormal chemical composition, with spectral lines of some rare elements enhanced by several order of magnitude (Morgan 1933, Adelman 1973, Preston 1974). The most abnormal lines concern the rare elements like Europium, Chromium, Strontium (7000 < Teff < 10000) and Lithium (10000-14000 K). The elements of the iron peak are in turn of solar abundances. These elements are concentrated in some spots (and even rings) at the surface of the star. This has been observed by Doppler imaging techniques L. Bigot and W. W. Weiss 27

(e.g. Piskunov & Rice 1993). All these features were unexplained during the first part of the 20th century. Babcock (1947) made a crucial discovery to explain the physical processes that occur in Ap stars. Indeed, he was the first to underline a Zeeman splitting of spectral lines, proof for the presence of a strong magnetic field. The typical strengths of the fields present in Ap stars are of the order of kilogauss, say at least 2 or 3 orders of magnitudes higher than the value of the global field of the Sun. These magnetic fields first appeared to be dipolar. The variability of the observed field was explained by the model of the oblique rotator (Babcock 1949, Stibbs 1950) which claimed that the axis of symmetry of the field is inclined to the rotation axis, so that as the star rotates the observer sees different aspects of the same field. The inclination, β, between the magnetic and rotation axes is discussed in Landstreet & Mathys (2000). They show that slowly rotating Ap stars are characterized by small values of β. There is now clear evidence for non-dipolar components of fields in Ap stars (e.g. Landstreet 1980, Mathys 1987, Landstreet 1992, Mathys & Hubrig 1997, Landolfi et al. 1998). Recently, Hubrig et al. (2000a) showed that the magnetic field in Ap stars becomes detectable after some time spent on the Main Sequence, roughly after 30%, which indicates that this abnormal intense magnetic field may be related to a special stage in the life of the star. The chemical peculiarities in Ap stars were explained by the presence of these huge fields. Indeed, the intense field is supposed to stabilize the external layers of the star against turbulent mixing. In the stable upper layers, the chemical elements are then subject to the balance between the gravitational settling and the radiative pressure (Michaud 1970). The elements which are concentrated in their line forming regions appear then overabundant. Kurtz (1978) made an important and unexpected discovery in the late sev- enties. In studying the effective temperature of Przybylski’s star, he was the first to detect a rapid light variation of about 12 min in an Ap star. This was clear evidence for the presence of acoustic oscillations. This discovery was un- expected because it was wrongly thought before that the huge magnetic field would suppress any type of oscillations. These Ap stars are the only stars of the Main Sequence to show the same short periods as the Solar-like stars, i.e. from about 5 to 15 min. Kurtz (1982) named them the rapidly oscillating Ap (roAp) stars. However, with typical amplitudes from 0.3 to about 8 mmag (Johnson B filter), the oscillations in the roAp stars are three orders of magni- tudes higher than those in the Sun. Their number has grown to 32 (Martinez et al. 2001). From the time of their discovery these stars showed multiplets with peaks exactly separated by the rotation frequency, Ω, of the star. Kurtz (1982) then proposed the model of the oblique pulsator which states that the axis of sym- 28 10 Aql, a new target for COROT

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Figure 1: Representation of the Hertzsprung-Russell diagram for variable stars. The roAp stars are located on the Main Sequence at the intersection with the δ Scuti instability strip. From Christensen-Dalsgaard & Dziembowski (2000). metry of the mode is aligned with the magnetic axis of the star, i.e. is inclined from the rotation axis. Then, exactly as in the model of the oblique rotator, the observer sees different aspects of the same mode during the rotation of the star. There is a double light variation, one with a short time scale due to the oscillations (few minutes) and one longer due to the rotation of the star (few days). In the frequency domain, the modulation by rotation makes each mode of degree ` to appear as a multiplet with (2` + 1) components separated by exactly the rotation frequency. This fact has been confirmed for several roAp stars by comparison with the rotation periods obtained from mean light vari- ations due the spots (e.g. Kurtz & Marang 1988, Kurtz et al. 1990, 1996) and is the best evidence in favor of the model of the oblique pulsator. An- other property of these oscillations is a phase jump by π radians exactly at the L. Bigot and W. W. Weiss 29 amplitude minima. This aspect is evidence for dipole (` = 1) modes. This variation indeed corresponds to the fact that the observer sees alternatively the two hemispheres of the mode with opposite phases (one in contraction, one in expansion). The source of these oscillations is still not clearly established. The short periods of the pulsations could lead us toward the stochastic excitation due to the turbulence in the outer layers of the stars, as for the Sun. However, the large amplitudes (∼ mmag) cannot be explain by this process. Dziembowski & Goode (1996) show that for large number of nodes, the oscillations are driven by the κ-mechanism of the hydrogen ionization zone, and not by the He II zone as in the δ Scuti stars. The frontier between the roAp stars and the Ap stars which do not show rapid light variations, the non oscillating Ap (noAp) stars, is not clear. Indeed, the two classes of stars have similar magnetic field strengths, similar color indexes, etc. However, Hubrig et al. (2000b) point out a significant difference in their masses, which are somewhat larger in the noAp stars. This may have important consequences in the excitation mechanism. Balmforth et al. (2000) proposed a kappa-mechanism depending on the latitude, as a consequence of the suppression of convective turbulence at the magnetic poles. They show that the excited modes can be in some cases aligned with the magnetic field. Reviews on roAp stars are available in Weiss (1986), Unno et al. (1989), Matthews (1991) and Kurtz (1990, 2000).

A model for roAp star pulsation

The model generally accepted to interpret pulsations in roAp stars is the model of the oblique pulsator (Kurtz 1982, Dziembowski & Goode 1985, Kurtz & Shibahashi 1986, Bigot & Dziembowski 2002). We are now going to describe the basic aspects of this model and its consequences on the interpretation of the seismic data. We consider only the main component of the magnetic field, say the dipole one. The non-perturbative treatment of the effects of a dipole magnetic field on acoustic modes has been done by Dziembowski & Goode (1996), Bigot et al. (2000) and Cunha & Gough (2000). These studies show shifts of frequencies of the order of several µHz for fields of about 1 kG. There is also a significant damping of the oscillations by dissipation of Alfv´enic waves. We will ignore here this effect. In roAp stars the axis of symmetry of the mag- netic field generally differs from the axis of rotation. Therefore, their combined effects on acoustic modes lead to a coupling of spherical harmonics, see Fig. 2 represented for dipole modes. In the context of roAp stars, only the coupling in the azimuthal order m matters. We may then write the general displacement 30 10 Aql, a new target for COROT

Mode axis Ω

B

β Observer i

A 1 A −1 A 0

ω−Ω ω ω+Ω

Figure 2: Representation of the model of the oblique pulsator in the special case of a dipole mode. This axis of symmetry of this mode generally differs from both rotation and magnetic axes. Then the inertial observer sees a variation of the aspect of this mode as the star rotates. This leads to the presence of a triplet in the spectrum. vector as a linear combination of the (2` + 1) unperturbed eigenvectors

` ξ~ = αmξ~m, (1) − mX= ` with αm coefficients to be determined. This degenerate perturbation formalism was first applied by Dziembowski & Goode (1985) in this context of roAp stars. It can be shown then that the equation of oscillation leads to the following eigenvalue problem

` 2 αm Ojm − ω δjm = 0 j = −`, ..., `, (2) m=−` X © ª where Ojm is the projection of the sum of the rotational, Lorentz and the usual adiabatic operators on unperturbed eigenvectors. The system (2) has (2` + 1) eigenvalues ω and (2` + 1) eigenvectors {αm} which are orthogonal. For the rotational operator we take into account both Coriolis and centrifugal forces. The latter is the dominant effect of the rotational shift of frequency and is typically two orders of magnitude higher than the Coriolis force. This is mainly due to the presence of large radial orders for which the modes are more sensitive to the centrifugal force. However, the effects of the Coriolis force are essential to keep an asymmetry in the problem and cannot be neglected. This asymmetry leads to unequal coefficients for the coupling, i.e. |αm| 6= |α−m|, see Bigot & Dziembowski (2002) for more details. L. Bigot and W. W. Weiss 31

Once the solutions of equation (2) are found, one has to express them in the reference frame of the observer in order to compare with observations. This reference frame is defined with the polar axis along the line-of-sight. It is well known (e.g. Kurtz 1982) that the fluctuations of luminosity of each eigenmode expressed in the observer’s system writes as the multiplet

δL ` ∝ Am cos(ω − mΩ)t, (3) L ` − mX= ` with consecutive peaks separated exactly by the rotation rate Ω. The amplitudes of the multiplet are given by

` m (`) (`) A` = dm0(i) αmdm0(β), (4) − jX= ` (`) where djm are expressed in term of the Jacobi polynomials (e.g. Edmonds 1960), and i is the angle between the line of sight and the axis of rotation. It can be shown that the inequality of side peaks Am 6= A−m, often observed in roAp stars, is only due to the asymmetry introduced by the Coriolis force. The Lorentz and the centrifugal forces have a mirror symmetry that cannot lead to this inequality.

The case of dipole modes

The case of dipole modes (` = 1) is of special importance to interpret pulsations in roAp stars, since these modes are the most observed ones. In that case, recently Bigot & Dziembowski (2002) have shown that the effects of rotation, mainly due to centrifugal force, are sufficient to tilt the axis of symmetry of this mode from the magnetic axis. The geometrical picture of the dipole mode is that during the pulsation cycle its axis of symmetry moves in a plane inclined from both the rotation and magnetic axes. In general, the maximum of the displacement vector describes an ellipse in that plane. In Fig. (3) we represent the fundamental parameters of the dipole mode, say the angle δ between the normal of the mode plane and the rotation axis, the ellipticity ψ defined as the arctan of the ratio between the major and semi axes of the ellipse and the mag 2 2 dimensionless frequency of the mode σ = (ω − ω1 )/DΩ + 1/3, where DΩ mag is the shift due to centrifugal force and ω1 is the magnetic shift without rotation for m = 1. These parameters are plotted as functions of the obliquity mag mag 2 angle β and for different values of the quantity µ = (ω0 − ω1 )/DΩ , which measures the strength of the magnetic field; |µ| increases with B. We show clearly on this figure that the dipole modes are very dependent on the 32 10 Aql, a new target for COROT

Figure 3: The fundamental parameters of the dipole modes σ, δ and ψ as function of the obliquity angle β and for different values of µ (see text). The three types of lines correspond to the three solutions of the system (2). The magnetic effects dominate over the rotational effects on the mode if |µ| > 1. The last quantity D measures the alignment of the mode axis with the magnetic axis: D = 0 means alignment and D = 1 means orthogonality with the magnetic axis. From Bigot & Dziembowski (2002). value of β and on the magnetic strength (through µ). We see that in general there is no preference for alignment of the dipole axis with the magnetic axis. The reason is that in roAp stars the centrifugal effects are in general comparable to the magnetic effects of the m components of the mode, and therefore are sufficient to affect the inclination of the mode. L. Bigot and W. W. Weiss 33

Amplitudes of the side lobes (` = 1)

For the moment, no rotational side lobes have been detected in 10 Aql. This might be due to the probable small rotational frequency Ω/2π. From v sin i < −1 5 km s (Hoffleit 1982), R = 2.42 R¯, and an arbitrary assumption i = 70 deg we get Ω/2π = 0.5 µHz, which is very small. That is probably the reason why we have not yet detected the rotational multiplets in 10 Aql. We hope that with the improved accuracy of COROT, we will resolve the rotational side lobes of the modes and get important constraints on this star. With this idea, we now illustrate the effects of the magnetic field on the am- plitudes of the dipole modes. We plot in Fig. (4) the amplitudes of the triplet as functions of the parameter µ. The magnetic effects are larger than the ro- tational effects on the mode, when |µ| À 1. We also consider two different values of β, the angle between the axes of the magnetic field and rotation. Fig. (4) shows that the magnetic field has a strong influence on the amplitudes of the triplet. For 10 Aql, we are certainly in the regime |µ| > 1. There is also an important dependence on β. Therefore, we see that a comparison between the observed values of the triplet and these calculations can lead to an estimate of B and β for this star.

Analysis of the frequencies in 10 Aql

The roAp stars are excellent candidates for asteroseismology. From their pulsa- tions we can potentially estimate their radii, luminosities, mean densities, ages, magnetic fields, rotation rates, inclinations i and β. These variable stars can then be used as tests for stellar structure models and for dynamo theories. One of the most important criteria to select a target for an asteroseismic space mission is the number of modes that we can expect for this star. The short periods of oscillations in roAp stars allow us to use the formalism of the asymptotic theory for p-modes (e.g. Tassoul 1980) to interpret their spectra. The non-magnetic and non-rotational frequencies are then given by ` ν = ∆ n + + ² , (5) 0 2 µ ¶ with ² a small constant depending on the stellar structure. The fundamental spacing ∆0 between two consecutive radial orders n and n + 1 is defined by −1 R dr ∆ = 2 ∼< ρ >∼ M 1/2R−3/2, (6) 0 c ( Z0 S ) with cS the sound speed. This quantity has been calculated for several stellar models by Saio & Shibahashi (1985). They found that these values are between 34 10 Aql, a new target for COROT

Figure 4: Amplitudes of the triplet (A0, A1, A−1) normalized to the unity as function of the magnetic field and for two different values of β = 10 deg (left) and 40 deg (right). The meaning of lines is A0 (full line) , A1 (dashed), A−1 (dot-dashed) and D (dotted).

20 and 100 µHz. From its value we can get an estimate on the mean density of the star, which is then a constraint on the stellar model. About half of the roAp stars are multiperiodic. One of them is in the field of COROT in the northern hemisphere. This star is 10 Aql (HD 176232, HR 7167). It is a F0 SrEu star whose recent spectroscopic investigation is available in Ryabchikova et al. (2000). The variability of this star has been discovered by Heller & Kramer (1990) who showed three distinct eigenmodes whose frequencies are ν1 = 1435.99, ν2 = 1385.37 and ν3 = 1239.26 µHz. Their amplitudes are less than 0.5 mmag (Johnson B filter). These amplitudes are much larger than the solar amplitudes but they are quite small for roAp stars which have generally ∆B > 1 mmag. These three modes are very close to the noise detection ∼ 0.2 L. Bigot and W. W. Weiss 35

Figure 5: The three detected frequencies in 10 Aql, ν1 = 1435.9, ν2 = 1385.4 and ν3 = 1239.2 µHz. Note also the low amplitudes (∆B < 0.5 mmag) of the eigenmodes. From Heller & Kramer (1990). mmag. The main source of this noise is the scintillation of the sky. Therefore it is a great advantage to observe these stars from space. There are others roAp stars which show small amplitudes like HR 1217 (Kurtz et al. 1989) or γ Equ (Martinez et al. 1996), and in these cases they exhibit more than one mode. In turn, the roAp stars which show large amplitudes, like HR 3831 (Kurtz et al. 1997) or α Cir (Kurtz et al. 1994) have only one single mode. Therefore since 10 Aql has similar amplitudes to HR 1217 and γ Equ we expect that there are many modes hidden in the noise of detection. The satellite COROT will considerably decrease this noise and will be able then to discover a large number of eigenmodes. The geometrical nature of the modes already detected in 10 Aql is not yet clearly established. In most roAp stars (e.g. Kurtz 1990, 1998) the modes responsible of the oscillations are dipole (` = 1) modes. We can expect the same nature of pulsation in 10 Aql. In order to interpret oscillations and to prepare the space mission, we have to choose a stellar evolutionary model for this star. The metallicity Z is cer- tainly one of the most undetermined parameters regarding the extreme chemical complexity of the outer layers. There is no reason to choose solar metallicity 36 10 Aql, a new target for COROT for an Ap star, since the Sun is known to be metal deficient compared with its neighborhood. Therefore, following Saio & Shibahashi (1985), we consider Z = 0.03. The effective temperature and luminosity are given by Matthews et al. (1999). The latter is obtained from Hipparcos measurements. They are: Teff = 8000§100 K and L = 21.4§0.6 L¯. The model is also constrained by the large frequency separation observed in 10 Aql, i.e. ∆12 = ν1 − ν2 = 50.6 µHz. The overshooting is supposed to be 0.2Hp, with Hp the pressure scale height. The stellar model is calculated with the CESAM evolutionary code developed at the Observatoire de la Cˆote d’Azur (Morel 1997). The model which reproduces well all these constraints is: M = 2.11 M¯, R = 2.42 R¯, X = 0.70, Z = 0.03 and a central fraction of hydrogen Xc = 0.4. We found a spacing of ∆0 = 50.15 µHz and a cut-off frequency equal to 1510 µHz (see Audard et al. 1998 for a discussion of the limit of frequencies for trapped acoustic modes in roAp stars). These results that we obtain are comparable with the observations of Heller & Kramer (1990). We note that the largest observed frequency is close to the cut-off frequency. In Figs (6), (7) and (8) we represent the three frequencies which are closest to the observed ones, assum- ing that they correspond to a same degree ` = 1. They have respectively the radial orders n = 24, 27, and 28. We note that the calculated non magnetic frequencies (full lines) differ from the observed ones (dot-dashed lines) by sev- eral µHz, up to about 8 µHz. However, it is now well known (Dziembowski & Goode 1996, Bigot et al. 2000, Cunha & Gough 2000) that the magnetic field is responsible for an important shift of frequency in roAp stars. It also breaks the degeneracy of frequencies which depend on |m|. In Fig. (6), (7) and (8) we place the frequencies corresponding to the radial orders n = 24, 27 and 28 affected by two different values of the magnetic field B = 0.5 and 0.8 kG. The best value of B that reproduces the data is 0.8 kG. We then get an agreement with the three observed frequencies within ∼ 0.3 µHz. This value of B agrees with the estimate given in Ryabchikova et al. (2000) from spectroscopy. The frequencies ν1 and ν2 correspond to the modes (n = 27, ` = 1, m = 0) and (n = 28, ` = 1, m = 0) and ν3 correspond to (n = 24, ` = 1, |m| = 1). We note that the shifts introduced by the field are quite large. This is due to the fact that these frequencies are close to the cut-off frequency.

obs obs We emphasize here that the non-obvious observed inequality ∆23 = ν2 − obs obs obs obs ν3 6= k∆12 = k(ν1 − ν2 ), with k integer, is well explain by the effect of the magnetic field and the shift that it produces. The magnetic spacings mag mag mag that we obtain for B = 0.8 kG are ∆12 = ν1 − ν2 = 50.4 µHz and mag mag mag ∆23 = ν2 − ν3 = 146.8 µHz which are close to the values observed by Heller & Kramer (1990). L. Bigot and W. W. Weiss 37

Figure 6: Positions of the frequency ν3 without a magnetic field (full line) and for B = 0.5 kG with m = 0 (dotted line) and for B = 0.8 kG with m = 0 (dashed line) and |m| = 1 (dot(×3)dashed line). The observed value is plotted with a dot-dashed line.

Conclusion

The roAp star 10 Aql is a very interesting target for the asteroseismic space mission COROT. The observations from the ground already lead to three eigen- modes with large radial orders. We have shown that the positions and ampli- tudes of the eigenmodes in the spectrum of oscillations are strongly influenced by the magnetic field, which by our estimation is close to 0.8 kG. The lack of accurate measurements leads to an absence of rotational sidelobes. The interest in COROT is to improve the resolution of observations, to resolve the multiplets and to discover a large number of modes hidden in the noise. As we have shown, the amplitudes will lead to a determination of the magnetic field 38 10 Aql, a new target for COROT

Figure 7: Same as Fig. (6) for the frequency ν2. and the obliquity angle between the rotation and magnetic axes, which then would be a constraint for dynamo theories for Ap stars.

Acknowledgments. L.B. thanks the ”Soci´et´e de Secours des Amis des Sciences” for financial support during the preparation of this article.

References Adelman, S. J. 1973, ApJ 183, 95 Audard N., Kupka F., Morel P., Provost J., Weiss W. W., 1998, A&A 335, 954 Babcock, H. W. 1947, ApJ 105, 105 Babcock, H. W. 1949, Obs 69, 191 Baglin, A., Foing, B. H. 1999, AdSpR 24, 145 Bagnulo, S., Landolfi, M., Landi Degl’Innocenti, M. 1999, A&A 343, 865 L. Bigot and W. W. Weiss 39

Figure 8: Same as Fig. (6) for the frequency ν1.

Balmforth, N. J., Cunha, M. S., Dolez, N., Gough, D. O., Vauclair, S. 2001, MNRAS 323, 362 Bigot, L., Provost, J., Berthomieu, G., Dziembowski, W. A., Goode, P. R. 1999, RoAJS 9, 129 Bigot, L., Provost, J., Berthomieu, G., Dziembowski, W. A., Goode, P. R. 2000, A&A 356, 218 Bigot, L., Dziembowski, W. A. 2002, A&A, accepted Christensen-Dalsgaard, J., Dziembowski, W. A. 2000, in Variable stars as essential astrophysical tools, ed. Cafer Ibanoglu, Kluwer Academic Publishers, Dordrecht, Boston, NATO science series. Series C, Mathematical and physical sciences; vol. 544, p. 1 Cunha, M. S., Gough, D. 2000, MNRAS 319, 1020 Dziembowski, W. A., Goode, P. R. 1985, ApJ 296, L27 40 10 Aql, a new target for COROT

Dziembowski, W. A., Goode, P. R. 1996, ApJ 458, 338 Edmonds, A. R. 1960, Angular Momentum in Quantum Mechanics, second edition, Princeton Univ. Press Heller, C. H., Kramer, K. S. 1990, MNRAS 246, 699 Hoffleit, D., 1982, Bright Star Catalogue, 4th edn, Yale University Observatory Hubrig, S., North, P., Mathys, G. 2000a, ApJ 539, 352 Hubrig, S., Kharchenko, N., Mathys, G., North, P. 2000b, A&A 355, 1031 Kurtz, D. W. 1982, MNRAS 200, 807 Kurtz, D. W., Shibahashi, H. 1986, MNRAS 223, 557 Kurtz, D. W., Marang, F. 1988, MNRAS 231, 565 Kurtz, D. W., Matthews, J. M., Martinez, P., Seeman, J., Cropper, M., Clemens, J. C., Kreidl, T. J., Sterken, C., Schneider, H., Weiss, W. W., Kawaler, S. D., Kepler, S. O. 1989, MNRAS 240, 881 Kurtz, D. W. 1990, ARA&A 28, 607 Kurtz, D. W., van Wyk, F., Marang, F. 1990, MNRAS 243, 289 Kurtz, D. W., Sullivan, D. J., Martinez, P., Tripe, P. 1994, MNRAS 270, 674 Kurtz, D. W., van Wyk, F., Roberts, G., Marang, F., Handler, G., Medupe, R. 1997, MNRAS 287, 69 Kurtz, D. W., Marang, F., van Wyk, F., Roberts, G. 1996, MNRAS 280, 1 Kurtz, D. W. 2000, in in Variable stars as essential astrophysical tools, ed. Cafer Ibanoglu, Kluwer Academic Publishers, Dordrecht, Boston, NATO science series. Series C, Mathematical and physical sciences, vol. 544, p.313 Landolfi, M., Bagnulo, S., Landi degl’Innocenti, M. 1998, A&A 338, 111 Landstreet, J. D. 1980, AJ 85, 611 Landstreet, J. D. 1992, A&ARv 4, 35 Landstreet, J. D., Mathys, G. 2000, A&A 359, 213 Martinez, P., Weiss, W. W., Nelson, M. J., Kreidl, T. J., Roberts, G. R., Mkrtichian, D. E., Dorokhov, N. I., Dorokhova, T. N., Birch, P. V., MNRAS 282, 243 Martinez, P., Kurtz, D. W., Ashoka, B. N., Chaubey, U. S. Girish, V., Gupta, S. K., Joshi, S., Kasturirangan, K., Sagar, R., Seetha, S. 2001, A&A 371, 1048 Mathys, G. 1989, The Observation of Magnetic Fields in Nondegenerate stars, Fundamentals of Cosmic Physics (Gordon and Breach Science Publishers, Inc.), vol. 13, 143 Mathys, G., Hubrig, S. 1997, A&AS 124, 475 Matthews, J. M. 1991, PASP 103, 5 Matthews, J. M., Kurtz, D. W., Martinez, P. 1999, ApJ 511, 422 Michaud, G. 1970, ApJ 160, 641 Morel, P. 1997, A&AS 124, 597 Morgan, W. W. 1933, ApJ 77, 330 Piskunov, N. E., Rice, J. B. 1993, PASP 105, 1415 Preston, G. W. 1974, ARA&A 12, 257 Ryabchikova, T. A., Savanov, I. S., Hatzes, A. P., Weiss, W. W., Handler, G. 2000, A&A 357, 981 L. Bigot and W. W. Weiss 41

Stibbs, D. W. N. 1950, MNRAS 110, 395 Tassoul, M. 1980, ApJS 43, 469 Unno W., Osaki Y., Ando H. & Shibahashi H. 1989 Non-radial oscillations of stars, second edition University of Tokyo Press Weiss, W. W. 1986, Upper Main Sequence star with Anomalous Abundances, Ed C. R. Cowley, M. M. Dworetsky et C. M´egessier, IAU coll. 90, D. Reidel Publ. Co., Dordrecht, Holland, p. 219 Comm. in Asteroseismology Vol. 141, 2002

Status of the COROT ground-based photometric activities

R. Garrido1, P. Amado1, A. Moya1, V. Costa1, A. Rolland1, I. Olivares1 and M. J. Goupil2

1.-Instituto de Astrofisica de Andalucia, Apdo. 3004, Granada, Spain 2.-DASGAL, UMR CNRS 8633, Observatoire de Paris-Meudon, France

Abstract

We report photometric observations made at the Observatorio de Sierra Nevada of the secondary targets in the input catalogue of COROT. These observations consist of Str¨omgren –uvby- and Crawford –Hβ– measurements of stars in Scenarios 1, 2 and 4. The cooler stars are being measured in the CaII H&K photometric system. We present here the present status of observations and reductions together with some preliminary results on the monitoring of photo- metric variability for some selected stars.

Introduction

In order to have an homogeneous data base of the secondary targets (up to 8 mag) it was decided to perform photometric and spectroscopic observations of the preliminary COROT selected fields. Here we present the photometric observations carried out in the four band –uvby- Str¨omgren system, where as usual: (b − y) is basically a temperature indicator m = (v − b) − (b − y) is a metallicity indicator and c = (u − v) − (v − b) is a luminosity indicator and in the Crawford –Hβ– system where Hβ is another independent temperature indicator for most of the star sam- ple, hence the reddening will be calculated using the Hβ index. The use of the CaII H&K band has been introduced since it is more than two times more sensitive to metallicity changes than m1, with sensitivity R. Garrido et al. 43

Figure 1: V versus (b − y) color diagram. (b − y) is unreddened. Stars fainter than 8 mag. are probably missidentifications of multiple systems.

increasing to lower temperatures, as shown by Twarog and Anthony-Twarog (1991).

Figure 2: c versus (b − y) color diagram. Colors are unreddened. 44 Status of the COROT ground-based photometric activities

Figure 3: m versus (b − y) color diagram. Colors are unreddened.

Instrumentation

All the Str¨omgren and Crawford observations were made with a photoelectric photometer able to work in two modes: in the Str¨omgren mode, it measures simultaneously the four uvby bands and in the Crawford mode, it measures simultaneously two bands nw (narrow and wide) centered on the Hβ line. This photometer is attached to a 0.9m telescope on Sierra Nevada Obser- vatory, property of the IAA (CSIC) situated near the city of Granada. The building contains two domes hosting two telescopes. CaII H&K observa- tions were made with a CCD camera attached to the other telescope with a diameter of 1.5m. Details of this instrumentation can be downloaded from http://www.iaa.es/osn/principal/html.

Catalogue observations

The observations we present here began in early 2000 when it was not yet known the final scenario chosen for COROT, at that time there was two possibilities: one starting at α = 5h37m (and 12h later) called scenario 1 and other starting at α = 7h19m (and 12h later) called scenario 2. We decided to observe both scenarios in order to have more physical information on the possible targets to help in the final selection. After some discussions during the meeting of the Scientific Committee held in Paris in spring 2001 an intermediate scenario was finally decided for COROT: scenario 4, starting at α = 6h03m and ending at 7h31m (and 12h later) , always within an equatorial band of §10 degrees. This R. Garrido et al. 45

Figure 4: Photometric variability of HD 49434. Stars are magnitude differences be- tween HD 49434 and HD 48922. Points are for HD 49933 minus HD 48922.

Figure 5: Power spectra of both photometric series shown in Figure 4.

final scenario has some stars in common with scenarios 1 and 2 but most of the stars were new. Before the decision was taken to select scenario 4 most of the stars in scenarios 1 and 2 were observed. In particular for scenario 1, 539 stars out of 661 were observed in the Str¨omgren and Crawford systems and 78 out of 144 in the CaII H&K system. For scenario 2 these numbers were: 442 out of 458 and 69 out of 128. We started the new observations for scenario 4 in summer 2001 and all the visible stars in the season were observed in Str¨omgren and 46 Status of the COROT ground-based photometric activities

Table 1: Fitted frequencies of HR 6534.

ν(µHz) Amp (mag) phase(rad) νi/ν0 νi/ν1 158.92 0.0022 § 0.0003 3.40 § 0.16 1 - 207.45 0.0039 § 0.0004 4.57 § 0.09 0.77 1 237.38 0.0018 § 0.0004 1.47 § 0.19 0.67 0.87 291.95 0.0022 § 0.0004 4.84 § 0.16 0.56 0.71 353.76 0.0029 § 0.0003 0.25 § 0.12 0.45 0.59

Crawford systems: a total of 155 objects. No observations were made in the CaII H&K system during this summer. In conclusion, 155 out of 419 stars have been observed in scenario 4 in the St¨omgren and Crawford systems and 105 out of 187 stars were observed in the CaII H&K system. We expect to finish all the observations in summer 2002, including the remaining stars from scenarios 1 and 2 in order to have an homogeneous catalogue to compare with spectroscopic observations made also for stars of these two scenarios. We present in Figure 1, 2 and 3 some preliminary and partial results of the photometric reductions now in progress.

Figure 6: (b+v) photometric variations of HR 6534. R. Garrido et al. 47

Table 2: Frequency of the fundamental radial mode and period ratios for radial modes of a 1.75M¯ model with no rotation (v = 0) and with a rotational veloc- ity of v = 100 km/s and v = 255 km/s.

v (km/s) log Teff ν0 (µHz) P1/P0 P2/P0 P3/P0 P4/P0 P5/P0 0 3.90 207.2 0.772 0.6305 0.533 0.458 0.399 100b 3.90 210.6 0.777 0.633 0.536 0.462 0.398 0 3.89 188.2 0.772 0.628 0.530 0.456 0.398 100b 3.89 191.2 0.782 0.627 0.534 0.460 0.397 250a 3.89 203.4 0.787 0.643 0.545 0.473 0.416 250b 3.89 221.2 0.884 0.736 0.635 0.553 0.488 a no degenerate coupling included tb degenerate coupling included

Photometric monitoring

Some bright stars from scenarios 1, 2 or 4 were observed to be slightly vari- able in spectroscopy by Claude Catala and then selected to be monitored in photometry. From these HD 171834, HD 164259, HD 46304, HD 174866 and HD 180868 were found to be “non-variable” during one night of observation (6 hours aproximately) at the level of 2 mmags rms (see Figure 5b for a typical power spectrum of a “non-variable” star). Only HD 49434 showed some pho- tometric variability as can be seen in Figure 4. This long term variability, see power spectrum in Figure 5, is typical for the newly discovered γ Dor stars as discussed in Bruntt et al. (2002). One of the comparison stars to test variability turned out to be a new δ Scuti star: HD 159170 = HR 6534. Light curves are plotted in Figure 6. The star is clearly multiperiodic and a preliminary Fourier solution is shown in Table 1. Period ratios seem to indicate that HR 6534 is a pure radial pulsator but the star has a rotational velocity of 255 kms−1 and when high rotation is included in the pulsation model the radial overtones do not show anymore the classical period ratios as shown in Table 2. More observations have been made during 2001 from OSN and from SAT in La Silla in order to confirm these results that will be discussed in a forthcoming paper by Garrido et al. 2002. Unfortunateley HR 6534 (or “Orianita” as MJG wants it to be renamed) is outside the scenario 4 (α = 17h33m30s and δ = -05d44m41s)

References Anthony-Twarog, B. J., Twarog, B. A. 1995, AJ 109, 2828 Bruntt, H., Catala, C., Garrido, R., Rodriguez, E. 2002, in preparation Garrido, R., Goupil, M. J., et al. 2002, in preparation Comm. in Asteroseismology Vol. 141, 2002

Mode identification using the exoplanetary camera

R. Garrido1, A. Moya1, M. J.Goupil2, C. Barban2, C. van’t Veer-Menneret2, F. Kupka3 and U. Heiter3

1.-Instituto de Astrofisica de Anadalucia, Apdo. 3004, Granada, Spain 2.-DASGAL, UMR CNRS 8633, Observatoire de Paris-Meudon, France 3.-Institut fur¨ Astronomie, Universit¨at Wien, Austria

Abstract

Numerical simulations are presented, which show that color information, as supplied by the exoplanetary camera of COROT, can be relevant to discriminate among several spherical degrees ` of the non-radial oscillations of δ Scuti stars.

Introduction

Multiperiodic δ Scuti stars show power spectra indicating that they pulsate in a mixing of radial and non-radial modes. Whereas period ratios can be used as an indication of the radial order n for radial modes they are not useful for the non-radial ones. Rotation and evolution –the so called “avoided crossing” phenomenon– complicate even more any mode identification. Spectroscopy – through line profile variations– and photometry –through the use of amplitude ratios and phase differences of different colors– have been widely utilized as techniques for mode identification. Here we will focus on the capability of the explanetary camera in COROT to supply color information which can be relevant for mode identification.

Numerical simulations

In front of the exoplanetary camera a bi-prism device will supply colored spotted images which will give three colors. These colors will depend on the spectral type of the considered star and on the position over the CCD plane. The full spot will be divided into three parts called “blue”, “green” and “red”, containing 20%, 20% and 60% respectively of the total flux. R. Garrido et al. 49

Figure 1: Amplitude ratio vs phase differences between “blue” and “green” bands of the exoplanetary camera of COROT for a model atmosphere of T = 7500 K, log g = 4.0 and periods close to a pulsation cosntant of Q = 0.035d. Regions for every ` value contain predictions for phase lags between 90 and 135 degrees and values of R between 0.25 and 1.00.

The Kurucz model atmospheres have been improved in the way explained in detail in Heiter et al. (2002), basically by a better treatment of the convection and a larger number of shells in order to obtain good derivatives with respect to temperature and gravity. The linearised formula given by Watson (1988) has been used to compute the diagrams shown in Figure 1 and 2. Regions for different ` have been evaluated for an improved model atmosphere of T = 7500 K, log g = 4, Pop I and a period region around the fundamental radial mode –pulsation constant Q = 0.035 d– following the paper by Garrido et al. (1990). Basically the predictions are based on calculations assuming a phase lag –angle between temperature and radius variations– between 90 and 135 degrees and an R value –a parameter measuring the non-adiabaticity of the atmosphere– between 0.25 and 1.00. From these figures we can see a clear separation for high ` orders, giving very different regions in the amplitude ratio vs phase differences plane. For the lowest ` orders this separation is less evident but we expect that the accuracy at which COROT will supply light curves allows a precision for amplitudes and 50 Mode identification using the exoplanetary camera

Figure 2: Same as figure 1 but for the central region. phases sufficient to discriminate among several ` orders and hence to help doing asteroseismology of these stars.

References Garrido, R., Garcia-Lobo, E., Rodriguez, E. 1990, A&A 234, 262 Heiter, H., et al. 2002, in preparation Watson, R. D. 1988, Ap&SS 140, 255 Comm. in Asteroseismology Vol. 141, 2002

COROT and the late stages of stellar evolution

T. Lebzelter, H. Pikall, F. Kerschbaum

Institut fur¨ Astronomie, Turk¨ enschanzstrasse 17, A-1180 Vienna, Austria

Abstract

Small amplitude variability is observed during the late stages of stellar evolution of low- to intermediate-mass stars. In this paper we discuss the occurrence of small-amplitude variations in the upper right part of the HRD, namely among the AGB and post-AGB stars, and we propose two observing programs for the COROT Additional Program.

AGB and post-AGB evolution

After exhaustion of central hydrogen, stars of low to intermediate mass evolve into red giants. The ascent on the first giant branch (RGB) is followed by a phase of helium core burning and a phase of helium and hydrogen shell burning, the latter being called the Asymptotic Giant Branch phase (AGB). This phase is characterized by cool, highly extended stellar atmospheres/envelopes. Pulsation is a critical aspect of the late stages of stellar evolution. Regular, semiregular and irregular variations have been found in the light and velocity changes of these stars. Based on regularity, period and amplitude three types of variables have been defined: Miras, Semiregular variables (SRVs) and Irregular variables (Lb). Pulsation affects the structure of the stellar atmosphere. Expansion leads to conditions better suited for the formation of dust. Mass loss and dust production are related and important tests for stellar evolution While variability is probably mainly due to pulsation, it may be affected by other sources. The attempt to understand the irregularities in the light change of these stars is one of the main aspects currently investigated in these objects (Percy 1997, Lebzelter et al. 2000, Kerschbaum et al. 2001, Percy et al. 2001). The AGB phase terminates when the star reaches a very high mass loss rate. While losing the remainder of its atmosphere, the degenerated core is uncovered and the effective temperature rises shifting the objects to the “left” in the HRD. 52 COROT and the late stages of stellar evolution

When all nuclear sources of energy are used up, the final (core-)contraction is inevitable and the objects arrive at the white dwarf (WD) cooling track, which is characterized by decreasing luminosity and temperature. A number of interesting phases are associated with these final steps in stellar evolution. Here we want to focus on three classes of objects: R CrB stars (RCB), Hydrogen deficient Carbon-stars (HdC) and Extreme Helium Stars (eHe). All three groups are characterized by the absence of Hydrogen. While for the first two groups H is underabundant by a factor of 10−4 to 10−8 (number density), some objects among the eHes have a larger fraction of H.

Following Iben et al. (1996) two evolutionary scenarios are thought to lead to the R CrB phenomenon: One possibility is the merging of two, close degenerated objects – like WD or even Neutron stars – in a binary system (double degenerate scenario, DD). Recent calculations of WD-evolution with He-accretion (Saio & Jeffery 2001) support this idea. The accreted material is heated and compressed by shocks which starts the nuclear fusion. He Shell burning releases enough luminosity to virtually move the star to the upper right corner of the HRD. Alternatively a final Helium flash can produce enough luminosity to expand the envelope so the object appears as a cool supergiant again. Helium flashes occur regularly during the final stages of AGB evolution, but in some cases a very late FINAL flash may happen after the star has lost its envelope and is already on the WD cooling track (final flash scenario, FF). Both scenarios lead again to a cool, extended envelope, so that these stars are found close to the AGB.

According to Clayton (1996) all RCB have an infrared excess, while for none of the HdC was an infrared excess observed. IRAS data imply dust temperatures in the range of 650–900 K (Walker 1986). There is evidence for old, fossil dust shells with large spatial dimensions, implying an age of 104 to 105 years for the shell around R CrB itself (Gillett et al. 1986), but without a connection to the recent mass loss, which is responsible for the declines. The emisson in the infrared is changing but not correlated to the big declines. Measurements of the polarisation support the existence of dust.

Beside the obvious large amplitude variations in RCB, the declines, which are due to phases of heavy mass loss and the emission of dust in puffs directly in the line of sight, small amplitude pulsation is observed in these stars. However, investigation of these variations has just begun and their origin and type of pulsation is not yet known. Typical pulsation periods are 30–110 days for RCB and HdC and between 0.1 and a few days for the eHe (Weiss et al. 1996, and references therein). Feast (1996) points out that the turnover timescales for a few large convection cells are of the same order for cool objects, but there is doubt if the change in brightness can be explained by these. T. Lebzelter, H. Pikall and F. Kerschbaum 53

Open questions

AGB stars

Our interest in COROT concerning AGB stars is motivated by two open ques- tions that require precise photometry with a high time resolution. At first glance, neither requirement seems necessary for investigations of AGB stars, as the variations in these objects typically occur on a time scale of several tens of days up to several years with amplitudes between 0.1 and 8 magnitudes. However, the definition of the variability region in the HRD, where the AGB stars are located, can be established by much smaller variations on much smaller timescales. AGB variables are not located within the classical instability strip but seem to form a separate instability region. The blue border of this region is not known yet. Several investigations detected small amplitude variability in late K- and early M-type giants, i.e. on the left side of the AGB. Jorissen et al. (1997) and Fekel et al. (2000) showed that small amplitude variability is common in this kind of stars. Semiregular and irregular variables with amplitudes of more than 0.1 mag are found also among early M-type giants (e.g. Lebzelter 1999a). Edmonds & Gilliland (1996) found variable K-giants in the globular cluster 47 Tuc with period lengths of a few days and amplitudes of about 10 mmag. Koen & Laney (2000) detected short time variations (periods of a few days) in M giants, although part of their results may be due to problems with the Hipparcos data they used (Kerschbaum et al. 2000). A more detailed investigation of the variability in this region of the HRD is necessary to derive the origin of the observed variability and to deduce the parameters for the onset of pulsation among cool giants. Also for the classical, large amplitude AGB variables, photometry with high time resolution is needed. It is suspected that these stars show also variations that do not origin from stellar pulsations. Opposite to what is found in Sun-like stars, the surface of a cool AGB giant should be covered only by a few large convective cells (Schwarzschild 1975; Freytag et al. 1997). Their occurrence and motion may be echoed in small amplitude changes of the star (e.g. Lebzelter 1999b). Beside that several investigators report the observation of short time outbursts in MIRAS (Maffei & Tosti 1995, de Laverny et al. 1998). For a few hours the star’s brightness changes by up to several 0.1 mag. The mechanism of these outbursts is not understood yet. Both kinds of variations, due to surface structures and due to outbursts, happen on a short time scale and may be detectable only with high precision photometry. Due to the lack of appropriate models an accurate prediction of the time scales and amplitudes of these variations cannot be made. 54 COROT and the late stages of stellar evolution

Post-AGB stars The small amplitudes of the pulsation detected or expected in the groups of post-AGB stars described above should predestine these objects for high preci- sion continuous photometry offered by COROT. Periods have been derived only for a small fraction (2) of the known objects, all with rather big uncertainties, so there is a clear demand for a detailed study of these stars. The determination of exact stellar parameters like effective temperature, luminosity, mass and the chemical composition are needed to calculate detailed models of these objects – both for evolutionary and pulsational aspects. The Hipparcos mission could not determine distances to any of the objects, because many of the objects were not bright enough, but the detection of RCB in the LMC during the MACHO-project (Alcock et al. 1996) supported high luminosi- ties. Estimates of the brightness were MV ≈ −4 to −5. This however doesn’t constrain evolutionary models (which could shed light on the object’s masses) sufficently. Atmospheric abundances, derived from detailed static modeling of the atmospheres, simultaneously yield values for temperatures and surface grav- ities (Asplund et al. 2000, Pandey et al. 2001). Pulsation theory relates the period of pulsation to the star’s mean density regarding the pulsation mode, which makes it important to measure the period with high accuracy, but pe- riods of the order of 30-110 days are hard to observe, due to seasonal effects. The temporal change of period is an indication if the star is contracting or expanding, but prior to that the period has to be established. High luminosity in stellar envelopes allows strange modes, which can grow on very short timescales, comparable to a few of their pulsation periods. Mak- ing it simple, one could say, that such envelopes tend to show instabilities of greater order and manifold than low-luminosity objects. The reason derives partly from the fact that the Kelvin-Helmholtz timescale becomes comparable to the dynamical timescale for very high luminosities. Linear theory showed that a density inversion enables mode trapping in the outer parts of the en- velopes [REFERENCE?]. Stellar winds are thought to be driven by different mechanisms, but high luminosities help anyway, if the matter is not heated, but accelerated. The connection between pulsation, wind and dust formation is still not understood. While dust formation and dust driven winds can be modeled simultaneously (H¨ofner 1999), the influence on the pulsation has not been in- vestigated so far. This also leads to the question if all RCB are pulsating, and what the correlation of systematic period variations to the evolution of these objects is. Another important question is: How does convection work in cool super- giants? As for AGB stars, which are in many ways similar to the RCB and HdC, short time variations with low amplitudes are expected. In the case of some eHe stars the question arises: Is the observed irregularity intrinsic or due T. Lebzelter, H. Pikall and F. Kerschbaum 55 to shortcomings in the observations.

Perspectives

Within the course of the Additional Program we propose two observing pro- grams. M giants and AGB stars are found all over the sky so that we aim to measure the light variability of both kinds of stars that are in the final fields of view of COROT. The data will contribute to both questions discussed for the AGB stars above. We plan to combine these observations with spectroscopic measurements and photometric follow up studies from the ground. The second observing proposal aims to measure the light change of at least one post-AGB object with high photometric precision. It turned out that a number of stars from the three groups described in this paper would be close to the planned field of view of COROT. One star, FH Sct – an RCB – will be in the Exoplanetary Field if one of the prime targets is observed in the asteroseismology field. No pulsational period has been derived for this star yet. The apparent brightness is V = 12.2 mag and therefore FH Sct fits perfectly in the sensitivity range of the instrument. For both proposals the Exoplanetary Field offers sufficient time resolution.

Acknowledgments. TL is supported by the Austrian Science Fund Project P14365-PHY. We wish to thank W.W.Weiss for organizing the COROT Science Week in Vienna.

References Alcock, C., Allsman, R. A., Alves, D. R., et al. 1996, ApJ 470, 583 Asplund, M., Gustafsson, B., Lambert, D. L., Rao, N. K. 2000, A&A353, 287 Clayton, G. C. 1996, PASP 108, 225 Clayton, G. C. 2001, AP&SS 275, 143 de Laverny P., Mennessier, M. O., Mignard, F., Mattei, J. A. 1998, A&A 330, 169 Edmonds, P. D., Gilliland, R. L. 1996, ApJ 464, L157 Feast, M. W. 1986, in Hydrogen Deficient Stars and Related Objects, ed. K. Hunger (Dordrecht, Reidel), p. 151 Fekel, F. C., Henry, G. W., Henry, S. M. 2000, in Garcia Lopez, R. J., Rebolo, R., Zapatero Osorio, M. R. (eds.), 11th Cambridge Workshop on Cool Stars, Stellar Systems and the Sun, ASP Conf. Ser. 223, 925 Freytag, B., Holweger, H., Steffen, M., Ludwig, H.G. 1997, in Paresce, F. (ed.), Science with the VLT interferometer”, ESO Astrophys.Symp., Springer, p. 316 Gillett, F. C., Backman, D. E., Beichman, C., Neugebauer, G. 1986, ApJ 310, 842 H¨ofner, S. 1999, A&A 346, L9 Iben, I., Tutukov, A. V. & Yungelson, L. R. 1996, ApJ 456, 750 Jorissen, A., Mowlavi, N., Sterken, C., Manfroid, J. 1997, A&A 324, 578 56 COROT and the late stages of stellar evolution

Kerschbaum, F., Lebzelter, T., Lazaro, C. 2001, A&A 375, 527 Koen, C., Laney, D. 2000, MNRAS 311, 636 Lawson, W. A. & Cottrell, P. L. 1997, MNRAS 285, 266 Lebzelter, T. 1999a, A&A 346, 537 Lebzelter, T. 1999b, A&A 351, 644 Lebzelter, T., Kiss, L. L., Hinkle, K. H. 2000, A&A 361, 167 Maffei, P., Tosti, G. 1995, AJ 109, 2652 Pandey, G., Kameswara Rao, N., Lambert, D. L., et al. 2001, MNRAS 324, 937 Percy, J. R. 1997, JAAVSO 25, 93 Percy, J. R., Wilson, J. B., Henry, G. W. 2001, PASP 113, 983 Saio, H. 1995, MNRAS 277, 1393 Saio, H. & Jeffery, C. S. 2000, MNRAS 313, 671 Schwarzschild, M. 1975, ApJ 195, 137 Walker, H. J. 1986, in Hydrogen Deficient Stars and Related Objects, ed. K. Hunger (Dordrecht, Reidel), p. 407 Walker, H. J. 1994, CCP7 Newslett. 21, 40 Weiss, A., Fried, R., Olson, C. E. 1996, A&AS 116, 31 Woitke, P., Goeres, A. & Sedlmayr, E. 1996, A&A 313, 217 Comm. in Asteroseismology Vol. 141, 2002

Pulsations of Luminous Blue Variables

E. A. Dorfi1, A. Gautschy2

1Institut fur¨ Astronomie, Turk¨ enschanzstr. 17, A-1180 Wien, Austria, dorfi@astro.univie.ac.at 2ETH-Bibliothek, R¨amistr. 101, Ch-8092 Zurich,¨ [email protected]

Abstract

From simulations of massive stars we find regular low-amplitude radial pulsa- tions can be excited. Such pulsations were encountered during core hydrogen burning as well as during the early core helium burning stage of evolution. For four selected models we present light curves in the V-band. The results are dis- cussed with the aim to guide observations to identify and monitor such regularly pulsating variable massive stars in nature.

Introduction

The COROT mission offers the possibility for high-precision measurements of photometric variations of stars. With increasing accuracy more and more do- mains with light variability are being identified on the Hertzsprung-Russell (HR) diagram. This applies in particular to the high-luminosity regime where the lu- minous blue variables (LBVs) are located. The results from the HIPPARCOS mission (Eyer & Grenon 1997) demonstrated that almost all stars with luminos- ity class lower than III show noticeable light variability (above some 10 mmag) of some as yet unknown origin. From the observational point of view the nature of LBV variability is still not well understood. On timescales of years and months, eruption-like outbursts have been reported but small amplitude, quasi-regular photometric variations have also been detected on timescales of several days (e.g. Nota & Lamers (1997) for a summary on LBV properties). Due to the small number of known LBVs (see Table 1) and the insufficient observational coverage, the physical nature of either the micro-variability or the outbursts is not clear up to date. Based on a search through the literature the latest discovery of a new LBV 58 Pulsations of Luminous Blue Variables

Table 1: Known luminous blue variables (after Bohannan 1997 and Shemmer et al. 2000)

Galaxy # LBVs Milky Way 5 η Car, P Cyg, AG Car, HR Car, HD160529 LMC 6 S Dor, R71, R127, HD269582, R110, R143 SMC 1 R40 M31 4 AE And, AF And, Var A-1, Var 15 M33 4 Var B, Var C, Var 2, Var 38, B416 NGC2403 5 V12, V22, V35, V37, V38 M81 3 I1, I2, I3 M101 3 V1, V2, V10 NGC1058 1 SN1961V

has been reported by Shemmer et al. (2000) for the star B416 located in M33. Hence, photometric satellite missions like COROT could offer a unique oppor- tunity to gain new insights into the physical mechanism operating in the most luminous stars. The theoretical results of the LBV-phase have been strongly debated and the suggested models range from strange mode instabilities (Kiriakidis et al. 1993) to an oscillatory dynamically unstable pulsation for stars near to the Eddington limit (Stothers & Chin 1993). However, none of the models presented are able to explain the micro-variability, the eruptive outbursts or the stellar winds also observed during the LBV phase. According to Leitherer (1997) the observation- ally deduced mass loss rates for LBVs not undergoing a major outburst range −4 −1 −7 −1 from 10 M¯ yr for η Car to 6·10 M¯ yr for R71 (see also Table 1). The final outflow velocities are typically around 200 km/s. In this short contribution we want to focus on the most luminous stars, the so-called LBVs (luminous blue variables) or S Dor variables. Since the observed number of such variables is rather small, theoretical predictions of the variability become more important. We have performed radial non-linear pulsation computations for these stars by applying the full set of the radiation hydrodynamics equations (RHD, e.g. Mihalas & Mihalas 1984). The initial models are computed from usual stellar evolution calculations where the corresponding evolutionary paths are plotted in Fig. 1. The pulsational models investigated belong to massive stars at the stage of either core hydrogen buring or early core helium burning. The numer- ical method used for the non-linear pulsation simulations is described in detail in Dorfi (1998). In the dynamical computations the radiative transfer is per- formed by using the Rosseland-mean OPAL92-opacity for the radiative flux and E. A. Dorfi and A. Gautschy 59

Table 2: Table of LBV model parameters and pulsation properties

Name M L Teff P ∆R/R ∆u us ∆mV 3 [M¯] [10 L¯] [K] [days] [km/s] [km/s] [mag] M20C 20 66 27 100 0.290 0.033 63 0.11 M45 45 537 33 890 0.575 0.080 248 48 0.11 M60H 60 933 34 680 1.611 0.368 598 211 0.27 M60C 60 900 18 000 4.086 0.281 356 114 0.26

the Planck-mean for the corresponding radiation energy equation. The calcu- lated profiles (density, temperature, etc.) are taken as input for a frequency- dependent radiative transfer. After folding the obtained spectra with the filter passbands we get e.g. the V-light curves presented in Fig. 1. To enhance the structural differences the areas below the observable light curves have been filled with grey colour.

Theoretical light curves of LBVs

The properties of the LBV models discussed in this paper are summarized in Table 2. The models exhibit rather regular pulsations (cf. Fig. 1) with pe- riods P shown in column 5. The relative changes of the photospheric radius ∆R/R are given in column 6 indicating that only the masses above 30 M¯ show pronounced radial oscillations with velocity variations ∆u (column 7). These motions go along with strong shock waves running through the atmosphere and the shock speed us is given in column 8. Finally, the last column states the changes in the visual magnitude ∆mV, which allows a direct comparison with observed LBVs. According to the results we emphasize the non-adiabatic nature of the pulsations, where large velocity amplitudes are accompanied by small luminosity variations of the order of about 0.1 mag. For light curves in different passbands we refer to Dorfi & Gautschy (2000). Due to the dom- inance of the radiation pressure these oscillations are mainly confined to the outer layers. From a comparison with linear stability analysis they seem to be rather regular radial modes (for a review e.g. Gautschy & Saio 1995). The temporal evolution of a pulsating star (model: M60H) is plotted in Fig. 2 through the motion of different mass shells over two pulsational cycles. The period of the outer (observable) shell is P = 1.611 days. In contrast to this periodicity the inner shells are oscillating with half of that period, namely P = 0.805 days. This short period corresponds exactly to the pulsation period obtained from a linear stability analysis which causes also the most prominent 60 Pulsations of Luminous Blue Variables

Figure 1: The HRD showing the location of luminous blue variables together with stellar evolution calculations. The grey shaded inserts exhibit the theoretical V-light curves of the LBV model calculations around their mean V-magnitude as a function of the pulsation phase, summarized in Table 2. Note that the magnitude scale of the upper panels is twice that of the lower ones. The upper models have masses of 60 M¯, the lower ones 45 M¯ (right) and 20 M¯ (left), respectively. E. A. Dorfi and A. Gautschy 61 peak in a power spectrum of the inner mass shells (Dorfi & Gautschy 2000). However, the non-linear reactions of the outermost layers induce a longer pe- riodicity visible for an external observer because the motions of the upper at- mosphere are dominated by two strong shock waves seen at phases 0.35 and 0.7. These outward running waves travel at speeds of us = 136 km/s and us = 211 km/s. The Mach numbers of these shocks vary between 5 and 8. After the acceleration event the mass shells are moving basically along ballistic trajectories characterized by the dynamical time scale of the stellar envelope. Clearly, such shock compressions heat the thermal gas and the corresponding features are imprinted on the light curves. Due to the low optical depth in the outer layers no phase shift of the features is introduced in the light curves if observed in different passbands. ] ú R [ s u i ad R

Time in pulsation periods

Figure 2: The paths of different mass shells as a function of the pulsation period show the ballistic behavior of the outermost atmosphere (M = 60 M¯, L = 933 000 L¯, Teff = 34 680 K). For an observer this motion pattern leads to periodic luminosity variations with P = 1.611 days whereas the internal pulsation period is close to 0.8 days (see text for more details).

Atmospheres of pulsating LBVs

The non-linearities of the pulsations increase with the mass of the star and the effects described in this section will occur for masses above 30 M¯. As already seen in the light curves (Fig. 1) and the mass shells (Fig. 2) the at- 62 Pulsations of Luminous Blue Variables mospheric motions are characterized by the shock waves leading to prominent perturbations of the outgoing stellar radiation. Since the waves deposit energy and momentum into the outermost stellar regions, the whole density and tem- perature stratification changes. Consequently, the location of the photosphere defined through 2 4 L = 4πσR Teff (1) can be shifted by large values. Taking the mean value over a pulsational cy- cle, the effective temperature changes e.g. for model M60C from the initial hydrostatic value of Teff = 18 000 K to 20 700 K. Caused by the large amount of radiation pressure, the density inversion of the hydrostatic model cannot be maintained in dynamical models, which allows one to see deeper into the stellar atmosphere. The non-monotonic gradients in the dynamic atmospheres also permit sudden frequency-dependent changes of the optical depth meaning that not all the radiation leaving the star has to come from the same atmospheric layer. Both effects, the compression of the gas by the shock waves as well as the non-monotonic behavior of the photosphere, are responsible for the compli- cated structure of the light curves (Fig. 1). However, as demonstrated in Fig. 2 the underlying mass motions can exhibit a much simpler pattern. Figure 3 illustrates these changes of the observable parameters comparing a hydrostatic with a dynamical atmosphere in the case of a pulsating LBV with 60 M¯. The thick solid line represents the initial model where the dominance of the radiative forces (Fig. 3, lower right panel) generate an unavoidable density and gas pressure inversion situated around 80 R¯ (Fig. 3, left panels). The grey shaded areas depict the upper and lower boundaries of physical quantities during a pulsational cycle. The resulting dynamic structure differs in many as- pects from the hydrostatic stratification since outwards of 30 R¯ shock waves are propagating through the stellar layers. Correspondingly, the density, tem- perature and pressure stratification leads to a different appearance of such a pulsating star relative to a hydrostatic configuration. This dynamic atmosphere (outside of 30 R¯ for model M60C) exhibits radial structures, which cannot be encompassed by the initially hydrostatic structure which is based on static equilibrium conditions. In particular, the aforementioned density inversion can never be supported in such a dynamical environment. The effective photo- spheric radius shrinks from R = 97.4 R¯ to a mean value oscillating around R¯dyn = 74.2 R¯. Since the mean luminosity L has to remain constant we get from Eq. (1) how the temperature of the dynamical atmospheres Tdyn relates to the initial hydrostatic value of Teff by

R Tdyn = Teff ¯ . (2) sRdyn E. A. Dorfi and A. Gautschy 63

Figure 3: The difference between hydrostatic and hydrodynamic models for the model M60C (M = 60 M¯, L = 900 000 L¯, Teff = 18 000 K). The thick solid line plots the initial hydrostatic structure, the grey shaded areas show the variation of the density, velocity as well as gas and radiation pressure during a fully developed pulsational cycle.

Inserting the aforementioned values for the radii we obtain from the last equa- tion Tdyn = 20 620 K in agreement with the averaged value of the non-linear simulations of 20 700 K. Due to the fact that the atmosphere is extended, thin and rather transparent the main effect comes from the inward shift of the effective photosphere allowing such a simple description of an overall bluer ap- pearance of the LBV. The shock waves running through the atmosphere with high Mach numbers (see us in Table 2) affect mainly the light curves as depicted in Fig. 1.

Conclusions

Summarizing a large number of non-linear pulsation computations (e.g. Dorfi et al. 2001) we emphasize that in general the observations of variable stars with large amplitude pulsations cannot be interpreted by simple hydrostatic atmo- spheric models. The stratification of the outer stellar layers is strongly modified by the shock waves, which heat and accelerate the atmospheric material. Al- though the small number of galactic LBVs makes it very unlikely that such 64 Pulsations of Luminous Blue Variables stars will be observed within the COROT mission it should be pointed out that the deduced stellar parameters of stars with large amplitude pulsations have to analyzed through dynamical atmospheres and not by adopting hydrostatic model atmospheres. The short life time of about 104 years (based on statistics) of the LBV phenomenon makes observational studies more difficult. Nevertheless, it is necessary to obtain regular sampled photometric data with accuracies of the order of 0.01 mag together with a time resolution of much less than one day. Since we have found very regular pulsating stars in the non-linear computations a careful monitoring would be essential to verify or disprove these predicted oscillations. If the observations of such stars are not done with an appropriate sampling the predicted pulsations can either be overlooked or interpreted as some kind of irregular flickering. To investigate the pulsational properties of LBVs well-sampled photometric data would be of great importance.

Acknowledgments. This project was supported by the Fonds zur F¨orderung der Wissenschaften (FWF) under project number S7305-AST.

References Bohannan, B. 1997, in Luminous Blue Variables: Massive Stars in Transition, Eds. A. Nota and H. J. G. L. M. Lamers, ASP Conf. Series Vol. 120, p.3 Dorfi, E. A. 1998, 27th Saas Fee Course, Springer, Berlin, p.263 Dorfi, E. A., Gautschy, A. 2000, ApJ 545, 982 Dorfi, E. A., Feuchtinger, M. U., H¨ofner, S. 2001, in Stellar Pulsation - Nonlinear Studies, Eds. D. Sasselov and M. Takeuti, ASSL, Kluwer, Dordrecht, p.137 Eyer, L., Grenon, M. 1997, in HIPPARCOS - Venice ‘97, ESA SP-402, p.467 Gautschy, A., Saio, H. 1995, ARA&A 33, 75 Kiriakidis, M., Fricke, K. J., Glatzel, W. 1993, MNRAS 264, 50 Leitherer, C. 1997, in Luminous Blue Variables: Massive Stars in Transition, Eds. A. Nota and H. J. G. L. M. Lamers, ASP Conf. Series Vol. 120, p.58 Mihalas, D., Mihalas, B. W. 1984, Foundations of Radiation Hydrodynamics, Oxford University Press, New York Nata, A., Lamers, H. J. G. L. M. (Eds.) 1997, Luminous Blue Variables: Massive Stars in Transition, ASP Conf. Series Vol. 120 Shemmer, O., Leibowitz, E. M., Szkody, P. 2000, MNRAS 311, 698 Stothers, R. B., Chin, C.-W. 1993, ApJ 408, L85 Comm. in Asteroseismology Vol. 141, 2002

By-product of the V1162 Ori multisite campaign: the new δ Scuti star GSC 4778 00324

Lampens, P.1, Van Cauteren, P.2,5, Niarchos, P.3, Gazeas, K.3, Manimanis, V.3, Arentoft, T.1,4, Wils, P.5, Bruch, A.6, Garrido, R.7, Shobbrook, R.8

1Koninklijke Sterrenwacht van Belgi¨e, Ringlaan 3, 1180 Brussels, Belgium 2Beersel Hills Observatory (BHO), Laarheidestraat, 3, 1650 Beersel, Belgium 3Section of Astrophysics, Astronomy and Mechanics, University of Athens, 15783 Zografos, Athens, Greece 4Vrije Universiteit Brussel, Pleinlaan 2, 1050 Brussels, Belgium 5Vereniging Voor Sterrenkunde (VVS), Belgium 6Laborat´orio Nacional de Astrof´ısica, C.P. 21, 37500-000, Itajub´a - MG - Brazil 7Instituto de Astrof´ısica de Andaluc´ıa, CSIC, Apdo. 3004, 18080 Granada, Spain 8Research School of Astronomy and Astrophysics, Australian National University, Weston Creek PO, ACT 2611, Australia

Introduction

Among the comparison stars used to monitor the intriguing intermediate am- plitude δ Scuti V1162 Ori two new variable stars have been discovered in the same field. One of them is GSC 4778 00324, for which short-period variability with an amplitude of about 0.01 mag with respect to the ”official” comparison star had been reported previously (Van Cauteren and Lampens, 2000). This variability was also the cause for the first claim of multiperiodicity for V1162 Ori (Hintz et al. 1998). The colour indices obtained by Handler (2000) put the new variable star in the middle of the δ Scuti instability strip. Therefore we decided to monitor GSC 4778 00324 taking as much as possible advantage from the international multisite campaign on V1162 Ori itself (Arentoft et al. 2001). In various cases the effective field-of-view was too small to allow simultaneous measuring with the main target and a few participants secured data by dedi- cating some extra time. However, the bulge of the data discussed here comes from those smaller telescopes with a sufficiently large chip that allow to simul- taneously obtain data on both δ Scuti stars (Belgium and Greece). Figure 1 shows the field with the identification of the known and the new variable stars. The δ Scuti candidate is called ”New 1”. The star called ”New 2” is a new 66 The new δ Scuti star GSC 4778 00324 eclipsing binary with an orbital period very close to an integer multiple of one day. The present contribution reflects the contents of a poster, presented at IAU Symp. 185 (Lampens et al. 2002), in more detail and reports on the final data analysis for the new δ Scuti star.

Figure 1: The field of V1162 Ori (8x12 arcmin2) and the newly detected variable stars.

Data and reduction

The observations were collected between October 1999 and March 2001. All observations used here were made using a standard V filter. Occasionally the R filter was used. The reduction was performed by each individual observing team. In some cases the packages MOMF (Kjeldsen and Frandsen 1992; by TA) or MIRA AP (Axiom Research Inc.; by PVC) were used. Table 1 lists the individual contributions. We mention the name of the observer, the location of the observatory, the number of measurements, an identifier for the comparison star and the mean data string length. In total we have made use of 2949 data points for the analysis of GSC 4778 00324 (V= 10.26; b-y = 0.20). In principle the same comparison star as used for the V1162 Ori campaign was adopted (GSC 4778 00019, V= 9.73; B-V = 1.55). Arentoft et al. (2001) have shown that this comparison star shows nightly zero-point shifts but only at the level of 0.01-0.02 mag such that this will not affect the frequency search in the range of interest (see below). In addition all differential data have been rescaled to the same mean difference in magnitude. The time base length of the complete data sets is 520 days. This allows a theoretical frequency resolution of 1.5/T or 0.0029 cycles/day. P. Lampens et al. 67

Table 1: Logbook of observations of GSC 4778-00324

Name Observatory Nr of data Comp.star Mean length Remark (S=GSC4778) (days)

Bruch OPD 432 S 285 0.158 C2 used Garrido Sierra Nevada 23 S 019 0.073 Lampens et al. Hoher List 53 S 019 0.080 Niarchos et al. Athens Univ. 2467 S 019 0.103 Shobbrook Siding Spring 13 S 019 0.120 Sterken ESO 105 S 019 0.093 Van Cauteren Beersel Hills 1054 S 019 0.106 All 4147 S 019 95.7 hrs T=520 d

Frequency analyses

The same criteria as for the final data selection of V1162 Ori were applied also in this case (Arentoft et al. 2001). The analysis was performed in three steps: (a) on the (northern) winter 1999-2000 data set, (b) on the (northern) winter 2000-2001 data set, and (c) on the complete data set. In all cases the data points were weighted in proportion to the inverse square of the residuals on a night-to-night basis. In each set up to three frequencies could be identified. We are confident that the two most dominant frequencies are secure as they are common to each individual analysis and have signal-to-noise ratios well above 4. Figure 2 illustrates the frequency search for the complete data set. Figure 2a illustrates the periodogram of the original data, the subsequent frames show the periodograms after prewhitening with the most dominant frequency of the previous frame. Whereas there is no doubt about the main frequency at 12.28 cycles/day, some ambiguity arises concerning the second frequency. In our previous analysis (Lampens et al. 2002) we adopted the frequency at 16.62 c/d as the most probable second frequency since it appeared to be very slightly more dominant in the yearly analyses in which all the data points were treated equally. However, Figure 2b now shows that the 1/day alias frequency at 15.62 cycles/day is slightly stronger (the corresponding amplitude is larger by 1 mmag). This, together with the symmetry of the alias pattern in Figure 2b (5 distinguishable peaks at each side of the adopted peak), causes us to identify the latter as the most probable second frequency. After simultaneous prewhitening with the two most dominant frequencies, the remaining standard deviation is 6.7 mmag. The nightly rms values of the data usually range from 3 to 10 mmag. Considering the possible contribution of slight zero-point shifts in the shortest data sets, this may reflect the overall precision of the differential data. But the periodogram in Figure 2c still exhibits additional power in the range 10- 15 cycles/day. Possible values for a third frequency together with the most 68 The new δ Scuti star GSC 4778 00324 dominant frequencies are summarized in Table 2, presenting the final results of the analysis. Depending on the choice for the second frequency (either 15.62 or 16.62 cycles/day) the third frequency assumes quite distinct values and its value is therefore uncertain even though it was identified with a signal-to-noise ratio larger than 4 in most cases. The ratio between the two most dominant frequencies is either 0.74 (16.62 c/d) or 0.79 (15.62 c/d).

0.004 0.004

0.0035 0.0035

0.003 0.003

0.0025 0.0025

0.002 0.002

0.0015 0.0015

0.001 0.001

0.0005 0.0005

0 0 0 5 10 15 20 25 0 5 10 15 20 25

0.004 0.004

0.0035 0.0035

0.003 0.003

0.0025 0.0025

0.002 0.002

0.0015 0.0015

0.001 0.001

0.0005 0.0005

0 0 0 5 10 15 20 25 0 5 10 15 20 25

Figure 2: Periodograms of successive frequency searches (a-d: from left to right).

Figure 3 illustrates two of the associated mean light curves. The V am- plitudes are of order of several millimags only. It must be remarked that the solutions presented here are incomplete as there exists at least one string of data for which the fit is not quite satisfactory (on JD 2451575). The best P. Lampens et al. 69

−0.05

−0.04

−0.03

−0.02

−0.01

0

0.01

0.02

0.03

0.04

0.05 0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 0.8 0.9 1 −0.05

−0.04

−0.03

−0.02

−0.01

0

0.01

0.02

0.03

0.04

0.05 0 0.1 0.2 0.3 0.4 0.5 0.6 0.7 0.8 0.9 1

Figure 3: Phase diagrams at the frequencies of 12.28 c/d (left) and 15.62 c/d (right) after prewhitening for the other frequency.

solution in terms of largest reduction of the standard deviation is flagged (1) in Table 2. The long, continuous data set obtained in Brazil (Table 1) was used as an independent test of the proposed frequency solution. Figure 4 shows the yet unused data as well as the two-frequency fit from solution (1). We readjusted the phases in order to have a perfect match between the calculated and the 70 The new δ Scuti star GSC 4778 00324

Table 2: Frequency analyses for GSC 4778 00324

Season Nr Ident. Freq. Ampl. Phase S/N Remark c/d mag 2πrad

1999-2000 1546 f1 12.2801 0.0041 0.960 7.8 weighted data f2 15.6157 0.0034 0.927 5.6 RSS = 6.3 mmag f3 11.7173 0.0023 0.317 4.3 uncertain 2 2000-2001 1403 f1 12.2802 0.0048 0.671 7.4 weighted data 2 f2 16.6215 0.0038 0.827 6.9 RSS = 6.7 mmag 2 f3 13.2568 0.0027 0.367 4.4 uncertain (1) 1999-2001 2949 f1 12.2810 0.0041 0.136 8.1 weighted data f2 15.6169 0.0036 0.135 6.8 strong alias at 16.62 c/d f3 10.5852 0.0021 0.421 4.7 RSS = 6.3 mmag

(2) 1999-2001 2949 f1 12.2810 0.0042 0.118 8.2 weighted data f2 16.6171 0.0027 0.611 4.8 strong alias at 15.62 c/d f3 12.9543 0.0021 0.802 3.9 RSS = 6.5 mmag 2 Phases corrected for the origin difference between data sets

observed light curve. The standard deviation of the residuals dropped by 24% to the level of 6.4 mmag. This shows that the quality of the presently avail- able data set is probably not sufficient to permit the identification of additional frequencies with amplitudes as small as 2-3 mmag.

−0.02

−0.015

−0.01

−0.005

0

0.005

0.01

0.015

0.02

0.025 514.68 514.7 514.72 514.74 514.76 514.78 514.8 514.82

Figure 4: Observations from OPD, Brazil, and two-frequency fit (JD 2451514). P. Lampens et al. 71

Conclusions

A by-product of the international multisite campaign on V1162 Ori is the serendipitous discovery of a pulsating star of very small amplitude and with variable light curve among the comparison stars in the field. GSC 4778 00324 has the colours of a typical δ Scuti star. This star was frequently observed during the 1999-2000 and 2000-2001 seasons, mainly from Greek and Belgian sites. The final frequency analysis, based on assigning weights to the data, confirms the δ Scuti type and the presence of at least two and probably three frequencies in the range 10-15 cycles/day. V amplitudes are of order 4 mmag or smaller. The proposed frequency solutions are incomplete but the presently available data set does not permit the identification of more frequencies. In it- self it is remarkable that observations from non-ideal sites in the vicinity of large European cities (Athens and Brussels) allow to detect and monitor such subtle variations and can effectively contribute to the analysis of these small-amplitude multiperiodic pulsating stars.

Acknowledgments. This research was supported by project G.0265.97 of the Belgian Fund for Scientific Research (FWO) - Flanders. T. Arentoft (TA) acknowledges financial support by the Flemish Ministry for Foreign Policy, European Affairs, Science and Technology, under contract BIL 98/11/52 and by the National Research Foundation of South Africa. We thank C. Sterken for providing us with observations of the new variable star. R.R. Shobbrook, Honorary Visiting Fellow at the Australian National University, acknowledges support from Mount Stromlo and Siding Spring Observatories. P. Lampens and P. Van Cauteren (PVC) thank Prof. Dr. W. Seggewiss for granting telescope time at the Hoher List Observatory (University of Bonn, Germany). We also thank M. Sperl for updated versions of the programme Period98.

References Arentoft, T. et al. 2001, A&A 374, 1056 Handler, G. 2000, IBVS 4817 Hintz, E., Joner, M. & Kim, C. 1998, PASP 110, 689 Kjeldsen, H. & Frandsen, S. 1992, PASP 104, 413 Lampens, P. et al. 2002, poster in Proceedings of IAU Coll. 185 on Radial and Nonradial Pulsations as Probes of Stellar Physics, Leuven, Belgium, 26-31 July, 2001. ASP Conference Series, eds. C. Aerts, T. Bedding & J. Christensen-Dalsgaard, in press. Van Cauteren, P. & Lampens, P. 2000, IBVS 4849 Comm. in Asteroseismology Vol. 141, 2002

Strømgren photometry of SX Phe = HD 223065

A. Stankov1, D. Sinachopoulos2, E. Elst3, M. Breger1

1 Institut fur¨ Astronomie, Turk¨ enschanzstrasse 17, 1180 Vienna, Austria 2 Institute of Astronomy and Astrophysics, National Observatory of Athens, I. Metaxa and Bas. Pavlou, 15236 Athens, Greece 3 Koninklijke Sterrenwacht van Belgi¨e, Ringlaan 3, 1180 Uccle, Belgium

Abstract

We present the analysis of more than 220 photometric measurements of the variable star SX Phe. The observations were made with the Johnson V and the four Strømgren (u, v, b, y) filters. An analysis of the parameters δc1 and δm1 is presented. The behavior of δc1 during the luminosity variations of the star corresponds to that expected from the standard luminosity calibrations by Crawford (1979) with δV = −10δc1. The behavior of the metallicity index, δm1, does not correspond to the expectations from the standard calibrations.

Introduction

SX Phe stars are a Pop. II subgroup of the δ Scuti stars. In order to shed light on the variability of SX Phe and related stars, many studies of SX Phe have been undertaken since the report on our previous observations (Elst 1978). Rodriguez et al. (1990) used Strømgren photometry to study the behavior of the δm1 index for SX Phe stars. In addition, Garrido et al. (1990) used uvby photometry for the identification of pulsation modes (modal discrimination) of the pulsating δ Scuti stars. The observations presented here were obtained to provide deeper understanding of the variability of SX Phe in the Strømgren filters.

Observations and data reduction

The ESO 50cm telescope was used to monitor SX Phe during the nights of 1992 October 18 and 31, as well as 1992 November 1. HD 223011 was used as A. Stankov et al. 73 SX Phoen icis

6.8

7.0

7.2

7.4

14.5 14.6 14.7 14.8

6.8

7.0 V

7.2

7.4

27.5 27.6 27.7 27.8

6.8

7.0

7.2

7.4

28.5 28.6 28.7 28.8 HJD [2448900+]

Figure 1: Photometry of SX Phe in V a comparison star. The photometer of the 50cm ESO telescope is only a single- channel photometer. Consequently, the measurements through the four filters were not obtained simultaneously, as was the case in Garrido et al. (1990). Due to technical problems, we had to use two different photomultipliers for our 74 Strømgren photometry of SX Phe = HD 223065 observations. During the first night we used a HAMAMATSU and during the other two nights a EMI 9789QB tube. In addition, we observed 16 bright photometric standard stars each night for transforming our observations into the standard uvby system. The SNOPY software on theHP 1000 computer was used for this transformation. The com- parison star did show any significant variability during this period. From the statistical fluctuations of its reduced magnitudes and indices we estimate that the accuracy of a single HD 223011 measurement is σV = 0.006, σ(b−y) =

0.005, σm1 = 0.008, σc1 = 0.007 mag.

0.12 0.100

0.095 0.10

0.090 0.08 1 1 c m 0.085

G G 0.06 0.080

0.04 0.075

0.02 0.070 7.5 7.4 7.3 7.2 7.1 7.0 6.9 6.8 0.06 0.08 0.10 0.12 0.14 0.16 0.18 0.20 V (b-y)

Figure 2: Left pannel: δc1 versus V ; dotted line: expected slope from calibrations from Crawford. Right pannel: δm1 versus (b − y).

Since SX Phe is about 0.9 magnitudes fainter than the comparison star, we expect that the accuracy of measurement is lower for SX Phe. More than 10000 counts per exposure (30 seconds each) of SX Phe were collected in each filter during the first night using the HAMAMATSU photo-multiplier. The efficiency of the EMI 9789QB tube was somewhat lower. Table 1 lists the measurements of SX Phe, while Figure 1 shows the light curves of SX Phe obtained during three nights.

Variations of the Stromgren parameters during the pulsational cycles

Crawford (1979) has presented calibrations of these parameters, based on a large number of different stars. A pulsating star such as SX Phe changes its luminosity and temperature during its pulsation cycle, but is not affected by a variable abundance. This provides an excellent opportunity to check these calibrations whether also apply the atmospheric changes during pulsation. A. Stankov et al. 75

We can now turn to the variations of the metallicity parameter for A/F stars, δm1, and the luminosity parameter, δc1. The observed variations are shown in Fig. 2. For δc1, we derive a slope of δV /δc1 = 9.8, in good agreement with the expected value of 9.0 derived for different stars by Crawford. The two slopes are also shown in Fig. 2, where the Crawford slope is displaced vertically for greater clarity. The right panel of Fig. 2 shows the relation between δm1 and color, b − y. A definite non-zero slope is found. A comparison with the standard values given by Crawford (1979), describing the variation of δm1 with temperature, shows very poor agreement. A logical, but incorrect, interpretation would be that the metal content changes during the cycle! We note here that SX Phe is extremely metal-poor and that the calibrations were set up for stars with near-solar abundances. The calibrations, therefore, cannot be expected to hold for SX Phe. The behavior of SX Phe confirms the observed behavior of two other (metal- poor) SX Phe stars, KZ Hya and CY Aqr (Rodriguez et al. 1990).

References Crawford, D. L. 1979, AJ 84, 1858 Elst E. 1978, A&AS 32, 161 Garrido, R., Garcia-Lobo, E., Rodriguez, E. 1990, A&A 234, 262 Rodriguez, E., Rolland, A., Lopez de Coca, R., et al. 1990, Rev. Mexic. Astron. Astrof. 21, 386 76 Strømgren photometry of SX Phe = HD 223065

Table 1: Measurements of SX Phe

HJD +244 8900 V (b-y) m1 c1 14.4976 7.400 .185 .076 .781 14.5134 7.417 .163 .096 .763 14.5280 7.097 .097 .123 .920 14.5332 7.069 .108 .123 .938 14.5376 7.151 .122 .126 .900

14.5426 7.263 .152 .098 .863 14.5468 7.337 .162 .097 .821 14.5518 7.397 .179 .086 .793 14.5560 7.439 .181 .089 .771 14.5605 7.465 .187 .082 .760

14.5650 7.471 .182 .080 .758 14.5691 7.441 .174 .081 .750 14.5731 7.331 .135 .093 .786 14.5771 6.989 .022 .138 .948 14.5820 6.792 .057 .138 1.036

14.5842 6.851 .074 .138 1.037 14.5882 7.033 .114 .126 .963 14.5922 7.181 .140 .109 .905 14.5962 7.301 .161 .096 .851 14.6001 7.371 .175 .082 .821

14.6040 7.417 .182 .084 .792 14.6079 7.450 .188 .080 .773 14.6118 7.467 .188 .077 .767 14.6166 7.465 .180 .088 .745 14.6206 7.433 .168 .085 .757

14.6244 7.344 .143 .107 .777 14.6282 7.166 .087 .118 .882 14.6320 6.994 .075 .137 .968 14.6360 7.007 .095 .126 .977 continued on next page A. Stankov et al. 77 continued from previous page HJD +244 8900 V (b-y) m1 c1 14.6399 7.090 .114 .123 .947

14.6438 7.186 .139 .125 .892 14.6476 7.273 .155 .115 .854 14.6515 7.335 .163 .116 .823 14.6564 7.387 .173 .102 .809 14.6604 7.411 .183 .096 .783

14.6643 7.429 .182 .092 .781 14.6682 7.431 .175 .111 .754 14.6720 7.428 .174 .105 .759 14.6759 7.412 .168 .107 .761 14.6798 7.383 .157 .114 .780

14.6839 7.327 .144 .115 .815 14.6880 7.251 .120 .130 .857 14.6918 7.192 .116 .131 .890 14.6958 7.173 .121 .127 .903 14.6997 7.193 .136 .115 .898

14.7037 7.238 .141 .126 .857 14.7076 7.291 .156 .115 .835 14.7115 7.339 .166 .109 .817 14.7153 7.380 .172 .106 .800 14.7209 7.433 .180 .098 .783

14.7249 7.453 .182 .098 .776 14.7289 7.461 .181 .101 .757 14.7342 7.442 .171 .098 .764 14.7391 7.353 .135 .113 .789 14.7436 7.084 .068 .134 .925

14.7481 6.899 .071 .140 1.012 14.7527 6.972 .100 .140 .998 14.7574 7.151 .135 .136 .905 14.7622 7.296 .161 .107 .860 continued on next page 78 Strømgren photometry of SX Phe = HD 223065

continued from previous page HJD +244 8900 V (b-y) m1 c1 14.7647 7.346 .171 .101 .829 14.7694 7.401 .183 .091 .812 14.7742 7.451 .186 .105 .771 14.7789 7.470 .191 .094 .758 14.7834 7.466 .184 .088 .763 27.5212 7.411 .200 .085 .783

27.5309 7.460 .197 .084 .765 27.5324 7.460 .190 .092 .757 27.5374 7.419 .166 .103 .753 27.5390 7.381 .166 .091 .771 27.5438 7.209 .101 .117 .868 27.5455 7.097 .072 .130 .936 27.5496 6.974 .085 .132 .999 27.5512 6.982 .095 .128 1.002 27.5560 7.078 .121 .126 .959 27.5575 7.119 .128 .127 .934

27.5592 7.167 .136 .119 .920 27.5640 7.276 .154 .117 .868 27.5655 7.306 .159 .114 .855 27.5672 7.329 .164 .117 .832 27.5719 7.379 .171 .117 .800

27.5734 7.401 .168 .112 .801 27.5751 7.404 .179 .104 .794 27.5799 7.421 .191 .093 .782 27.5818 7.427 .181 .106 .771 27.5834 7.434 .177 .113 .755

27.5884 7.433 .171 .109 .775 27.5900 7.425 .174 .106 .770 27.5916 7.423 .169 .111 .762 27.5963 7.399 .156 .116 .775 27.5979 7.378 .146 .126 .787

continued on next page A. Stankov et al. 79 continued from previous page HJD +244 8900 V (b-y) m1 c1 27.6007 7.337 .134 .129 .818 27.6046 7.254 .121 .133 .859 27.6062 7.228 .119 .132 .877 27.6078 7.208 .119 .134 .886 27.6124 7.188 .125 .130 .895

27.6140 7.194 .133 .120 .906 27.6157 7.207 .131 .134 .881 27.6199 7.251 .142 .121 .877 27.6215 7.270 .143 .127 .864 27.6232 7.289 .155 .119 .850

27.6280 7.346 .164 .106 .825 27.6295 7.360 .170 .104 .811 27.6312 7.382 .169 .105 .810 27.6350 7.414 .173 .110 .783 27.6366 7.426 .181 .103 .773 27.6382 7.443 .180 .103 .768 27.6400 7.445 .180 .107 .771 27.6416 7.448 .189 .095 .765 27.6433 7.461 .180 .103 .759 27.6451 7.471 .174 .105 .766

27.6468 7.467 .175 .112 .741 27.6484 7.459 .178 .101 .759 27.6533 7.410 .155 .105 .783 27.6548 7.377 .134 .126 .772 27.6564 7.323 .127 .118 .801

27.6604 7.093 .079 .114 .935 27.6620 6.990 .061 .138 .970 27.6636 6.943 .065 .136 1.001 27.6684 6.960 .089 .139 1.005 27.6700 7.004 .107 .130 .991

27.6716 7.069 .111 .141 .945 continued on next page 80 Strømgren photometry of SX Phe = HD 223065

continued from previous page HJD +244 8900 V (b-y) m1 c1 27.6762 7.214 .151 .112 .894 27.6778 7.255 .155 .121 .863 27.6794 7.296 .160 .120 .836 27.6842 7.374 .178 .105 .808

27.6858 7.389 .183 .099 .806 27.6874 7.403 .188 .103 .787 27.6923 7.456 .188 .094 .784 27.6939 7.464 .188 .104 .760 27.6978 7.464 .193 .091 .758

27.6994 7.459 .192 .088 .765 27.7033 7.439 .169 .099 .759 27.7050 7.406 .164 .105 .755 27.7089 7.258 .120 .109 .807 27.7107 7.119 .068 .113 .895

27.7146 6.804 .051 .143 1.021 27.7163 6.808 .065 .139 1.035 27.7201 6.916 .094 .141 1.022 27.7218 6.982 .118 .132 .978 27.7257 7.133 .142 .115 .937 27.7296 7.249 .156 .113 .869 27.7335 7.332 .173 .104 .823 27.7374 7.377 .185 .097 .811 27.7414 7.411 .189 .094 .773 27.7454 7.428 .190 .099 .759

27.7493 7.426 .191 .085 .767 27.7533 7.411 .183 .089 .761 27.7573 7.376 .167 .106 .757 28.5088 7.373 .181 .109 .769 28.5104 7.388 .190 .094 .774

28.5121 7.406 .181 .106 .755 28.5171 7.418 .193 .096 .733 continued on next page A. Stankov et al. 81 continued from previous page HJD +244 8900 V (b-y) m1 c1 28.5188 7.425 .186 .094 .742 28.5204 7.414 .189 .096 .724 28.5254 7.391 .169 .099 .747

28.5270 7.369 .182 .078 .761 28.5286 7.356 .157 .101 .765 28.5337 7.240 .132 .098 .845 28.5353 7.202 .111 .118 .869 28.5369 7.157 .118 .105 .902

28.5410 7.120 .125 .102 .930 28.5426 7.125 .122 .110 .922 28.5442 7.144 .116 .126 .897 28.5489 7.192 .136 .111 .886 28.5505 7.216 .141 .109 .881

28.5521 7.232 .149 .106 .860 28.5537 7.256 .151 .107 .843 28.5579 7.304 .163 .102 .813 28.5595 7.319 .167 .100 .805 28.5611 7.332 .173 .094 .801

28.5656 7.374 .174 .096 .785 28.5672 7.381 .178 .099 .765 28.5691 7.398 .180 .088 .782 28.5729 7.412 .184 .087 .774 28.5745 7.420 .177 .097 .760 28.5761 7.421 .179 .099 .750 28.5777 7.431 .176 .095 .756 28.5824 7.429 .168 .094 .763 28.5841 7.416 .173 .090 .759 28.5857 7.405 .163 .099 .761

28.5900 7.343 .146 .098 .791 28.5916 7.302 .138 .100 .810 28.5934 7.257 .122 .109 .837 continued on next page 82 Strømgren photometry of SX Phe = HD 223065

continued from previous page HJD +244 8900 V (b-y) m1 c1 28.5980 7.118 .100 .118 .913 28.5996 7.091 .095 .122 .927

28.6012 7.072 .099 .121 .939 28.6050 7.092 .108 .125 .933 28.6067 7.110 .124 .118 .915 28.6084 7.157 .117 .125 .911 28.6130 7.250 .149 .099 .857

28.6146 7.281 .150 .108 .823 28.6162 7.309 .158 .094 .826 28.6202 7.353 .175 .090 .789 28.6218 7.384 .165 .095 .797 28.6234 7.394 .175 .089 .785

28.6280 7.432 .185 .074 .778 28.6296 7.439 .189 .075 .763 28.6312 7.453 .182 .079 .756 28.6351 7.459 .170 .089 .744 28.6367 7.464 .175 .096 .736

28.6383 7.456 .172 .087 .750 28.6429 7.377 .144 .099 .753 28.6445 7.316 .125 .097 .791 28.6461 7.219 .087 .107 .839 28.6499 6.847 .039 .126 1.005

28.6516 6.805 .051 .120 1.042 28.6532 6.809 .066 .116 1.054 28.6548 6.854 .065 .137 1.040 28.6596 7.057 .117 .117 .959 28.6613 7.110 .134 .111 .929 28.6629 7.175 .134 .115 .897 28.6672 7.299 .164 .099 .838 28.6688 7.338 .167 .094 .827 28.6740 7.411 .176 .083 .810 continued on next page A. Stankov et al. 83 continued from previous page HJD +244 8900 V (b-y) m1 c1 28.6757 7.417 .187 .083 .788

28.6796 7.456 .184 .088 .762 28.6812 7.469 .186 .089 .761 28.6857 7.473 .195 .070 .775 28.6874 7.471 .179 .091 .752 28.6913 7.446 .164 .090 .757

28.6930 7.407 .167 .079 .773 28.6968 7.290 .123 .109 .812 28.6984 7.191 .103 .104 .872 28.7030 6.953 .073 .132 .979 28.7046 6.957 .076 .138 .984

28.7092 7.042 .100 .132 .970 28.7108 7.091 .112 .129 .940 28.7155 7.212 .149 .102 .887 28.7172 7.252 .147 .113 .855 28.7213 7.327 .168 .114 .834

28.7259 7.375 .188 .101 .802 28.7299 7.408 .192 .098 .781 28.7339 7.423 .192 .094 .786 28.7380 7.433 .174 .121 .753 28.7420 7.429 .179 .107 .757 28.7459 7.415 .170 .114 .771 28.7499 7.377 .160 .106 .796 28.7538 7.323 .142 .126 .816