5 Results

5.1 General

Most of the photometrically identified emission line objects can be confirmed to show H$\alpha$ in emission. K799 shows [NII] in emission and can therefore no longer be classified as a Be star.

W4:31 and W2:79 show traces of a rotationally broadened stellar absorption profile without any signature of emission, although these stars were classified as emission line objects by photometry. Taking the bright Be star A22 (= G5 = K223 = W1:15) as a reference with a color index of F555W-F656N = 0 $.\!\!^{\rm m}$236 (Keller et al. 2000) the colors of 0 $.\!\!^{\rm m}$36 and -0 $.\!\!^{\rm m}$25 for W4:31 and W2:79 respectively, would indicate strong emission at least for W4:31.

G108 and G113 also show traces of a rotationally broadened stellar absorption profile without any signature of emission. The R-H$\alpha$ color index of A22 is given as 0 $.\!\!^{\rm m}$46 (Grebel 1995). The colors of 0 $.\!\!^{\rm m}$22 and 0 $.\!\!^{\rm m}$10 for G108 and G113 would indicate a 80% lower H$\alpha$ emission strength with respect to A22. Note that G108 did not show emission also a year earlier (Hummel et al. 1999, their Fig.1e) when observed with FORS1.

The status of K255 and G96 remains unclear due to the large fraction of the interstellar contribution to H$\alpha$. After subtraction of a homogeneous background for G96 one third of the raw H$\alpha$ emission was found to be of stellar origin for the FORS1 observations (Fig. 3). The FORS2 observations showed a rather inhomogeneous background. Moreover the box shape of the residual H$\alpha$ profile indicates an interstellar origin. Spectra of G28, G64 and G111 suffer from low S/N.

Measured FWHM of H$\alpha$ emission range from 590 kms^-1 (G84) to 150 kms^-1 (K1830). Note that the FWHM of the instrumental profile of 2.7 Å corresponds to 120 kms^-1. H$\alpha$ emission strengths and line shapes (Fig. 2, Table 2) are very similar to those of Galactic Be stars, in particular when compared with surveys of similar spectroscopic resolution (Dachs et al. 1986). A simple preliminary classification with respect to line profile asymmetry is given in Table 2.

Figure 2: Spectra of 62 targets in the region of H$\alpha$ normalized to the local stellar continuum $F_{\rm c}$. Spectra are offset in ordinate by multiples of 4.0 $F_{\rm c}$. The abscissa gives geocentric velocities. Prefixes: G after (Grebel 1995), K after (Keller at al. 1999), W after (Keller et al. 2000), A, B after (Roberts 1974), and F after (Hummel et al. 1999). The number after the "-'' in the name indicates an intensity scaling factor used in the plot for better visibility. Some spectra (G12, G38, G64, G28, G84, W2:79) still contain a bad column residual, which could not be corrected by the flat fielding due to non-linearities

Figure 3: Variablity of H$\alpha$ emission line profiles of Be stars in the field of NGC 330. Bold: profiles of the present study collected in Nov./1999 with a spectral resolution of $\Delta V =120$ km s-1 with FORS2 at UT2 in MXU mode. Thin: Profiles of our first study (Hummel et al. 1999) collected in Oct./1998 with a spectral resolution of $\Delta V =225$ km s-1 with FORS1 at UT1 in MOS mode. Profiles of the present study were shifted in V to match the observed velocity of the profiles obtained earlier. Target numbers after Grebel (1995)

5.2 Line profile variability

In Fig. 3 we show the new observations presented in this study together with line profiles obtained a year earlier with FORS1. The spectral resolution of the present H$\alpha$ line profiles is a factor 2 higher than of those profiles obtained a year earlier. No significant variability is detected for G1, G32, G41, G55, G77, and G78. F1, G109, and G115 show signatures of a change in the profile shape. Decreasing emission equivalent width is found in G9, G59, and G68. The only candidate for a long-term V/R-ratio variability in Fig. 3 is G109, however a time scale of one year is not long enough to detect V/R variability in Be stars.

5.3 Emission line statistics

33 of the 48 line profiles in emission (G96, K799, K4151 and A30a excluded) showed single peaks. Among the 15 double peak profiles six are clearly asymmetric (d/aa) and two further ones show asymmetry with a V/R-ratio close to 1 (d/a). This means that N_V/R amounts to 40%-53% for the sufficiently resolved double peak profiles. We found six marginally asymmetric (s/a) and two strongly asymmetric (s/aa) single peak profiles meaning 6% to 24% (2-8 of 33) of the single peak profiles.

We assumed that the number of long-term V/R-variables which are accidently caught during a symmetric phase does bias the number statistics for single peak and double profiles in the same way. Furthermore we assumed that asymmetric profiles with very low $v \sin i$ cannot be detected, also not in high-resolution spectroscopy, since the kinematical broadening is less than the non-kinematical broadening. This selection effect, being independent of the spectral resolution, biases all statistics in the same manner and cancels out when compared to the statistics of other studies.

A further selection effect is that the empirical minimum H$\alpha$ equivalent width for GDO is $W_\alpha = -10$ Å (Hanuschik et al. 1995). Circumstellar disks with lower $W_\alpha$ are probably physically not able to establish a GDO at all. This empirical result most probably reflects a lower limit for the viscosity which provides the required interaction between different particle trajectories to establish a GDO. We assume that the fraction of Be stars with densities below the critical density for GDOs is independent of inclination and spectral type.

The lower number of N_V/R for the single peak profiles with respect to the double peak profiles is therefore due to the finite resolution alone. Taking the statistics for the double peak profiles as representative, we find:

$$N_{V/R}^{{\rm NGC~330}}=0.47\pm0.13$$ (3)

where 0.47 is the mean value of both statistics (counting the d/a profiles as Class2 profiles or not) and $\pm0.13$ is given by the standard error for the binomial probability distribution:

$$\sigma = \sqrt{ n_{\rm I} N_{V/R} N{\sc i + ii} }.$$ (4)

Thereby we assume the sample amounts to N I + II, and the parent population, the number of all Be stars in NGC 330, is more than twice as large as the sample. $n_{\rm I}$ is the probability to find a Be star of Class1 in the sample, hence

$$n_{\rm I} = \frac{ N{\sc i} }{ N{\sc i+ii} }$$ (5)

and

$$n_{\rm I} + n_{V/R} =1.$$ (6)

 

 

Table 2: H$\alpha$ emission line parameters. From left to right: MXU slit number, target names (prefixes: A, B from Roberts 1974, K from Keller et al. 1999, G from Grebel 1995, F from Hummel et al. 1999, W from Keller et al. 2000 and LIN from Lindsay 1961), emission equivalent width of H$\alpha$ in Å, peak radial velocity (p) if single peak (comment = s) or velocity of the violet peak (V) the central depression (c) and red peak (R) in case of double peak (comment = d), FWHM in km s^-1, intensity of the of profiles in units of the local stellar continuum, comment: s) symmetric single-peak, d) symmetric double-peak, d/aa) asymmetric double peak, s/aa) asymmetric single peak, s/a) marginally asymmetric single peak, d/a) marginally asymmetric double peak, neb) [N II] line emission, abs) no emission, stellar absorption profile, -) low S/N , no classification possibles, $^{{\rm a)}}$) only V and R given, $^{{\rm b)}}$) blend with neighbor star, $^{{\rm c)}}$) 2 stellar components

slit	name	$W_\alpha$	$V_{\rm p}$ or VV/$V_{\rm c}$/VR	FWHM	$I_{\rm p}$ or IV/$I_{\rm c}$/IR	comment
		Å	km s-1	km s-1	$I_{\rm c}$
72	A20	0	-	-	-	abs
38	F1	32	-70/20/110	500	3.6/3.6/3.9	d/a
20	G1, K80, B21	27	110	260	5.2	s
44	G3, K242, B17	18	150	220	4.2	s
70	G5, K223, A22	24	210:	230	4.9	s
25	G6, K238, B35	43	130	340	5.8	s
15	G7, K213, B12	29	150	200	6.6	s
56	G8, A39	20	130	260	4.1	s $^{{\rm b)}}$
08	G9, K258	28	160	270	5.4	s
66	G9a, K203, B6	26	120	250	5.0	s
62	G10, K228, B14	18	160	300	3.4	s
49	G11a, K206, B5	33	130	290	5.9	s
22	G12, K229, B36	28	140	320	5.2	s $^{{\rm b)}}$
21	G15, K480, A24	43	110	360	5.8	s/a
52	G19, K211, B41	52	140	300	7.0	s $^{{\rm b)}}$
71	G20, A30 $^{{\rm c)}}$	20	140	350	3.3	s
	(G20):, A30a	6	80	400	1.4	s
23	G23, K239, B34	28	120	280	5.4	s
55	G28, K441	-	-	-	-	-
24	G31, K471	29	120	280	5.2	s
13	G32, K419, B9	18	120	190	4.6	s
69	G38, K439	24	30/160/260	380	2.8/2.1/3.9	d/aa
06	G41, K874	26	110	320	4.4	s $^{{\rm b)}}$
50	G44, K875	37	140	290	6.3	s
54	G47	19	140	280	4.4	s
53	G51	9	30/160/230	470	2.0/1.7/1.8	d/aa
07	G55, K857	34	160	360	5.0	s
36	G56	20	70	460	3.1	s/aa
09	G59, K991	21	160	230	4.5	s
12	G64, K870	0	-	-	-	abs
65	G68	8	110	440	1.8	s/a
68	G71	24	30/160/250	500	3.4/2.6/3.7	d
51	G72, K1845	37	120	380	4.9	s/a
19	G73	17	110	270	3.8	s/a $^{{\rm b)}}$
47	G77, K1918	10	30/160/250	500	1.9/1.6/1.9	d
46	G78, K2030	17	20/150/250	510	2.6/1.7/2.6	d
34	G84	8	20/120/250	590	1.8/1.5/1.7	d
43	G96	-	-	-	-	f
59	G98	40	70/260/300	500	4.5/2.5/4.2	d/a
37	G99	11	-30/140/250	570	2.2/1.6/2.2	d
35	G104	35	120	360	5.0	s/aa
26	G108	0	-	-	-	abs

a) only V and R given.
b) blend with neighbor star.
c) 2 stellar components.

 

Table 2: continued. H$\alpha$ emission line parameters

slit	name	$W_\alpha$	$V_{\rm p}$ or VV/$V_{\rm c}$/VR	FWHM	$I_{\rm p}$ or IV/$I_{\rm c}$/IR	comment
		Å	km s-1	km s-1	$I_{\rm c}$
57	G109	10	10/120/200	280	2.4/1.2/3.7:	d/aa
42	G110	0	-	-	-	abs
58	G111	-	-	-	-	-
17	G112	14	70/250 $^{{\rm a)}}$	370	3.2/2.5 $^{{\rm a)}}$	d/aa
45	G113	0	-	-	-	abs
05	G114	20	100/190/240	310	4.1/2.7/2.9	d/aa $^{{\rm b)}}$
29	G115	10	70	320	2.6	s/a
18	K44, B31	8	$-70/240^{{\rm a)}}$	460	1.6/2.1 $^{{\rm a)}}$	d/aa
30	K111	17	140	270	3.4	s
67	K233	26	120	330	3.6	s $^{{\rm b)}}$
64	K255, B18	-	-	-	-	f
63	K459	18	140	370	2.9	s
31	K754	32	70/120/170	360	4.8/4.3/4.6	d $^{{\rm b)}}$
28	K798	23	95	250	5.1	s/a
61	K799	-	-	-	-	neb
60	K821	22	140	340	3.6	s
48	K1830	10	80	150	3.7	s
14	K4154, (LIN305)	-	-	-	-	neb
11	W2:79	-1	160:	-	0.5:	abs
10	W3:16	8	-20/160/320	500	2.0/0.8/1.9	d
73	W4:31	0	-	-	-	abs

 

 

Table 3: Number statistics of N_V/R for Galactic samples

		${\rm MK}=[$O9-B9]			${\rm MK}=[$O9-B4]			${\rm MK}=[$O9-B4]
								$\Delta V_{{\rm p}}^{\alpha} < 120$ km s-1
Ref.	epoch	N I+II	N II	NV/R	N I+II	N II	NV/R	N I+II	N II	NV/R
Copeland & Heard (1963)	1938-1962	54	36	0.67	--	--	--	--	--	--
Andrillat & Fehrenbach (1982)	1980-1981	61	24	0.40	33	4	0.12	13	6	0.46
Andrillat (1983)	1981	52	13	0.25	29	5	0.17	9	2	0.22
Hanuschik et al. (1988)	1982, 1985	36	5	0.14	30	12	0.40	10	3	0.30
Doazan $^{{\rm a)}}$ et al. (1991)	1978-1988	96 $^{{\rm a)}}$	33 $^{{\rm a)}}$	0.34 $^{{\rm a)}}$	47	11	0.23	23	5	0.21
Slettebak et al. (1992)	1989	40	15	0.38	25	17	0.68	9	2	0.22
Hanuschik $^{{\rm b)}}$ et al. (1996)	1982-1993	69	25	0.36	46	14	0.30	20	5	0.25

a) H$\beta$ is used in case H$\alpha$ is not available.
b) Including the study of Slettebak et al. (1992).