A&A 448, 351-366 (2006)
DOI: 10.1051/0004-6361:20053066
S. Simón-Díaz1,2 - A. Herrero1,3 - C. Esteban1,3 - F. Najarro4
1 - Instituto de Astrofísica de Canarias, 38200 La Laguna,
Tenerife, Spain
2 -
Isaac Newton Group of Telescopes, Apartado de Correos 321,
38700 Santa Cruz de la Palma, Tenerife, Spain
3 -
Departamento de Astrofísica, Universidad de La Laguna,
Avda. Astrofísico Francisco Sánchez, s/n, 38071
La Laguna, Spain
4 -
Instituto de Estructura de la Materia, CSIC, C/ Serrano 121,
28006 Madrid, Spain
Received 15 March 2005 / Accepted 1 October 2005
Abstract
We present the results of a spectroscopic analysis of the
Trapezium cluster stars inside the Orion nebula. The rotational velocities
were obtained using the Fourier analysis method, and we found agreement with
values derived by the usual method based on linewidth measurements.
The rotational velocity derived for Ori C by this method
is consistent with the variability of some of its spectral features that
have a period of 15.42 days.
By means of the fit of H , He I, and He II observed
profiles with F ASTWIND synthetic profiles, stellar parameters and wind
characteristics were derived. This methodology let us estimate the
errors associated with these parameters, and we found that macroturbulence
effects have to be included for a good
fit to the He I-II lines in the spectrum of
Ori C.
By means of a very accurate study, oxygen abundances were derived
for the three B0.5V stars Ori A, D, and
Ori B. Final abundances are consistent with the nebular gas-phase results presented in Esteban et al. (2004) and are lower than
those given by Cunha & Lambert (1994). Our results suggest
a lower dust depletion factor of oxygen than previous estimations for the
Orion nebula.
Key words: stars: abundances - stars: early-type - stars: fundamental parameters - ISM: abundances - ISM: H II regions - ISM: individual objects: Orion nebula
New developments in massive-star model-atmosphere codes have led to interesting new possibilities for stellar spectroscopic studies. Improvements in computational methods, as well as an increase in the efficiency of computers, have made it possible to model the atmospheres of hot luminous stars, taking into account not only strong NLTE effects and hundreds of thousands of metallic lines that produce the so-called line-blanketing (Hubeny & Lanz 1995), but also winds with expanding spherical geometries (Santolaya-Rey et al. 1997; Hillier & Miller 1998; Pauldrach et al. 2001; Puls et al. 2005).
The new improvements included in these latest generation models call
for a review of previous results. For example, the papers by Herrero et al.
(2002), Crowther et al. (2002), and Martins et al. (2002) show
that the SpT -
calibrations used previously (Vacca et al. 1996)
needed to be revised to lower effective
temperatures for a given spectral type. Recent analyses by Repolust et al.
(2004) and Martins et al. (2005) in the Milky Way, and by
Massey et al. (2004, 2005) at SMC and LMC metallicities, reinforce this
conclusion and together imply a need to revisit the ionizing flux distribution
that is used for the study of H II regions and starbursts.
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Figure 1:
Atlas of the INT+ IDS spectra in the 4000-5000 Å region used for this study. The ![]() ![]() ![]() ![]() ![]() ![]() ![]() |
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Recent work by Trundle et al. (2002) and Urbaneja et al. (2005)
have shown that abundance gradients in some spiral galaxies derived from stellar
and nebular studies tend to be coherent but very dependent on the calibration
used in the strong line nebular methods. However, until now, there have
been no detailed systematic studies that compare results from nebular
and stellar studies. This is the first in a series of papers aimed at this
subject. We selected some
galactic H II regions for a detailed study of the interaction
between massive stars and the surrounding ISM, looking for consistency in the
derived parameters (
,
luminosities, and ionizing flux distribution of
the stellar population), as well as the abundances of C, N, O, Si, and Mg.
For this first study, we selected the Orion nebula, a well-studied and
resolved Galactic H II region powered by a cluster of a few massive stars,
the Trapezium cluster.
The Orion complex contains the massive on going star-forming region closest
to Earth at about only 450 pc. The Orion nebula, M 42, is part of this
complex. It is a well-known H II region (e.g. O' Dell 2001; Ferland
2001) ionized by the Trapezium cluster stars ( Ori), a group of early type stars located in the core of the nebula. Together with
Ori C (HD 37022, O7V), the main ionizing source, we find other
B0.5V stars that are perfect targets for a stellar abundance study.
The most recent study of the chemical composition of the Orion nebula has been presented by Esteban et al. (2004) who re-analyzed the emission line spectrum of the nebula to determine the physical conditions and abundances of the ionized gas-phase. Cunha & Lambert (1992, 1994) included some of the Trapezium cluster stars in a survey of B-type stars in the Orion OB1 association. They presented a spectroscopic analysis of these stars to determine C, N, O, Si, and Fe stellar abundances.
This paper focuses on a spectroscopic analysis of the Trapezium cluster stars to derive their
stellar parameters and oxygen abundances. The stellar
parameters obtained for Ori C will be used in future papers
as input for modeling the Orion nebula with photoionization codes. The derived
stellar abundances are compared to those obtained by Esteban et al.
(2004) through nebular studies.
Our paper is structured as follows: in Sect. 2 we present
the observations. In Sects. 3 and 4 we obtain
the v sin i and the stellar parameters of our targets. The study of
Ori C is presented in Sect. 5,
and the oxygen abundance analyses in Sect. 6.
A discussion of the results and the conclusions of this work are
presented in Sect. 7.
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Figure 2:
Atlas of the INT+ IDS spectra in the H![]() ![]() ![]() ![]() ![]() ![]() ![]() |
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The bulk of the observations used here were carried out with the Isaac
Newton 2.5 m. Telescope ( INT) at the Roque de los Muchachos Observatory in La Palma on 21 December 2002. The Intermediate Dispersion Spectrograph
( IDS) was used with the 235 mm camera and two different gratings. We observed
the spectral region between 4000 and 5050 Å using the H2400B grating,
which resulted in an effective spectral resolution R
7500
(equivalent to a 0.23 Å/pixel resolution and
2.6 pixel FWHM
arc lines). The H1800V grating was used for the H
region,
resulting in a similar spectral resolution (0.3 Å/pixel, R
8000).
With these configurations, three exposures were needed to cover the whole
range. A large number of flat fields and arcs for the reduction process were
then obtained.
The reduction and normalization of the spectra was made following standard
techniques, with IRAF, and with our own software developed in IDL. The
signal-to-noise ratio (SNR) of the reduced spectra depends on the
spectral range, but is usually about 200-250 in the blue region and 250 in the H region (see Table 2).
We found some problem when rectifying the INT observations in the H region. In the blue wing of the H
line we observed a feature that is independent of rectification.
No known line should be present at this wavelength, so we argue
that it must be cosmetic.
Special care has to be taken over the nebular contamination of the stellar spectra. The stars we studied are located inside H II regions, so the stellar spectra are contaminated by the nebular emission spectrum. It would be desirable to remove this nebular contribution, as it fills the cores of the Balmer H and He I lines. However, this is not easy even though we have long slit observations; the nebular emission has a spatial structure that complicates its subtraction, especially for the Balmer H lines, the most important nebular contribution. If this is not correctly done, then an over or under subtraction will appear. After trying different possibilities, we concluded that the best solution for this problem is not to subtract the emission lines and to ignore these regions in the final spectrum. For the Balmer H and He I lines, this is satisfactory, as emission lines are narrower than absorption lines. Nebular contributions could be more difficult to separate for metal lines; however, the contamination of the stellar metal lines used for the abundance analysis due to nebular lines is negligible.
Table 1:
Identification, spectral type, and photometric visual data of the studied objects. The
and
values for Orion stars are from Preibisch et al. (1999). The
values for these stars were calculated considering a distance
450
50 pc
to the Orion nebula. Data for HD 214680 and HD 47839 are from Herrero et al. (1992). Photometric data for HD 149438 are from Humphreys (1978). Uncertainties in
,
,
and
are 0.01, 0.03, and 0.3, respectively.
Table 2: SNR achieved for the different spectra for the three ranges observed with the INT+ IDS.
The INT observations consist of the brightest three stars
in the Trapezium cluster ( Ori A, C, D), together with the two nearby stars
Ori A and B. Two standard stars were included
in this set, 10 Lac and 15 Mon (O9V and O7V, respectively).
The other standard star,
Sco, was kindly provided by Dr. Gehren.
This is a slow rotating B0.2V star that is perfect for a preliminary abundance
analysis. The spectrum was obtained with CASPEC, attached
to the ESO 3.6 m telescope. The SNR of this spectrum is
200-300 in the blue region and
150 in the H
region.
For the study of the spectral variability of Ori C, we
used FEROS spectra, some of them downloaded from the FEROS database and
others kindly provided by Dr. Stahl. These observations
were carried out with FEROS at the ESO 1.52 m telescope in La Silla. The instrument is designed for high-dispersion spectroscopy with
48 000 in the spectral range 3700-9200 Å. The achieved SNR is 300 at about 4500-5000 Å. Different phase observations were used for the variability study (see Table 7), and is presented in Sect. 5.
The analysis of stellar spectra makes use of a number of free parameters like the micro and macroturbulent velocities and the projected rotational velocity, v sin i. The last one has acquired particular importance in recent times because of the mixing that rotation may induce in the interior of massive stars (e.g. Maeder & Meynet 2000; Villamariz et al. 2002). However, some methods for determining the rotational velocities do not distinguish between rotation and other surface-broadening mechanisms, like macroturbulence.
Conventionally, v sin i values are based on linewidth measurements of individual features, mainly metal lines that are apparently free of blends. As the principal broadening mechanism of these lines is stellar rotation, it is possible with sufficient resolution to determine v sin i from the FWHM of the line. Usually metal lines are used; however, in cases of high rotational velocities or high temperatures, metal lines appear blended or are very weak. Therefore, in these cases, the whole He spectrum is used if the v sin i is high; if v sin i is not extremely high, only He I lines are used, as these lines are less affected by pressure broadening than are He II lines. However, the v sin i derived must be tested with some metal lines (if available), as we are not completely sure that rotation broadening will dominate over pressure broadening.
The Fourier method of determining v sin i is based on how in
Fourier space, convolutions transform into products and how, of the rotation,
macroturbulence, natural, and instrumental profiles (turbulence and
instrumental assumed Gaussian), only the rotation function has zeros in
its Fourier transform. These zeros will appear in the total transform
function, while Carroll (1933) showed that the position of the zeros
are related to the v sin i. Actually, the frequency of the first zero (
)
is related to the rotational velocity through
![]() |
(1) |
The main problems in applying of the Fourier method to early type
stars are related to the quality of the observed spectra (i.e. spectral
resolution and SNR). The lowest v sin i limit that can be determined
is given by the spectral resolution (
in Å/pixel), as the
sampling of the computational Fourier transform cannot be extended beyond 0.5/
.
The noise in the observed spectra transforms as white noise that obscure
the first zero.
The advantages of the Fourier method are that rotational broadening can be separated from other broadening mechanisms and therefore that metal lines, as well as He lines, can be used for the v sin i determination (even for low values of v sin i). This is very useful for fast-rotating stars and spectral types earlier than O9 that show blended or very weak metal lines.
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Figure 3:
Four different synthetic lines generated with F ASTWIND were convolved with a v sin i of 60 km s-1 and degraded to an SNR of 200. The different lines show the Fourier transform of the lines on a
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We tested the Fourier method with theoretical and observational cases, and it works well for massive hot stars. A paper with these results, as well as determinations for a number of O star rotational velocities, is in preparation (Simón-Díaz & Herrero 2006); Fig. 3 shows a typical example, and for a recent application to A-stars see Royer et al. (2002).
The v sin i of our sample of stars was determined through the
Fourier and FWHM methods. Results are presented in Table 3,
together with some values found in the literature for comparison.
The resolution in the IDS spectra is 0.23 Å/pixel, so the lowest
v sin i that could be detected with the Fourier method is 20 km s-1. For
Ori C FEROS spectra, the resolution is
0.03 Å/pixel, so v sin i
2 km s-1 for detection. For
the
Sco spectrum, the resolution is 0.1 Å/pixel, so the
lowest detectable v sin i is 8 km s-1.
All the papers refered to in Table 3 use the FWHM method
applied to the optical spectra of the stars except those by Howarth
et al. (1997), who use a cross correlation technique for
IUE spectra, and Schönberner et al. (1988), who
compare the observed spectrum of Sco with synthetic profiles.
Comments on the individual stars' v sin i determination, as well as the
comparison between the values derived through Fourier and FWHM methods
are presented in Sect. 4.1. Agreement between both
methodologies is very good; however, there are some interesting cases
(see the study of Ori C in Sect. 5).
Table 3: Projected rotational velocity derived from the Fourier and FWHM analyses. References from the literature are: a Simón-Díaz et al. (2003), b Howarth et al. (1997), c McNamara & Larsson (1962), d Schönberner et al. (1988), e Killian et al. (1991).
Table 4:
Stellar parameters derived from F ASTWIND analysis. Only an upper limit for log Q can be derived for these stars. The microturbulences considered for the HHe analysis in each star are shown in the corresponding fitting plots. A normal value for the He abundance was considered for all the stars ( = 0.09).
The analyses were performed using the latest version of F ASTWIND (an acronym for Fast Analysis of STellar atmospheres with WINDs), a code originally described by Santolaya-Rey et al. (1997). See Puls et al. (2005) for the newest description of the code along with a discussion of comparisons with previous models and other spherical mass-losing codes. The latest version uses more complete line-blanketing and a temperature correction method based on the energy balance of electrons. The technique used for deriving the stellar parameters is already standard and has been described elsewhere (Herrero et al. 2002; Repolust et al. 2004), so we only give the main points. The analyses are based on visual fitting of hydrogen Balmer lines and He I and He II lines. Through the He I/He II ionization equilibrium, the effective temperature can be estimated; the wings of the Balmer lines are useful for determining the gravity and can give us some information about the stellar wind.
The code also needs other parameters, such as microturbulence,
He abundance, and wind properties (mass loss, terminal velocity
and parameter). Actually, wind properties are related
through the Q parameter (Q =
/ (
R)1.5).
Once the observed lines are fitted with the modeled ones, effective temperature,
gravity, He abundance, microturbulence, and log Q are determined. The
low density in the winds of the studied objects makes the spectrum insensitive
to changes in Q, so that we can only determine upper values in most cases.
Microturbulence has no effect on the H/He spectrum of early type stars with large
gravity, as has been shown by Villamariz & Herrero (2002). Therefore
only effective temperatures, gravities, and He abundances were determined for this
step of the analysis. Of course, microturbulence is important for the derivation
of metallic abundances and will be determined in the corresponding section.
The code also provides the emergent flux distribution, so mass,
radius, and luminosity can be calculated if
is known (see Herrero
et al. 1992).
Errors in
and log g can be established by generating a microgrid
around central values; visual comparisons between modeled lines and
observations allow us to determine the range of possible values for
these parameters. For further comments on the effects of varying
the various physical parameters used in the analyses, their mutual
interplay and their error bars, see viz. Repolust et al. (2004),
Herrero et al. (2002), Villamariz & Herrero (2002), and
Villamariz et al. (2002). Errors in R, M and L are calculated considering the
propagations of the uncertainties in
,
log g, and
.
The fits of the synthetic F ASTWIND H and He I-II profiles to the observed ones are shown in Figs. 4 to 10. The derived parameters for our sample of stars are shown in Table 4, corresponding to the best fits. Some comments on the individual analyses and the comparison between spectroscopic and evolutionary results are given in Sects. 4.1 and 4.3.
The Fourier method was applied to some O II, Si III-IV,
N II, and He I lines, deriving a v sin i = 55
0.6 km s-1. Metal and He I lines are in agreement. The v sin i derived
through the linewidth measurement method is consistent with this value (see
Table 3).
Figure 4 shows the good fit of the F ASTWIND profiles to
the observed spectra for the parameters given in Table 4, except
for the forbidden component of He I 4471 which is not
reproduced well throughout our analyses. Note also that the apparently bad fit of
He II
4200 is due to the blend with the N III line
at the same wavelength. We are very close to the applicability limit of the He I-II ionization equilibrium for deriving the
,
as He I
4200 and
He I
4541 lines are very faint; however, these lines are still
sufficiently sensitive to changes in
and log g for deriving the
stellar parameters accurately. The stellar parameters obtained here very well agree with those
obtained by Cunha & Lambert (1992) using the Strömgren index c0 and the wings of H
from the Kurucz's (1979) LTE model atmospheres, though it is not the case for the other two stars in common with these authors. A comparison of values is given
in Table 5.
Table 5: Comparison of stellar parameters for HD 37020, HD 37023, HD 37042. The first values refer to the Cunha & Lambert (1992) determinations, the second values to this work. We see that there is excellent agreement for HD 37020, but poor agreement (specially for log g) for the other two stars.
Figure 5 shows the fitting of the HHe lines. Observed
He I lines are slightly broader than the theoretical ones.
Table 5 compares the stellar
parameters we derived with those by Cunha & Lambert (1992).
In this case the agreement is not as good as for HD 37020, although
the
are compatible, the log g they derived is very high.
Comparing the spectrum of this O9V star with that
of the standard star 10 Lac (also classified as O9V), we found that
there are no unblended metal lines due to its high rotational velocity
(except Si IV 4089, but it is in the blue wing of H
). A good v sin i determination was possible using the Fourier method with the He I lines. A v sin i = 131
4 km s-1 was derived. We used the Si IV line to check the reliability of this value; a v sin i = 136
5 km s-1 was obtained. He II does not give good results.
A v sin i = 140
11 km s-1 was derived using the FWHM method for the He I
5015 line.
Figure 6 shows the fitting of the synthetic profiles with the observed ones. A v sin i = 131 km s-1 was considered for the H-He analysis. The wings of the He I lines cannot be well fitted, which might be explained by the presence of a companion.
A very good fit of the observed and synthetic F ASTWIND profiles was obtained (see Fig. 7). Again the
forbidden component of He I 4471 was too weak. For
this star, the He I lines fit better if a microturbulence
of 10 km s-1 was considered.
This is the third star in common with Cunha & Lambert (see Table 5); we also find for this star (as for HD 37023) that the stellar parameters derived by these authors are very different from ours (they obtain a very high log g and a higher
3000 K).
The projected rotational velocity of this star was determined
accurately by means of the FWHM method (v sin i = 37
4 km s-1).
The Fourier method applied to the INT+ IDS spectrum gives
v sin i = 30
0.8 km s-1. This larger difference could be due
to the fact that the v sin i is close to
the computational Fourier transform limit (
20 km s-1 for this
spectrum), or because the microturbulence is affecting the determination
of the rotational first zero in the Fourier transform (Gray 1973).
The whole set of HHe lines is perfectly fitted with the F ASTWIND synthetic profiles (Fig. 8). The parameters derived by Herrero et al. (2002) are
= 35 500 K, log g = 3.95, and
= 0.09. Our results agree with those values.
A very good fit to the synthetic F ASTWIND profiles was obtained for this
star (see Fig. 9). In this case the problem with the
forbidden component of the He I 4471 can be seen clearly.
Table 11 (Sect. 6.4) summarized the stellar
parameters obtained in this and previous work.
We used the spectrum of HD 214680 (O9V) convolved with a high
v sin i (350 km s-1) for recognizing those lines in the spectrum of
HD 47839 not contaminated by the secondary star; three metal lines
were found. Using these lines (Si IV
4212,
4654 and N III
4379) a v sin i = 67
4 km s-1
was determined. A similar value was derived by applying the
FWHM method to the same lines (66
6 km s-1).
The fitting of the H and He lines for the stellar parameters
shows how the He I are contaminated by the secondary star lines. The
spectroscopically derived mass (36
9 M
)
agrees very well
with the dynamical mass derived by Gies (1993). Herrero
et al. (1992) derived a
= 39 500 K, log g = 3.70, and
= 0.07 for this star. Although we expect to obtain a lower
due to including of line-blanketing effects, the value we obtained is slightly higher because there is also a large difference between the log g we derived (4.0) and the one by Herrero et al. (3.7). There is also another difference, as we do not need a lower He abundance for fitting the He lines. This could be due to a binarity effect; when a composite spectrum is considered in a binary system, the lines can appear diluted or magnified due to the combination of the fluxes of the primary and the companion. If the system is out of eclipse, the total flux will be
higher than when the primary is observed isolated, so when the spectrum is
normalized all the lines will appear diluted; and then a lower He abundance is needed to fit the He lines and a lower gravity is derived.
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Figure 11: HR diagram with the studied stars. Evolutionary tracks from Meynet & Maeder (2003). Isocrones from Schaller et al. (1992), corresponding to 2, 2.5, and 3 Myr. |
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From the optical spectra of the Orion stars, only upper limits for the
Q parameter can be achieved. These estimations are based on
the effect of the wind on the He II 4686 and
H
lines, with the second contaminated by the nebular
contribution. Some tests have shown that the other H and
He lines are not affected by the uncertainties in log Qfor the range of values considered, so the derived parameters
will not be affected.
Masses, radii, and luminosities were derived for all the
targets, as indicated in Table 4 together with their
uncertainties. The main source of uncertainty for these parameters
is the one associated with the absolute magnitude, except for very large
uncertainties in log g. An error in
0.3 propagates
to the mass, radius, and logarithmic luminosity, giving uncertainties
of
37%, 13%, and 3%, respectively.
The stars were plotted on the HR diagram in Fig. 11. The evolutionary tracks from Meynet & Maeder (2003), corresponding to initial masses ranging from 9 to 120 Mand initial rotational velocities of 0 km s-1 were also plotted. All stars are found in the Main Sequence close to the ZAMS, as is expected because of their youth. Nevertheless, we can
see the separation from the ZAMS increasing with luminosity, as pointed
out by Herrero et al. 2004. The loci of the Orion stars is consistent with an isochrone at about 2.5
0.5 Myr, derived from the tracks with zero initial rotational
velocity, which is slightly older than the upper limit given by
Palla & Stahler (1999, 2 Myr) and somewhat larger than
other Trapezium age determinations found in the literature (e.g.
Hillenbrand 1997,
1 Myr). However, it has to be considered
that, at large initial rotational velocities, the ZAMS is
slightly shifted to the right and modifies the derived ages. Hence,
until the role of the initial rotational velocities is properly
understood (for example, the distribution of initial rotational
velocities in clusters), the use of isochones for massive stars in
very young clusters should be regarded with special caution.
A good agreement between gravities derived from the evolutionary tracks and those obtained from quantitative analysis of the spectra is found (see Table 6). There is a trend for the most massive stars to have larger spectroscopic than evolutionary masses, but the number of objects is too small to draw any general conclusion.
Table 6:
Comparison of masses and gravities derived from the evolutionary tracks and from quantitative analysis of the spectra. The quoted log
values are given corrected to two decimal places to be consistent with the corresponding evolutionary masses. Note, however, that these are not an indication of the precision of these values, which we consider to be 0.1 dex.
HD 37022 (
Ori C, O7V) is the main ionizing source of the Orion
nebula. We want to derive its stellar parameters as a first step in
determining the effect of its ionizing flux on the photoionization of the
surrounding nebula in a consistent way. Once these parameters are known,
the spectral energy distribution can be modeled by using model atmosphere
codes. In this way one of the inputs used in the photoionization codes will
be fixed consistently.
This star is known to have variable spectral features that vary
in phase or antiphase with a period of 15.422
0.002 d (Stahl et al.
1993; Walborn & Nichols 1994; Stahl et al. 1996).
These variable features were discovered after Conti (1972) showed
for the first time that
Ori C has a variable inverted P-Cygni
profile in the He II 4686 line. Among
them are H
emission, variability in the equivalent
width of some atmospheric and wind lines, and X-ray emission (Caillault
et al. 1994; Gagné et al. 1997).
Different explanations for this variability were postulated.
The possibility of
Ori C being a binary and this binarity
explaining the spectral variability has been dismissed (Stahl et al.
1996). The variability has been associated with the rotation
of the star. Stahl et al. (1996) proposed the presence of a dipolar magnetic
field in
Ori C, with the magnetic pole inclined 45
from the rotation axis (inclined 45
from the line of
sight). The geometry of this system would imply alignment between
magnetic pole and the line of sight at phase 0.5, and
they would be perpendicular at phase 0.0 (when maximum H
emission is found). Babel & Montmerle (1997) propose the magnetically confined wind-shock model (MCWS) for explaining the variability in the star. According to this model, the radiatively line-driven wind is
confined by a dipolar magnetic field towards the magnetic equator of the
system, thereby generating both a cold, dense disk due to the collision of material
coming from both hemispheres and a hot post-shock region.
The wind characteristics of
Ori C were determined by Howard & Prinja
(1989) and Stahl et al. (1996) through UV spectrum studies.
The former derived a mass loss rate of 4
10-7 M
yr-1,
the latter determined a terminal velocity somewhat greater than 2500 km s-1
through the absorption in C IV lines. It is common for O7V stars
to have stellar winds; what is not so common is the detection of magnetic
fields in O stars. Donati et al. (2002) succeeded in detecting
Zeeman features in the spectrum of
Ori C through
spectropolarimetric observations with the Anglo-Australian Telescope. They
detected variability in the Stokes V profiles of some photospheric metal lines. This variability
has a coherent modulation with the period derived from other variable
features. However, the geometry derived contradictes the one from
Stahl et al. (1996), with the magnetic pole aligned with the line
of sight at phase 0, where they found a maximum in the longitudinal component of the
magnetic field.
Once spectral variability is known, it is very important to understand the
cause of this variability and to determine which lines are reliable
for the stellar atmosphere modeling before comparing synthetic and
observed H -He profiles. Preliminary work with the INT+ IDS spectrum showed that a better spectral resolution was needed to apply the Fourier method to obtain
the v sin i. This spectrum allowed us neither to have a good enough
sampling (
)
nor to carry on a study of the variability,
so we decided to use some FEROS spectra with better quality that
covered all variability phases (see Table 7), which are available
in the ESO archive.
Through the study of the FEROS spectra observed in the different phases, we have found all the variable spectral features described in Stahl et al. (1996) and some new ones:
Table 7:
FEROS spectra used for the study of the spectral variability of Ori C. All spectra were downloaded from the ESO- FEROS database except f85221, kindly provided by O. Stahl. The different phases were calculated from
-2.400.000,5 = 48832,5 (Stahl et al. 1996), and P = 15.422 days.
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Figure 12:
The most representative phases were selected to show the variability of the He II ![]() ![]() ![]() ![]() ![]() ![]() |
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This variability can be easily explained considering the model
proposed by Stahl et al. (1996) and developed by Babel
& Montmerle (1997). According to this model we would
have an O7V star with a disc. The disc is produced by the confinement
of the wind by a dipolar magnetic field through the magnetic equator.
We will have a cool disc with material falling back to the stellar surface.
If we consider that at phase 0 the cool disc is seen edge-on, the
blue-shifted emission appearing in He II 4686,
H
and the other hydrogen Balmer lines can be explained as
stellar photons absorbed and reemited with a doppler shift corresponding
to the velocity of the disc material falling to the surface of the
star (in a process similar to what occurs in a stellar wind but with blue-shifted emission
and red-shifted absorption). As density in the disc is very high,
then a strong blue-shifted emission will appear. At phase 0.5,
when the disc is seen face-on, the blue-shifted emission disappears.
The emission appearing in the red wings of the former lines could be
explained as the effect of the scattering of stellar photons by the
wind material confined by the magnetic field and accreting to the disc
(see Fig. 12).
The disc will also have continuum emission that will affect the
total continuum flux received from the star. This effect will be at its
maximum when the disc is face-on because the emitting
region is larger at this phase. The variability observed in
He I, He II, and metal lines (except for the emission
in He I 4686) is only a consequence
of this effect. As the total continuum flux is varying with the
phase, the normalized spectrum will be affected. All absorption lines
will be artificially weaker when the disc continuum emission is at
maximum. This variability in the lines can allow us to estimate
the amount of continuum flux that comes from the disc, and then
the visual magnitude variability. By assuming that at phase 0.0 the
lines are not affected by the continuum emission from the disc, we
scaled the spectra at the other phases to fit in the former spectrum;
the scaling factor will be related to the ratio of visual fluxes (i.e.
the difference in magnitudes) between the stellar component and the
stellar+disc contribution. This can be seen in Fig. 13. From this study we would expect a variability of
0.16 mag. In their catalogue of
suspected variable stars, Kukarkin et al. (1981) found a change in
of 0.06 mag
(5.10-5.16) for
Ori C. Hipparcos has also classified this star as variable; although Hipparcos data do not show a clear pattern, the median magnitude
in Hipparcos system is 4.61 mag and the variability of this magnitude
varies between 4.56 and 4.70, which is agreement with our study.
![]() |
Figure 13:
Magnitude variability expected from the study of the He II ![]() ![]() |
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It is very important to have a good determination of the projected
rotational velocity of this star, as it is supposed that the spectral
variability of Ori C is related to its rotation. The
derived v sin i should be independent of the phase and should be
coherent with a period of
15.4 days. It is shown in Sect. 5.2 that the profiles of metal, as well as He I and He II lines are dependent on the phase; however, this is only an artificial dependence due to the presence of the disk. Once the spectra
of different phases are corrected for the effect of dilution by the disc
continuum emission, all the metal, He I and He II lines are
independent of the phase, except those related to the disc, see
Fig. 12.
The Fourier method allows us to separate pure rotational broadening from
other broadening mechanisms that affects the shape of the lines. We used this
method with some metal lines at phase
0 (see Table 8).
A v sin i = 24
3 km s-1 was derived. Figure 14
shows an example of the application of the Fourier method in
determination of the v sin i of
Ori C.
Table 7 offers a comparison of the v sin i values obtained with the Fourier and FWHM methods. We see that the Fourier method gives more consistent values for all these lines. Differences within the FWHM method may reach a factor of 2 and, in fact, we had problems when trying to fit the profile of some of the lines with a Gaussian profile for measuring their linewidth.
The derived value for v sin i through the Fourier method is also more
coherent with an O7V star rotating with a 15.422 days period.
Considering R
11 R
,
the upper limit for
v sin i is
35 km s-1, so that the inclination
of the rotational pole is i
45, in agreement with previous
results (see next section).
Table 8:
Projected rotational velocities derived through Fourier and FWHM methods for some metal lines present in the spectrum of Ori C. Values were derived at phase 0.972 (see explanation in text).
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Figure 14:
Fourier analysis of the N IV ![]() ![]() |
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Once the spectral variability of Ori C is understood and
its v sin i has been derived, we can proceed to model its stellar
atmosphere and wind through the observed spectrum of the star.
The lines used for this analysis are shown in Fig. 15;
basically, they are the ones used in the other analyses plus
He I 5875
.
Some of the lines are contaminated
(see Sect. 5.2), so this must be taken into account.
The H
and H
lines were selected
as the most reliable lines for deriving log g (less
contaminated than H
and H
). The whole
set of He I-II lines was considered except
He II
4686, but it must be taken into account
that the strength of all these lines vary with the phase. Phase 0.972 will be used for determining the stellar parameters, since the effect of the continuum emission of the disc is smaller at this
phase (see Sect. 5.2).
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Figure 15:
HHe analysis of HD 37022 (![]() ![]() ![]() ![]() ![]() ![]() |
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Although our study has shown that the rotational velocity
(derived from Fourier analysis) is 24 km s-1, when this
broadening is considered all synthetic lines appear narrower
than the observed ones. We tried to solve the problem by means of a different
v sin i value; however, it does not work, because then the shape of
the synthetic profiles do not fit with the observed ones, as the
cores of the modeled profiles are too wide when the FWHM of
the He lines is fitted. An extra-broadening mechanism
has to be included. When a Gaussian macroturbulent broadening (Gray 1973) is used, the fit clearly improves; however, in this case the H I and He II lines cannot be fitted
simultaneously with the He I ones, because when the former are
fitted (for a
= 39 000 K), some of the synthetic lines
in the latter appear stronger and narrower than observed. A better
fitting for the He I lines is obtained if a higher
if considered, but then the synthetic He II lines appear too strong.
There is no way to solve this problem in this region of the
parameter space either by changing the rotational
velocity, the macroturbulence, or the microturbulence.
In Sect. 4.2, Fig. 10 shows that the fitting of the He I-II lines could follow similar behavior in the case of the O7V star HD 47839 (selected as reference star). A variation of 1000 K in the effective temperature strongly changes the strength of the He I lines. For this spectrum it also occurs that when the He I lines appear fitted, the He II lines are slightly stronger than observed, and if the He II lines are fitted, the He I lines are stronger than observed.
Puls et al. (2005) have shown that there is a discrepancy for the He I singlets between the synthetic F ASTWIND and CMFGEN lines for effective temperatures
between 36 000 and 41 000 K (being the CMFGEN profiles shallower). Therefore the He I triplet system can be considered more reliable (i.e. the He I 4471 line). Knowing this discrepancy, we considered the He I
4471 for the fitting with F ASTWIND synthetic profiles. Our best model corresponds to
= 39 000
1000 K and log g = 4.1 dex. A stellar radius
R = 10.6
1.5 R
was derived, which
implies an inclination of the rotational axis of i = 44
12
.
This value agrees with previous independent studies
(Stahl et al. 1996; Donati et al. 2002), although
our derived v sin i is more reliable and our radius is not obtained
from an SpT - R calibration, but is the result of the spectral
analysis of the star.
Table 9:
Equivalent widths and derived line abundances for the set of O II lines used in our analysis. Line abundances refer to the microturbulence given in brackets for each star (
in km s-1). Uncertainties in the line abundances come from the propagation of the uncertainties of the equivalent width measurements (see text). Some of the O II lines of the Orion stars
have not been used, as they appear blended. O II
4072, 4076, and 4078 lines were ruled out in the analysis of
Sco due to the poor quality of the CASPEC spectrum in this region. Final oxygen abundances for each star were calculated through a weighted mean of the linear values. Errors represent the statistical deviation for these mean values.
Three of the stars studied in Orion are perfect targets for a stellar
abundance analysis, as they have many narrow unblended lines.
These are HD 37020, HD 37023, and HD 37042, all three B0.5V stars. Cunha
& Lambert (1992, 1994) present carbon, nitrogen,
oxygen, and silicon abundances from LTE and NLTE analyses
for these stars. For comparison purposes, a similar analysis has been
done for Sco, a B0.2V star (Walborn & Fitzpatrick 1990) with very
low v sin i. Stellar abundances for this star have been derived elsewhere in the literature (Hardorp & Scholz 1971; Kane et al. 1980; Peters & Polidan 1985;
Schönberner et al. 1988; Becker & Butler 1988; Kilian et al. 1991; Martin 2004, see Table 11). The other two Orion stars have been ruled
out: HD 37041 has a relatively high projected rotational velocity, so
metallic lines are broadened and then appear blended; HD 37022, being
an O7V star, does not have enough oxygen lines for an accurate abundance analysis.
We therefore derived oxygen abundances for the Trapezium Cluster B0.5V stars for comparing them with the M 42 nebular abundances obtained by Esteban et al. (2004).
We made use of the curve of growth classical method to determine
the oxygen abundances. When this methodology is used, it is important
to remove all lines that appear blended. The
spectrum of Sco (a star with similar spectral type to our
targets in Orion and a very low v sin i) was used for an identification of the O II lines present in the spectra of the Orion stars. The whole set of lines is shown in Tables 12 and 13
divided into multiplets, together with their log gf values (basically
taken from the NIST database).
The equivalent widths of all the O II lines listed in Tables 12 and 13 were measured for the four stars; however, only our set of suitable lines (see below) is shown in Table 9. To measure the equivalent widths we use our own software developed in IDL. A least
squares profile-fitting procedure was used, with Gaussian profiles
fitting the line and polynomials of degree one or two to fit the local
continuum. Errors in the measurements due to uncertainty in the
position of the local continuum (estimated as 1/SNR, Villamariz
et al. 2002) were also considered. The estimated value of
the uncertainty in the measurement of the EWs is
5 mÅ
and
10 mÅ for some problematic lines.
Some of the lines that appear unblended in the spectrum of Sco
cannot be used in the analyses of the other stars, since they
have larger rotational broadening and then appear blended (or lie in the
wings of H or He lines). Once a first set of unblended
lines was selected, a preliminary abundance analysis was done separately
for each multiplet. In this way the dispersion in the line abundances for
the zero slope value of the microturbulence are minimized as all the
lines in a multiplet are formed in the same region in the stellar
atmosphere. Problematic lines, errors in the measurement of the equivalent
widths or artificial trends can then be identified; such lines will be
removed in the global analysis; e.g. this is the case for the
O II
4414 line, an isolated line that systematically gives
lower abundances.
The set of suitable lines finally used in the abundance analyses is
presented in Table 9. They are lines coming from
transitions between configurations 2p2 (3P) 3s-2p2 (3P) 3p
(NIST multiplets 64, 65, and 72) and 2p2 (3P) 3p-2p2 (3P) 3d
(NIST multiplets 90, 148, and 130). We ruled out the lines that do not follow the general trend and found that lines from multiplets 99, 118, 161, 188, 172 give systematically lower abundances. This effect can be due to the definition of the oxygen model atom we used, or can be associated
with the log gf values. Some preliminary comparisons with TLUSTY analyses have shown that the difference in the line abundances is also present for the case of Sco.
For each star we proceed as follows. A grid of 16 F ASTWIND models
combining four abundances and four values of microturbulence is
calculated. In this way, the curves of growth for each line can be
constructed by plotting the theoretical equivalent width for each
abundance and microturbulence versus the abundance (see Fig. 16 for the case of the O II 4414 line in HD 37042).
Through the observed equivalent width and its error, one abundance
(and its derived uncertainty) can be derived for each line and
microturbulence. The individual line abundances are dependent on the
microturbulence. The microturbulence value that minimizes the
dependence of the line abundances on the line strength in the
log - EW diagrams will be the microturbulence we
are looking for. These diagrams can also be used as a diagnostic tool
to check the reliability of the different lines for the abundance
determination (see Sect. 6.2). Figure 17 shows the log
- EW diagrams for two different microturbulences in the study of HD 37023.
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Figure 16:
Example of the curve of growth for the line O II ![]() |
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![]() |
Figure 17:
Example of log ![]() |
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The microturbulence derived from the zero slope for each star is presented
in Table 10. Uncertainties in the microturbulence are obtained
considering the errors derived for the slope in a linear fit of the data,
due to errors in the individual line abundances. This step also allows
us to estimate the contribution of the uncertainty in the microturbulence
to the total oxygen abundance. This uncertainty depends mainly
on the quality of the spectra, in this case 0.06 dex.
Abundance values for each line, as well as their uncertainties, are obtained
for that microturbulence (see Table 9). The final abundance
value is calculated through a weighted mean of the linear individual line
abundances (10
), and its uncertainty (
)
is the one associated with this mean.
Final values for the total oxygen abundances for each star are shown
in Table 10. The final uncertainty in the oxygen abundance
takes four different sources of errors into account: those associated
with the statistical analysis, those derived from the error in the
determined microturbulence and finally those caused by the uncertainty both
in the stellar parameters and in the atomic data. All these sources of
error are added quadratically to derive the final abundance
uncertainty (see Villamariz et al. 2002).
Table 10:
Oxygen abundances for the three B0.5V stars inside Orion nebula and the reference star Sco. Oxygen LTE and NLTE abundances derived by Cunha & Lambert
(1994) for the Orion stars as well as those calculated by Esteban et al. (2004) for the nebula are also presented for comparison.
Table 11:
Comparison of stellar parameters and abundances derived for Sco in previous studies found in the literature and in this work.
Table 12:
Preliminary set of O II lines selected for the analysis, divided by multiplets. The spectrum of the low v sin i star Sco was used to identify the lines. The log gf values are from the NIST database.
Table 13:
(Continued) Preliminary set of O II lines selected for the analysis, divided by multiplets. The spectrum of the low v sin i star Sco was used to identify the lines. The log gf values are from the NIST database.
Table 11 summarizes the oxygen abundances appearing
in the literature for Sco. Our derived value is
compatible with previous results but a little higher
than most of them. This difference can be easily explained by
taking into account that, for this range of stellar parameters,
the oxygen abundance derived from O II lines is very
sensitive to a change in
and log g: the lines become fainter when
a higher
is considered and then the derived oxygen
abundance is higher. The uncertainty in the oxygen abundance
due to a change of
1000 K in
is
0.08 dex.
This effect is considered in the uncertainty for the given values;
however, the central value will slightly depend on the
derived stellar parameters.
We considered two F ASTWIND models with different
(32 000,
32 500 K) and the same log g, thereby obtaining oxygen abundances of 8.70 and 8.74 dex, respectively.
The derived abundance is also very dependent on the chosen microturbulence, especially if lines with high equivalent width are used. We have taken these dependencies into account in our uncertainties.
The oxygen abundances derived for the Orion stars are
compatible within the errors (see Table 10). HD 37023
has a slightly lower abundance, but one still compatible with
the other abundances. In Sect. 4.1 we have seen that
the fitting of the H and He lines is not as good as for the
other B0.5V stars, since the observed lines appear slightly broader.
The oxygen abundances in the Orion stars are systematically
lower than that derived for Sco.
Esteban et al. (2004) have recently published a reappraisal
of the chemical composition of the Orion nebula. They derived a total oxygen gas-phase abundance (O) = 8.65
0.03. However, some oxygen is expected to be depleted onto dust grains in
ionized nebula, so the total gas+dust oxygen abundance should take this depletion into account.
In a previous work (Esteban et al. 1998), these authors
estimate the depletion onto dust grains by comparing Si and Fe nebular
abundances with those obtained by Cunha & Lambert (1994)
for B stars in the Orion association, assuming a certain composition
for the main dust molecules. Taking this correction
into account, the final gas+dust oxygen abundance that
Esteban et al. (2004) propose is
(O) = 8.73
0.03, where an oxygen abundance correction for
dust
0.08 is applied.
Our stellar results are compatible with those obtained by Esteban et al. (2004) for the gas phase; however, the dust+gas corrected abundance seems to be too high compared with our derived stellar abundances, although still marginally consistent inside the uncertainties.
Our oxygen abundances are systematically lower than the NLTE abundances
by Cunha & Lambert (1994). In that paper the former authors
comment that the LTE abundances are slightly more reliable than
the NLTE abundances they present. The difference between their LTE abundances and our results are even higher. These differences can be associated in
part with the differences in the derived stellar parameters
for these stars. The
and log g obtained by these authors are
higher than ours (see Table 5), so the oxygen abundances
they derive are obviously higher.
This difference in the effective temperatures may also affect the
derived stellar silicon abundances used by Esteban et al. 1998 for
estimating the oxygen depletion. Cunha & Lambert (1994), derive
their silicon abundances by using 3 Si III lines. Preliminary
silicon analysis by our group has shown that a difference of 1000 K
in
can shift the Si III abundances up to 0.2 dex, deriving a
higher abundance for the lower
.
Alternatively, our result could suggest that the molecules that Esteban et al. use to estimate the O dust depletion in Orion cannot be present in large amounts in this nebula. Consequently, Si, Mg, and Fe (the refractory elements being the main constituents of those molecules) have to form other molecules without oxygen.
By means of a detailed spectroscopic analysis of the optical
spectra of the Trapezium cluster stars, we derived stellar parameters
and oxygen abundances. Projected rotational velocities were obtained through Fourier method. This
method has been extensively used for late type stars, but not widely applied to early type stars. Our results show this method is very useful for distinguishing between rotational
broadening and another broadening mechanisms that can be present in
early type stars (e.g. macroturbulence). The agreement is very good when comparing with results from the line-width method. The Fourier method applied to the high resolution Ori C
FEROS spectra allow us to derive a very accurate v sin i that agrees
with the period of variability of some spectral features in
Ori C.
Stellar parameters and their uncertainties were derived
for these stars using H , He I, and He II lines and the F ASTWIND code.
The presence of many O II lines in the optical spectrum of three B0.5V Orion stars allowed us to work on a very detailed abundance analysis using the curve of growth method. This analysis was performed through a careful selection of suitable lines from a previous study of the different O II multiplets. In this way, the dispersion in the line abundances is reduced, and the final abundance value derived is very precise.
The derived oxygen abundances in the Orion stars agree with the nebular gas-phase abundances obtained by Esteban et al. (2004), and 0.2 dex lower than the NLTE abundances derived by Cunha & Lambert (1994). The gas+dust corrected oxygen abundances estimated by
Esteban et al. (1998, 2004), using the Cunha & Lambert stellar abundances in the Orion association, seem to be too high compared with our derived abundances, although still marginally consistent within the uncertainties. This result suggests a lower dust depletion factor of oxygen than previous estimations for the Orion nebula. A revision of the silicon, magnesium,
and iron stellar abundances in the Trapezium cluster stars is presently under way in our group to confirm this result.
Acknowledgements
We want to thank M. A. Urbaneja for the original procedures for deriving the stellar abundances and his invaluable help with F ASTWIND, D. Lennon for his comments, and also P. Dufton and R. Ryans for calculating some TLUSTY models for comparison. This work was partially funded by the Spanish Ministerio de Educación y Ciencia under project AYA2004-08271-C02-01. We are very grateful to T. Gehren and O. Stahl for lending us the spectra ofSco and
Ori C. This research made use of the ESO-FEROS database.
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Figure 4:
Analysis of HD 37020 (![]() ![]() ![]() |
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Figure 5:
As in Fig. 4 for HD 37023 (![]() |
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Figure 6:
As in Fig. 4 for HD 37041 (![]() ![]() |
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Figure 7:
As in Fig. 4 for HD 37042 (![]() |
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Figure 8: As in Fig. 4 for HD 214680 (10 Lac, O9V). A variation of 0.1 dex in log g was considered in this case. |
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Figure 9:
As in Fig. 4 for HD 149438 (![]() ![]() ![]() ![]() |
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Figure 10:
As in Fig. 4 for HD 47839 (15 Mon, O7V). A variation of 1000 K in
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