5 Relations between emission line strengths and rotational velocity

Briot (1986) searched for a correlation between projected rotational velocity and emission characteristics for a large sample of Be stars, dividing early (B0-B5) stars into 2 classes showing strong (Fe II emission, strong IR excess) and weak (absence of Fe II emission, little or no IR excess) emission characteristics. It was found that the Be stars showing strong emission characteristics had a larger mean v sin i than the group of weak Be stars. Late (B6-B9) stars were found to have a mean v sin i comparable to the strongly emitting early stars. Briot (1986) interpreted this as suggesting that only hot, rapidly rotating stars can develop strong emission characteristics.

We characterised stars of spectral type B0-B4 as either strong or weak emitters based on the presence of Fe II emission (in practice a division at EW $_{\rm Br{\gamma}}=8$ Å; Sect. 4.3). The division into early and late stars at B4 was made based on the finding that neither Mg II or He I emission was seen in any star of a later spectral type. A further group comprising of stars with no evidence of emission (essentially appearing to be normal B stars - Sect. 4.1) was also defined ("non-e'' stars). As we noted in Sect. 4.1 these may be misclassified B stars, or Be stars which are currently in a non-emission phase. In that section we showed that a comparison of the spectral class distribution of the "non-e'' and "e'' stars indicates that they appear to be from different populations.

We choose to express the projected rotational velocities in terms of the critical breakup velocity of the individual star, $v_{\rm crit}$, such that $\omega \sin i = v \sin i/v_{\rm crit}$. The critical breakup velocity was calculated according to Porter (1996)

$$v_{\rm crit}=\sqrt{0.67{\times}GM/R}$$ (1)

where M and R are the stellar mass and radius taken from Schmidt-Kaler (1982).

The results of the analysis can be found in Table 9. Unlike Briot (1986) we find no statistically significant difference in mean $\omega \sin i$ between the strong and weak emitters. Likewise, no difference between the projected breakup velocities of the early and late stars is observed, in this case agreeing with the result of Briot (1986). However, the mean $\omega \sin i$ of the "non-e'' stars is found to be lower than those of the "e'' stars at the 4$\sigma$ level. Two explanations are possible for this result. Firstly, the historical identification of the stars we identify as normal ("non-e'') B stars as Be stars may have been incorrect; consequently we are simply observing the result that Be stars rotate more rapidly than B stars (e.g. Slettebak 1966). Alternatively, since a relationship between rotational velocity and the Be phenomenon clearly exists, stars with lower rotational velocities may be more prone to phase changes between B and Be states. The result that the B ("non-e'') and Be ("e'') stars are from different populations can then be explained by the result of Briot (1986) that there are no slowly rotating late Be stars, so the slowly rotating Be stars that have undergone a phase change and now appear as normal ("non-e'') B stars should preferentially be of an early spectral type. This is the distribution found in Sect. 4.1 for the "non-e'' objects. Future monitoring of the sample to see if the "non-e'' objects undergo a phase transition back to "e'', and hence truly are Be stars, will be necessary to differentiate between the two cases.

 

Table 9: Mean breakup velocities of the subsets of the sample based on emission characteristics

	$\omega \sin i$
B0-B4(EW > 8 Å) "e''	0.42 $\pm$ 0.03
B0-B4(EW < 8 Å) "e''	0.37 $\pm$ 0.03
B5-B9 "e''	0.42 $\pm$ 0.04
B0-B9 "non-e'' (groups 2&4)	0.21 $\pm$ 0.04

Figure 5: Plot of Br$\gamma$ FWHM against projected rotational velocity. Solid line indicates $FWHM=v \sin i$

Figure 6: Plot of He I 2.058 $\mu$m FWHM against projected rotational velocity. Solid line indicates $FWHM=v \sin i$

Figure 7: Stars belonging to Group 1: I. Wavelength given in microns, flux normalised and with offset applied

Figure 8: Stars belonging to Group 1: II. Wavelength given in microns, flux normalised and with offset applied

Figure 9: Stars belonging to Groups 2 & 4. Wavelength given in microns, flux normalised and with offset applied

Figure 10: Stars belonging to Group 3. Wavelength given in microns, flux normalised and with offset applied

Figure 11: Stars belonging to Group 5: I. Wavelength given in microns, flux normalised and with offset applied

Figure 12: Stars belonging to Group 5: II. Wavelength given in microns, flux normalised and with offset applied