A&A 371, 932-942 (2001)
W. Hummel1,6 - W. Gässler1,7 - B. Muschielok1 - H. Schink2 - H. Nicklas2 - G. Conti3 - E. Mattaini3 - S. Keller4 - K.-H. Mantel1 - I. Appenzeller5 - G. Rupprecht6 - W. Seifert5 - O. Stahl5 - K. Tarantik1,8
1 - Institut für Astronomie und Astrophysik und Universitäts-Sternwarte München, Scheinerstr. 1, 81679 München, Germany
2 - Universitäts-Sternwarte Göttingen, Geismarlandstr. 11, 37083 Göttingen, Germany
3 - CNR - Istituto Fisica Cosmica, Via Bassini 15, 20133 Milano, Italy
4 - Mount Stromlo Observatory, Private Bag, Weston Creek, Weston, ACT 2611, Australia
5 - Landessternwarte Heidelberg, Königstuhl 12, 69117 Heidelberg, Germany
6 - European Southern Observatory, Karl-Schwarzschildstr. 2, 85748 Garching, Germany
7 - National Astronomical Observatory of Japan, Saburo Telescope, 650 North A'ohoku Place, Hilo,
Hawaii, HI 96720 USA
8 - Max-Planck-Institut für extraterrestrische Physik, Postfach 1312, 85741 Garching, Germany
Received 4 October 2000 / Accepted 1 March 2001
We perform an observational test on global oscillations in Be star circumstellar disks in the metal deficient environment of the SMC. According to the hybrid model of disk oscillations early-type Be stars require an optically thin line force to establish a density wave. The low metallicity in the SMC should therefore diminish or prevent the formation of disk oscillations in early-type Be stars. We present short wavelength range spectra around H of 48 Be stars in the young open cluster NGC 330 in the SMC. We find that the fraction of early-type Be stars in NGC 330 which host a global disk oscillation does not differ from the known fraction of Galactic field Be stars. This observational result is in contradiction to the theoretical prediction. We discuss several interpretations and propose a further observational test.
Key words: line: formation, profiles - stars: circumstellar matter, emission-line, Be
A considerable fraction of all Galactic field Be stars exhibit a long term variable so-called V/R asymmetry (V/R means the ratio of the violet to the red continuum-subtracted peak intensity in a double peak line profile) in their emission lines with a mean cycle of 6.8 years (Copeland & Heard 1963) while the remaining fraction show symmetric emission lines. Hanuschik (1988) extended the terminology of long V/R-variables and called them Class2 line profiles to account for asymmetric and long-term variable single-peak profiles which share the same physical phenomenon seen at low inclination (see Hanuschik et al. 1995), while symmetric line profiles without V/R variability are called Class1. The cyclic behavior of Class2 emission lines was for a long time a matter of debate (for a discussion see Ballereau & Chauville 1989) but is now interpreted as due to one-armed global oscillations in a nearly Keplerian circumstellar disk (Okazaki 1991, 1997; Papaloizou et al. 1992; Savonije & Heemskerk 1993). The model predictions were confirmed by observations (Hanuschik et al. 1995; Telting et al. 1994; Hummel & Hanuschik 1997; Mennickent et al. 1997).
These waves are confined to the inner region of the circumstellar disk and the confinement radius is predicted to amount to 10-15 stellar radii (Savonije 1999; Okazaki 1997, 2000) in agreement with results based on line profile modeling (Hummel 2000a).
In the hybrid model for global disk oscillations (hereafter GDOs) for viscous decretion disks (Okazaki 1997, hereafter O97) two different mechanisms are proposed to confine a GDO. For the late-type Be stars a quadrupole of the potential induced by the rotationally flattened star provides a deviation from the purely Keplerian flow to establish the necessary confinement constraint as proposed by Papaloizou et al. (1992). Since early-type Be stars (MK B0-B4) do not rotate as close to the critical break-up velocity as late type Be stars (MK = B5-B9) (Fukuda 1982) the quadrupole of the potential as due to rotational flattening is no longer efficient for MK=B0-B4. For early-type Be stars the optically thin line force of the stellar radiation is proposed to perform the wave confinement to establish GDOs (O97).
In this study we present an observational test for the hybrid scenario. The efficiency of the optically thin line force depends on the metallicity of the central star. The terminal velocity in winds of O stars in the SMC is usually about 1000 kms-1 lower with respect to their Galactic counterparts (Garmany et al. 1985; Haser et al. 1998). Kudritzki et al. (1987) interpreted the lower terminal velocity as due to a lower line force and a four times lower wind momentum assuming = 0.1 , ( = 0.02, Cox 2000). We selected the young open cluster NGC 330 in the SMC, since it is one of the best studied SMC clusters and shows a large number fraction of Be/(B+Be) stars (Feast 1972; Grebel et al. 1992; Keller et al. 2000). Furthermore the metallicity of this cluster is supposed to be even lower than in the field around NGC 330, hence enhancing the metallicity effect to be studied (Grebel & Richtler 1992). Hill (1999) found Z=0.004 for NGC 330 and Z=0.006 for the field around NGC 330, meaning = 0.2-0.3 .
The relation between the optically thin line force F and
the metallicity Z is not well known. For B supergiants and O-type stars
F scales something between linear and square root with the
metallicity (Pauldrach & Puls, priv. comm.):
As a consequence the confinement constraint for GDOs around early-type Be stars in the SMC is predicted to be considerably reduced with respect to Galactic Be stars. The prediction of the hybrid model would be a larger confinement radius and a larger oscillation period, or even a complete lack of GDOs in early-type Be stars.
The hybrid scenario can therefore easily be tested by a comparison between
the number fraction of Be stars with long-term V/R-variability
(Class2 Be stars, N II) to the total number of all Be stars
|Figure 1: Central part (2.3 2.3 ) of the preparation image (6.8 6.8 ) obtained with FORS1 at Antu (former UT1) and position of photometrically identified Be stars. North is top and East is left. Prefixes: A, B from Roberts (1974), K from Keller et al. (1999) and G from Grebel (1995)|
|Open with DEXTER|
For the present study we selected Be stars mainly from two catalogs: Grebel (1995) and Keller et al. (1999). The distance modulus of the SMC is (Grebel 1995) and Keller et al. (2000) used . The 116 Be stars detected by Grebel (1995) range from to which corresponds to spectral types B0 to B6 using the calibration of Zorec & Briot (1991). Be stars detected by Keller at al. (1999) range from to corresponding to spectral types B0 to B3. This means that all targets of Keller et al. (1999) and those targets of Grebel (1995) with 18 25 (G1-G103) belong to the spectral classes B0-B4 under consideration.
The completeness of this study can be estimated to be around 40% for the early-type Be stars B0-B4 in NGC 330.
The spectroscopic observations were collected
as two 900 s exposures with FORS2 in MXU mode
at VLT Kueyen (former UT2) on Paranal as part of the first commissioning
period of the instrument.
The instrument is described by Böhnhardt (2000),
and Seifert et al. (2000), while the mask exchange
unit (MXU) of FORS2 is described by Schink et al. (2000).
The nominal pointing was RA = 00:56:19.043, DEC = -72:27:59.81
for the instrument setup. Up to 7 reference stars were
used to align the telescope pointing to the correct position.
Additionally a reference slit was specified to correct
the image flexure of the instrument.
The wavelength range coverage is 6533Å
Most of the slits were longer than 8 arcsec, providing
sufficient sky background.
Some slits were occupied with two or
more targets. Four slits were inclined with respect to the
dispersion direction to optimize the sky background region
and the spectra required an additional row-by-row
alignment handling using the wavelength calibration.
The log of observations is given in Table 1.
Thanks to the negligible image flexure of FORS2 calibration frames could
be collected a day earlier during day-time with the telescope in zenith position.
Most of the photometrically identified emission line objects can be confirmed to show H 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 = 0 236 (Keller et al. 2000) the colors of 0 36 and -0 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 color index of A22 is given as 0 46 (Grebel 1995). The colors of 0 22 and 0 10 for G108 and G113 would indicate a 80% lower H 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. After subtraction of a homogeneous background for G96 one third of the raw H 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 profile indicates an interstellar origin. Spectra of G28, G64 and G111 suffer from low S/N.
Measured FWHM of H 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 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 normalized to the local stellar continuum . Spectra are offset in ordinate by multiples of 4.0 . 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|
|Open with DEXTER|
|Figure 3: Variablity of H 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 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 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)|
|Open with DEXTER|
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 NV/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 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 equivalent width for GDO is Å (Hanuschik et al. 1995). Circumstellar disks with lower 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 NV/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:
|slit||name||or VV//VR||FWHM||or IV//IR||comment|
|Å||km s-1||km s-1|
|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|
|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|
|21||G15, K480, A24||43||110||360||5.8||s/a|
|52||G19, K211, B41||52||140||300||7.0||s|
|23||G23, K239, B34||28||120||280||5.4||s|
|13||G32, K419, B9||18||120||190||4.6||s|
|slit||name||or VV//VR||FWHM||or IV//IR||comment|
|Å||km s-1||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|
|Hanuschik et al. (1988)||1982, 1985||36||5||0.14||30||12||0.40||10||3||0.30|
|Doazan et al. (1991)||1978-1988||96||33||0.34||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 et al. (1996)||1982-1993||69||25||0.36||46||14||0.30||20||5||0.25|
In Cols. 3-5 of Table 3 we collected
NV/R of different emission line surveys
of the Galaxy.
Only Copeland & Heard (1963) and
Hanuschik et al. (1988) give numbers for
NV/R. The values of the other references and
N I + II and N I
were derived by simply counting the number of
asymmetric line profiles.
Known binary stars such as Per, Tau and HR 2142
From Table 3 (Cols. 3-5) we find
The number statistics NV/R for Galactic Be stars
as given in Col. 5 of Table 3 is biased in two ways with
respect to NV/R for NGC 330. First
contains also late-type Be stars. To account for this
difference we derived
again, but now only for the early type stars
(see Cols. 6-8 of Table3).
The second difference between the Galactic and the NGC 330 number statistics
in Eq. (3)
is for the resolved double-peak profiles alone, which selects mostly
Be stars in NGC 330 (more precisely: the large
Be stars in
NGC 330, where
is the H
Assuming that the observed
Be stars is a product of the geometric effect of
random axis orientation and of the distribution of
relative equatorial rotation velocity
Nevertheless we tune the Galactic sample of Be stars for the
conditions under which
again, but take only Be stars
into account for which
O9-B4] and for which the H
peak separation of the double peak profile
So far we adapt the instrumental resolution of FORS2 to the
Galactic emission line sample.
We find (see Cols. 9-11 of Table3):
Our detailed number statistics for the Galactic Be stars show that:
We see three possible explanations for the discrepancy in the NV/R number statistics.
The results of Kudritzki et al. (1987) show that the ratio of weak lines to strong lines (= wind parameter) is mostly independent on the metallicity for meaning that the weak line force is locked with the metallicity.
If the metallicity of the SMC is not sufficiently low to diminish the line force the conditions to establish disk oscillations in NGC 330 are similar to those in the Galaxy, hence should equal ;
A weak point in this study is that is based on field Be stars and is therefore based on an age-averaged sample while NV/R of NGC 330 is a representative number for a fixed stellar age of T=19 My (Keller et al. 1999). This becomes important when considering that, at least in the Galaxy, the Be phenomenon occurs in the second half of the main sequence life-time (Fabregat & Torrejón 2000). The relation between w and spectral type of NGC 330 might therefore be different from that of the Galactic sample (Fukuda 1982). This relation depends on the evolution of w during the main sequence life time of stars.
Recent stellar evolution models for rotating massive stars
(Heger et al. 2000; Heger & Langer 2000) indicate
that the surface rotation velocity during the main sequence
lifetime is reduced by about 10% to 20% of the initial
surface rotation velocity:
The lower metallicity could affect the period of the density waves more strongly than the general formation of GDOs. In this case NV/R might be similar to that for Be stars in the Galaxy, but the periods should be much longer in the early-type Be stars than in the later type Be stars.
This second test for the hybrid model, the predicted larger period of V/R variability, cannot be performed with the present data set since our observations suffer from the short time coverage of only one year. This means also that our asymmetric line profiles match only one criterion to be classified as a Class2 profile, namely the asymmetry; the second necessary condition, the long term variability, is not proven yet. The only Galactic Be star known to the authors showing an asymmetric H emission line without a cyclical variability is CMa (Hanuschik et al. 1996). All other Galactic Be stars with Class2 profiles are by definition V/R-variable. Given the low spectral resolution we expect V/R variability to be first detected in Be stars with the highest (those with comment d/aa in Table 2): G38, G51, G109, G112, and G114. Four further candidates are G40, G103, G109 and G78 (Hummel et al. 1999).
An observational test would be the continued observation of Be stars in NGC 330 known to exhibit asymmetric emission lines on a once-a-year basis, with a higher spectral resolution. One could derive the mean V/R-ratio periods for the early-type and late-type Be stars in NGC 330 separately and compare the values to the known mean V/R-ratio period for Galactic Be stars of years (Copeland & Heard 1963).
The funding through "Verbundforschung Astronomie'' by the Bundesministerium für Forschung und Technologie BMBF grants 05AV9WM12, 05AV9MGA and 05AV9V0A6 is gratefully acknowledged. We thank the referee John Telting for his valuable comments on the selection effects. We thank the VIRMOS Consortium for cooperation in all matters regarding the Mask Manufacturing Unit.