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Astron. Astrophys. 355, L35-L38 (2000)

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4. Evolutionary state of HD 35929 and V351 Ori

Although more systematic observations are needed to define the precise pulsational periods of HD 35929 and V351 Ori, the available data already provide some preliminary constraints on their position in the H-R diagram, stellar mass and evolutionary state. As remarked by the referee, the discussion presented in this section is useful as an illustration of how well-defined periods for these stars could indeed tightly constrain their evolutionary states.

Fig. 3 shows the location of HD 35929 in the H-R diagram. The dotted box accounts for the uncertainty in the spectral type (A5 to F0: Malfait et al. 1998, Miroshnichenko et al. 1997, van den Ancker et al. 1998) and distance ([FORMULA]360 to 430 pc: van den Ancker et al. 1998). The instability strip for the first three radial modes, as predicted by Marconi & Palla (1998), is also shown together with the PMS evolutionary tracks computed by Palla & Stahler (1993) and the post-MS evolutionary tracks for 2.5 and 3.0 [FORMULA] of Castellani et al. (1999). The two circles indicate the best combination of the stellar parameters (M, L, [FORMULA]) that yield a period equal to the observed one, [FORMULA] d. These values are listed in Table 2. The two solutions indicate a mass of 3.4 or 3.8 [FORMULA], pulsating in the first overtone (FO) and second overtone (SO) respectively: in both cases, HD 35929 can be considered a PMS pulsator, as expected.

[FIGURE] Fig. 3. The position of HD 35929 in the H-R diagram according to the estimates of spectral type and distance found in the literature (dotted box). The shaded region is the instability strip predicted by Marconi & Palla (1998). The PMS and post-MS evolutionary tracks are shown as thick and thin solid lines, respectively. The birth-line for the PMS evolutionary tracks is displayed by the dashed line.


[TABLE]

Table 2. Derived stellar parameters


A combination of parameters for a post-MS stellar mass can also reproduce the pulsational period observed in HD 35929. In this case, the best choice would be a post-MS model of a 2.7 [FORMULA] star with [FORMULA] and [FORMULA] K, pulsating in the SO mode. The location of this solution is quite close to the lower circle shown in Fig. 3. Although only few specific studies exist on HD 35929, some evidence supports the fact that HD 35929 is a young star associated with the Ori OB-1c association. For example, Malfait et al. (1998) have discussed the infrared excess observed toward this star, whereas Miroshnichenko et al. (1997) have suggested that the star might be in a transition phase between a PMS Herbig Ae star and a [FORMULA] Pictoris-type object. We also note that the pulsational character (period and H-R diagram location) of HD 35929 is similar in many respects to that of the well known Herbig Ae star HR 5999. From these considerations, we favor the conclusion that HD 35929 is a PMS star with a mass in the narrow range 3.4-3.8 [FORMULA], pulsating in the FO or SO.

Finally we note that if the Strömgren indices [c1] and [m1] measured for this star are taken into account, together with published values for the [FORMULA] index, one derives an effective temperature [FORMULA] and a luminosity varying between 26 and 36 [FORMULA]. This means that according to the Strömgren photometry, HD35929 should be located on a [FORMULA] evolutionary track and the period based on present data would be too long even for a pulsation in the fundamental mode. A possible explanation for this inconsistency could be that HD35929 is also a rapid rotator ([FORMULA]150 Km s-1) so that the assumption of radial pulsation could be not completely correct. Future observations and numerical simulations are needed in order to properly address this problem.

Fig. 4 shows the location of V351 Ori in the H-R diagram. Here, the uncertainty on the luminosity is larger than for HD 35929, because of the distance ambiguity. The lower value corresponds to the minimum distance of 260 pc given in the Hipparcos catalogue (van den Ancker et al. 1998). The upper limit (inverted triangle) assumes that V351 Ori is located in the Orion molecular cloud at a distance of 460 pc. The dashed box corresponds to an uncertainty of [FORMULA] dex in [FORMULA]. Finally, the filled circle marks the position estimated by van den Ancker et al. (1996) with the associated error bar. As already pointed out, present data for V351 Ori suggest a pulsation period of [FORMULA].

[FIGURE] Fig. 4. Same as Fig. 3 but for V351 Ori. The filled triangles refer to the upper and lower limits of the distance. The open symbols are the best pulsation models reproducing an oscillation consistent with the observed period.

Using the constraints provided by this preliminary period and by the topology of the instability strip, we have computed linear nonadiabatic models to find the best set of stellar parameters that reproduce the pulsation of V351 Ori. The results are shown in Fig. 4 and the stellar parameters are given in Table 2. The solutions yield a stellar mass of 1.85 [FORMULA] or 2.15 [FORMULA], respectively, pulsating in the SO (open square in Fig. 4), or the third overtone (TO) mode (open circle in Fig. 4. For lower modes, the luminosity of the model would be lower than the estimated lower limit for V351 Ori, whereas higher modes are probably excluded by the closeness of the observational box to the second overtone blue boundary and we did not consider their occurrence. These solutions would tend to favor a distance of V351 Ori, smaller than that of the young stellar population of the Orion complex.

Recently, Koval'chuk & Pugach (1998) have argued on the basis of several peculiar photospheric properties that V351 Ori is in fact more evolved than previously thought and conclude that this star does not belong to the group of Herbig stars. From Fig. 4, we see that the the post-MS track of a 2 [FORMULA] star intersects the corresponding PMS track at about the position of the best pulsational models. The degeneracy of the tracks does not allow to use the pulsational analysis to discriminate between the two evolutionary phases. However, this uncertainty would disappear if the distance of V351 Ori were the same as that of the Orion population stars. Then, the difference between the pre- and post-MS tracks would be large enough that the period analysis would rule out one of the two solutions. Future studies of V351 Ori should address this important aspect.

As in the case of HD35929, we used the measured Strömgren indices [c1] and [m1] to provide an independent evaluation of the location in the HR diagram of this star. The results, [FORMULA], [FORMULA], now suggest a stellar mass of [FORMULA] for this pulsator, in which case the observed periods would be consistent with an oscillation in the TO pulsation mode.

In conclusion, the present observations yield compelling evidence for the occurrence of [FORMULA] Scuti-type pulsation in two Herbig Ae stars, HD 35929 and V351 Ori. The comparison with evolutionary and pulsational models provides independent, even if quite preliminary, constraints on the mass and evolutionary state of these stars.

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© European Southern Observatory (ESO) 2000

Online publication: March 21, 2000
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