Astron. Astrophys. 336, 960-965 (1998)

3. Discussion

3.1. Spectral classification

The spectral class of V635 Cas was determined by comparison with Hiltner 102 and the standards published by Walborn & Fitzpatrick (1990). The Walborn (1971) scheme hinges on the ratios of neutral and singly-ionised helium and the first three ions of silicon.

The comparison with the standard star Hiltner 102 which is identified as a B0 III in the Simbad database led us to reclassify Hiltner 102 as a O9.7 II star. The spectral classification was based on the He II 4541 Å / He I 4387 Å and the He II 4200 Å / He I 4144 Å ratios. For an O 9.7 star the strength of the He II 4541 Å Si III 4552 Å. The luminosity class was determined from the Si V 4089 Å / He I 4026, 4121, 4144 Å and the Si IV 4116 Å / He I 4121 Å ratios. Hiltner 102 may have a slight nitrogen enhancement.

V635 Cas is harder to classify due to two factors. It is fainter than Hiltner 102 (V = 15.5 c.f. Hiltner 102, V = 10.42) and the disk emission causes the filling in of the bluer singly ionized hydrogen and helium lines. For example the filling in is evident in H 4340 Å. In addition the HeII lines appear to be filled in. The strong He II 4200 Å absorption indicates that the star is earlier that previously assumed based on its optical colours: it is an 09e star.

3.2. Probing the circumstellar disk

No far red optical spectra have been previously published for V635 Cas. Dramatic spectral variability occurred. The H line changed from emission to absorption on a timescale of four months or less (Unger 1993). This is the first time that a phase change has been seen in this system. The phase change was also reflected in the Paschen and HeI lines and by the low in the lightcurve. We interpreted this as a disk loss event which was discussed by Unger (1993) and in Paper I. Here we will interpret the emission line data.

With sufficient resolution many of the emission lines in Be stars are double peaked. Huang (1972) interpreted this in terms of a simple model consisting of a disc rotating about the star. The outer radius of the emission region can be estimated using the ratio of the peak separation to the star's rotational velocity.

The spectral resolution of our observations is not sufficient to show the expected double peak structure of the lines. Previous H spectra of higher resolution (0.8 Å dispersion, 1.6 Å resolution) indicate that the peak separation remains fairly constant and that it was 462 km s^-1 in 1991 August 28 (Unger 1993). The unresolved spectra are consistent with either this or a narrower value.

3.2.1. Hei lines

The HeI emission is clearly seen at 6678 Å and 7065 Å on 1991 October 16 and 27. Hence the line at 7281 Å is probably HeI and not an artifact introduced by the removal of the variable telluric H₂O absorption band.

3.2.2. Paschen lines

Generally, Be stars with earlier spectral type have stronger emission in Paschen lines (Andrillat et al. 1990). We observe P19-P11 in emission on 1991 October 16 and 20, and weaker emission from P18-P11 on 1991 October 27. The emission lines disappear, with P13, P12 and P11 in absorption on 1992 February 19.

We hoped to be able to constrain the electron density in the disk by investigating the relative Paschen line strengths. The line fluxes were dereddened using a standard Galactic law, (Rieke & Lebofsky 1985; Howarth 1983) and assuming E(B-V) = 1.5 for V635 Cas (Hutchings & Crampton 1981). It is difficult to fix the continuum level between the various Paschen lines due to the broad emission wings and blending with other lines (see Table 1).

Table 1. The observed lines in the far red CCD spectra.
Notes: Line identification from Meinel et al. (1975); the line centres can vary due to blending; the presence of OI 8446 Å and CaII triplet can only be inferred by an increase in the relative strengths of the Paschen lines. The NI and KI are only suspected to be present, see text.

Line ratios were calculated for the unblended lines (i.e. P19, P17, P14, P12, P11). We estimate that the errors in the line fluxes are 30%. These ratios were compared with case B optically thin recombination line strengths for a = K, = cm^-3 and = K, cm^-3 plasma (Hummer & Storey 1987; Storey & Hummer 1995). The line ratios were also compared with the optically thick line ratios based on the simple assumption that the disk is a K blackbody with the same line widths and emitting areas for all the observed Paschen lines.

General results for Be star systems indicate that P and P line ratios are well away from the case B values (Sellgren & Smith 1992) and that P19 and higher Paschen lines are consistent with optically thin emission (Briot 1981). However, these data have a turnover in the relative line strengths at both P17 and P11. For example on 1991 October 16 the relative line fluxes for P17, P14, P12, P11, normalised to P17 are 1.0:2.4:1.9:1.5. This means that we cannot constrain the disk density. For optically thin emission we would expect that the relative line strengths increase as you descend the Paschen series, the opposite is true for optically thick emission from a blackbody. To constrain these models we require observations of additional lines in the series and at a higher resolution so that the lines are not blended.

3.2.3. Oi lines

The OI 8446 Å emission is more frequent in early type stars (Andrillat et al. 1990) and if seen is always present in emission (Andrillat 1986). This line is blended with P18 and we would expect it to be in emission when the OI 7772-74-75 Å lines are in emission. The OI 8446 Å line has a greater tendency to go into emission than the O I 7772-74-75 Å line due to Bowen fluorescence (Bowen 1947). Assuming the lines are optically thin and adopting the P19 equivalent width for P18, which would under estimate the strength of P18, we obtain equivalent widths of -1.2 Å and -2.0 Å for OI 8446 Å on 1991 October 16 and 20, respectively. The measured equivalent widths of of OI 7772-74-75 Å on 1991 October 16 and 20 are -0.5 Å and -1.6 Å, respectively. We may have underestimated the strength of the OI 8446 Å line but our values are well below the observed ratio of 4 which has been reported for these lines in other Be stars (Jaschek et al 1993).

3.2.4. Other possible features

There are two absorption lines present in all the spectra at 7665 and 7699 Å which could be KI lines. These lines have low excitation potentials of 1.6 eV but more importantly the ionization potential of potassium is 4.3 eV, so any emission must be shielded from the strong UV flux, i.e. the line is probably formed in the outer regions of the disk. Alternatively the KI lines could be due to interstellar absorption.

There are two possible FeII emission lines at 6480 Å and 6508 Å. The FeII has a low excitation potential, 4-5 eV, and we may expect the line to be formed in the outer disk if we assume the excitation potential is correlated with the disk size as in the Balmer series. However, the FeII line can be subject to Ly fluorescence and hence we could also see it in the inner parts of the disk (Slettebak et al 1992). Higher resolution spectra to determine the positions of all of these lines and the individual peak separations are needed.

The emission line at 7235 Å is present in all the 1991 October spectra. Assuming it is not an artefact due to the removal of the telluric H₂O absorption band we tentatively identify it as CII 7234 Å. This line is excitable by resonance fluorescence from the UV continuum and we would expect the line to be formed near the star (Williams & Ferguson 1983).

There are possibly some NI 8629, 8680-83-86, 8703-12-19 Å emission lines. If neutral nitrogen is seen it is always in emission and it is more frequent in early type spectral classes (Jaschek et al 1992; Andrillat et al 1990).

Finally there is a possible emission feature at 8810 Å on 1991 October 27. Higher resolution spectra are needed to confirm the presence of these lines.

© European Southern Observatory (ESO) 1998

Online publication: July 27, 1998
helpdesk.link@springer.de