A&A 448, 697-701 (2006)
E. Chapellier - D. Le Contel - J. M. Le Contel - P. Mathias - J.-C. Valtier
Observatoire de la Côte d'Azur, Département GEMINI - UMR 6203, BP 4229, 06304 Nice Cedex 4, France
Received 12 July 2005 / Accepted 24 October 2005
High resolution and high signal to noise ratio spectroscopic observations of the classical Cephei star Peg were obtained between 1991 and 2005. The analysis of these data combined with previously published results shows that Peg is a spectroscopic binary with an orbital period of 370.5 d. We discovered three new frequencies in addition to the well-known 6.5897 d-1 (0.15175 d) one. That at 6.01 d-1 is a typical Cephei frequency. The two others at 0.68 and 0.87 d-1 are similar to the high degree g-mode frequencies found in SPB stars. Thus, Peg is a hybrid Cephei-SPB star. Its position in the HR diagram is compatible with such a status.
In addition, a small increase of the main period has been detected between the 1995 and 2005 observations.
Key words: stars: binaries: spectroscopic - stars: oscillations - stars: individual: Cephei - stars: variables: general
Peg (HD 886, B2 IV, mv = 2.83) is one of the so-called "classical'' Cephei stars. It has been assigned to this group by McNamara (1953) who demonstrated its variability in radial velocity with a single period of 0.1517 d and an amplitude of 3.5 km s-1. Its photometric variability, with the same period as the spectroscopic one and corresponding to an amplitude of 15 mmag in yellow light, was detected by Williams (1954). He showed that, as in the other stars of the group, the light curve lags about behind the radial velocity one. Jerzykiewicz (1970), observing the star through UBV Johnson's filters, derived the amplitudes U=27 mmag, B=18 mmag and V=17 mmag. He also observed that the U-B colour index was smallest around light maxima. A more precise determination of the period was obtained by Sareyan et al. (1975) and its constancy was confirmed by Valtier et al. (1985). Smith & McCall (1978) and Cugier et al. (1994) attributed this pulsation to a radial mode. From high time and high spectral resolution observations obtained with an electronographic camera, Duchesne et al. (1967) showed that line profile and equivalent width variations were associated with the light and radial velocity period.
Harmanec et al. (1979) proposed a binary period of 6.83 d which has not been confirmed yet.
We have obtained new spectroscopic observations from which new results could be derived both on the binary nature and on the pulsation of Peg. Section 2 presents the observations, Sect. 3 examines the binarity and the pulsation characteristics of the star. The results are discussed in Sect. 4.
|Figure 1: Fit of heliocentric radial velocities (HRV) with the 370.5 d orbital period. Open circles represent old observations; dots correspond to our new observations.|
Our data set consists of 5 series of observations, i.e., 2 consecutive nights in 1991,
5 nights spread over 36 days in 1992, 8 consecutive nights in 1993, 8 nights
spread over 12 days in 1995 and 2 nights in 2005 (Table 1). The radial velocity curves show that the variations cannot be accounted for by a single frequency contrary to what has
been claimed up to now. We searched for new frequencies in the two most abundant series of data obtained
in 1993 and 1995 using the Period04 software (Lenz & Breger 2005).
We have not prewhitened the data for orbital motion because the uncertainties
on the shape of the curve are still too large.
We first analyzed the data year by year.
In 1993, four frequencies are detected at respectively 6.591, 0.659, 0.840 and 5.99 d-1
Figure 2 shows the corresponding Fourier spectra.
|Figure 2: Amplitude spectra for the 1993 data. The successive panels show the periodograms after different stages of prewhitening. The dotted line indicates the position of the detected frequency given in the upper right corner. The bottom one shows the periodogram for the residuals.|
In 1995 our data lie on the steep ascending branch of the orbital curve and the velocity increases by more than 10 km s-1 in 12 days. Consequently in our periodogram a low non-pulsational frequency appears with a large amplitude. After prewhitening these data for this frequency, we are able to compensate for the orbital movement and then look for pulsation frequencies. The same four frequencies are again present at 6.589, 0.670, 0.872 and 6.02 d-1 (Table 3). Figure 3 presents the corresponding Fourier spectra. To make the figure clearer, only the periodograms corresponding to the pulsations frequencies are presented. Using a straight line to prewhiten the 1995 data for the orbital motion, we find the same four frequencies but with small differences for the two low ones: f2 = 0.701 d-1, K2=1.65 km s-1 and f3=0.892 d-1, K3=0.82 km s-1. Note that we also tried to prewhiten the data for the orbital curve calculated in Sect. 3.1. Unfortunately, this curve is not accurate enough and a slight slope remains so the data had to be prewhitened again for a low frequency.
If we analyze together the two sets of data of 1993 and 1995 after prewhitening for a low frequency, we find again the four frequencies at f1=6.5898, f2=0.686, f3=0.866 and f4=6.01 d-1, with respective amplitudes K1=3.46, K2=1.41, K3=1.12 and K4=0.43 km s-1. The number of cycles between the two sets is so large that large uncertainties in the cycle count occur.
Considering the results of the two data sets and the different methods used to remove the orbital movement, we finally adopt the main values d-1, d-1 and d-1 for the pulsation frequencies. For f1, we retained the 6.5897763 d-1 value given by Valtier et al. (1985). Figures 4 and 5 display the adjustment of the observational data using the values provided in Table 3.
The variability of equivalent widths and line profiles discovered by Duchesne et al. (1967) were confirmed by Le Contel (1968, 1975) and by Smith & McCall (1978). Le Contel & Morel (1985) related the existence of weak components in the wings of the Si III lines to the development of shock waves in the atmosphere of Peg. In our new data, weak variations of the equivalent widths (at a level of 2%) and of the Full Widths at Half Maximum are observed, correctly folded with the f1 frequency. However, the spectral resolution is not sufficient to allow a direct comparison with Le Contel's observations.
As in other Cephei variables, the constancy of the main pulsation period (6.5897 d-1; 0.15175 d) of Peg has been checked by different authors (Van Hoof 1970; Sareyan et al. 1975; Smith & McCall 1978; Lane & Percy 1979; Valtier et al. 1985). They all confirmed that the pulsation period remained constant from 1931 to 1985. Conversely, Koubsky et al. (1981), including in their analysis old data from 1899 to 1931, suggested that the period might have increased at a rate of 0.06 s century-1. Our present data obtained between 1991 and 1995 agree with the Valtier et al. (1985) ephemeris. However, our two 2005 data sets show a phase lag of +0.021 and +0.024 d with respect to this last ephemeris which corresponds to an increase in the main period of at least 0.08 s between 1995 and 2005.
|Figure 4: Adjustment of the 1993 observational data with the frequencies, amplitudes and phases given in Table 3.|
Such abrupt variations have been detected in several other Cephei stars (Chapellier 1985). In some of these stars such a phenomenon has been explained by the light-time effect associated with a long period (several tens of years) companion (Mendenhall 1930; Odell 1985; Chapellier 1990; Pigulski 1992; Pigulski & Boratyn 1992; Pigulski 1993). In the case of Peg, since no other companion was observed by McAlister et al. (1989) and since no observation has been obtained between 1995 and 2005, no interpretation of this period variation can be given.
The first result arising from our study is the discovery of the binarity of Peg. The length of the orbital period (about one year) and the shape of the orbital curve explain why it had not been detected earlier. The value of the orbital elements should be improved by new data obtained during the descending branch of the curve.
The second important result is that Peg is a multi-periodic variable: in addition to the well-known f1 frequency, three new frequencies were detected in two independent data sets as well as in the whole data set. Moreover, we cannot exclude the existence of other frequencies associated with smaller amplitudes. The two high frequencies, i.e. f1 and f4, are related to the classical Cephei instability while f2 and f3 are too low, but are well in the range of the SPBs typical frequencies.
The Cephei type variations in Peg were detected more
than fifty years ago while the SPBs variations are detected in our data sets
separated by about 2 years.
They are also observed in our 1992 data set although there are not enough data points
to perform a separate detailed analysis.
If we reanalyze all the previously published data between 1952 and 1978 and
if we prewhiten them for the well-known f1 frequency, a dominant peak appears
at 0.69 d-1 corresponding to f2.
Therefore we consider that the two types of variations are always present in Peg.
|Figure 5: Adjustment of the 1995 observational data with the frequencies, amplitudes and phases given in Table 3.|
Since the mechanism responsible for the instability of Cephei and SPB stars has been discovered, theoretical models indicate the existence of a small region in the HR Diagram where the two types of variations can coexist (Pamyatnykh 1999). This region coincides with the extreme cool limit of the Cephei instability zone of slightly evolved stars. From its classification (B2IV) and its temperature, Peg is among the coldest Cephei stars and is very close to the hotter limit of the SPBs instability zone as determined by Pamyatnykh (1999). Thus, it is not surprising that Peg appears to be a hybrid star. Such SPBs pulsations have been observed in another Cephei star, Eri, by Handler et al. (2005). Using the data from a very large observational campaign in 2002-2003, they detected a low amplitude frequency at 0.432 d-1 in photometry. This frequency has not been detected in spectroscopy (Aerts et al. 2004), but was found again in the 2003-2004 photometric campaign (Jerzykiewicz et al. 2005). On the other hand, high frequency variations have also been observed in two SPB stars, 53 Psc (Wolf 1987; Mathias & Waelkens 1995; Le Contel et al. 2001) and Her (Chapellier et al. 1987, 2000) but the phenomenon is spurious.
From a theoretical point of view, the existence of excited p and g modes should allow the simultaneous study of both the external and the internal zones of the star. It may also help to refine the limits of the SPBs instability zone. Our data do not allow mode identification but if f1 is actually radial, f2 is necessarily non-radial. Simultaneous photometric and high resolution spectroscopic observations are requested to give more precise frequency values, mode parameters and a better determination of the orbital elements.
In addition, a small increase in the main period was detected between the 1995 and 2005 observations.
The authors are grateful to the staff of the Haute Provence Observatory for their technical assistance.
|2 400 000+||exposure||length||spectra||line|
|48 538.4071||8.0-15.0||0.08||9||SiIII 4567 Å|
|48 539.5095||15.0||0.15||11||SiIII 4567 Å|
|48 851.5500||1.0||0.16||103||SiIII 4567 Å|
|48 854.7287||2.5-3.3||0.18||55||SiIII 4567 Å|
|48 857.5272||2.0-8.3||0.22||57||SiIII 4567 Å|
|48 884.5563||3.0-6.0||0.23||56||SiIII 4567 Å|
|48 887.4727||3.3-10.0||0.11||5||SiIII 4567 Å|
|49 201.5383||1.7-10.0||0.16||40||SiIII 4567 Å|
|49 202.5371||0.8-10.0||0.17||45||SiIII 4828 Å|
|49 203.5476||1.2-7.5||0.17||73||SII 5320 Å|
|49 204.5271||2.8-10.0||0.17||19||SiIII 4567 Å|
|49 205.5616||1.7-10.0||0.14||29||SiIII 4828 Å|
|49 206.5616||1.7-10.0||0.15||34||SiIII 4828 Å|
|49 207.5367||2.0-6.7||0.18||44||SII 5453 Å|
|49 208.5383||3.2-10.0||0.19||32||SiIII 4567 Å|
|49 969.4953||4.5-10.8||0.32||59||SiIII 4567 Å|
|49 970.4938||3.3-8.0||0.32||64||SiIII 4567 Å|
|49 974.5765||9.0-15.0||0.33||16||SiIII 4567 Å|
|49 975.5270||3.8-15.0||0.25||40||SiIII 4567 Å|
|49 976.5134||3.7-7.2||0.25||56||SiIII 4567 Å|
|49 977.5239||3.3-8.3||0.27||47||SiIII 4567 Å|
|49 978.5221||3.3-11.7||0.25||50||SiIII 4567 Å|
|49 981.4696||5.4-11.7||0.29||56||SiIII 4567 Å|
|53 418.2773||5.0-6.0||0.05||12||SiIII 4567 Å|
|53 536.5505||1.7||0.08||66||SiIII 4567 Å|