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Astron. Astrophys. 363, 1019-1025 (2000)

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4. Discussion

We classified WDS 00550+2338 as a K0 IV star while it would be more correct to fix its position between G8-G9 and KO-K1 which, indeed, implies a rather small difference in temperature between real components, moreover when photometric and spectral methods agree in its value determination.

It is noteworthy that without knowledge of the star binarity, the spectral classification as well as the physical parameter calculation would be a complicated task because of certain irreconciliable details in the composite spectrum. To what extent this circumstance may have an influence on the derived astrophysical parameters can be illustrated by the following. Spinrad & Taylor (1969, 1971) included HD 5286 in their list of super metal rich (SMR) stars which they generally define as stars with metal abundance [FORMULA]. Later, Faber et al. (1985), on the basis of the new data, suggest almost normal metallicity for this star. In both cases, the authors did not mention its binarity.

Obviously, any astrophysical parameter calculated for a close binary system does not represent any real star and may therefore contribute to the confusion when describing its physical properties, since it is mostly unclear what component they must be attributed to. It is especially important in the cases where the brightness difference between components is less than [FORMULA], which means that fainter component contribution may be more significant and clearly seen on the main background spectrum of the brighter one (Jaschek & Jaschek 1987).

On the basis of spectroscopic study, McWilliam (1990) obtained some model atmosphere parameters, in particular [FORMULA]K for this star. Practically the same value ([FORMULA]K) has recently been obtained by Taylor (1999) who used the results of near infrared photometry, in particular V-K color index. In spite of different methods used to determine the temperature, both authors treated the star as a single object. Although our classification of K0 IV practically coincides with earlier classifications of Abt (1981) and Keenan & McNail (1989), who classified this star as K1 IV, we stress the uncertainty of the temperature and chemical abundance assignments, since they are attributed rather to a composite virtual object than to any of the real components (especially if photometric methods are used).

An exhaustive spectroscopic study of the WDS 00550+2338 chemical abundance must be noted (Luck & Challener 1995) not least because they first noticed the very important presence of the Li I 6708 doublet and derived its equivalent width.

Recent results of the differential photometry performed separately for the components A and B suggest they are dwarfs of spectral types K3 V and K5 V, respectively (Brummelaar et al. 1996). However, the absolute magnitudes M(A) = +3.1 and M(B) = +3.7 calculated using Hipparcos parallax differ significantly from those of MS dwarfs belonging to the same spectral types (for which [FORMULA]; Schmidt-Kaler 1982) clearly suggesting a PMS location. Because of relatively small difference in brightness between components ([FORMULA]) it would be reasonable to suggest that they are both luminous subgiants, since these G and K type stars are located on the HR diagram within the same range ([FORMULA]) of absolute magnitudes (Jennens & Helfer 1975; Schmidt-Kaler 1982). Thus, all aforementioned data support the suggestion that both components belong to the subgiants of luminosity class IV.

It must also be noted that the location of WDS 00550+2338 on the HR diagram coincides well with that of T Tau stars situated above the main sequence in a strip of G-K spectral types (Bertout 1989; Schatzman & Praderie 1993). In general, they are found 2 or 3 orders of magnitude above the main sequence at absolute magnitudes between +7 and +3 (Herczeg & Drechsel 1994).

Further strong evidence for the PMS status of a newly formed system is the presence of a Li I 6708 doublet in its spectrum (Luck & Challener 1995) as well as the inclusion of WDS 00550+2338 in the catalogue of stars with Ca II H and K emission cores (Glebocki et al. 1980). Unfortunately, in our spectrum, the K emission is situated on the limit of the spectral sensitivity and was therefore not reliably measured. Ca II doublet photoelectric flux measurements for this star are given by Duncan et al. (1991).

A very important clue to the age of WDS 00550+2338 A is the Li I 6708 absorption line which is usually only found in young objects, as it is depleted very rapidly in the stellar atmosphere, though no emission features are seen in its spectrum (Sterken & Jaschek 1996).

It is well known that one of the most characteristic properties of young PMS stars is their brightness variation which can be either irregular (and whose nature remains unclear) or regular (due to hot spot induced variations related to rotational period) (Bertout 1989; Bouvier & Bertout 1989; Herbst 1990; Simon et al. 1990; Sterken & Jaschek 1996). The time-scales of these variations differ significantly but we adopt the definition of flare-like events given by Gahm (1990) which refers to those variations similar to flares on flare stars over time-scales of several hours.

Photometric data given in Table 1 and shown in Fig. 1 not only confirm the variability of WDS 00550+2338 but also allow one to make certain suggestions regarding its nature. An overview of these data shows that variations in brightness and B-V color index were insignificant for the observational period but November 27, 1997, a sudden increase ([FORMULA], [FORMULA]) was registered over the course of about 2 hours, which is evidence of its variability.

Apart from this sudden and relatively strong increase, another variation of smaller amplitude ([FORMULA] in B and [FORMULA] in V) was registered on November 26, 1997. This increase was higher than the 3[FORMULA] level and may therefore be taken into account as a real fluctuation. Notice that both variations occured on the same time scale (of about two hours) on two consecutive nights (see data in Table 1). The strong flare occurred on November 27, 1997 and was preceeded by a weaker flare just a day before. It is unclear though, to what component the activity must be attributed.

As it seen from Fig. 1, the B-V color index becomes more blue during the strong event but remained unchanged throughout the course of the weaker one.

A second observation was carried out in 1999 to follow WDS 00550+2338 photometric behavior and detect new changes in brightness. The results showed no significant brightness or color index variations during that period.

Taking into account the spectral type of the star as well as the seemingly irregular character of the detected variations, one may suggest that these changes are similar to the flare-like events (Gahm 1990) observed in young PMS stars. However, additional long-term monitoring seems necessary to clarify the nature of the variability.

Finally, the total A+B mass of the pair, 2.1[FORMULA], has been derived using semimajor axis and period values obtained by Docobo & Costa (1990) and the Hipparcos parallax value. The use of better orbit and precise parallax value allowed us to improve significantly the total mass value, 3.85[FORMULA] obtained by Stefensson & Sanwall (1969) on the basis of the orbit calculated by Muller (1957).

According to Schmidt-Kaler (1982), a K0 III-IV star mass is about 1.0[FORMULA] while a K5 dwarf has a mass less than 0.8[FORMULA]. Under the assumption that both components are dwarfs (as suggested by Brummelaar et al. 1996), the total mass of the system should be about 1.5[FORMULA], discordant (within the estimated error bar of our mass determination equal to 0.3[FORMULA]) with 2.1[FORMULA] value obtained in the present study.

Thus, both luminosity and total mass estimations support the above main sequence location of the star, with a variability similar to the flare activity of young PMS stars.

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© European Southern Observatory (ESO) 2000

Online publication: December 5, 2000
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