_lines.ps
Astronomy & Astrophysics manuscript no. 6489 c© ESO 2018
October 28, 2018
The VLT-FLAMES survey of massive stars:
Wind properties and evolution of hot massive stars in the LMC
M. R. Mokiem1, A. de Koter1, C. J. Evans2, J. Puls3, S. J. Smartt4, P. A. Crowther5, A. Herrero6,7, N. Langer8,
D. J. Lennon9,6, F. Najarro10, M. R. Villamariz11,6, and J. S. Vink12
1 Astronomical Institute Anton Pannekoek, University of Amsterdam, Kruislaan 403, 1098 SJ Amsterdam, The Netherlands
2 UK Astronomy Technology Centre, Royal Observatory, Blackford Hill, Edinburgh, EH9 3HJ, UK
3 Universitäts-Sternwarte München, Scheinerstr. 1, D-81679 München, Germany
4 The Department of Pure and Applied Physics, The Queen’s University of Belfast, Belfast BT7 1NN, Northern Ireland, UK
5 Department of Physics and Astronomy, University of Sheffield, Hicks Building, Hounsfield Rd, Shefffield, S3 7RH, UK
6 Instituto de Astrofı́sica de Canarias, E-38200, La Laguna, Tenerife, Spain
7 Departamento de Astrofı́sica, Universidad de La Laguna, Avda. Astrofı́sico Francisco Sánchez, s/n, E-38071 La Laguna, Spain
8 Astronomical Institute, Utrecht University, Princetonplein 5, 3584 CC Utrecht, The Netherlands
9 The Isaac Newton Group of Telescopes, Apartado de Correos 321, E-38700, Santa Cruz de La Palma, Canary Islands, Spain
10 Instituto de Estructura de la Materia, Consejo Superior de Investigaciones Cientı́ficas, CSIC, Serrano 121, E-28006 Madrid, Spain
11 Grantecan S.A., E-38200, La Laguna, Tenerife, Spain
12 Astrophysics Group, Lennard-Jones Laboratories, Keele University, Staffordshire, ST55BG, UK
Accepted: 24 January 2007
ABSTRACT
We have studied the optical spectra of a sample of 28 O- and early B-type stars in the Large Magellanic Cloud, 22 of which are associated with
the young star forming region N11. Our observations sample the central associations of LH9 and LH10, and the surrounding regions. Stellar
parameters are determined using an automated fitting method (Mokiem et al. 2005), which combines the stellar atmosphere code fastwind (Puls
et al. 2005) with the genetic algorithm based optimisation routine pikaia (Charbonneau 1995). We derive an age of 7.0 ± 1.0 and 3.0 ± 1.0 Myr
for LH9 and LH10, respectively. The age difference and relative distance of the associations are consistent with a sequential star formation
scenario in which stellar activity in LH9 triggered the formation of LH10. Our sample contains four stars of spectral type O2. From helium and
hydrogen line fitting we find the hottest three of these stars to be ∼49−54 kK (compared to ∼45−46 kK for O3 stars). Detailed determination of
the helium mass fraction reveals that the masses of helium enriched dwarfs and giants derived in our spectroscopic analysis are systematically
lower than those implied by non-rotating evolutionary tracks. We interpret this as evidence for efficient rotationally enhanced mixing leading to
the surfacing of primary helium and to an increase of the stellar luminosity. This result is consistent with findings for SMC stars by Mokiem et
al. (2006). For bright giants and supergiants no such mass discrepancy is found; these stars therefore appear to follow tracks of modestly or non-
rotating objects. The set of programme stars was sufficiently large to establish the mass loss rates of OB stars in this Z ∼ 1/2 Z⊙ environment
sufficiently accurate to allow for a quantitative comparison with similar objects in the Galaxy and the SMC. The mass loss properties are found
to be intermediate to massive stars in the Galaxy and SMC. Comparing the derived modified wind momenta Dmom as a function of luminosity
with predictions for LMC metallicities by Vink et al. (2001) yields good agreement in the entire luminosity range that was investigated, i.e.
5.0 < log L/L⊙ < 6.1.
Key words. Magellanic Clouds – stars:atmospheres – stars: early-type – stars: fundamental parameters – stars: mass loss
1. Introduction
Massive stars play an intricate role in the evolution of galax-
ies. Because of the large energies associated with their stel-
lar winds, ionising radiation, and life-ending supernova ex-
plosions, they dictate galactic structuring processes such as
star formation and the creation and evolution of supperbub-
bles (e.g. Oey 1999). Mounting evidence also points to a direct
link between massive stars and exotic phenomena such as γ-ray
bursts (e.g. Hjorth et al. 2003) and the reionisation of the early
universe (Bromm et al. 2001). Accordingly, understanding the
properties of these stars, both in terms of their fundamental pa-
rameters as well as their evolution, is fundamental.
The initial metal composition (Z) of the gas out of which
massive stars form has a strong impact on their global prop-
erties and characteristics. Many studies have shown that pa-
rameters such as the effective temperature and ionising fluxes
are strong functions of Z (e.g. Kudritzki 2002; Mokiem et al.
2004; Massey et al. 2005; Mokiem et al. 2006), adding an ex-
tra dimension to the conversion of morphological properties
http://arxiv.org/abs/0704.1113v1
2 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
such as spectral type to physical quantities. Theoretical and ob-
servational arguments (e.g. Kudritzki & Puls 2000; Vink et al.
2001) also point to a relation between the strengths of the stel-
lar winds of these objects and metallicity. As wind mass loss
leads to partial evaporation of the star (e.g. Chiosi & Maeder
1986) – with possible consequences for the nature of the com-
pact object that is left behind after the final supernova explosion
– and to loss of angular momentum (e.g. Meynet & Maeder
2000), it also dictates to a large extent its evolutionary path
and fate (e.g. Yoon & Langer 2005; Woosley & Heger 2006).
Quantifying the mass loss versus metallicity dependence Ṁ(Z),
therefore, is an important quest in astrophysics.
Due to their proximity and low metal content the
Magellanic Clouds provide us with unparallelled laboratories
to test and enlarge our knowledge of massive stars. These
galaxies, therefore, have been in the focal point of many stud-
ies analysing their massive star content. Early studies (e.g.
Conti et al. 1986; Garmany et al. 1987; Massey et al. 1989;
Parker et al. 1992; Walborn et al. 1999) predominantly relied
on photometric data and spectral type calibrations. Only rel-
atively recent the advent of large telescopes and the develop-
ment of sophisticated stellar atmosphere models has allowed
for more detailed analyses of individual stars (e.g. Puls et al.
1996; Hillier & Miller 1999; Crowther et al. 2002; Bouret et al.
2003; Martins et al. 2004). Though all these studies have con-
tributed enormously to our understanding of massive stars, the
samples analysed so far have been rather limited in size (a few
objects at a time) and have been focused predominantly on ob-
jects in the Small Magellanic Cloud (SMC) – because its metal
deficiency is more extreme than that of the Large Magellanic
Cloud (LMC) (but see Massey et al. 2004, 2005).
The limited sizes of the samples that have so far been stud-
ied are at the root cause of the perhaps somewhat disheartening
conclusion that – in spite of all the progress that has been made
– we still cannot provide robust and sound answers to the ques-
tion: what is the role of metal content, stellar winds and rota-
tion in the evolution of massive stars? To help attack this prob-
lem, our research group has conducted a VLT-FLAMES Survey
of Massive Stars (see Evans et al. 2005). In this ESO Large
Program the Fibre Large Array Multi-Element Spectrograph
at the Very Large Telescope was used to obtain optical spec-
tra of more than 50 O- and early B-type stars in the Magellanic
Clouds.
Here we present the homogeneous analysis, employing au-
tomated spectral fitting methods, of a sample of 28 O-type and
early B-type stars in the Large Magellanic Cloud; 22 targets
from the FLAMES survey and 6 from other sources. This is
so far the largest sample of massive stars studied in the LMC
and it almost doubles the amount of massive objects in this
galaxy for which parameters have been derived from quantita-
tive spectroscopy. Specifically, we will try to establish the mass
loss rates of OB stars in this Z ∼ 1/2 Z⊙ environment to a level
of precision that allows for a quantitative comparison with sim-
ilar objects in the SMC and our Galaxy. This will provide a new
Z-point in testing the fundamental prediction provided by radi-
ation driven wind theory for the mass-loss – metallicity depen-
dence: Ṁ(Z) ∝ Z0.5−0.7 (e.g. Kudritzki et al. 1987; Puls et al.
2000; Vink et al. 2001).
The majority of our LMC sample is associated with
the spectacular star forming region N11 (Henize 1956). It
has a Hα luminosity only surpassed by that of 30 Doradus
(Kennicutt & Hodge 1986), ranking it as the second largest
H ii region in the Magellanic Clouds. N11 is host to several
OB associations of apparently different ages, the formation of
which is believed to have been triggered by stellar activity in
the central OB cluster (Parker et al. 1992; Walborn & Parker
1992; Walborn et al. 1999). Our observations sample both the
central cluster LH9 as well as the younger cluster LH10, allow-
ing for an investigation of a possible sequential star formation
scenario. We will use our N11 stars to test predictions of mas-
sive star evolution, including the role of rotation, and the star
formation history.
This paper is organised as follows: in Sect. 2 we describe
the LMC data set that was analysed using our automated ge-
netic algorithm based fitting method. A short description of this
method is given in Sect. 3 and the results obtained are presented
in Sect. 4, with fits and comments on individual objects given
in the appendix. In Sect. 5 we investigate the discrepancy be-
tween spectroscopically determined masses and those derived
from evolutionary tracks. The evolutionary status of N11 is dis-
cussed in Sect. 6. Finally, Sect. 7 summarises and lists our most
important findings.
2. Data description
Our OB-type star sample is mainly drawn from the targets ob-
served in the LMC within the context of the VLT-FLAMES
survey of massive stars (see Evans et al. 2005). Two fields in
the LMC centred on the clusters N11 and NGC 2004 were
observed in the survey. Here we analyse a subset of the ob-
jects observed in the N11 field. This set consists of all O-type
spectra obtained, excluding those that correspond to confirmed
binaries, and five early B-type spectra of luminous giant and
supergiant stars.
To improve the sampling in luminosity and temperature, we
supplemented the FLAMES targets with six relatively bright O-
type field stars. These objects are part of the Sanduleak (1970,
hereafter Sk) and Brunet et al. (1975, hereafter BI) catalogues,
and were observed as part of the programs 67.D-0238, 70.D-
0164 and 074.D-0109 (P.I. Crowther) using the Ultraviolet and
Visual Echelle Spectrograph (UVES) at the VLT.
The observations of the FLAMES targets are described ex-
tensively by Evans et al. (2006), to which we refer for full de-
tails. Here we only summarise the most important observa-
tional parameters. Basic observational properties of the pro-
gramme stars together with common aliases are given in Tab. 1.
Note that N11-031, BI 237, BI 253 and Sk−67 166 were stud-
ied recently using line blanketed stellar atmosphere models. In
the appendix a comparison with these analyses is provided. The
FLAMES targets were observed with the Giraffe spectrograph
mounted at UT2. For six wavelength settings a spectrum was
acquired six times for each object with an effective resolving
power of R ≃ 20 000. These multiple exposures, often at differ-
ent epochs, allowed for the detection of variable radial veloc-
ities. As a result, a considerable number of binaries could be
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 3
Table 1. Basic parameters. Primary identification numbers for N11 are from Evans et al. (2006). Identifications starting with
“Sk”, “PGMW” and “BI” are from, respectively, Sanduleak (1970), Parker et al. (1992) and Brunet et al. (1975). Photometric
data for these objects are from Evans et al. (2006) and Parker et al. (1992), the latter are flagged with an asterisk. For non-N11
objects these are from Ardeberg et al. (1972), Issersted (1979) and Massey (2002). Wind velocities given without brackets are
from Crowther et al. (2002), Massa et al. (2003) and Massey et al. (2005). For Sk −66 18 the wind velocity was measured from
O iv (1031-1037 Å) in its FUSE spectrum. Values between brackets are calculated from the escape velocity at the stellar surface.
For N11-031 the value of v∞ is from Walborn et al. (2004), though this is actually based on the value obtained for Sk-68 137 by
Prinja & Crowther (1998).
Primary ID Cross-IDs Spectral V AV MV v∞
Type [km s−1]
N11-004 Sk −66 16 OC9.7 Ib 12.56 0.74 −6.68 [2387]
N11-008 Sk −66 15 B0.7 Ia 12.77 0.84 −6.57 [1619]
N11-026 ... O2 III(f*) 13.51 0.47 −5.46 [3116]
N11-029 ... O9.7 Ib 13.63 0.56 −5.43 [1576]
N11-031 PGMW 3061 ON2 III(f*) 13.68∗ 0.96 −5.78 3200
N11-032 PGMW 3168 O7 II(f) 13.68∗ 0.65 −5.47 [1917]
N11-033 PGMW 1005 B0 IIIn 13.68 0.43 −5.25 [1536]
N11-036 ... B0.5 Ib 13.72 0.40 −5.18 [1714]
N11-038 PGMW 3100 O5 II(f+) 13.81∗ 0.99 −5.68 [2601]
N11-042 PGMW 1017 B0 III 13.93 0.22 −4.79 [2307]
N11-045 ... O9 III 13.97 0.50 −5.03 [1548]
N11-048 PGMW 3204 O6.5 V((f)) 14.02∗ 0.47 −4.95 [3790]
N11-051 ... O5 Vn((f)) 14.03 0.19 −4.66 [2108]
N11-058 ... O5.5 V((f)) 14.16 0.28 −4.62 [2472]
N11-060 PGMW 3058 O3 V((f*)) 14.24∗ 0.81 −5.07 [2738]
N11-061 ... O9 V 14.24 0.78 −5.04 [1898]
N11-065 PGMW 1027 O6.5 V((f)) 14.40 0.25 −4.35 [2319]
N11-066 ... O7 V((f)) 14.40 0.25 −4.35 [2315]
N11-068 ... O7 V((f)) 14.55 0.28 −4.23 [3030]
N11-072 ... B0.2 III 14.61 0.09 −3.98 [2098]
N11-087 PGMW 3042 O9.5 Vn 14.76∗ 0.62 −4.36 [3025]
N11-123 ... O9.5 V 15.29 0.16 −3.37 [2890]
BI 237 ... O2 V((f*)) 13.89 0.62 −5.23 3400
BI 253 ... O2 V((f*)) 13.76 0.71 −5.45 3180
Sk −66 18 ... O6 V((f)) 13.50 0.37 −5.37 2200
Sk −66 100 ... O6 II(f) 13.26 0.34 −5.58 2075
Sk −67 166 HD 269698 O4 Iaf+ 12.27 0.31 −6.54 1750
Sk −70 69 ... O5 V 13.95 0.28 −4.83 2750
detected (Evans et al. 2006), which we subsequently excluded
from our analysis.
To allow for a sky subtraction a master sky-spectrum was
created from combining the sky fibres in the Giraffe spectro-
graph (typically 15), individually scaled by their relative fibre
throughput. Even though in general the sky background is low
and this approach successfully removes the background con-
tribution, in crowded regions such as N11 accurate subtraction
of nebular features remains very difficult. As a result of this,
the line profiles of many of our programme stars still suffer
from nebular contamination. This in principle does not hamper
our analysis. For most stars the core nebular emission is well-
resolved and we simply disregard this contaminated part of the
profile in the automated line fits. Mokiem et al. (2006) showed
by performing tests using synthetic data that with this reduced
amount of information, the automated method can still accu-
rately determine the correct fit parameters. Also tests assessing
the impact of possible residual nebular contamination in the
4 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
line wings or over-subtraction of sky components showed that
its effect is negligible.
For each wavelength range the individual sky-subtracted
spectra were co-added and then normalised using a cubic-
spline fit to the continuum. A final spectrum covering 3850-
4750 and 6300-6700 Å was obtained by merging the nor-
malised data. Depending on the magnitude of the target these
combined spectra have typical signal-to-noise ratios of 100-
The first four field stars listed in Tab. 1 were observed with
the VLT-UVES spectrograph in service mode on 29 and 30
November 2004 under program 74.D-0109. UVES is a two
armed cross-dispersed echelle spectrograph allowing for simul-
taneous observations in the blue and red part of the spectrum.
In the blue arm standard settings with central wavelengths of
390 and 437 nm were used to observe the spectral ranges 3300-
4500 and 3730-5000 Å. For the red arm standard settings with
central wavelengths of 564 and 860 nm provided coverage be-
tween 4620-5600 and 6600-10400 Å. A 1.2′′ wide slit was
used, providing a spectral resolution of 0.1 Å at Hγ, corre-
sponding to an effective resolving power of R ≃ 40 000, a value
which applies to all UVES setups. Individual exposure times
ranged from 1500 to 2200 seconds.
Sk −67 166 was observed on 27 and 29 September 2001
under program 67.D-0238 with UVES using a 1′′ slit. A stan-
dard blue setting with central wavelength 437 nm provided con-
tinuous coverage between 3730-5000 Å. A non-standard red
setup with central wavelength 830 nm was used to observe
the range between 6370-10250 Å. Three exposures, each of
1000 s were obtained for each setup. Note that the data for
this target was previously presented by Crowther et al. (2002).
Finally, Sk −70 69 was observed on 1 and 2 December 2002
using UVES under program 70.D-0164. Standard blue and red
settings with central wavelengths 390 and 564 nm were used
in simultaneous 2400 s exposures using dichroic. A second
non-standard red setup with central wavelength 520 nm (4170-
6210 Å) without dichroic was used in a 1500 s exposure. The
typical two pixel S/N ratios obtained for all spectra are 80 at
Hγ and 60 at Hα.
Spectral types for FLAMES stars were determined by vi-
sual inspection of the spectra, using published standards. In
particular the atlas of Walborn & Fitzpatrick (1990) was used
while considering the lower metallicity environment of the
LMC. The classifications given in Tab. 1 are from Evans et al.
(2006). A comparison with previously published spectral types
is also given in the latter paper. For the field stars we adopted
the classifications given by Walborn (1977), Walborn et al.
(1995, 2002a) and Massey et al. (1995, 2005).
Photometric data for the FLAMES targets was obtained
predominantly from B and V images of the N11 field
taken with the Wide Field Imager at the 2.2-m Max Planck
Gesellschaft/ESO telescope on 2003 April 2004 (Evans et al.
2006). The photometry for stars flagged with an asterisk was
adopted from Parker et al. (1992). The caption of Tab. 1 lists
the references to the various sources of the photometric data of
the field stars.
To calculate the interstellar extinction (AV ) given in Tab. 1
we adopted intrinsic colours from Johnson (1966, and refer-
ences therein) and a ratio of total to selective extinction of
RV = 3.1. With these AV values, extinction corrected visual
magnitudes (V0) were calculated from the observed V-band
magnitudes. Finally, we calculated the absolute visual magni-
tude MV , while adopting a distance modulus of 18.5 for the
LMC (Panagia et al. 1991; Mitchell et al. 2002).
3. Analysis method
All optical spectra are analysed using the automated fitting
method developed by Mokiem et al. (2005, hereafter referred
to as Paper I). Here we will suffice with a short description
of the method and refer to the before mentioned paper and to
Mokiem et al. (2006, hereafter Paper II) for the details.
In short the automated fitting method uses the genetic algo-
rithm based optimisation routine pikaia (Charbonneau 1995) to
determine the set of input parameters for the stellar atmosphere
code fastwind (Puls et al. 2005) which fit an observed spec-
trum the best. This best fitting model is constructed by evolving
a population of fastwind models over a course of generations.
At the end of every generation the parameters of the models
which relatively fit the observed spectrum the best are used to
construct a new population of models. By repeating this proce-
dure a natural optimisation is obtained and after a number of
generations (see below) the best fitting model, i.e. global opti-
mum in parameter space, is found.
Using the concept of a unified model atmosphere the fast
performance code fastwind incorporates non-LTE and an ap-
proximate approach to line blanketing to synthesise hydrogen
and helium line profiles. Consequently, given the spectral range
of the observed data set in this study we will focus on the mod-
elling of the optical hydrogen and helium lines. To account for
the accuracy with which each individual line can be reproduced
by fastwind we adopt the line weighting scheme as described
in Paper I.
In the fitting of the spectra we allow for seven free parame-
ters. These are the effective temperature Teff, the surface gravity
g, the helium number density defined as YHe ≡ N(He)/N(H),
the microturbulent velocity vturb, the projected rotational veloc-
ity vr sin i, the mass loss rate Ṁ and the exponent of the beta-
type velocity law describing the supersonic regime of the stellar
wind. For the terminal velocity of the wind v∞, which cannot
be accurately determined from the optical spectrum, we adopt
values determined from the analysis of ultraviolet (UV) wind
lines. If no UV determination of v∞ is available a scaling re-
lation of v∞ with the escape velocity (vesc) defined at the stel-
lar surface is used throughout the fitting process. For our pro-
gramme stars we adopt the ratio as determined by Lamers et al.
(1995) of v∞/vesc = 2.6. Similar as in Paper II we adopt fixed
values for the atmospheric abundances of the background met-
als. These were scaled with respect to mass ratios based on
the Solar abundances of Grevesse & Sauval (1998, and ref-
erences therein). The metallicity scaling factor was set equal
to the mean metal deficiency of 0.5 as found for the LMC
(Russell & Bessell 1989; Rolleston et al. 2002).
For our current data set the spectral range and quality, with
exception of the signal-to-noise ratio, is similar to the set anal-
ysed in Paper II. Consequently, we adopt the same minimum
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 5
number of generations that have to be calculated to determine
the best fit, i.e. to assure that the global optimum in parameter
space is found. In Paper II, using formal tests that accounted for
nebular contamination of the line profiles and a signal-to-noise
ratio of 50, this number was determined to be 150.
The uncertainties of the fit parameters are determined us-
ing so-called optimum width based error estimates. In Paper I
we have argued that an error estimate for a parameter can be
defined as the maximum variation of this parameter within the
global optimum in parameter space. By measuring the width
of this optimum in terms of fit quality one can determine which
models, based on their fit quality, are associated with the global
optimum. The maximum variation of the individual fit param-
eters within this group of models then gives the error estimates
(see Sect. 4 in Paper I).
In Tab. 3 the optimum width based error estimates of the fit
parameters for our programme stars are listed. The uncertain-
ties in R⋆, L⋆, Ms and Mev, that are derived from the fit param-
eters, were calculated using the same approach as in Paper II.
The adopted uncertainty in the visual magnitude is ∆MV = 0.13
(Panagia et al. 1991) for all our programme stars.
3.1. The assumption of spherical symmetry
fastwind assumes a spherically symmetric star and wind. For
very high rotational speeds, vrot, this may potentially be a prob-
lematic assumption. As it is well known, a high rotation rate
leads to a distortion of the stellar surface and, via the von Zeipel
theorem, to a decrease in flux and effective temperature from
the pole to the equator (von Zeipel 1924, (slightly) modified
by Maeder (1999) for the relevant case of shellular, i.e., radi-
ally dependent, differential rotation, ω = ω(r)). This so-called
“gravity” darkening does not only affect the stellar parameters,
but also the wind (e.g. Cranmer & Owocki 1995; Owocki et al.
1996; Petrenz & Puls 2000). Thus, it might be questioned in
how far the derived properties (which then depend on incli-
nation angle and have a somewhat local character) are repre-
sentative for the global quantities (e.g., mass, luminosity and
mass-loss rate) which refer to integral quantities.
As shown by various simulations (Cranmer & Owocki
1995; Petrenz & Puls 1996, 2000), the difference between local
and global quantities remains small unless the star rotates faster
than >∼ 60. . .70 % of its critical speed. Unfortunately, however,
we cannot directly access the actual rotational speed, vrot, but
only its projected value, vr sin i.
Accounting for the average value of < sin i >= π/4, we
have calculated the ratio of average to critical rotational speed,
Ω ≈ vr sin i/(< sin i > vcrit), and found that only two objects
lie above this value, namely N11-033 and N11-051, both with
Ω ≈ 0.7. All other objects lie well below Ω = 0.3 (except for
N11-087 withΩ ≈ 0.4). Of course, we cannot exclude that also
the latter objects lie above the decisive threshold (if observed
pole-on), but the rather large sample size implies that such a
possibility should be actually present only for a minority of ob-
jects. In conclusion, we suggest that the majority of our objects
is, if at all, only weakly affected by a distortion of surface and
wind. Note, e.g., that Ω = 0.3 leads to a difference in Teff and
R⋆ between pole and equator of less than 5%, i.e., of the same
order or less than our error estimates (cf. Tab. 3). Regarding the
derived gravities and masses, finally, we have applied a consis-
tent centrifugal correction anyhow (cf. Sect. 4.3).
For the two objects with large rotational speeds, on the
other hand, it is rather possible that we observe them al-
most equator-on, i.e., the observed profiles contain an intrin-
sic averaging over the complete stellar disk and thus corre-
spond, at least in part, to the global, polar angle averaged val-
ues. Howarth & Smith (2001) analyzed two galactic fast rota-
tors, HD 149757 (ζ Oph) and HD 191423, accounting for ef-
fects of non-sphericity and von Zeipel’s theorem. Their study
showed that for these objects, rotating at Ω = 0.9, differ-
ences in effective temperature and radius between pole and
equator can amount to 20–30 %. Interestingly, analysis of the
same objects by Herrero et al. (2002), Villamariz et al. (2002),
and Villamariz & Herrero (2005) using both hydrostatic and
spherically symmetric expanding atmospheres yielded (aver-
age) parameters that agreed to within the standard deviation
(5–10%) with the results obtained by Howarth & Smith. For
stellar parameters it thus appears justified to conclude that non-
sphericity has a very small impact even on stars with extreme
vr sin i. But note also that with respect to global mass-loss rates
the situation is much more unsecure and more-D simulations
would be required to constrain their actual values. As shown,
e.g., by Petrenz & Puls (2000), the modified wind-momentum
rate from a rapidly rotating (Ω = 0.85) B-supergiant might be
underestimated up to one magnitude if seen equator on. As the
wind properties of N11-031 and N11-051 are rather normal it
seems however unlikely that in these cases rotational effects
have this type of dramatic impact.
4. Fundamental parameters
4.1. Effective temperatures
In Fig. 1 the distribution of the effective temperatures of our
programme stars as a function of spectral type is shown. The
different luminosity classes are denoted using circles, trian-
gles and squares, respectively, for class V, III and I-II objects.
Similar as in Paper II we see that for a given spectral type
the dwarfs are systematically hotter than the giants and su-
pergiants. This separation can be interpreted as the result of
the reduced surface gravities of the more evolved objects. A
lower surface gravity results in an increased helium ionisation
(e.g. Kudritzki et al. 1983), reducing the Teff needed for a given
spectral type (e.g. Mokiem et al. 2004). A second reason is that
the more evolved objects have stronger winds. These denser
winds induce an increased line blanketing effect, further reduc-
ing the required temperature (e.g. Schaerer & de Koter 1997).
Massey et al. (2005), who also analysed a sizeable sample of
LMC stars, do not find evidence for the dwarfs being hotter
than the giants. Though, we note that their analysis only con-
tained two giant Teff determinations for spectral type later than
For comparison we also show in Fig. 1 as a dashed line
the observed Teff vs. spectral calibration for Galactic O-type
dwarfs from Martins et al. (2005a). With a dotted line the av-
6 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
Table 2. Fundamental parameters of the LMC sample determined using GA optimised spectral fits, with Teff in kK, log g and
log gc in cm s
−2, R⋆ in R⊙, L⋆ in L⊙, vturb and vr sin i in km s
−1, Ṁ in M⊙yr
−1and Ms and Mev in M⊙. Results were obtained
using a population of 72 fastwind models evolved over a minimum of 150 generations. Gravities corrected for centrifugal
acceleration (log gc) were used to calculate the spectroscopic masses (Ms). Evolutionary masses (Mev) were derived from the
tracks of Schaerer et al. (1993).
ID ST Teff log g log gc R⋆ log L⋆ YHe vturb vr sin i Ṁ β Ms Mev
N11-004 OC9.7 Ib 31.6 3.36 3.37 26.5 5.80 0.10 5.7 81 1.78·10−6 1.18 59.9 47.5
N11-008 B0.7 Ia 26.0 2.98 2.99 29.6 5.55 0.10 18.0 83 4.96·10−7 1.87 31.4 31.7
N11-026 O2 III(f*) 53.3 4.00 4.00 10.7 5.92 0.11 19.0 109 1.81·10−6 1.08 42.4 81.9
N11-029 O9.7 Ib 29.4 3.23 3.24 15.7 5.21 0.07 19.1 77 1.73·10−7 1.63 15.5 24.9
N11-031 ON2 III(f*) 45.0 3.85 3.86 13.7 5.84 0.10 20.0 116 3.88·10−6 0.89 49.3 60.8
N11-032 O7 II(f) 35.2 3.45 3.46 14.0 5.43 0.09 11.3 96 8.06·10−7 1.03 20.6 33.8
N11-033 B0 IIIn 27.2 3.21 3.35 15.6 5.07 0.08 17.9 256 2.44·10−7 1.03 19.8 21.0
N11-036 B0.5 Ib 26.3 3.31 3.32 15.6 5.02 0.08 13.6 54 1.06·10−7 0.80a) 18.4 20.0
N11-038 O5 II(f+) 41.0 3.72 3.74 14.0 5.69 0.10 9.2 145 1.52·10−6 0.98 38.8 48.3
N11-042 B0 III 30.2 3.69 3.70 11.8 5.01 0.10 4.0 42 1.89·10−7 1.19 25.1 21.4
N11-045 O9 III 32.3 3.32 3.35 12.0 5.15 0.07 16.8 105 5.48·10−7 0.80a) 11.8 24.6
N11-048 O6.5 V((f) 40.7 4.19 4.20 9.9 5.38 0.06 1.0 130 1.67·10−7 0.80a) 56.6 36.6
N11-051 O5 Vn((f)) 42.4 3.75 3.88 8.4 5.31 0.08 19.7 333 1.01·10−6 0.60 19.5 36.4
N11-058 O5.5 V((f) 41.3 3.89 3.90 8.4 5.27 0.10 14.5 85 1.52·10−7 1.42 20.3 34.4
N11-060 O3 V((f*)) 45.7 3.92 3.93 9.7 5.57 0.12 19.3 106 5.22·10−7 1.26 29.2 49.4
N11-061 O9 V 33.6 3.51 3.52 11.7 5.20 0.09 19.8 87 2.14·10−7 1.80 16.7 26.6
N11-065 O6.5 V((f) 41.7 3.89 3.90 7.4 5.17 0.17 9.4 83 3.63·10−7 0.80a) 15.8 32.9
N11-066 O7 V((f)) 39.3 3.87 3.88 7.7 5.10 0.11 4.8 71 4.08·10−7 0.80a) 16.2 29.5
N11-068 O7 V((f)) 39.9 4.13 4.13 7.1 5.06 0.10 15.6 54 3.43·10−7 1.12 25.2 29.2
N11-072 B0.2 III 30.8 3.78 3.78 7.9 4.70 0.12 7.6 14 2.35·10−7 0.84 13.8 17.7
N11-087 O9.5 Vn 32.7 4.04 4.09 8.9 4.91 0.10 15.5 276 1.38·10−7 0.80a) 35.6 20.9
N11-123 O9.5 V 34.8 4.22 4.23 5.4 4.58 0.09 9.0 110 7.62·10−8 0.80a) 17.8 18.8
BI 237 O2 V((f*)) 53.2 4.11 4.11 9.7 5.83 0.10 12.8 126 7.81·10−7 1.26 44.6 75.0
BI 253 O2 V((f*)) 53.8 4.18 4.19 10.7 5.93 0.09 18.6 191 1.92·10−6 1.21 64.6 84.1
Sk −66 18 O6 V((f)) 40.2 3.76 3.76 12.2 5.55 0.14 10.8 82 1.07·10−6 0.94 31.5 40.7
Sk −66 100 O6 II(f) 39.0 3.70 3.71 13.6 5.58 0.19 8.7 84 8.81·10−7 1.27 34.7 41.4
Sk −67 166 O4 Iaf+ 40.3 3.65 3.66 21.3 6.03 0.28 20.0 97 9.28·10−6 0.94 75.0 70.4
Sk −70 69 O5 V 43.2 3.87 3.88 9.0 5.41 0.17 16.1 131 1.03·10−6 0.78 22.7 39.7
a) assumed fixed value
erage temperature of the SMC dwarfs studied in Paper II is
shown. The LMC dwarfs are found to occupy the region in
between these two average temperature scales, with an average
behaviour intermediate to that of the Galactic and SMC dwarfs.
We interpret this as the result of the metallicity of the LMC that
is lower than the Galactic value and higher than in the SMC. As
the amount of line blanketing in a stellar atmosphere depends
on metallicity the LMC objects have temperatures in between
that of objects in the other two galaxies.
4.2. The Teff scale of O2 stars
Our sample contains four O2 type stars. This spectral type was
introduced by Walborn et al. (2002b) and is assigned based pri-
marily on the ratios of selective emission lines of N iv and N iii.
By modelling these lines Walborn et al. (2004) have shown that
indeed, as had been hypothesised, the O2 spectra correspond to
higher effective temperatures. However, the correct treatment
of nitrogen lines in stellar atmosphere models is notoriously
difficult. The relevant ionisation stages of this atom represent
much more complex ion models compared to the relatively
simple hydrogen and helium ions. Moreover, for the higher ions
of nitrogen the ionisation depends on the extreme-UV radiation
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 7
Table 3. Optimum width based error estimates for the seven fit parameters. The ND entries correspond to error in vturb that reach
up to the maximum allowed value of vturb and, therefore, are formally not defined. See text for details on the calculation of the
uncertainties in the derived parameters. Units: Teff in kK, log g and log gc in cm s
−2, R⋆ in R⊙, L⋆ in L⊙, vturb and vr sin i in
km s−1, Ṁ in M⊙yr
−1and Ms and Mev in M⊙.
ID ∆Teff ∆log gc ∆R⋆ ∆ log L⋆ ∆YHe ∆vturb ∆vr sin i log∆Ṁ ∆β ∆Ms ∆Mev
N11-004 −0.5
−0.06
+0.06 ±1.7 ±0.06
−0.01
+0.02
−0.14
+0.05
−0.00
+0.19
N11-008 −0.6
−0.06
+0.05 ±1.9 ±0.07
−0.01
+0.04
−0.20
+0.16
−0.31
+1.08
N11-026 −3.9
−0.05
+0.05 ±0.8 ±0.14
−0.01
+0.02
−0.05
+0.10
−0.12
+0.03
N11-029 −0.6
−0.05
+0.06 ±1.0 ±0.07
−0.01
+0.02
−0.13
+0.16
−0.42
+0.31
N11-031 −1.6
−0.06
+0.07 ±0.9 ±0.10
−0.02
+0.02
−0.05
+0.06
−0.04
+0.06
N11-032 −0.7
−0.07
+0.06 ±0.9 ±0.06
−0.02
+0.02
−0.16
+0.09
−0.11
+0.28
N11-033 −0.9
−0.06
+0.09 ±1.0 ±0.08
−0.01
+0.03
−0.26
+0.10
−0.20
+0.74
N11-036 −0.3
−0.09
+0.06 ±1.0 ±0.09
−0.01
+0.02
−1.99
+0.30 −
N11-038 −1.0
−0.06
+0.07 ±0.9 ±0.07
−0.01
+0.03
−0.16
+0.10
−0.13
+0.22
N11-042 −0.8
−0.07
+0.09 ±0.7 ±0.07
−0.01
+0.03
−0.47
+0.40
−0.47
+0.77
N11-045 −1.2
−0.09
+0.21 ±0.8 ±0.08
−0.02
+0.03
−0.20
+0.14 −
N11-048 −2.0
−0.21
+0.06 ±0.7 ±0.10
−0.01
+0.02
+13.2
−1.48
+0.74 −
N11-051 −1.9
−0.06
+0.06 ±0.6 ±0.09
−0.02
+0.03
−0.28
+0.11
−0.09
+0.52
N11-058 −0.5
−0.06
+0.05 ±0.5 ±0.06
−0.01
+0.01
−0.36
+0.08
−0.20
+0.60
N11-060 −1.0
−0.05
+0.09 ±0.7 ±0.10
−0.02
+0.03
−13.1
−0.22
+0.06
−0.21
+0.20
N11-061 −0.6
−0.09
+0.07 ±0.7 ±0.06
−0.01
+0.03
−0.21
+0.14
−0.49
+0.38
N11-065 −0.9
−0.07
+0.15 ±0.5 ±0.06
−0.03
+0.04
−1.06
+0.14 −
N11-066 −1.3
−0.15
+0.19 ±0.5 ±0.09
−0.03
+0.05
−0.49
+0.22 −
N11-068 −1.0
−0.19
+0.07 ±0.4 ±0.07
−0.02
+0.03
−14.8
−0.79
+0.16
−0.24
+0.36
N11-072 −0.6
−0.07
+0.09 ±0.5 ±0.06
−0.02
+0.02
−1.29
+0.24
−0.32
+0.31
N11-087 −0.6
−0.09
+0.11 ±0.6 ±0.07
−0.03
+0.02
−2.82
+0.27 −
N11-123 −0.7
−0.11
+0.00 ±0.3 ±0.06
−0.01
+0.02
−2.03
+0.35 −
BI 237 −3.8
−0.08
+0.15 ±0.8 ±0.18
−0.01
+0.03
−12.6
−0.16
+0.12
−0.17
+0.16
BI 253 −5.5
−0.15
+0.14 ±0.9 ±0.19
−0.02
+0.03
−18.0
−0.05
+0.10
−0.21
+0.07
Sk −66 18 −1.1
−0.13
+0.09 ±0.8 ±0.07
−0.03
+0.04
−10.4
−0.05
+0.06
−0.10
+0.11
Sk −66 100 −1.0
−0.09
+0.08 ±0.9 ±0.07
−0.03
+0.05
−0.12
+0.12
−0.19
+0.18
Sk −67 166 −0.8
−0.08
+0.00 ±1.3 ±0.07
−0.04
+0.06
−0.05
+0.02
−0.04
+0.09
Sk −70 69 −1.4
−0.14
+0.13 ±0.6 ±0.08
−0.03
+0.05
−13.1
−0.11
+0.14
−0.25
+0.15
field, which may be affected by non-thermal processes, such
as shocks. Adding to the complexity is the fact that the nitro-
gen abundance has to be treated as a free parameter, as many
early type stars show evidence of atmospheric abundance en-
hancements (e.g. Crowther et al. 2002; Bouret et al. 2003). In
our analysis method we solely model the hydrogen and helium
lines and self consistently allow for abundance enhancements
by treating the helium abundance as a free parameter. In prin-
ciple our temperature determination should, therefore, not be
affected by such problems. As a result of this, our analysis can
provide an independent confirmation for the hot nature of O2
stars.
In Fig. 1 we see that the objects with an O2 spectral type
indeed correspond to the hottest stars in our sample. With ex-
ception of the giant N11-031, which based on its helium lines
we find to be cooler (see below), they have temperatures in ex-
cess of 50 kK and, therefore, are significantly hotter than the
O3 star at Teff = 46 kK. The error bars, though, are consider-
able. Compared to the average error of ∼3 percent the O2 ef-
fective temperatures have an uncertainty of 5 up to 11 percent.
This large uncertainty can be explained by the weak or even
absent neutral helium lines in the spectra of these objects. As a
result of this, our fitting method has to predominately rely on
the He ii line profiles to determine the correct helium ionisation
equilibrium. Based on a single ionisation stage the determina-
tion of this equilibrium is more uncertain, which explains the
larger error bars. One should also be wary for a possible degen-
eracy effect between Teff and YHe that can occur as a result of
8 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
O2 O3 O4 O5 O6 O7 O8 O9 B0 B1
spectral type 
Fig. 1. Effective temperatures as a function of spectral type for
the investigated LMC sample. Different luminosity classes are
denoted using circles, triangles and squares for class V, III and
I-II objects, respectively. Shown as a dashed line is the Teff cal-
ibration of Martins et al. (2005a) for Galactic O-type dwarfs.
The dotted line corresponds to the average effective tempera-
ture of the SMC dwarfs studied in Paper II. The LMC dwarfs
are found to lie in between these two average scales, with a
typical ∼2 kK offset from both the average Galactic and SMC
relations.
the missing neutral helium lines (see e.g. Paper I). This, how-
ever, is not an issue as the helium abundances derived for all
O2 stars correspond to normal values close to 0.10.
It is also possible to question whether the fact that the He i
lines are so weak could result in a systematically higher Teff
determination by the automated method compared to “by eye”
fits. In other words, would a “by eye” fit prefer a lower tem-
perature solution? For the O2 giant N11-026 this is a relevant
question, as the weak He i λ4471 line, shown in Fig. A.1 in the
appendix, is not fitted perfectly. Note, however, that due to the
relative plot scale the discrepancy is exaggerated and that the
fit quality is good and is comparible to that of the He ii lines.
Still, to assess whether the fit of this particular line could be
improved, i.e. to search for a solution predominantly using the
He i diagnostic, we ran test fits with increasing relative weight
of He i λ4471. We found that a solution with an improved fit
could be obtained for an increase of the relative line weight
with a factor of five. This solution has a Teff lower by 3.7 kK
and all other fit parameters approximately equal to the solution
with the low He i λ4471 weight. The overal fit quality of the
other lines has somewhat decreased. As these lines are rela-
tively strong, this is difficult to discern by eye. Also note that
within the lower error estimate for Teff of 3.9 kK the two solu-
tions agree. Similar results were also obtained for the O2 stars
BI 237 and BI 253, where a reduction of Teff by 3.8 kK and
5.4 kK was found, respectively, for an increase in the relative
weight of He i λ4471 by a factor of six and two, respectively.
Consequently, based on our analysis we can only give a range
for the Teff scale of the O2 stars of 49 − 54 kK, the bounds
being set by the temperatures based on the He i λ4471 and hy-
drogen/He ii diagnostics.
The upper end of our hydrogen/helium based O2 effec-
tive temperature range is in good agreement compared to
the average effective temperature of 54 kK as determined by
Walborn et al. (2004). As we already mentioned the giant N11-
031 with a Teff lower by approximately 8 kK compared to the
other O2 stars forms an exception to this. Its effective tempera-
ture compares better to the temperature of the O3 star N11-060.
A close inspection of the spectra of the two stars reveals why
this is so. Both stars have exactly the same equivalent width
ratio of the He i λ4471 and He ii λ4541 lines, which results in
similar values for Teff . To also test whether the relatively strong
He i λ4471 line could be dominating the fit, forcing a relatively
low Teff, we refitted the spectrum of N11-031 ignoring the neu-
tral helium lines. This again resulted in an effective temperature
of 45 kK.
Interestingly, the nitrogen line analysis of Walborn et al.
for N11-031 did result in a higher effective temperature of
55 kK. We are not sure whether this discrepancy is the re-
sult of a systematic offset between the N v λ4603-20/N iv and
He i λ4471/He ii temperature scales. A more recent compari-
son by one of us (P.A.C.) of the spectral fit of Walborn et al. to
new VLT data, however, showed relatively large discrepancies
in the helium line fits compared to the nitrogen line fits, indi-
cating that this could be the case. The discrepancy could also
be due to differences in fitting assumptions and approaches.
Walborn et al. adopted a surface gravity of log g = 4.0 and
estimated a mass loss rate of Ṁ = 1.0 × 10−6 M⊙yr
−1 from
the He ii λ4686 line. Our fit indicates that the former param-
eter should be lower by 0.15 dex, consequently lowering Teff.
With respect to the mass loss rate we find a value higher by
a factor four. An increase in Ṁ of this magnitude can have a
significant effect on the strength of different nitrogen lines (e.g.
Crowther et al. 2002) and, therefore, on the derived effective
temperature.
Massey et al. (2005) analysed a total of 11 O2 stars. They
could determine the effective temperature for three dwarfs and
one giant, with 47.0 kK . Teff . 54.5 kK. For the supergiants
only lower limits of Teff & 42 kK and higher could be derived.
These results are in agreement with our findings. However,
these authors note that the correlation with Teff for the O2–3.5
spectral types is not tight. In particular no good agreement was
found between the ratios of the N iii and N iv emission lines
and the He i and He ii lines. Consequently, a more thorough
investigation of the O2 stars based on both the nitrogen and
helium spectrum is necessary to resolve this degenerate class
adequately.
4.3. Gravities
The distribution of our programme stars in the log Teff – log gc
plane is presented in Fig. 2. To calculate the surface gravity
corrected for centrifugal acceleration (log gc) the method dis-
cussed by Herrero et al. (1992) and Repolust et al. (2004) was
adopted. Different luminosity classes are denoted using circles,
triangles and squares for, respectively, class V, III and I-II ob-
jects. In this figure we see that the dwarfs, with exception of
two objects, form a group clearly separated from the latter two
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 9
4.40 4.50 4.60 4.70 4.80
log Teff
Fig. 2. The log Teff – log gc plane for the analysed LMC objects.
Different luminosity classes are denoted using circles, triangles
and squares for dwarfs, giants, and bright giants and super-
giants, respectively. With exception of two objects the dwarfs
are clearly separated from the luminosity class I-II objects. The
giants seem to overlap with both the luminosity class V and I-II
objects. Shown as a dotted line is the average log Teff – log gc
relation of the SMC O-type stars of luminosity class I-II-III
from Paper II.
groups. In contrast to our findings in Paper II we do not find
a clear separation between the dwarfs and giants. Instead the
latter group of objects shows an overlap with both luminosity
class V and I-II objects.
Shown in Fig. 2 as a dotted line is the average log Teff –
log gc relation of the SMC O-type stars of luminosity class I-
II-III from Paper II. The majority of the evolved LMC objects
seem to agree with this trend, that illustrates the evolutionary
correlation between effective temperature and surface gravity.
However, note that some evolved objects are at considerable
distance from the average relation. Thus, a calibration of the
two parameters for a given luminosity class should be taken
with care (see also Repolust et al. 2004).
The comparison of spectroscopically determined masses
(Ms) with masses obtained from evolutionary calculations is
presented in Fig. 3. Using the same symbols as in Fig. 2 dwarfs,
giant and supergiants are distinguished. Evolutionary masses
(Mev) were derived from the evolutionary tracks calculated for
a metallicity of Z = 0.4 Z⊙ from Schaerer et al. (1993). The er-
rors on these masses correspond to the maximum mass interval
in the error box spanned by the uncertainties in luminosity and
effective temperature.
The tracks from Schaerer et al. do not include the effects
of rotation. Consequently, this additional source of error is not
included. Calculations including vrot show that in some cases
one can no longer assign an unambiguous M(L, Teff). This is
a result of rotationally enhanced mixing or the unknown incli-
nation angle, if the star has a non-spherically symmetric dis-
tribution of R⋆ and Teff , causing complicated tracks including
loops during the secular redward evolution (Meynet & Maeder
2005). Assessing the impact of vrot on the Mev determination we
showed in Paper II from a comparison of masses derived from
 0  20  40  60  80  100  120
Spectroscopic Mass [Msun]
Fig. 3. Comparison of spectroscopic masses with masses de-
rived from the evolutionary tracks of Schaerer et al. (1993).
The one-to-one correlation between the two mass scales is
given by the dashed line. Symbols have the same meaning as
in Fig. 2. The objects with the highest evolutionary masses ex-
hibiting a mass discrepancy correspond to the four O2 stars in
our sample.
non rotating tracks to those obtained from tracks calculated for
vrot = 300 km s
−1 that the error in Mev will not increase by more
than approximately ten percent.
In Fig. 3 we see that the majority of the objects are lo-
cated left of the one-to-one correlation, given by the dashed
line. The error bars of twelve objects do not even touch
this correlation. Consequently, we find a significant mass dis-
crepancy. Similar mass discrepancy problems have been dis-
cussed by e.g. Herrero et al. (1992), de Koter & Vink (2003),
and Repolust et al. (2004). Most of these classical problems
were attributed to limitations in the stellar atmosphere mod-
els (Herrero 1993) and to potential biases in the fitting process
(see Paper I). Here we cannot explain the found discrepancy in
such a manner. We will provide a more thorough investigation
and discussion in Sect. 5.
4.4. Helium abundances
For the SMC sample that we analysed in Paper II we found a
correlation between the helium surface abundance and surface
gravity. This could partly be explained by evolutionary effects.
As the surface gravity decreases when a star evolves away from
the ZAMS, objects with lower gravities would correspond to
more evolved objects and are more likely to have atmospheres
enriched with helium.
To investigate whether the scenario discussed above also
applies to the current LMC sample, we plot the helium abun-
dance as determined with the automated method as a function
of log gc in Fig. 4. Also shown as a dashed line is a measure
of the initial helium abundance. This value of YHe = 0.09 was
calculated by averaging the surface helium abundances of the
dwarf type objects with a helium abundance smaller than the
total sample average. Compared to this measure for the ini-
tial helium abundance a correlation between the surface gravity
10 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
2.80 3.20 3.60 4.00 4.40
log gc
Fig. 4. Helium abundances as a function of surface gravity.
Symbols have the same meaning as in Fig. 2. The dashed line
at YHe = 0.09 is a measure for the initial helium abundance
and corresponds to the mean of the helium abundances of the
dwarfs with YHe smaller than the sample average. Starting at
the highest gravities an increase of the helium fraction is seen
down to log gc ≈ 3.6.
and helium enrichment can be observed. Starting at the high-
est gravities, we find an increase in the average helium abun-
dance towards lower log gc. Note that the two objects exhibit-
ing the largest helium fractions are supergiants. The increase
can again be partially explained as an evolutionary effect. In
Paper II a similar correlation between average helium abun-
dance and gravity was found to exist down to the lowest gravi-
ties investigated.
Interestingly, in Fig. 4 we see that for our LMC sample no
helium enrichment is found at log g . 3.6. Not even the super-
giants show evidence of enrichment below this gravity. Why is
this so? In Fig. 5, where we present the HR-diagram for our
sample, the answer to this question is given. Shown as a grey
area in this figure is the region in which the rotating evolu-
tionary models of Meynet & Maeder (2005) predict a helium
surface enrichment of at least ten percent. The majority of the
evolved objects are located outside this region. Consequently,
based on their specific evolutionary phase these objects are not
expected to show any enrichment. Also important is the fact
that we have selected the hottest objects from the N11 cluster.
As a result of this our sample is biased and does not contain
the low gravity objects with a high luminosity, i.e. those more
likely to be enriched, as these have evolved into cool B-type
stars.
Apart from the supergiants also four dwarfs are found to be
enriched. Figure 5, in which we have highlighted stars with a
helium abundance of at least 0.12 using open symbols, shows
that one these dwarfs, Sk −66 18, is located relatively close to
the region in which enrichment is predicted. Consequently, ro-
tationally enhanced mixing is a possible explanation. The three
remaining dwarfs, N11-60, N11-065 and Sk −70 69, in contrast
lie relatively close to the ZAMS. Therefore, “normal” mixing
cannot explain their enrichment. Instead, a possible explana-
tion is given by chemically homogeneous evolution. As this
4.24.34.44.54.64.74.8
log Teff
Fig. 5. Hertzsprung-Russell diagram for the LMC sample.
Symbols have the same meaning as in Fig. 2. Over plotted as
grey lines are the evolutionary tracks of Schaerer et al. (1993)
for Z = 0.4 Z⊙, with a black line representing the ZAMS. Open
symbols indicate objects with YHe ≥ 0.12. The grey area corre-
spond to the region in which the rotating evolutionary models
of Meynet & Maeder (2005) predict a relative helium enhance-
ment of at least 0.01.
may also be linked to the mass discrepancy, we will return to it
in Sect. 5.
4.5. Microturbulence
The microturbulent velocities determined with the automated
fitting method are also given in Tab. 2. Though the error es-
timates are relatively large (see Tab. 3), we find that for the
current data set the vturb determinations were sufficiently accu-
rate to reveal a weak correlation between this parameter and
the surface gravity. This is shown in Fig. 6, where it can be
seen that for log g . 3.6 the average microturbulence recovered
from the line profiles increases systematically. The situation for
log g & 3.6 is less clear, as the error bars are on average larger
and the values for vturb are more or less randomly distributed
between 0 and 20 km s−1. The reason for this is that for larger
values of the surface gravity the line profiles become intrinsi-
cally broader due to the increased Stark broadening, making it
more difficult to accurately recover vturb from the line profiles
only.
Based on samples predominantly consisting of unevolved
early B-type Galactic stars other authors have also found
a relation between microturbulence and surface gravity, e.g.
Kilian et al. (1991), Gies & Lambert (1992) and Daflon et al.
(2004). More recently Hunter et al. (2006) analysed a sample
of early B-type stars in the Magellanic clouds and also found
a trend of increasing vturb for decreasing log g. To derive the
values for vturb all these authors relied on curve-of-growth tech-
niques, which were applied to metal lines such as those of Si iii
and O iii calculated using plane parallel models. Consequently,
our line profile based analysis is an independent confirmation
of the existence (or requirement by lack of a physical explana-
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 11
2.8 3.2 3.6 4.0 4.4
log g [cm/s2]
Fig. 6. Microturbulent velocities determined using line profile
fits as a function of surface gravity. Symbols have the same
meaning as in Fig. 2. For log g . 3.6 a trend of increasing
vturb with decreasing log g is visible. For larger gravities the
uncertainties in the vturb determinations are too large to discern
any possible relation.
tion or failures in the line broadening mechanisms) of micro-
turbulence in the atmospheres of these type of stars.
The physical mechanism explaining the observed microtur-
bulence or its relation to the surface gravity is still poorly un-
derstood. Kudritzki (1992) and Lamers & Achmad (1994) have
argued that the observed microturbulence might be the result of
a stellar outflow, implying the existence of a velocity gradient
in the photospheric layers, which can mimic microturbulence-
like desaturation effects. Smith & Howarth (1998), however,
showed by applying a simple core-halo model to the Galactic
O9.7 supergiant HD 152003 that this effect would not be suffi-
cient to explain the observed vturb. Indeed, our analysis employ-
ing a unified photosphere and wind model confirms that the mi-
croturbulence cannot be explained as an artifact of a transonic
velocity field.
A possible explanation for the gravity dependence could
be related to instabilities in the wind. These instabilities are
reflected in the large turbulent velocities (∼100–200 km s−1)
necessary to fit the wind lines in the intermediate and outer
wind (e.g. Groenewegen & Lamers 1989; Haser et al. 1998;
Evans et al. 2004b) and could be related to shocks due to the in-
trinsic line-driven instability (Lucy 1983; Owocki et al. 1988;
Owocki & Puls 1999). For a low surface gravity the wind starts
at larger Rosseland optical depth compared to the high grav-
ity case. Therefore, the line forming region of low gravity ob-
jects contains a relatively large contribution originating from
(the base of) the wind. Consequently, this region could be af-
fected by the onset of the line-driven instability introducing
wind-turbulence into the line profiles.
Tentative support for the above described scenario may be
implied by Fig. 7, where we show vturb as function of the line
forming region of He i λ4471 in units of the stellar radius. The
location of the line forming region is defined as the position
at which the radial optical depth in the line core reaches a
value of 2/3. This figure shows that when the radial distance
1.00 1.02 1.04 1.06 1.08 1.10
R(τline=2/3) [Rstar]
Fig. 7. Microturbulent velocity as a function of the location of
the line forming region of He i λ4471, which is defined as the
location where the radial optical depth in the line core reaches
a value of τ = 2/3. A weak trend is visible that suggests that
for increasing extension of the atmosphere larger values of vturb
are necessary to fit the line profiles.
to this position increases also the average microturbulent ve-
locity increases. Though the trend is weak, it does seem to
indicate that when the atmosphere becomes more extended,
higher values of vturb are necessary to reproduce the line pro-
files. It may appear that this implies that the line forming re-
gion enters in to the regime where wind turbulence develops.
However, for this statement no compelling evidence is avail-
able, as we do not find any correlation between vturb and the
distance between the line forming region and (for instance) the
sonic point. Moreover, for most objects at R(τline = 2/3) > 1.04
values of vturb close to the maximum allowed value are found.
Consequently, these values should be interpreted as lower lim-
its. Therefore, we can only conclude that our analysis points
towards a gradient in the turbulent velocity, possibly connected
to a link between microturbulence and wind instabilities, and
suggest further investigation in this direction.
4.6. Wind parameters
The wind parameters and their uncertainties determined using
the automated method are listed in Tabs. 2 and 3. Compared
to our SMC analysis we find that we were able to accurately
determine these parameters for a significantly larger number
of objects. This is mainly the result of an on average higher
signal-to-noise ratio of the spectra as well as of denser winds
for the LMC objects compared to their SMC counterparts. In
total we determined 22 mass loss rates and 6 upper limits. The
upper limits are defined (and can be identified in Tab. 3) by an
error bar − log Ṁ > 1.0 dex.
To provide a meaningful comparison of the mass loss rates
we place the LMC objects in the modified wind momentum
luminosity diagram. This diagram shows as function of stellar
luminosity the distribution of the so-called modified wind mo-
mentum, which is defined as Dmom ≡ Ṁv∞R
⋆ . Not only does
this allow for an assessment of the behaviour of Ṁ within our
12 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
4.5 5.0 5.5 6.0 6.5
log(L/Lsun)
Fig. 8. Modified wind momentum (Dmom) in units of
g cm s−2 R⊙
1/2 vs. luminosity. Symbols have the same mean-
ing as in Fig. 2. Upper limits are shown as inverted triangles.
Dashed lines correspond to the predicted wind-momentum lu-
minosity relations (WLR) from Vink et al. (2000, 2001). The
upper, middle and lower relation, respectively, correspond to
predictions for Galactic, LMC and SMC metallicity. The ob-
served modified wind momenta show a strong correlation with
luminosity with an average relation that lies in between the pre-
dicted Galactic and SMC WLR. This is quantified by the em-
pirical WLR that was constructed for the LMC objects and that
is shown by the dotted line. The open square corresponds to
the wind momentum corrected for clumping of the supergiant
Sk −67 166. Shown as a dashed-dotted is the empirical WLR
obtained using this corrected Dmom.
sample, it also provides a convenient method to compare the
observed wind strengths to the predictions of line driven wind
theory. According to this theory Dmom is predicted to behave as
log Dmom = x log
L⋆/L⊙
+ log D◦ , (1)
where L⋆ is the stellar luminosity (Kudritzki et al. 1995;
Puls et al. 1996). In this equation x is the inverse of the slope of
the line-strength distribution function corrected for ionisation
effects (Puls et al. 2000). The vertical offset D◦ is a measure
for the effective number of lines contributing to the accelera-
tion of the outflow.
In Fig. 8 the distribution of the modified wind-momenta
for our programme stars are presented. Indicated using circles,
triangles and squares are objects of, respectively, luminosity
class V, III and I-II. Upper limits are shown as grey inverted
triangles. A clear correlation between L⋆ and Dmom can be ob-
served in this figure. Over an order of magnitude in L⋆ the av-
erage modified wind-momentum decreases by approximately
1.5 dex. A comparison of the behaviour of Dmom with pre-
dictions is facilitated by the theoretical WLRs calculated by
Vink et al. (2000, 2001) that are shown as a set of dashed lines.
The upper, middle and lower of these predicted power laws
were calculated for, respectively, Solar, LMC and SMC metal-
licity. Compared to these predictions we find that the LMC
wind-momenta are approximately confined between the theo-
retical WLR for Galactic and SMC metallicity. In other words,
compared to the Galactic and SMC case, stars in the LMC have
intermediate wind strengths.
To further quantify the behaviour of LMC winds relative to
that of Galactic and SMC outflows we have fitted a power law
to the observed modified wind-momentum distribution, consis-
tently accounting for both the symmetric errors in L⋆ and the
asymmetric errors in Dmom, This yielded the following empiri-
cal WLR
log Dmom = (1.81 ± 0.18) log
L⋆/L⊙
+ (18.67 ± 1.01) . (2)
The theoretical LMC relation from Vink et al. is given by
x = 1.83 and Dmom = 18.43. The error bars of the theoretical
and empirical relations are in agreement. More importantly, in
Fig. 8 the empirical relation, shown as a dotted line, is found to
lie between the predicted Galactic and SMC relations. These
predictions have been found to be in good agreement with
the observed Galactic WLR (Repolust et al. 2004, Paper I) and
observed SMC WLR (Paper II). Consequently, our empirical
LMC WLR is quantitative evidence for the fact that massive
stars in this system have mass loss rates intermediate between
those of massive stars in the Galaxy and SMC.
The differences between the empirical and theoretical LMC
WLR at the start and end of the observed luminosity range
are, respectively, 0.17 and 0.16 dex. In Fig. 8 this seems to
imply a systematic offset between the two relations. However,
we note that these differences are still smaller than the typi-
cal uncertainty in Dmom of 0.2 dex. More importantly, no cor-
rection was applied for the possibility that the winds of our
sample stars are, in contrast to our assumption, not smooth but
structured. In recent years evidence has been mounting for this
so-called clumping in the winds of O- and early B-type stars.
In particular spectroscopic modelling of UV (resonance) lines
(Crowther et al. 2002; Hillier et al. 2003; Bouret et al. 2003;
Massa et al. 2003; Martins et al. 2004, 2005b; Bouret et al.
2005; Fullerton et al. 2006) seems to suggest the existence of
clumping factors in the range of 10-100, implying correspond-
ing reductions of Ṁ by factors 3 up to 10. More recently,
Puls et al. (2006) also found clumping factors of the order 5
to 10 (normalized to the unknown clumping properties in the
outermost, radio-emitting wind) from the analysis of Hα, in-
frared, millimetre and radio fluxes. In the present study we try
to account for possible wind clumping effects by correcting the
mass loss rates of stars with Hα in emission. Markova et al.
(2004) and Repolust et al. (2004) argue that the mass loss rates
of these stars could be overestimated as a result of the fact that
Hα emission lines are formed over a relatively large volume
where clumping might have set in. In contrast, for stars with
Hα in absorption the line is formed relatively close to the stel-
lar surface, where clumping effects are negligible. Based on the
comparison of dwarfs and supergiants in their Galactic sample
Repolust et al. derived a numerical correction factor of 0.44 for
the mass loss rates of supergiants with Hα in emission.
We have applied the clumping correction to the super giant
Sk −67 166, which has a Hα emission profile. In Fig. 8 its new
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 13
wind momentum is indicated using an open symbol. Using this
value the following empirical WLR is obtained1
log Dmom = (1.43 ± 0.17) log
L⋆/L⊙
+ (20.77± 0.97) . (3)
In Fig. 8 we see that for log L⋆/L⊙ & 5.3 the new WLR com-
pares better to the Vink et al. relation. For lower luminosities
the situation is less clear. Due to the large uncertainties, this
range has a relatively low weight in the fit. Consequently, a
discrepancy for low L⋆ is less significant than the good agree-
ment obtained for the higher luminosities. For this reason it
is difficult to investigate the existence of a “weak wind prob-
lem” for stars at log L⋆/L⊙ . 5.3, first reported by Bouret et al.
(2003). This study, as well as later studies (Hillier et al. 2003;
Evans et al. 2004a; Martins et al. 2004, 2005b) report a steep-
ening in the WLR relation relative to predictions starting at
about the above mentioned L⋆, leading to an over prediction of
the wind strength by up to a factor 100 at log L⋆/L⊙ ∼ 4.5.
Our LMC results do not appear to confirm this break be-
tween observations and seem to follow the predictions down
to log L⋆/L⊙ ≈ 5.0. In a forthcoming paper we will present a
comprehensive overview of the observed WLR relations in our
Galaxy and the Magellanic Clouds as well as a thorough dis-
cussion of the successes and failures of the theory of radiation
driven winds in predicting the WLR, including possible causes
for the weak wind problem (Mokiem et al. in preparation).
5. The mass discrepancy
In Fig. 9 we investigate the mass discrepancy of our LMC tar-
gets as a function of the helium surface abundance. On the ver-
tical axis a measure for this discrepancy is shown, which is
scaled to the mean of the evolutionary and spectroscopic mass.
This ensures that positive and negative discrepancies follow
the same linear scale. For non-enriched stars, i.e. YHe . 0.10,
approximately three times as many objects lie above the one-
to-one correlation than below it. Therefore, in contrast to our
finding for the equivalent SMC case (Paper II) a significant
mass discrepancy is found for our sample of non-enriched LMC
stars. The reason for this is unclear. As we stated before our
analysis employs state-of-the-art atmosphere models, and is in
principle not hampered, as were previous studies, by poten-
tially unoptimised fits. As stated, in our SMC dataset, analysed
in an identical manner, no evidence was found for a mass dis-
crepancy for YHe < 0.10.
Massey et al. (2005) also study the Ms vs. Mev problem in
a set of LMC stars. For objects hotter than 45 kK they find
a mass discrepancy that is even stronger than what we find.
This behaviour could be the result of an underestimate of the
photospheric line pressure in this high temperature regime. In
Fig. 9 we have highlighted the objects with Teff ≥ 45 kK using
open symbols. Although their mass discrepancy is considerable
1 The small errors in the relevant parameters of Sk −67 166, which
dominates the WLR at very high luminosity, cause the significant dif-
ference between Eqs. 2 and 3. Note that our least square fitting does
not account for the correlation of both quantities (due to R⋆). In view
of the well known distance these effects are likely small compared to
the Galactic case (see Markova et al. 2004; Repolust et al. 2004).
-1.00
-0.50
0.05 0.10 0.15 0.20 0.25 0.30
Fig. 9. Mass discrepancy as a function of helium abundance
for the LMC sample. Evolutionary masses used to calculate
the discrepancy were derived from the non-rotating tracks of
Schaerer et al. (1993). Symbols have the same meaning as in
Fig. 2. The open symbols correspond to stars with Teff ≥ 45 kK.
they do not stand out as a separate group.To further illustrate
this, note that the hottest supergiant in our sample Sk −67 166
(Teff = 40 kK) at YHe = 0.28 has a Ms of 75 M⊙ which is
in very good agreement with its Mev of 70 M⊙. Consequently,
though we can not explain the mass discrepancy we do not an-
ticipate that it is connected to a flawed treatment of the photo-
spheric radiation pressure that manifests itself at metallicities
as high as that of the LMC environment (but not yet at values
typical for the SMC).
If we accept the analysis of stellar and photospheric pa-
rameters, the presence of a mass discrepancy may point to an
oversimplified picture of the evolution of massive stars used
to determine Mev and/or, possibly, to a breakdown of the as-
sumption of a spherically symmetric atmosphere. One obvious
simplification may be that we have used tracks for non-rotating
stars. One of the effects introduced by rotation is a wide bi-
furcation in the evolutionary tracks (Maeder 1987). Stars rotat-
ing faster than roughly half the surface break-up velocity will
essentially follow tracks representative of homogeneous evo-
lution. Langer (1992) showed that as a result of this the M/L-
ratio will be a monotonically decreasing function for increasing
helium enrichment. Consequently, stars evolving along homo-
geneous tracks are expected to be increasingly under-massive
for increasing helium abundance. One would therefore derive a
positive mass discrepancy if the evolution of a rapidly rotating
star was incorrectly described using non-rotating or modestly
rotating model tracks.
So, can rapid rotation be used to explain the observed mass
discrepancy? Let us first focus on the supergiants. For these
stars the mass discrepancy problem appears absent, though we
note two exceptions at YHe = 0.07 and YHe = 0.09. This ab-
sence of signs of a distinct mass discrepancy was also found
in our study of SMC bright giants and supergiants (Paper II).
The conclusion seems to be that the supergiants follow non-
rotating or modestly rotating evolutionary tracks. This implies
that, apparently, all supergiant targets that we have selected in
14 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
both the SMC (Paper II) and LMC happen to have started out
their evolution with low or modest initial rotational velocities.
As the initial rotational velocity distribution derived in Paper II
for the NGC 346 cluster implies that only some 5 − 15 percent
of stars are expected to evolve along homogeneous tracks this
should not be alarming. One supergiant, however, might not fit
this picture: Sk -67 166. This source does not feature a mass
discrepancy, but does show an increased helium abundance.
Given its position in the HRD this is quite difficult to explain in
term of evolution without rotation. A scenario accounting for a
more delicate interplay between mass loss and rotation may be
required to explain its properties (see Herrero & Lennon 2004).
Now let us turn to the dwarfs and giants. For our SMC
sample we found a correlation between the helium surface en-
richment and the mass discrepancy for class V and III stars,
with Ms being systematically smaller than Mev for YHe & 0.11.
Interestingly, this dependence was only found for dwarfs and
giants (see above). It was suggested that this behaviour was
the result of rotationally enhanced mixing, enriching the at-
mospheres with primary helium. The dwarfs and giants at
YHe > 0.11 displayed in Fig. 9 also systematically suffer from
a positive mass discrepancy. Consequently, this scenario also
seems to be a good explanation for the situation in the LMC
sample. However, we realise that these seven objects are a rel-
atively small sample and that the typical uncertainties in YHe
are 0.03. This statement should therefore be seen as a working
hypothesis.
For helium enriched dwarfs and giants the above scenario
seems a logical one, as helium enrichment early on in the evo-
lution must imply efficient mixing. For the dwarfs and giants
that do not show excess helium in their surface layers such
a clue or indication is not present (unless they evolve left of
the main sequence as do chemically homogeneous stars, see
Sect. 5.1). In principle these non YHe-enriched stars could ro-
tate sufficiently rapid to cause a mass discrepancy, but not to
the extent of chemically homogeneous evolution. However, if
so, a very significant fraction of the stars should have started
their evolution at super-critical rotation. At least for the stars
studied in the SMC cluster NGC 346 (Paper II) this is not the
case.
5.1. Chemically homogeneously evolving O2 stars
Three out of the four O2 stars are found to lie to the left of the
ZAMS in Fig. 5 and exhibit relatively large mass discrepan-
cies (see open symbols in Fig. 9). These are N11-026, BI 237
and BI 253. For a discussion of the effective temperature of the
fourth O2 object, N11-031 (Teff ∼ 45 kK), we refer to Sect. 4.1.
Even though the hottest three (Teff ∼ 53 − 54 kK) are not sig-
nificantly enriched in helium it is tempting to speculate on a
possible (near) homogeneous evolution of these objects. This
would not only provide an explanation for their mass discrep-
ancy, it would also explain their peculiar location in the HRD.
The possibility that these are true ZAMS stars is very exciting.
However, due to the short evolutionary time scales in this part
of the HR-diagram, it would also be unlikely. Based on the non-
rotating model tracks of Schaerer et al. (1993) we estimate that
a 60 M⊙ star will within 2 Myr evolve away from the ZAMS
to a location in the HRD at Teff ≈ 45 kK. This is well beyond
the error bars on the parameters of these objects. Given the fact
that N11-026 and BI 253 are thought to be associated with the
LH 10 and 30 Doradus clusters, respectively, which have ages
of approximately 3 (see Sect. 6.2) and 2 Myr (de Koter et al.
1998), a normal evolutionary scenario appears unlikely. Based
on its large radial velocity the field star BI 237 probably is a
runaway star (Massey et al. 2005), therefore it is also likely to
be relatively old.
The problems with a reconciliation of the hottest three O2
stars with fully homogeneous evolution and an age of at least
2 Myr, are i) that their surface helium abundances are not sig-
nificantly enriched, and ii) that their rotational velocities are
not extreme. The latter issue need not be “a smoking gun” con-
sidering the possibility that we may see them relatively pole-
on and the fact that their vr sin i values, ranging between 110-
190 km s−1, are above the sample average. Concerning the first
point, it appears that we must concede to the possibility that
stars may evolve along tracks similar to those for homogeneous
evolution while in fact they are not fully homogeneous – i.e.
the near surface layers are not yet strongly affected by the mix-
ing that must occur deeper in. Having said this, we do point to
the fact that the error bars on YHe do allow for relatively large
ages even within the hypothesis of fully homogeneous evolu-
tion. Within the error bars the maximum YHe is ∼ 0.13 for all
three stars. It takes fully homogeneously evolving stars of 60
and 40 M⊙ about ∼ 1.5 and ∼2.0 Myr, respectively, to build up
this amount of helium enrichment (S. -C. Yoon, private com-
munication). This is close to the derived cluster ages of LH 10
and 30 Doradus. For reference, if the star would evolve along
near-homogeneous tracks they would be older.
We tentatively conclude that the HRD position, helium
abundance, and rotational velocity of the three hottest O2 stars
in our sample could be consistent with (near-)homogeneous
evolution. Evidence for efficient rotation-induced mixing dur-
ing the main sequence phase of O stars has also been presented
by Lamers et al. (2001), based on the chemical abundance pat-
tern of ejected circumstellar nebulae.
5.2. Large sample trends
To firmly establish our findings with respect to the mass dis-
crepancy problems, we compare in Fig. 10 spectroscopic and
evolutionary masses for all stars that have been analysed using
our automated fitting method so far. Presented in this figure are
the combined Galactic (Paper I), SMC (Paper II) and the cur-
rent LMC samples, corresponding to a total of 71 O- and early
B-type stars.
The figure confirms our two main findings. First, the good
agreement between Ms and Mev for the supergiants is clearly
visible. This is an encouraging finding and we believe that this
means that the improvements in the stellar atmosphere mod-
els, evolutionary calculations and spectral analysis techniques
have finally resolved the long standing mass discrepancy as
found by Herrero et al. (1992). The correspondence between
the spectroscopic and evolutionary mass scale, in particular for
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 15
Fig. 11. The FLAMES field for the star forming region N11 with the associations LH9 (south of centre) and LH10 (north of
centre). The FLAMES targets are identified using circles and their star identification is given. From Evans et al. (2006). [Included
as separate JPG file in astro-ph submission.]
-1.00
-0.50
0.05 0.10 0.15 0.20 0.25 0.30
Fig. 10. Mass discrepancy as a function of helium abundance
for the combined samples from Paper I (Galactic), Paper II
(SMC) and the current LMC sample. Symbols have the same
meaning as in Fig. 2. Note the good correspondence between
Ms and Mev for the bright giants and supergiants (square sym-
bols).
increasing helium enrichment is striking. The reason for this is
probably related to the fact that the class I-II objects are found
in the region of the HR-diagram in which stars are not expected
to undergo extreme evolutionary phases, such as homogeneous
evolution. Consequently, the enriched bright giants and super-
giants have evolved along relatively simple evolutionary tracks,
which (in this respect) appear well understood.
Our second finding, i.e. that of a correlation of mass dis-
crepancy with the helium abundance, is also corroborated by
the large sample in Fig. 10. For YHe > 0.10 all dwarfs and gi-
ants, except for one dwarf, are found above the Ms = Mev line.
In all fairness, given the typical uncertainty of 0.03 in YHe, from
a statistical point of view the region 0.09 < YHe < 0.12 should
be regarded with care. For larger helium abundance, however,
the correlation can be regarded to be statistically significant. A
total of 11 objects show a positive mass discrepancy, with only
a single counter example. Moreover, the magnitude of the dis-
crepancy seems to be related to the amount of enrichment. This
strongly points to efficient mixing in the main sequence phase,
leading to (near-)chemically homogeneous evolution.
6. The evolutionary status of N11
In this section we explore the evolutionary status of the N11
field. We will first briefly outline the current understanding of
N11 with respect to the OB associations LH9 and LH10 and
in particular the sequential evolutionary link between the two.
Based on the analysis of the 22 stars in our sample associated
with these clusters we will then estimate their ages and discuss
whether they are compatible with a sequential star formation
scenario.
6.1. LH9 and LH10 in N11
N11 is an intricate giant H ii region containing several massive
star forming complexes. The largest of these are the OB associ-
ations LH9 and LH10 (Lucke & Hodge 1970). An image of the
region, including the adopted star identifications, is presented
in Fig. 11 (Evans et al. 2006). Several studies have suggested
that the star forming activity in these associations are linked.
In particular the study by Parker et al. (1992) has provided a
key to understanding of the structure and formation of LH9
and LH10. Their analysis of the stellar content of N11 revealed
the presence of several O3-O5 stars and possible ZAMS stars
in LH10. In contrast the earliest spectral type associated with
LH9 was found to be O6. They also found the slope of the ini-
tial mass function of LH10 to be significantly flatter than that
of LH9, indicating that the former contains a higher ratio of
high mass to low mass stars. Combined with the fact that the
reddening of LH10 is larger than that of LH9, it was concluded
that LH10 is the younger of the two clusters. Based on these
findings the authors propose an evolutionary link between the
two, where star formation in LH10 could possibly be triggered
by the stellar winds and supernovae of massive stars in LH9.
Further evidence for a sequential star formation sce-
nario was given by Walborn & Parker (1992), who found a
dual structural morphology of N11 analogous to that of the
30 Doradus H ii region. In the latter substantial evidence sug-
gests the presence of current star formation in regions sur-
rounding the central star cluster (e.g. Walborn & Blades 1987;
Hyland et al. 1992, also see Walborn & Blades 1997). This sec-
ondary burst seems to be set off approximately 2 Myr after the
initial star formation took place, possibly initiated by the en-
ergetic activity of the evolving cluster core. Walborn & Parker
argue that this is very similar to N11 where the stellar content of
LH9 and LH10 also suggests an age difference of ∼2 Myr (also
see Walborn et al. 1999). The process in N11, however, would
be advanced by∼2 Myr, classifying it as an evolved 30 Doradus
analogue, though less massive.
More recently, several bright IR sources showing charac-
teristics of young stellar objects were discovered in the N11B
nebula surrounding LH10 by Barbá et al. (2003). These objects
are probably intermediate mass (pre-)main-sequence Herbig
Ae/Be stars belonging to the same generation as do the LH10
objects (the pre-main-sequence evolutionary timescales of in-
termediate mass stars being longer than that of massive OB
stars). Barbá et al. also found that the massive stars in LH10
have blown away their ambient molecular material and are cur-
rently disrupting the surface of the parental molecular cloud
material surrounding LH10 (i.e. the material in N11B).
16 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
4.44.54.64.7
log Teff
65 66
Fig. 12. Comparison of the programme stars located in N11 with isochrones derived from the evolutionary tracks of Schaerer et al.
(1993). The symbols used to denote dwarf, giants and bright- and supergiants are, respectively, circles, triangles and squares.
Filled and open symbols are used to distinguish stars associated with LH9 and LH10, respectively. Black symbols refer to objects
located within a radius of two arc minutes of the LH9 and LH10 core, whereas the grey symbols correspond to stars found outside
this radius. Labels correspond to the N11 identification numbers given in Tab. 1. Isochrones (dashed lines) are shown from one
up to ten million years with 1 Myr intervals. For reference evolutionary tracks from Schaerer et al. (1993) are drawn as grey lines,
where for clarity purpose only the blueward part of the evolution for the most massive tracks is shown. The ZAMS corresponding
to these tracks is given by the black solid line.
6.2. The ages of LH9 and LH10
To estimate the ages of our programme stars we compare their
location in the HR-diagram to theoretical isochrones in Fig. 12.
The luminosity classes V, III and I-II are represented using, re-
spectively, circles, triangles and squares. To differentiate be-
tween stars associated with LH9 and LH10, respectively, filled
and open symbols are used. Membership is defined on the ba-
sis of minimum distance to either the LH9 or LH10 cluster
core. This, therefore, should be taken with some care: the cores
are about 4 arcmin or 60 parsec apart, which can be traversed
in 2 Myr if the proper motion of the star is some 30 km s−1.
Runaway O- and B-type stars have typical velocities of 50–
100 km s−1, therefore, it is entirely possibly that (a few) indi-
vidual objects are assigned to the wrong association. For posi-
tion reference see Fig. 11. Isochrones in Fig. 12 are shown as
dashed lines for one up to ten million years with 1 Myr intervals
and were derived from the evolutionary tracks of Schaerer et al.
(1993). Note that these tracks do not account for the effects of
rotation.
The distribution of the stars in Fig. 12 is such that they can
be separated in a group of objects younger than three million
years and a group of objects older than this age. We also see that
the oldest objects are predominantly found in LH9, whereas
LH10 contains the largest fraction of young objects, suggest-
ing that the clusters can indeed be separated in terms of age. At
first sight it seems that they can be characterised by an age of
∼7 Myr and ∼2 Myr, respectively. However, a significant num-
ber of objects in both clusters are found to be several million
years younger and older than these preliminary ages. A possi-
ble explanation for this large age scatter could be “contamina-
tion” by field stars. To assess this possibility we have assigned
grey symbols to all objects outside a radius of two arc minutes
from the cluster cores. Disregarding these objects reduces the
age scatter significantly; still a number of stars appear to con-
tradict with the notion of two coeval populations.
To investigate the age distributions in more detail the indi-
vidual age estimates are shown in Fig. 13. LH9 and LH10 ob-
jects, which are shown using the identical symbols as in Fig. 12,
are placed in, respectively, the left and right part of the diagram
and are separated using a dashed line. First concentrating on the
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 17
8 29 33 36 42 45 58 65 66 123 4 26 31 32 38 48 51 60 61 68 72 87
N11-#
Fig. 13. Individual age estimates for stars in N11 based on the
non-rotating tracks of Schaerer et al. (1993). Symbols have the
same meaning as in Fig. 12. The left and right part of the
diagram, respectively, contain stars associated with LH9 and
LH10 and are separated using a dashed line. The upward point-
ing arrow denotes a lower age limit for N11-065 determined
from its surface helium abundance using chemically homoge-
neous evolutionary models (Yoon et al. 2006; Yoon & Langer
2005).
LH9 objects we see that they are near-coeval with exception of
four dwarfs. Of these the two non-core members are located
at a distance of approximately six arc minutes from the central
concentration. Therefore, it is probable that they are either spa-
tially not related to N11 or that possibly their formation was
triggered more recently by the stellar activity in LH9. The third
dwarf N11-065 occupies a location very close to the ZAMS and
in principle would be the youngest member of LH9. However,
we also find a considerable helium enrichment for this object
(YHe= 0.17). Combined with the fact that N11-065 has a large
mass discrepancy, this suggests that it might be evolving chem-
ically homogeneously. Consequently, a more appropriate age
estimate should be derived from tracks appropriate for this kind
of evolution. Adopting such tracks from Yoon et al. (2006, also
see Yoon & Langer 2005) we derive a lower limit of 6 Myr for
the age of N11-065 based on the surface helium abundance.
In Fig. 13 this estimate is indicated using an upward pointing
arrow, and is in good agreement with the bulk of the LH9 stars.
The dwarf N11-123 exhibits no chemical peculiarities and
its location on the sky places it in the central concentration of
LH9, suggesting that it should have been formed in the burst
of star formation that formed the cluster. However, given the
fact that N11-123 is in principle the only star deviating from
the coeval nature of LH9 we suspect that it is like N11-058
and N11-066 part of the periphery of the cluster and that its
relatively close position to the core is due to a projection ef-
fect. Consequently, we conclude that the central concentration
of LH9 is coeval with an age of ∼7.0 ±1.0 Myr. The error bar
is of the same order of magnitude as the 0.5–1.0 Myr intro-
duced by uncertainties in the evolutionary tracks due to effects
of (relatively modest) rotation (see Paper II).
The right part of Fig. 13 shows that the stars in the cen-
tral concentration of LH10 have ages ranging from one up to
approximately six million years. Despite this large scatter it is
clear that the majority of the stars are younger than 4.5 Myr
and that only one object (N11-087) in the cluster core is older
than this age. An explanation for the large age of the latter ob-
ject might be that it formed in LH9 and, over time, migrated
towards LH10. As explained at the start of this section, this
is a possibility given that the distance from the centre of LH9
to the current position of N11-087 is only three arc minutes.
Considering the error bars on the age determinations and the
additional uncertainty of 0.5–1.0 Myr introduced by rotation
(see Paper II), we finally estimate an age best describing LH10
of ∼3.0 ±1.0 Myr.
6.3. Sequential star formation?
If the starbursts in N11 are sequential the formation of LH10
should be induced by supernova explosions and/or stellar winds
in LH9. Given the age difference and distance between the
two clusters we find that both means of triggering are possi-
ble. In the first case the age difference of approximately four
million years is compatible with the time of ∼3 Myr it takes
before the most massive stars end their life in a supernova (e.g.
Schaller et al. 1992). The time needed for the supernova shock
to cross the distance of∼3 arc minutes between LH9 and LH10,
corresponding to approximately 40 parsec at the distance of
the LMC, is only ∼105 years (see e.g. Falle 1981); hence it
is a possible scenario. If we consider triggering through stellar
winds, the time scales are also compatible. Garcia-Segura et al.
(1996), for instance, using hydrodynamical simulations have
shown that a 60 M⊙ star can create a wind-driven bubble in the
interstellar medium of ∼50 pc during its main sequence lifetime
of 3 Myr. Consequently, this scenario seems to be appropriate
as well and in agreement with the determined time scales.
In view of the above we conclude that a sequential scenario
for LH9 and LH10 seems very likely. A combination of super-
novae and stellar winds from stars in LH9 may have initiated
star formation in LH10.
7. Summary and conclusions
We have analysed a sample of 28 massive OB-type stars lo-
cated in the LMC. For a homogeneous and consistent treatment
of the data we employed the automated fitting method devel-
oped by Mokiem et al. (2006), which combines the genetic al-
gorithm based optimisation routine pikaia (Charbonneau 1995)
with the fast non-LTE unified stellar atmosphere code fastwind
(Puls et al. 2005). The sample is mostly drawn from the targets
observed within the context of the VLT-FLAMES survey of
massive stars (Evans et al. 2005). In total 22 of these stars are
located in the LH9 and LH10 clusters within the giant H ii re-
gion N11. This region is believed to have been the scene of
sequential star formation, with the stellar activity in LH9 ignit-
ing secondary starbursts in different associations in N11. Our
main findings are summarised below.
18 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
i) The effective temperature per spectral sub-type of the LMC
stars is found to be intermediate between that of Galactic
and SMC O- and early B-type stars, with the LMC objects
being, respectively, cooler and hotter by typically ∼2 kK.
ii) Based on the helium and hydrogen lines it was possible to
determine the effective temperatures, though with relatively
large error bars, of the four O2 stars in our sample. Three
of these are found to be hotter by more than 3 to 7 kK (see
Sect. 4.2) compared to the O3 star in our sample, suggesting
that O2 stars indeed represent a hotter subgroup within the
O-type class. However, we note that for the one ON2 star
a relatively low Teff of 45 kK was obtained, indicating that
the N iv and N iii classification lines (Walborn et al. 2002b)
are not fully compatible with the helium lines traditionally
used for classification.
iii) The spectroscopically determined masses of the dwarf and
giant stars in our set of programme stars are found to be
systematically smaller than those derived from non-rotating
evolutionary tracks. For helium enriched dwarfs and giants,
i.e. those having YHe > 0.11, we find that all show this mass
discrepancy. The same was found in an analysis of SMC
stars using the same methods (Mokiem et al. 2006). We in-
terpret this as evidence for efficient rotationally enhanced
mixing leading to the surfacing of primary helium and to
an increase of the stellar luminosity.
iv) The bright giants and supergiants do not show any mass
discrepancy, regardless of the surface helium abundance.
This also is consistent with the finding for Galactic and
SMC class I-II objects studied with the same methodology
(Mokiem et al. 2005, 2006). This suggests that shortly after
birth all these stars must have rotated at less than about 30
to 40 percent of the surface break-up velocity.
v) A weak correlation is found between microturbulent veloc-
ity and surface gravity. More extended atmospheres (i.e.
lower gravity stars) require a relatively large vturb to fit the
lines. The reason for this relation is unclear, however, it
does not seem to be connected to the lines being formed
closer to the sonic point of the wind flow in low gravity
stars.
vi) From a comparison of modified wind momenta Dmom we
find that the wind strengths of LMC stars are weaker com-
pared to Galactic stars, and stronger compared to SMC
stars. Comparing the derived Dmom as a function of lumi-
nosity with predictions for LMC metallicities by Vink et al.
(2001) yields good agreement in the entire luminosity range
that was investigated (5.0 < L/L⊙ < 6.1).
vii) We have determined an age of, respectively, ∼7.0±1.0 Myr
and ∼3.0±1.0 Myr for the clusters LH9 and LH10. The age
difference and relative distances are in good agreement with
a sequential star formation scenario, in which stellar activ-
ity in LH9 triggered the formation of LH10.
Acknowledgements. We would like to thank Sung-Chul Yoon and
Alexander van der Horst for constructive discussions, and George
Meynet for providing us with the evolutionary models of the Geneva
group. M.R.M. acknowledges financial support from the NWO
Council for Physical Sciences. S.J.S. acknowledges the European
Heads of Research Councils and European Science Foundation
EURYI (European Young Investigator) Awards scheme, supported by
funds from the Participating Organisations of EURYI and the EC Sixth
Framework Programme. A.H. and F.N. acknowledge support from the
Spanish MEC through project AYA2004-08271-C02. JSV acknowl-
edges RCUK for his Fellowship. Spectral fits were calculated us-
ing the LISA computer cluster at SARA Computing & Networking
Services.
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20 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
Appendix A: Fits and comments on individual
objects
The observed spectra shown in this section were corrected for
radial velocities. If not noted differently the lines that were fit-
ted are the hydrogen Balmer lines Hα, Hγ and Hδ; the He i
singlet line at 4387 Å; the He i triplet lines at 4026, 4471 and
4713 Å; and the He ii lines at 4200, 4541 and 4686 Å. Over
plotted are the best fit spectra, unless noted differently. For a
discussion of the line weighting scheme adopted in our fitting
procedure we refer to Paper I (especially to their Table 3.).
N11-004 (Fig. A.1) With exception of He i λ4471 good fits
for all lines of this OC9.7 supergiant were obtained. The under
prediction of the He i λ4471 line is possibly connected to the
so-called generalised dilution effect (Voels et al. 1989, see also
Repolust et al. 2004).
N11-008 (Fig. A.1) All lines are reproduced well. Note that
given the weakness of He ii λ4686 and lower weight assigned
to this line, the reproduction of its profile can be considered
good.
N11-026 (Fig. A.1) A high effective temperature of 53 kK
was needed to fit the spectrum of this O2 giant. Note that even
though on its relative plot scale the very weak He i λ4471 line
is not fitted perfectly, the quantitative fit quality is good and
is comparible to that of the He ii lines. Still, trying to assess
whether the fit could be improved we ran test fits with increas-
ing relative weight of the He i λ4471 line. It turned out that an
increase of a factor of five was necessary to obtain a solution
with an improved He i λ4471 line fit, which is shown in Fig. A.1
as a set of dotted lines. This best fit has a Teff lower by 3.7 kK
and all other fit parameters approximately equal compared to
the regular fit. The overal fit quality of the other diagnostic lines
is somewhat less, though this is difficult to discern by eye be-
cause these lines are relatively strong (one has to zoom in to the
0.5 to 1 percent level, as we do for He i λ4471). Note, however,
that the error estimates for Teff are relative large and within the
lower error estimate of 3.9 kK (cf. Tab. 3) the solutions agree.
Also note that even for the lower Teff solution this O2 star is
still considerably hotter than an O3 star (see Sect. 4.1).
N11-029 (Fig. A.1) A good fit was obtained. The slight un-
derestimation of the core strength of the He ii λ4200 and
He ii λ4541 lines is, given the relative weakness of these lines,
not significant.
N11-031 (Fig. A.2) To fit the spectrum of this O2 star a
relatively low effective temperature of 45 kK was necessary.
Walborn et al. (2004) derived a Teff of 55 kK for this object.
However, their determination of this parameter was based on
the analysis of the N iv-v ionisation adopting a fixed gravity of
log g = 4.0. Based on the analysis of the helium ionisation we
exclude a Teff higher than ∼47 kK. This is illustrated by the dot-
ted lines, which correspond to a model calculated for an effec-
tive temperature higher by the upper error estimate we derived
for Teff. For even higher temperatures the He i λ4471 would, in
contradiction to the observations, disappear completely (also
see Sect. 4.1).
N11-032 (Fig. A.2) The He i and He ii blend at 4026 Å was
not observed for this object. In the final fit, therefore, also the
Hδ Balmer line is shown.
N11-033 (Fig. A.2) This B0 giant is fast rotator. To repro-
duce the observed line profiles a projected rotational velocity
of 256 km s−1 was necessary. In the presented fit the strength
of He ii λ4541 seems to be under predicted. However, given the
relative weakness of this line this small discrepancy is negligi-
ble. Note that the sharp change in the central part of the Hα
profile is the result of an over subtraction of the core nebular
feature and was not taken into account in the fitting procedure.
N11-036 (Fig. A.2) The line profiles of N11-036 could be re-
produced quite accurately because of its relatively slow rota-
tion.
N11-038 (Fig. A.3) The relatively poor reproduction of
He i λ4387 in this O5 bright giant is the result of the low weight
assigned to this line based on the spectral type of N11-038. A
good fit to the He i λ4471 could not be obtained due to its pecu-
liar line profile shape. Possibly this triangular profile is related
to macroturbulence.
N11-042 (Fig. A.3) A relatively low projected rotational ve-
locity was found for this B0 giant allowing for a nearly perfect
N11-045 (Fig. A.3) A good fit was obtained for this O9 giant.
Despite the good fit quality a problem is apparent when Ms and
Mev are compared in Tab. 2. The spectroscopic mass is found
to be less than half its evolutionary equivalent.
N11-048 (Fig. A.3) Based on the fact that the He ii λ4686 is
the strongest line in its spectrum Parker et al. (1992) classified
this star as a Vz star. The fit parameters, however, place this
object at a considerable distance from the theoretical ZAMS in
Fig. 5. We also find that the spectrum can not be reproduced ac-
curately. The width of the neutral helium lines is systematically
unpredicted, whereas the width of the He ii is over predicted.
This and the fact that the helium abundance that is recovered
from the spectrum is rather low (YHe = 0.06) are indicative of
a possible binary nature.
N11-051 (Fig. A.4) The spectrum of the O5 dwarf N11-051
shows very broad lines, indicating fast stellar rotation. To ob-
tain the final fit a projected rotational velocity of 333 km s−1
was required.
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 21
N11-004   OC9.7 Ib
 4330  4340  4350
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4192  4200  4208
HeII 4200
 4537  4541  4545
HeII 4541
 4683  4686  4689
HeII 4686
N11-008   B0.7 Ia
 4335  4340  4345
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4192  4200  4208
HeII 4200
 4537  4541  4545
HeII 4541
 4684  4686  4688
HeII 4686
N11-026   O2 III(f*)
 4335  4340  4345
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4681  4686  4691
HeII 4686
N11-029   O9.7 Ib
 4335  4340  4345
 6552  6562  6572
 4023  4026  4029
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4194  4200  4206
HeII 4200
 4537  4541  4545
HeII 4541
 4684  4686  4688
HeII 4686
Fig. A.1. Comparison of the observed line profiles of N11-004, -008, -026 and -028 with best fitting synthetic line profiles
obtained using the automated fitting method (grey lines). Wavelengths are given on the horizontal axis in Å. The vertical axis
gives the normalised flux. Note that this axis is scaled differently for each line. The dotted line profiles for N11-026 correspond
to the best fit obtained adopting a five times larger relative weight for the He i λ4471 line. This best fit has a Teff lower by 3.7 kK
and all other fit parameters approximately equal compared to the fit for the nominal He i λ4471 weight.
N11-058 (Fig. A.4) A good fit was obtained for all lines of
this O5.5 dwarf. The small under prediction of the He i4387
line strength is the result of the relatively low weight assigned
to this line for the spectral type of this star.
N11-060 (Fig. A.4) Note that the He i λ4471 line of this O3
dwarf has a strength that is comparable to the strength of this
line in the spectrum of the O2 giant N11-031. As the effective
temperature we find for these two objects are also compara-
ble, the helium spectrum suggest a spectral type of O3 for both
objects.
N11-061 (Fig. A.4) All lines of this O9 dwarf are reproduced
correctly. The mass loss could be reliably determined using the
automated method at a rate of 2.1 × 10−7 M⊙. Note that the
22 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
N11-031   ON2 III(f*)
 4333  4340  4347
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4465  4471  4477
HeI 4471
 4710  4713  4716
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4680  4686  4692
HeII 4686
N11-032   O7 II(f)
 4096  4102  4108
 4335  4340  4345
 6547  6562  6577
 4386  4388  4390
HeI 4387
 4469  4471  4473
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-033   B0 IIIn
 4330  4340  4350
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4384  4388  4392
HeI 4387
 4467  4471  4475
HeI 4471
 4709  4713  4717
HeI 4713
 4192  4200  4208
HeII 4200
 4537  4541  4545
HeII 4541
 4681  4686  4691
HeII 4686
N11-036   B0.5 Ib
 4330  4340  4350
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4192  4200  4208
HeII 4200
 4537  4541  4545
HeII 4541
 4684  4686  4688
HeII 4686
Fig. A.2. Same as Fig. A.1, however, for N11-031, -032, -033 and -036. Shown as dotted profiles for N11-031 is the effect of and
a 2.2 kK increase in Teff .
value of 1.8 obtained for the wind acceleration parameter β is
rather high for a dwarf type object.
N11-065 (Fig. A.5) A relatively high helium abundance of
0.17 was recovered from the spectrum. The good fit quality
obtained for all lines, however, does not indicate an overesti-
mation of this parameter. A strong mass discrepancy is found
with Ms being smaller than Mev by approximately a factor of
N11-066 (Fig. A.5) Apart from a slight under prediction of the
width of He ii λ4686 all lines of this O7 dwarf could be fitted
with good accuracy. Like the previously discussed object, again
a large mass discrepancy is found with a ratio of Mev to Ms of
N11-068 (Fig. A.5) The final fit reproduces the observed line
profiles very accurately. No further comments are required.
N11-072 (Fig. A.5) The star has a very sharp lined spectrum.
The best fit required a vr sin i of 14 km s
−1. Note that given the
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 23
N11-038   O5 II(f+)
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4466  4471  4476
HeI 4471
 4709  4713  4717
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-042   B0 III
 4330  4340  4350
 6547  6562  6577
 4023  4026  4029
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4712  4713  4714
HeI 4713
 4192  4200  4208
HeII 4200
 4536  4541  4546
HeII 4541
 4684  4686  4688
HeII 4686
N11-045   O9 III
 4335  4340  4345
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4683  4686  4689
HeII 4686
N11-048   O6.5 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4465  4471  4477
HeI 4471
 4710  4713  4716
HeI 4713
 4194  4200  4206
HeII 4200
 4535  4541  4547
HeII 4541
 4682  4686  4690
HeII 4686
Fig. A.3. Same as Fig. A.1, however, for N11-038, -042, -045 and -048.
lower error estimate, given in Tab. 3, and the spectral resolution
of the data this value can also be interpreted as an upper limit.
N11-087 (Fig. A.6) To fit the spectrum of N11-087 a vr sin i
of 276 km s−1 was needed. The final fit shows a good repro-
duction of the line profiles with a slight under prediction of the
He ii λ4541 line. Given the signal-to-noise ratio of the spectrum
and the relative strength of this line, we do not believe this to
be significant.
N11-123 (Fig. A.6) Like the previously discussed object N11-
123 is also classified as an O9.5 dwarf. The effective tempera-
ture needed to fit the spectrum is, however, higher by more than
2 kK. This is the result of a log g that is higher by ∼0.2 dex.
BI 237 (Fig. A.7) A good fit was obtained for this O2 dwarf
star. The He i λ4471 line is hardly visible. Fortunately, the hy-
drogen and He ii lines provide the automated fitting method
with enough information to determine an effective temperature.
The error estimates on this parameter, however, are consider-
able (see Tab. 3). The effect of lowering Teff to the lower error
estimate of 49 kK is illustrated by the dotted profiles.
BI 237 was recently also analysed by Massey et al. (2005).
Compared to this study we find large differences for Teff
24 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
N11-051   O5 Vn((f))
 4330  4340  4350
 6547  6562  6577
 4020  4026  4032
HeI + HeII 4026
 4382  4388  4394
HeI 4387
 4466  4471  4476
HeI 4471
 4706  4713  4720
HeI 4713
 4192  4200  4208
HeII 4200
 4533  4541  4549
HeII 4541
 4678  4686  4694
HeII 4686
N11-058   O5.5 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4468  4471  4474
HeI 4471
 4709  4713  4717
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-060   O3 V((f*))
 4333  4340  4347
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4466  4471  4476
HeI 4471
 4710  4713  4716
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4680  4686  4692
HeII 4686
N11-061   O9 V
 4335  4340  4345
 6547  6562  6577
 4023  4026  4029
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4194  4200  4206
HeII 4200
 4536  4541  4546
HeII 4541
 4683  4686  4689
HeII 4686
Fig. A.4. Same as Fig. A.1, however, for N11-051, -058, -060 and -061.
(+5 kK), log g (+0.2 dex) and Ṁ (−0.4 dex). Our solution
for higher effective temperature and surface gravity can be ex-
plained by the fact that the ionisation structure of the atmo-
sphere is set by both these parameters (also see Paper I). The
reduced mass loss rate we find is the result of the increased
value for β (+0.4) obtained by the automated method.
BI 253 (Fig. A.7) The star has as an identical spectral type
as the previously discussed object. The photospheric parame-
ters determined from the spectrum are also very similar to the
parameters of BI 237. Note that in contrast to this agreement
BI 253 is found to have a much denser wind. Its mass loss rate
is larger by a factor 2.5 and reflects the fact that compared to
BI 237 the wind lines Hα and He ii λ4686 are more filled in.
In their analysis Massey et al. (2005) could only determine a
lower estimate for Teff of 48 kK, which is consistent with our
lower error estimate. The other parameters determined by these
authors are in fair agreement with our analysis. An exception
to this is the surface gravity for which Massey et al. estimate
a value lower by 0.3 dex, which is the result of the lower Teff
used for their fit.
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 25
N11-065   O6.5 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-066   O7 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-068   O7 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-072   B0.2 III
 4335  4340  4345
 6547  6562  6577
 4024  4026  4028
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4712  4713  4714
HeI 4713
 4192  4200  4208
HeII 4200
 4536  4541  4546
HeII 4541
 4684  4686  4688
HeII 4686
Fig. A.5. Same as Fig. A.1, however, for N11-065, -066, -068 and -072.
Sk −66 18 (Fig. A.7) With exception of the He ii lines at
4541 Å and 4686 Å we obtained a near perfect fit. The helium
abundance YHe = 0.14 indicates an enriched atmosphere.
Sk −66 100 (Fig. A.7) was previously analysed by Puls et al.
(1996). Compared to this study we find similar parameters with
exception of a lower Teff which is the result of the inclusion of
blanketing and an increased helium abundance of YHe = 0.19.
The final fit shows that with this large value all lines could
be reproduced quite accurately. The slight mismatch of the
He ii λ4686 line can be improved by increasing Ṁ to its up-
per error estimate. The effect of this increase of 0.12 dex on the
line profiles is shown using dotted lines.
Sk −67 166 (Fig. A.8) This supergiant has both Hα and
He ii λ4686 in emission indicating the presence of a dense stel-
lar outflow. To correctly reproduce the line profiles a very high
mass loss rate of 9.3×10−6 M⊙yr
−1 was required. This value is
in good agreement with the findings of Crowther et al. (2002)
who studied this objects using the model atmosphere code cm-
fgen (Hillier & Miller 1998) in the optical, UV and far-UV.
Compared to these authors we also find that the other fit pa-
rameters are in good agreement. An exception to this is the he-
26 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
N11-087   O9.5 Vn
 4325  4340  4355
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4465  4471  4477
HeI 4471
 4707  4713  4719
HeI 4713
 4192  4200  4208
HeII 4200
 4534  4541  4548
HeII 4541
 4680  4686  4692
HeII 4686
N11-123   O9.5 V
 4330  4340  4350
 6542  6562  6582
 4022  4026  4030
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4192  4200  4208
HeII 4200
 4536  4541  4546
HeII 4541
 4683  4686  4689
HeII 4686
Fig. A.6. Same as Fig. A.1, however, for N11-087 and -123.
lium abundance. Crowther et al. (2002) adopted a fixed value
of YHe = 0.2 whereas our automated fitting method was able to
self consistently determine a value of YHe = 0.28.
Sk −70 69 (Fig. A.8) To fit the spectrum of this O5 dwarf
Sk −70 69 a helium abundance of YHe = 0.17 was required.
Note that, similar to N11-065, which is also strongly enriched
with helium, the spectroscopic mass is found to be much
smaller than the evolutionary mass.
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 27
BI237   O2 V((f*))
 4333  4340  4347
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4465  4471  4477
HeI 4471
 4710  4713  4716
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4680  4686  4692
HeII 4686
BI253   O2 V((f*))
 4333  4340  4347
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4465  4471  4477
HeI 4471
 4710  4713  4716
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4680  4686  4692
HeII 4686
Sk -66 18   O6 V((f))
 4335  4340  4345
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
Sk -66 100   O6 II(f)
 4335  4340  4345
 6552  6562  6572
 4021  4025  4029
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4469  4471  4473
HeI 4471
 4711  4713  4715
HeI 4713
 4196  4200  4204
HeII 4200
 4537  4541  4545
HeII 4541
 4676  4686  4696
HeII 4686
Fig. A.7. Same as Fig. A.1, however, for BI 237, BI 253, Sk −66 18 and Sk −66 100. The dotted profiles for BI 237 correspond
to a model calculated with a Teff reduced by 3.8 kK compared to the best fit value. For Sk −66 100 the dotted profiles show the
effect of an increase in Ṁ by 0.12.
28 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
Sk -67 166   O4 If+
 4335  4340  4345
 6547  6562  6577
 4021  4025  4029
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4469  4471  4473
HeI 4471
 4711  4713  4715
HeI 4713
 4196  4200  4204
HeII 4200
 4537  4541  4545
HeII 4541
 4677  4686  4695
HeII 4686
Sk -70 69   O5 V
 4335  4340  4345
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4468  4471  4474
HeI 4471
 4709  4713  4717
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
Fig. A.8. Same as Fig. A.1, however, for Sk −67 166 and Sk −70 69.
This figure "6489fig11.jpg" is available in "jpg"
 format from:
http://arxiv.org/ps/0704.1113v1
http://arxiv.org/ps/0704.1113v1
	Introduction
	Data description
	Analysis method
	The assumption of spherical symmetry
	Fundamental parameters
	Effective temperatures
	The Teff scale of O2 stars
	Gravities
	Helium abundances
	Microturbulence
	Wind parameters
	The mass discrepancy
	Chemically homogeneously evolving O2 stars
	Large sample trends
	The evolutionary status of N11
	LH9 and LH10 in N11
	The ages of LH9 and LH10
	Sequential star formation?
	Summary and conclusions
	Fits and comments on individual objects

  [Abridged] We have studied the optical spectra of 28 O- and early B-type
stars in the Large Magellanic Cloud, 22 of which are associated with the young
star-forming region N11. Stellar parameters are determined using an automated
fitting method, combining the stellar atmosphere code FASTWIND with the
genetic-algorithm optimisation routine PIKAIA. Results for stars in the LH9 and
LH10 associations of N11 are consistent with a sequential star formation
scenario, in which activity in LH9 triggered the formation of LH10. Our sample
contains four stars of spectral type O2, of which the hottest is found to be
~49-54 kK (cf. ~45-46 kK for O3 stars). The masses of helium-enriched dwarfs
and giants are systematically lower than those implied by non-rotating
evolutionary tracks. We interpret this as evidence for efficient
rotationally-enhanced mixing, leading to the surfacing of primary helium and to
an increase of the stellar luminosity. This result is consistent with findings
for SMC stars by Mokiem et al. For bright giants and supergiants no such
mass-discrepancy is found, implying that these stars follow tracks of modestly
(or non-)rotating objects. Stellar mass-loss properties were found to be
intermediate to those found in massive stars in the Galaxy and the SMC, and
comparisons with theoretical predictions at LMC metallicity yielded good
agreement over the luminosity range of our targets, i.e. 5.0 < log L/L(sun) <
6.1.

Introduction
Massive stars play an intricate role in the evolution of galax-
ies. Because of the large energies associated with their stel-
lar winds, ionising radiation, and life-ending supernova ex-
plosions, they dictate galactic structuring processes such as
star formation and the creation and evolution of supperbub-
bles (e.g. Oey 1999). Mounting evidence also points to a direct
link between massive stars and exotic phenomena such as γ-ray
bursts (e.g. Hjorth et al. 2003) and the reionisation of the early
universe (Bromm et al. 2001). Accordingly, understanding the
properties of these stars, both in terms of their fundamental pa-
rameters as well as their evolution, is fundamental.
The initial metal composition (Z) of the gas out of which
massive stars form has a strong impact on their global prop-
erties and characteristics. Many studies have shown that pa-
rameters such as the effective temperature and ionising fluxes
are strong functions of Z (e.g. Kudritzki 2002; Mokiem et al.
2004; Massey et al. 2005; Mokiem et al. 2006), adding an ex-
tra dimension to the conversion of morphological properties
http://arxiv.org/abs/0704.1113v1
2 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
such as spectral type to physical quantities. Theoretical and ob-
servational arguments (e.g. Kudritzki & Puls 2000; Vink et al.
2001) also point to a relation between the strengths of the stel-
lar winds of these objects and metallicity. As wind mass loss
leads to partial evaporation of the star (e.g. Chiosi & Maeder
1986) – with possible consequences for the nature of the com-
pact object that is left behind after the final supernova explosion
– and to loss of angular momentum (e.g. Meynet & Maeder
2000), it also dictates to a large extent its evolutionary path
and fate (e.g. Yoon & Langer 2005; Woosley & Heger 2006).
Quantifying the mass loss versus metallicity dependence Ṁ(Z),
therefore, is an important quest in astrophysics.
Due to their proximity and low metal content the
Magellanic Clouds provide us with unparallelled laboratories
to test and enlarge our knowledge of massive stars. These
galaxies, therefore, have been in the focal point of many stud-
ies analysing their massive star content. Early studies (e.g.
Conti et al. 1986; Garmany et al. 1987; Massey et al. 1989;
Parker et al. 1992; Walborn et al. 1999) predominantly relied
on photometric data and spectral type calibrations. Only rel-
atively recent the advent of large telescopes and the develop-
ment of sophisticated stellar atmosphere models has allowed
for more detailed analyses of individual stars (e.g. Puls et al.
1996; Hillier & Miller 1999; Crowther et al. 2002; Bouret et al.
2003; Martins et al. 2004). Though all these studies have con-
tributed enormously to our understanding of massive stars, the
samples analysed so far have been rather limited in size (a few
objects at a time) and have been focused predominantly on ob-
jects in the Small Magellanic Cloud (SMC) – because its metal
deficiency is more extreme than that of the Large Magellanic
Cloud (LMC) (but see Massey et al. 2004, 2005).
The limited sizes of the samples that have so far been stud-
ied are at the root cause of the perhaps somewhat disheartening
conclusion that – in spite of all the progress that has been made
– we still cannot provide robust and sound answers to the ques-
tion: what is the role of metal content, stellar winds and rota-
tion in the evolution of massive stars? To help attack this prob-
lem, our research group has conducted a VLT-FLAMES Survey
of Massive Stars (see Evans et al. 2005). In this ESO Large
Program the Fibre Large Array Multi-Element Spectrograph
at the Very Large Telescope was used to obtain optical spec-
tra of more than 50 O- and early B-type stars in the Magellanic
Clouds.
Here we present the homogeneous analysis, employing au-
tomated spectral fitting methods, of a sample of 28 O-type and
early B-type stars in the Large Magellanic Cloud; 22 targets
from the FLAMES survey and 6 from other sources. This is
so far the largest sample of massive stars studied in the LMC
and it almost doubles the amount of massive objects in this
galaxy for which parameters have been derived from quantita-
tive spectroscopy. Specifically, we will try to establish the mass
loss rates of OB stars in this Z ∼ 1/2 Z⊙ environment to a level
of precision that allows for a quantitative comparison with sim-
ilar objects in the SMC and our Galaxy. This will provide a new
Z-point in testing the fundamental prediction provided by radi-
ation driven wind theory for the mass-loss – metallicity depen-
dence: Ṁ(Z) ∝ Z0.5−0.7 (e.g. Kudritzki et al. 1987; Puls et al.
2000; Vink et al. 2001).
The majority of our LMC sample is associated with
the spectacular star forming region N11 (Henize 1956). It
has a Hα luminosity only surpassed by that of 30 Doradus
(Kennicutt & Hodge 1986), ranking it as the second largest
H ii region in the Magellanic Clouds. N11 is host to several
OB associations of apparently different ages, the formation of
which is believed to have been triggered by stellar activity in
the central OB cluster (Parker et al. 1992; Walborn & Parker
1992; Walborn et al. 1999). Our observations sample both the
central cluster LH9 as well as the younger cluster LH10, allow-
ing for an investigation of a possible sequential star formation
scenario. We will use our N11 stars to test predictions of mas-
sive star evolution, including the role of rotation, and the star
formation history.
This paper is organised as follows: in Sect. 2 we describe
the LMC data set that was analysed using our automated ge-
netic algorithm based fitting method. A short description of this
method is given in Sect. 3 and the results obtained are presented
in Sect. 4, with fits and comments on individual objects given
in the appendix. In Sect. 5 we investigate the discrepancy be-
tween spectroscopically determined masses and those derived
from evolutionary tracks. The evolutionary status of N11 is dis-
cussed in Sect. 6. Finally, Sect. 7 summarises and lists our most
important findings.
2. Data description
Our OB-type star sample is mainly drawn from the targets ob-
served in the LMC within the context of the VLT-FLAMES
survey of massive stars (see Evans et al. 2005). Two fields in
the LMC centred on the clusters N11 and NGC 2004 were
observed in the survey. Here we analyse a subset of the ob-
jects observed in the N11 field. This set consists of all O-type
spectra obtained, excluding those that correspond to confirmed
binaries, and five early B-type spectra of luminous giant and
supergiant stars.
To improve the sampling in luminosity and temperature, we
supplemented the FLAMES targets with six relatively bright O-
type field stars. These objects are part of the Sanduleak (1970,
hereafter Sk) and Brunet et al. (1975, hereafter BI) catalogues,
and were observed as part of the programs 67.D-0238, 70.D-
0164 and 074.D-0109 (P.I. Crowther) using the Ultraviolet and
Visual Echelle Spectrograph (UVES) at the VLT.
The observations of the FLAMES targets are described ex-
tensively by Evans et al. (2006), to which we refer for full de-
tails. Here we only summarise the most important observa-
tional parameters. Basic observational properties of the pro-
gramme stars together with common aliases are given in Tab. 1.
Note that N11-031, BI 237, BI 253 and Sk−67 166 were stud-
ied recently using line blanketed stellar atmosphere models. In
the appendix a comparison with these analyses is provided. The
FLAMES targets were observed with the Giraffe spectrograph
mounted at UT2. For six wavelength settings a spectrum was
acquired six times for each object with an effective resolving
power of R ≃ 20 000. These multiple exposures, often at differ-
ent epochs, allowed for the detection of variable radial veloc-
ities. As a result, a considerable number of binaries could be
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 3
Table 1. Basic parameters. Primary identification numbers for N11 are from Evans et al. (2006). Identifications starting with
“Sk”, “PGMW” and “BI” are from, respectively, Sanduleak (1970), Parker et al. (1992) and Brunet et al. (1975). Photometric
data for these objects are from Evans et al. (2006) and Parker et al. (1992), the latter are flagged with an asterisk. For non-N11
objects these are from Ardeberg et al. (1972), Issersted (1979) and Massey (2002). Wind velocities given without brackets are
from Crowther et al. (2002), Massa et al. (2003) and Massey et al. (2005). For Sk −66 18 the wind velocity was measured from
O iv (1031-1037 Å) in its FUSE spectrum. Values between brackets are calculated from the escape velocity at the stellar surface.
For N11-031 the value of v∞ is from Walborn et al. (2004), though this is actually based on the value obtained for Sk-68 137 by
Prinja & Crowther (1998).
Primary ID Cross-IDs Spectral V AV MV v∞
Type [km s−1]
N11-004 Sk −66 16 OC9.7 Ib 12.56 0.74 −6.68 [2387]
N11-008 Sk −66 15 B0.7 Ia 12.77 0.84 −6.57 [1619]
N11-026 ... O2 III(f*) 13.51 0.47 −5.46 [3116]
N11-029 ... O9.7 Ib 13.63 0.56 −5.43 [1576]
N11-031 PGMW 3061 ON2 III(f*) 13.68∗ 0.96 −5.78 3200
N11-032 PGMW 3168 O7 II(f) 13.68∗ 0.65 −5.47 [1917]
N11-033 PGMW 1005 B0 IIIn 13.68 0.43 −5.25 [1536]
N11-036 ... B0.5 Ib 13.72 0.40 −5.18 [1714]
N11-038 PGMW 3100 O5 II(f+) 13.81∗ 0.99 −5.68 [2601]
N11-042 PGMW 1017 B0 III 13.93 0.22 −4.79 [2307]
N11-045 ... O9 III 13.97 0.50 −5.03 [1548]
N11-048 PGMW 3204 O6.5 V((f)) 14.02∗ 0.47 −4.95 [3790]
N11-051 ... O5 Vn((f)) 14.03 0.19 −4.66 [2108]
N11-058 ... O5.5 V((f)) 14.16 0.28 −4.62 [2472]
N11-060 PGMW 3058 O3 V((f*)) 14.24∗ 0.81 −5.07 [2738]
N11-061 ... O9 V 14.24 0.78 −5.04 [1898]
N11-065 PGMW 1027 O6.5 V((f)) 14.40 0.25 −4.35 [2319]
N11-066 ... O7 V((f)) 14.40 0.25 −4.35 [2315]
N11-068 ... O7 V((f)) 14.55 0.28 −4.23 [3030]
N11-072 ... B0.2 III 14.61 0.09 −3.98 [2098]
N11-087 PGMW 3042 O9.5 Vn 14.76∗ 0.62 −4.36 [3025]
N11-123 ... O9.5 V 15.29 0.16 −3.37 [2890]
BI 237 ... O2 V((f*)) 13.89 0.62 −5.23 3400
BI 253 ... O2 V((f*)) 13.76 0.71 −5.45 3180
Sk −66 18 ... O6 V((f)) 13.50 0.37 −5.37 2200
Sk −66 100 ... O6 II(f) 13.26 0.34 −5.58 2075
Sk −67 166 HD 269698 O4 Iaf+ 12.27 0.31 −6.54 1750
Sk −70 69 ... O5 V 13.95 0.28 −4.83 2750
detected (Evans et al. 2006), which we subsequently excluded
from our analysis.
To allow for a sky subtraction a master sky-spectrum was
created from combining the sky fibres in the Giraffe spectro-
graph (typically 15), individually scaled by their relative fibre
throughput. Even though in general the sky background is low
and this approach successfully removes the background con-
tribution, in crowded regions such as N11 accurate subtraction
of nebular features remains very difficult. As a result of this,
the line profiles of many of our programme stars still suffer
from nebular contamination. This in principle does not hamper
our analysis. For most stars the core nebular emission is well-
resolved and we simply disregard this contaminated part of the
profile in the automated line fits. Mokiem et al. (2006) showed
by performing tests using synthetic data that with this reduced
amount of information, the automated method can still accu-
rately determine the correct fit parameters. Also tests assessing
the impact of possible residual nebular contamination in the
4 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
line wings or over-subtraction of sky components showed that
its effect is negligible.
For each wavelength range the individual sky-subtracted
spectra were co-added and then normalised using a cubic-
spline fit to the continuum. A final spectrum covering 3850-
4750 and 6300-6700 Å was obtained by merging the nor-
malised data. Depending on the magnitude of the target these
combined spectra have typical signal-to-noise ratios of 100-
The first four field stars listed in Tab. 1 were observed with
the VLT-UVES spectrograph in service mode on 29 and 30
November 2004 under program 74.D-0109. UVES is a two
armed cross-dispersed echelle spectrograph allowing for simul-
taneous observations in the blue and red part of the spectrum.
In the blue arm standard settings with central wavelengths of
390 and 437 nm were used to observe the spectral ranges 3300-
4500 and 3730-5000 Å. For the red arm standard settings with
central wavelengths of 564 and 860 nm provided coverage be-
tween 4620-5600 and 6600-10400 Å. A 1.2′′ wide slit was
used, providing a spectral resolution of 0.1 Å at Hγ, corre-
sponding to an effective resolving power of R ≃ 40 000, a value
which applies to all UVES setups. Individual exposure times
ranged from 1500 to 2200 seconds.
Sk −67 166 was observed on 27 and 29 September 2001
under program 67.D-0238 with UVES using a 1′′ slit. A stan-
dard blue setting with central wavelength 437 nm provided con-
tinuous coverage between 3730-5000 Å. A non-standard red
setup with central wavelength 830 nm was used to observe
the range between 6370-10250 Å. Three exposures, each of
1000 s were obtained for each setup. Note that the data for
this target was previously presented by Crowther et al. (2002).
Finally, Sk −70 69 was observed on 1 and 2 December 2002
using UVES under program 70.D-0164. Standard blue and red
settings with central wavelengths 390 and 564 nm were used
in simultaneous 2400 s exposures using dichroic. A second
non-standard red setup with central wavelength 520 nm (4170-
6210 Å) without dichroic was used in a 1500 s exposure. The
typical two pixel S/N ratios obtained for all spectra are 80 at
Hγ and 60 at Hα.
Spectral types for FLAMES stars were determined by vi-
sual inspection of the spectra, using published standards. In
particular the atlas of Walborn & Fitzpatrick (1990) was used
while considering the lower metallicity environment of the
LMC. The classifications given in Tab. 1 are from Evans et al.
(2006). A comparison with previously published spectral types
is also given in the latter paper. For the field stars we adopted
the classifications given by Walborn (1977), Walborn et al.
(1995, 2002a) and Massey et al. (1995, 2005).
Photometric data for the FLAMES targets was obtained
predominantly from B and V images of the N11 field
taken with the Wide Field Imager at the 2.2-m Max Planck
Gesellschaft/ESO telescope on 2003 April 2004 (Evans et al.
2006). The photometry for stars flagged with an asterisk was
adopted from Parker et al. (1992). The caption of Tab. 1 lists
the references to the various sources of the photometric data of
the field stars.
To calculate the interstellar extinction (AV ) given in Tab. 1
we adopted intrinsic colours from Johnson (1966, and refer-
ences therein) and a ratio of total to selective extinction of
RV = 3.1. With these AV values, extinction corrected visual
magnitudes (V0) were calculated from the observed V-band
magnitudes. Finally, we calculated the absolute visual magni-
tude MV , while adopting a distance modulus of 18.5 for the
LMC (Panagia et al. 1991; Mitchell et al. 2002).
3. Analysis method
All optical spectra are analysed using the automated fitting
method developed by Mokiem et al. (2005, hereafter referred
to as Paper I). Here we will suffice with a short description
of the method and refer to the before mentioned paper and to
Mokiem et al. (2006, hereafter Paper II) for the details.
In short the automated fitting method uses the genetic algo-
rithm based optimisation routine pikaia (Charbonneau 1995) to
determine the set of input parameters for the stellar atmosphere
code fastwind (Puls et al. 2005) which fit an observed spec-
trum the best. This best fitting model is constructed by evolving
a population of fastwind models over a course of generations.
At the end of every generation the parameters of the models
which relatively fit the observed spectrum the best are used to
construct a new population of models. By repeating this proce-
dure a natural optimisation is obtained and after a number of
generations (see below) the best fitting model, i.e. global opti-
mum in parameter space, is found.
Using the concept of a unified model atmosphere the fast
performance code fastwind incorporates non-LTE and an ap-
proximate approach to line blanketing to synthesise hydrogen
and helium line profiles. Consequently, given the spectral range
of the observed data set in this study we will focus on the mod-
elling of the optical hydrogen and helium lines. To account for
the accuracy with which each individual line can be reproduced
by fastwind we adopt the line weighting scheme as described
in Paper I.
In the fitting of the spectra we allow for seven free parame-
ters. These are the effective temperature Teff, the surface gravity
g, the helium number density defined as YHe ≡ N(He)/N(H),
the microturbulent velocity vturb, the projected rotational veloc-
ity vr sin i, the mass loss rate Ṁ and the exponent of the beta-
type velocity law describing the supersonic regime of the stellar
wind. For the terminal velocity of the wind v∞, which cannot
be accurately determined from the optical spectrum, we adopt
values determined from the analysis of ultraviolet (UV) wind
lines. If no UV determination of v∞ is available a scaling re-
lation of v∞ with the escape velocity (vesc) defined at the stel-
lar surface is used throughout the fitting process. For our pro-
gramme stars we adopt the ratio as determined by Lamers et al.
(1995) of v∞/vesc = 2.6. Similar as in Paper II we adopt fixed
values for the atmospheric abundances of the background met-
als. These were scaled with respect to mass ratios based on
the Solar abundances of Grevesse & Sauval (1998, and ref-
erences therein). The metallicity scaling factor was set equal
to the mean metal deficiency of 0.5 as found for the LMC
(Russell & Bessell 1989; Rolleston et al. 2002).
For our current data set the spectral range and quality, with
exception of the signal-to-noise ratio, is similar to the set anal-
ysed in Paper II. Consequently, we adopt the same minimum
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 5
number of generations that have to be calculated to determine
the best fit, i.e. to assure that the global optimum in parameter
space is found. In Paper II, using formal tests that accounted for
nebular contamination of the line profiles and a signal-to-noise
ratio of 50, this number was determined to be 150.
The uncertainties of the fit parameters are determined us-
ing so-called optimum width based error estimates. In Paper I
we have argued that an error estimate for a parameter can be
defined as the maximum variation of this parameter within the
global optimum in parameter space. By measuring the width
of this optimum in terms of fit quality one can determine which
models, based on their fit quality, are associated with the global
optimum. The maximum variation of the individual fit param-
eters within this group of models then gives the error estimates
(see Sect. 4 in Paper I).
In Tab. 3 the optimum width based error estimates of the fit
parameters for our programme stars are listed. The uncertain-
ties in R⋆, L⋆, Ms and Mev, that are derived from the fit param-
eters, were calculated using the same approach as in Paper II.
The adopted uncertainty in the visual magnitude is ∆MV = 0.13
(Panagia et al. 1991) for all our programme stars.
3.1. The assumption of spherical symmetry
fastwind assumes a spherically symmetric star and wind. For
very high rotational speeds, vrot, this may potentially be a prob-
lematic assumption. As it is well known, a high rotation rate
leads to a distortion of the stellar surface and, via the von Zeipel
theorem, to a decrease in flux and effective temperature from
the pole to the equator (von Zeipel 1924, (slightly) modified
by Maeder (1999) for the relevant case of shellular, i.e., radi-
ally dependent, differential rotation, ω = ω(r)). This so-called
“gravity” darkening does not only affect the stellar parameters,
but also the wind (e.g. Cranmer & Owocki 1995; Owocki et al.
1996; Petrenz & Puls 2000). Thus, it might be questioned in
how far the derived properties (which then depend on incli-
nation angle and have a somewhat local character) are repre-
sentative for the global quantities (e.g., mass, luminosity and
mass-loss rate) which refer to integral quantities.
As shown by various simulations (Cranmer & Owocki
1995; Petrenz & Puls 1996, 2000), the difference between local
and global quantities remains small unless the star rotates faster
than >∼ 60. . .70 % of its critical speed. Unfortunately, however,
we cannot directly access the actual rotational speed, vrot, but
only its projected value, vr sin i.
Accounting for the average value of < sin i >= π/4, we
have calculated the ratio of average to critical rotational speed,
Ω ≈ vr sin i/(< sin i > vcrit), and found that only two objects
lie above this value, namely N11-033 and N11-051, both with
Ω ≈ 0.7. All other objects lie well below Ω = 0.3 (except for
N11-087 withΩ ≈ 0.4). Of course, we cannot exclude that also
the latter objects lie above the decisive threshold (if observed
pole-on), but the rather large sample size implies that such a
possibility should be actually present only for a minority of ob-
jects. In conclusion, we suggest that the majority of our objects
is, if at all, only weakly affected by a distortion of surface and
wind. Note, e.g., that Ω = 0.3 leads to a difference in Teff and
R⋆ between pole and equator of less than 5%, i.e., of the same
order or less than our error estimates (cf. Tab. 3). Regarding the
derived gravities and masses, finally, we have applied a consis-
tent centrifugal correction anyhow (cf. Sect. 4.3).
For the two objects with large rotational speeds, on the
other hand, it is rather possible that we observe them al-
most equator-on, i.e., the observed profiles contain an intrin-
sic averaging over the complete stellar disk and thus corre-
spond, at least in part, to the global, polar angle averaged val-
ues. Howarth & Smith (2001) analyzed two galactic fast rota-
tors, HD 149757 (ζ Oph) and HD 191423, accounting for ef-
fects of non-sphericity and von Zeipel’s theorem. Their study
showed that for these objects, rotating at Ω = 0.9, differ-
ences in effective temperature and radius between pole and
equator can amount to 20–30 %. Interestingly, analysis of the
same objects by Herrero et al. (2002), Villamariz et al. (2002),
and Villamariz & Herrero (2005) using both hydrostatic and
spherically symmetric expanding atmospheres yielded (aver-
age) parameters that agreed to within the standard deviation
(5–10%) with the results obtained by Howarth & Smith. For
stellar parameters it thus appears justified to conclude that non-
sphericity has a very small impact even on stars with extreme
vr sin i. But note also that with respect to global mass-loss rates
the situation is much more unsecure and more-D simulations
would be required to constrain their actual values. As shown,
e.g., by Petrenz & Puls (2000), the modified wind-momentum
rate from a rapidly rotating (Ω = 0.85) B-supergiant might be
underestimated up to one magnitude if seen equator on. As the
wind properties of N11-031 and N11-051 are rather normal it
seems however unlikely that in these cases rotational effects
have this type of dramatic impact.
4. Fundamental parameters
4.1. Effective temperatures
In Fig. 1 the distribution of the effective temperatures of our
programme stars as a function of spectral type is shown. The
different luminosity classes are denoted using circles, trian-
gles and squares, respectively, for class V, III and I-II objects.
Similar as in Paper II we see that for a given spectral type
the dwarfs are systematically hotter than the giants and su-
pergiants. This separation can be interpreted as the result of
the reduced surface gravities of the more evolved objects. A
lower surface gravity results in an increased helium ionisation
(e.g. Kudritzki et al. 1983), reducing the Teff needed for a given
spectral type (e.g. Mokiem et al. 2004). A second reason is that
the more evolved objects have stronger winds. These denser
winds induce an increased line blanketing effect, further reduc-
ing the required temperature (e.g. Schaerer & de Koter 1997).
Massey et al. (2005), who also analysed a sizeable sample of
LMC stars, do not find evidence for the dwarfs being hotter
than the giants. Though, we note that their analysis only con-
tained two giant Teff determinations for spectral type later than
For comparison we also show in Fig. 1 as a dashed line
the observed Teff vs. spectral calibration for Galactic O-type
dwarfs from Martins et al. (2005a). With a dotted line the av-
6 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
Table 2. Fundamental parameters of the LMC sample determined using GA optimised spectral fits, with Teff in kK, log g and
log gc in cm s
−2, R⋆ in R⊙, L⋆ in L⊙, vturb and vr sin i in km s
−1, Ṁ in M⊙yr
−1and Ms and Mev in M⊙. Results were obtained
using a population of 72 fastwind models evolved over a minimum of 150 generations. Gravities corrected for centrifugal
acceleration (log gc) were used to calculate the spectroscopic masses (Ms). Evolutionary masses (Mev) were derived from the
tracks of Schaerer et al. (1993).
ID ST Teff log g log gc R⋆ log L⋆ YHe vturb vr sin i Ṁ β Ms Mev
N11-004 OC9.7 Ib 31.6 3.36 3.37 26.5 5.80 0.10 5.7 81 1.78·10−6 1.18 59.9 47.5
N11-008 B0.7 Ia 26.0 2.98 2.99 29.6 5.55 0.10 18.0 83 4.96·10−7 1.87 31.4 31.7
N11-026 O2 III(f*) 53.3 4.00 4.00 10.7 5.92 0.11 19.0 109 1.81·10−6 1.08 42.4 81.9
N11-029 O9.7 Ib 29.4 3.23 3.24 15.7 5.21 0.07 19.1 77 1.73·10−7 1.63 15.5 24.9
N11-031 ON2 III(f*) 45.0 3.85 3.86 13.7 5.84 0.10 20.0 116 3.88·10−6 0.89 49.3 60.8
N11-032 O7 II(f) 35.2 3.45 3.46 14.0 5.43 0.09 11.3 96 8.06·10−7 1.03 20.6 33.8
N11-033 B0 IIIn 27.2 3.21 3.35 15.6 5.07 0.08 17.9 256 2.44·10−7 1.03 19.8 21.0
N11-036 B0.5 Ib 26.3 3.31 3.32 15.6 5.02 0.08 13.6 54 1.06·10−7 0.80a) 18.4 20.0
N11-038 O5 II(f+) 41.0 3.72 3.74 14.0 5.69 0.10 9.2 145 1.52·10−6 0.98 38.8 48.3
N11-042 B0 III 30.2 3.69 3.70 11.8 5.01 0.10 4.0 42 1.89·10−7 1.19 25.1 21.4
N11-045 O9 III 32.3 3.32 3.35 12.0 5.15 0.07 16.8 105 5.48·10−7 0.80a) 11.8 24.6
N11-048 O6.5 V((f) 40.7 4.19 4.20 9.9 5.38 0.06 1.0 130 1.67·10−7 0.80a) 56.6 36.6
N11-051 O5 Vn((f)) 42.4 3.75 3.88 8.4 5.31 0.08 19.7 333 1.01·10−6 0.60 19.5 36.4
N11-058 O5.5 V((f) 41.3 3.89 3.90 8.4 5.27 0.10 14.5 85 1.52·10−7 1.42 20.3 34.4
N11-060 O3 V((f*)) 45.7 3.92 3.93 9.7 5.57 0.12 19.3 106 5.22·10−7 1.26 29.2 49.4
N11-061 O9 V 33.6 3.51 3.52 11.7 5.20 0.09 19.8 87 2.14·10−7 1.80 16.7 26.6
N11-065 O6.5 V((f) 41.7 3.89 3.90 7.4 5.17 0.17 9.4 83 3.63·10−7 0.80a) 15.8 32.9
N11-066 O7 V((f)) 39.3 3.87 3.88 7.7 5.10 0.11 4.8 71 4.08·10−7 0.80a) 16.2 29.5
N11-068 O7 V((f)) 39.9 4.13 4.13 7.1 5.06 0.10 15.6 54 3.43·10−7 1.12 25.2 29.2
N11-072 B0.2 III 30.8 3.78 3.78 7.9 4.70 0.12 7.6 14 2.35·10−7 0.84 13.8 17.7
N11-087 O9.5 Vn 32.7 4.04 4.09 8.9 4.91 0.10 15.5 276 1.38·10−7 0.80a) 35.6 20.9
N11-123 O9.5 V 34.8 4.22 4.23 5.4 4.58 0.09 9.0 110 7.62·10−8 0.80a) 17.8 18.8
BI 237 O2 V((f*)) 53.2 4.11 4.11 9.7 5.83 0.10 12.8 126 7.81·10−7 1.26 44.6 75.0
BI 253 O2 V((f*)) 53.8 4.18 4.19 10.7 5.93 0.09 18.6 191 1.92·10−6 1.21 64.6 84.1
Sk −66 18 O6 V((f)) 40.2 3.76 3.76 12.2 5.55 0.14 10.8 82 1.07·10−6 0.94 31.5 40.7
Sk −66 100 O6 II(f) 39.0 3.70 3.71 13.6 5.58 0.19 8.7 84 8.81·10−7 1.27 34.7 41.4
Sk −67 166 O4 Iaf+ 40.3 3.65 3.66 21.3 6.03 0.28 20.0 97 9.28·10−6 0.94 75.0 70.4
Sk −70 69 O5 V 43.2 3.87 3.88 9.0 5.41 0.17 16.1 131 1.03·10−6 0.78 22.7 39.7
a) assumed fixed value
erage temperature of the SMC dwarfs studied in Paper II is
shown. The LMC dwarfs are found to occupy the region in
between these two average temperature scales, with an average
behaviour intermediate to that of the Galactic and SMC dwarfs.
We interpret this as the result of the metallicity of the LMC that
is lower than the Galactic value and higher than in the SMC. As
the amount of line blanketing in a stellar atmosphere depends
on metallicity the LMC objects have temperatures in between
that of objects in the other two galaxies.
4.2. The Teff scale of O2 stars
Our sample contains four O2 type stars. This spectral type was
introduced by Walborn et al. (2002b) and is assigned based pri-
marily on the ratios of selective emission lines of N iv and N iii.
By modelling these lines Walborn et al. (2004) have shown that
indeed, as had been hypothesised, the O2 spectra correspond to
higher effective temperatures. However, the correct treatment
of nitrogen lines in stellar atmosphere models is notoriously
difficult. The relevant ionisation stages of this atom represent
much more complex ion models compared to the relatively
simple hydrogen and helium ions. Moreover, for the higher ions
of nitrogen the ionisation depends on the extreme-UV radiation
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 7
Table 3. Optimum width based error estimates for the seven fit parameters. The ND entries correspond to error in vturb that reach
up to the maximum allowed value of vturb and, therefore, are formally not defined. See text for details on the calculation of the
uncertainties in the derived parameters. Units: Teff in kK, log g and log gc in cm s
−2, R⋆ in R⊙, L⋆ in L⊙, vturb and vr sin i in
km s−1, Ṁ in M⊙yr
−1and Ms and Mev in M⊙.
ID ∆Teff ∆log gc ∆R⋆ ∆ log L⋆ ∆YHe ∆vturb ∆vr sin i log∆Ṁ ∆β ∆Ms ∆Mev
N11-004 −0.5
−0.06
+0.06 ±1.7 ±0.06
−0.01
+0.02
−0.14
+0.05
−0.00
+0.19
N11-008 −0.6
−0.06
+0.05 ±1.9 ±0.07
−0.01
+0.04
−0.20
+0.16
−0.31
+1.08
N11-026 −3.9
−0.05
+0.05 ±0.8 ±0.14
−0.01
+0.02
−0.05
+0.10
−0.12
+0.03
N11-029 −0.6
−0.05
+0.06 ±1.0 ±0.07
−0.01
+0.02
−0.13
+0.16
−0.42
+0.31
N11-031 −1.6
−0.06
+0.07 ±0.9 ±0.10
−0.02
+0.02
−0.05
+0.06
−0.04
+0.06
N11-032 −0.7
−0.07
+0.06 ±0.9 ±0.06
−0.02
+0.02
−0.16
+0.09
−0.11
+0.28
N11-033 −0.9
−0.06
+0.09 ±1.0 ±0.08
−0.01
+0.03
−0.26
+0.10
−0.20
+0.74
N11-036 −0.3
−0.09
+0.06 ±1.0 ±0.09
−0.01
+0.02
−1.99
+0.30 −
N11-038 −1.0
−0.06
+0.07 ±0.9 ±0.07
−0.01
+0.03
−0.16
+0.10
−0.13
+0.22
N11-042 −0.8
−0.07
+0.09 ±0.7 ±0.07
−0.01
+0.03
−0.47
+0.40
−0.47
+0.77
N11-045 −1.2
−0.09
+0.21 ±0.8 ±0.08
−0.02
+0.03
−0.20
+0.14 −
N11-048 −2.0
−0.21
+0.06 ±0.7 ±0.10
−0.01
+0.02
+13.2
−1.48
+0.74 −
N11-051 −1.9
−0.06
+0.06 ±0.6 ±0.09
−0.02
+0.03
−0.28
+0.11
−0.09
+0.52
N11-058 −0.5
−0.06
+0.05 ±0.5 ±0.06
−0.01
+0.01
−0.36
+0.08
−0.20
+0.60
N11-060 −1.0
−0.05
+0.09 ±0.7 ±0.10
−0.02
+0.03
−13.1
−0.22
+0.06
−0.21
+0.20
N11-061 −0.6
−0.09
+0.07 ±0.7 ±0.06
−0.01
+0.03
−0.21
+0.14
−0.49
+0.38
N11-065 −0.9
−0.07
+0.15 ±0.5 ±0.06
−0.03
+0.04
−1.06
+0.14 −
N11-066 −1.3
−0.15
+0.19 ±0.5 ±0.09
−0.03
+0.05
−0.49
+0.22 −
N11-068 −1.0
−0.19
+0.07 ±0.4 ±0.07
−0.02
+0.03
−14.8
−0.79
+0.16
−0.24
+0.36
N11-072 −0.6
−0.07
+0.09 ±0.5 ±0.06
−0.02
+0.02
−1.29
+0.24
−0.32
+0.31
N11-087 −0.6
−0.09
+0.11 ±0.6 ±0.07
−0.03
+0.02
−2.82
+0.27 −
N11-123 −0.7
−0.11
+0.00 ±0.3 ±0.06
−0.01
+0.02
−2.03
+0.35 −
BI 237 −3.8
−0.08
+0.15 ±0.8 ±0.18
−0.01
+0.03
−12.6
−0.16
+0.12
−0.17
+0.16
BI 253 −5.5
−0.15
+0.14 ±0.9 ±0.19
−0.02
+0.03
−18.0
−0.05
+0.10
−0.21
+0.07
Sk −66 18 −1.1
−0.13
+0.09 ±0.8 ±0.07
−0.03
+0.04
−10.4
−0.05
+0.06
−0.10
+0.11
Sk −66 100 −1.0
−0.09
+0.08 ±0.9 ±0.07
−0.03
+0.05
−0.12
+0.12
−0.19
+0.18
Sk −67 166 −0.8
−0.08
+0.00 ±1.3 ±0.07
−0.04
+0.06
−0.05
+0.02
−0.04
+0.09
Sk −70 69 −1.4
−0.14
+0.13 ±0.6 ±0.08
−0.03
+0.05
−13.1
−0.11
+0.14
−0.25
+0.15
field, which may be affected by non-thermal processes, such
as shocks. Adding to the complexity is the fact that the nitro-
gen abundance has to be treated as a free parameter, as many
early type stars show evidence of atmospheric abundance en-
hancements (e.g. Crowther et al. 2002; Bouret et al. 2003). In
our analysis method we solely model the hydrogen and helium
lines and self consistently allow for abundance enhancements
by treating the helium abundance as a free parameter. In prin-
ciple our temperature determination should, therefore, not be
affected by such problems. As a result of this, our analysis can
provide an independent confirmation for the hot nature of O2
stars.
In Fig. 1 we see that the objects with an O2 spectral type
indeed correspond to the hottest stars in our sample. With ex-
ception of the giant N11-031, which based on its helium lines
we find to be cooler (see below), they have temperatures in ex-
cess of 50 kK and, therefore, are significantly hotter than the
O3 star at Teff = 46 kK. The error bars, though, are consider-
able. Compared to the average error of ∼3 percent the O2 ef-
fective temperatures have an uncertainty of 5 up to 11 percent.
This large uncertainty can be explained by the weak or even
absent neutral helium lines in the spectra of these objects. As a
result of this, our fitting method has to predominately rely on
the He ii line profiles to determine the correct helium ionisation
equilibrium. Based on a single ionisation stage the determina-
tion of this equilibrium is more uncertain, which explains the
larger error bars. One should also be wary for a possible degen-
eracy effect between Teff and YHe that can occur as a result of
8 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
O2 O3 O4 O5 O6 O7 O8 O9 B0 B1
spectral type 
Fig. 1. Effective temperatures as a function of spectral type for
the investigated LMC sample. Different luminosity classes are
denoted using circles, triangles and squares for class V, III and
I-II objects, respectively. Shown as a dashed line is the Teff cal-
ibration of Martins et al. (2005a) for Galactic O-type dwarfs.
The dotted line corresponds to the average effective tempera-
ture of the SMC dwarfs studied in Paper II. The LMC dwarfs
are found to lie in between these two average scales, with a
typical ∼2 kK offset from both the average Galactic and SMC
relations.
the missing neutral helium lines (see e.g. Paper I). This, how-
ever, is not an issue as the helium abundances derived for all
O2 stars correspond to normal values close to 0.10.
It is also possible to question whether the fact that the He i
lines are so weak could result in a systematically higher Teff
determination by the automated method compared to “by eye”
fits. In other words, would a “by eye” fit prefer a lower tem-
perature solution? For the O2 giant N11-026 this is a relevant
question, as the weak He i λ4471 line, shown in Fig. A.1 in the
appendix, is not fitted perfectly. Note, however, that due to the
relative plot scale the discrepancy is exaggerated and that the
fit quality is good and is comparible to that of the He ii lines.
Still, to assess whether the fit of this particular line could be
improved, i.e. to search for a solution predominantly using the
He i diagnostic, we ran test fits with increasing relative weight
of He i λ4471. We found that a solution with an improved fit
could be obtained for an increase of the relative line weight
with a factor of five. This solution has a Teff lower by 3.7 kK
and all other fit parameters approximately equal to the solution
with the low He i λ4471 weight. The overal fit quality of the
other lines has somewhat decreased. As these lines are rela-
tively strong, this is difficult to discern by eye. Also note that
within the lower error estimate for Teff of 3.9 kK the two solu-
tions agree. Similar results were also obtained for the O2 stars
BI 237 and BI 253, where a reduction of Teff by 3.8 kK and
5.4 kK was found, respectively, for an increase in the relative
weight of He i λ4471 by a factor of six and two, respectively.
Consequently, based on our analysis we can only give a range
for the Teff scale of the O2 stars of 49 − 54 kK, the bounds
being set by the temperatures based on the He i λ4471 and hy-
drogen/He ii diagnostics.
The upper end of our hydrogen/helium based O2 effec-
tive temperature range is in good agreement compared to
the average effective temperature of 54 kK as determined by
Walborn et al. (2004). As we already mentioned the giant N11-
031 with a Teff lower by approximately 8 kK compared to the
other O2 stars forms an exception to this. Its effective tempera-
ture compares better to the temperature of the O3 star N11-060.
A close inspection of the spectra of the two stars reveals why
this is so. Both stars have exactly the same equivalent width
ratio of the He i λ4471 and He ii λ4541 lines, which results in
similar values for Teff . To also test whether the relatively strong
He i λ4471 line could be dominating the fit, forcing a relatively
low Teff, we refitted the spectrum of N11-031 ignoring the neu-
tral helium lines. This again resulted in an effective temperature
of 45 kK.
Interestingly, the nitrogen line analysis of Walborn et al.
for N11-031 did result in a higher effective temperature of
55 kK. We are not sure whether this discrepancy is the re-
sult of a systematic offset between the N v λ4603-20/N iv and
He i λ4471/He ii temperature scales. A more recent compari-
son by one of us (P.A.C.) of the spectral fit of Walborn et al. to
new VLT data, however, showed relatively large discrepancies
in the helium line fits compared to the nitrogen line fits, indi-
cating that this could be the case. The discrepancy could also
be due to differences in fitting assumptions and approaches.
Walborn et al. adopted a surface gravity of log g = 4.0 and
estimated a mass loss rate of Ṁ = 1.0 × 10−6 M⊙yr
−1 from
the He ii λ4686 line. Our fit indicates that the former param-
eter should be lower by 0.15 dex, consequently lowering Teff.
With respect to the mass loss rate we find a value higher by
a factor four. An increase in Ṁ of this magnitude can have a
significant effect on the strength of different nitrogen lines (e.g.
Crowther et al. 2002) and, therefore, on the derived effective
temperature.
Massey et al. (2005) analysed a total of 11 O2 stars. They
could determine the effective temperature for three dwarfs and
one giant, with 47.0 kK . Teff . 54.5 kK. For the supergiants
only lower limits of Teff & 42 kK and higher could be derived.
These results are in agreement with our findings. However,
these authors note that the correlation with Teff for the O2–3.5
spectral types is not tight. In particular no good agreement was
found between the ratios of the N iii and N iv emission lines
and the He i and He ii lines. Consequently, a more thorough
investigation of the O2 stars based on both the nitrogen and
helium spectrum is necessary to resolve this degenerate class
adequately.
4.3. Gravities
The distribution of our programme stars in the log Teff – log gc
plane is presented in Fig. 2. To calculate the surface gravity
corrected for centrifugal acceleration (log gc) the method dis-
cussed by Herrero et al. (1992) and Repolust et al. (2004) was
adopted. Different luminosity classes are denoted using circles,
triangles and squares for, respectively, class V, III and I-II ob-
jects. In this figure we see that the dwarfs, with exception of
two objects, form a group clearly separated from the latter two
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 9
4.40 4.50 4.60 4.70 4.80
log Teff
Fig. 2. The log Teff – log gc plane for the analysed LMC objects.
Different luminosity classes are denoted using circles, triangles
and squares for dwarfs, giants, and bright giants and super-
giants, respectively. With exception of two objects the dwarfs
are clearly separated from the luminosity class I-II objects. The
giants seem to overlap with both the luminosity class V and I-II
objects. Shown as a dotted line is the average log Teff – log gc
relation of the SMC O-type stars of luminosity class I-II-III
from Paper II.
groups. In contrast to our findings in Paper II we do not find
a clear separation between the dwarfs and giants. Instead the
latter group of objects shows an overlap with both luminosity
class V and I-II objects.
Shown in Fig. 2 as a dotted line is the average log Teff –
log gc relation of the SMC O-type stars of luminosity class I-
II-III from Paper II. The majority of the evolved LMC objects
seem to agree with this trend, that illustrates the evolutionary
correlation between effective temperature and surface gravity.
However, note that some evolved objects are at considerable
distance from the average relation. Thus, a calibration of the
two parameters for a given luminosity class should be taken
with care (see also Repolust et al. 2004).
The comparison of spectroscopically determined masses
(Ms) with masses obtained from evolutionary calculations is
presented in Fig. 3. Using the same symbols as in Fig. 2 dwarfs,
giant and supergiants are distinguished. Evolutionary masses
(Mev) were derived from the evolutionary tracks calculated for
a metallicity of Z = 0.4 Z⊙ from Schaerer et al. (1993). The er-
rors on these masses correspond to the maximum mass interval
in the error box spanned by the uncertainties in luminosity and
effective temperature.
The tracks from Schaerer et al. do not include the effects
of rotation. Consequently, this additional source of error is not
included. Calculations including vrot show that in some cases
one can no longer assign an unambiguous M(L, Teff). This is
a result of rotationally enhanced mixing or the unknown incli-
nation angle, if the star has a non-spherically symmetric dis-
tribution of R⋆ and Teff , causing complicated tracks including
loops during the secular redward evolution (Meynet & Maeder
2005). Assessing the impact of vrot on the Mev determination we
showed in Paper II from a comparison of masses derived from
 0  20  40  60  80  100  120
Spectroscopic Mass [Msun]
Fig. 3. Comparison of spectroscopic masses with masses de-
rived from the evolutionary tracks of Schaerer et al. (1993).
The one-to-one correlation between the two mass scales is
given by the dashed line. Symbols have the same meaning as
in Fig. 2. The objects with the highest evolutionary masses ex-
hibiting a mass discrepancy correspond to the four O2 stars in
our sample.
non rotating tracks to those obtained from tracks calculated for
vrot = 300 km s
−1 that the error in Mev will not increase by more
than approximately ten percent.
In Fig. 3 we see that the majority of the objects are lo-
cated left of the one-to-one correlation, given by the dashed
line. The error bars of twelve objects do not even touch
this correlation. Consequently, we find a significant mass dis-
crepancy. Similar mass discrepancy problems have been dis-
cussed by e.g. Herrero et al. (1992), de Koter & Vink (2003),
and Repolust et al. (2004). Most of these classical problems
were attributed to limitations in the stellar atmosphere mod-
els (Herrero 1993) and to potential biases in the fitting process
(see Paper I). Here we cannot explain the found discrepancy in
such a manner. We will provide a more thorough investigation
and discussion in Sect. 5.
4.4. Helium abundances
For the SMC sample that we analysed in Paper II we found a
correlation between the helium surface abundance and surface
gravity. This could partly be explained by evolutionary effects.
As the surface gravity decreases when a star evolves away from
the ZAMS, objects with lower gravities would correspond to
more evolved objects and are more likely to have atmospheres
enriched with helium.
To investigate whether the scenario discussed above also
applies to the current LMC sample, we plot the helium abun-
dance as determined with the automated method as a function
of log gc in Fig. 4. Also shown as a dashed line is a measure
of the initial helium abundance. This value of YHe = 0.09 was
calculated by averaging the surface helium abundances of the
dwarf type objects with a helium abundance smaller than the
total sample average. Compared to this measure for the ini-
tial helium abundance a correlation between the surface gravity
10 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
2.80 3.20 3.60 4.00 4.40
log gc
Fig. 4. Helium abundances as a function of surface gravity.
Symbols have the same meaning as in Fig. 2. The dashed line
at YHe = 0.09 is a measure for the initial helium abundance
and corresponds to the mean of the helium abundances of the
dwarfs with YHe smaller than the sample average. Starting at
the highest gravities an increase of the helium fraction is seen
down to log gc ≈ 3.6.
and helium enrichment can be observed. Starting at the high-
est gravities, we find an increase in the average helium abun-
dance towards lower log gc. Note that the two objects exhibit-
ing the largest helium fractions are supergiants. The increase
can again be partially explained as an evolutionary effect. In
Paper II a similar correlation between average helium abun-
dance and gravity was found to exist down to the lowest gravi-
ties investigated.
Interestingly, in Fig. 4 we see that for our LMC sample no
helium enrichment is found at log g . 3.6. Not even the super-
giants show evidence of enrichment below this gravity. Why is
this so? In Fig. 5, where we present the HR-diagram for our
sample, the answer to this question is given. Shown as a grey
area in this figure is the region in which the rotating evolu-
tionary models of Meynet & Maeder (2005) predict a helium
surface enrichment of at least ten percent. The majority of the
evolved objects are located outside this region. Consequently,
based on their specific evolutionary phase these objects are not
expected to show any enrichment. Also important is the fact
that we have selected the hottest objects from the N11 cluster.
As a result of this our sample is biased and does not contain
the low gravity objects with a high luminosity, i.e. those more
likely to be enriched, as these have evolved into cool B-type
stars.
Apart from the supergiants also four dwarfs are found to be
enriched. Figure 5, in which we have highlighted stars with a
helium abundance of at least 0.12 using open symbols, shows
that one these dwarfs, Sk −66 18, is located relatively close to
the region in which enrichment is predicted. Consequently, ro-
tationally enhanced mixing is a possible explanation. The three
remaining dwarfs, N11-60, N11-065 and Sk −70 69, in contrast
lie relatively close to the ZAMS. Therefore, “normal” mixing
cannot explain their enrichment. Instead, a possible explana-
tion is given by chemically homogeneous evolution. As this
4.24.34.44.54.64.74.8
log Teff
Fig. 5. Hertzsprung-Russell diagram for the LMC sample.
Symbols have the same meaning as in Fig. 2. Over plotted as
grey lines are the evolutionary tracks of Schaerer et al. (1993)
for Z = 0.4 Z⊙, with a black line representing the ZAMS. Open
symbols indicate objects with YHe ≥ 0.12. The grey area corre-
spond to the region in which the rotating evolutionary models
of Meynet & Maeder (2005) predict a relative helium enhance-
ment of at least 0.01.
may also be linked to the mass discrepancy, we will return to it
in Sect. 5.
4.5. Microturbulence
The microturbulent velocities determined with the automated
fitting method are also given in Tab. 2. Though the error es-
timates are relatively large (see Tab. 3), we find that for the
current data set the vturb determinations were sufficiently accu-
rate to reveal a weak correlation between this parameter and
the surface gravity. This is shown in Fig. 6, where it can be
seen that for log g . 3.6 the average microturbulence recovered
from the line profiles increases systematically. The situation for
log g & 3.6 is less clear, as the error bars are on average larger
and the values for vturb are more or less randomly distributed
between 0 and 20 km s−1. The reason for this is that for larger
values of the surface gravity the line profiles become intrinsi-
cally broader due to the increased Stark broadening, making it
more difficult to accurately recover vturb from the line profiles
only.
Based on samples predominantly consisting of unevolved
early B-type Galactic stars other authors have also found
a relation between microturbulence and surface gravity, e.g.
Kilian et al. (1991), Gies & Lambert (1992) and Daflon et al.
(2004). More recently Hunter et al. (2006) analysed a sample
of early B-type stars in the Magellanic clouds and also found
a trend of increasing vturb for decreasing log g. To derive the
values for vturb all these authors relied on curve-of-growth tech-
niques, which were applied to metal lines such as those of Si iii
and O iii calculated using plane parallel models. Consequently,
our line profile based analysis is an independent confirmation
of the existence (or requirement by lack of a physical explana-
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 11
2.8 3.2 3.6 4.0 4.4
log g [cm/s2]
Fig. 6. Microturbulent velocities determined using line profile
fits as a function of surface gravity. Symbols have the same
meaning as in Fig. 2. For log g . 3.6 a trend of increasing
vturb with decreasing log g is visible. For larger gravities the
uncertainties in the vturb determinations are too large to discern
any possible relation.
tion or failures in the line broadening mechanisms) of micro-
turbulence in the atmospheres of these type of stars.
The physical mechanism explaining the observed microtur-
bulence or its relation to the surface gravity is still poorly un-
derstood. Kudritzki (1992) and Lamers & Achmad (1994) have
argued that the observed microturbulence might be the result of
a stellar outflow, implying the existence of a velocity gradient
in the photospheric layers, which can mimic microturbulence-
like desaturation effects. Smith & Howarth (1998), however,
showed by applying a simple core-halo model to the Galactic
O9.7 supergiant HD 152003 that this effect would not be suffi-
cient to explain the observed vturb. Indeed, our analysis employ-
ing a unified photosphere and wind model confirms that the mi-
croturbulence cannot be explained as an artifact of a transonic
velocity field.
A possible explanation for the gravity dependence could
be related to instabilities in the wind. These instabilities are
reflected in the large turbulent velocities (∼100–200 km s−1)
necessary to fit the wind lines in the intermediate and outer
wind (e.g. Groenewegen & Lamers 1989; Haser et al. 1998;
Evans et al. 2004b) and could be related to shocks due to the in-
trinsic line-driven instability (Lucy 1983; Owocki et al. 1988;
Owocki & Puls 1999). For a low surface gravity the wind starts
at larger Rosseland optical depth compared to the high grav-
ity case. Therefore, the line forming region of low gravity ob-
jects contains a relatively large contribution originating from
(the base of) the wind. Consequently, this region could be af-
fected by the onset of the line-driven instability introducing
wind-turbulence into the line profiles.
Tentative support for the above described scenario may be
implied by Fig. 7, where we show vturb as function of the line
forming region of He i λ4471 in units of the stellar radius. The
location of the line forming region is defined as the position
at which the radial optical depth in the line core reaches a
value of 2/3. This figure shows that when the radial distance
1.00 1.02 1.04 1.06 1.08 1.10
R(τline=2/3) [Rstar]
Fig. 7. Microturbulent velocity as a function of the location of
the line forming region of He i λ4471, which is defined as the
location where the radial optical depth in the line core reaches
a value of τ = 2/3. A weak trend is visible that suggests that
for increasing extension of the atmosphere larger values of vturb
are necessary to fit the line profiles.
to this position increases also the average microturbulent ve-
locity increases. Though the trend is weak, it does seem to
indicate that when the atmosphere becomes more extended,
higher values of vturb are necessary to reproduce the line pro-
files. It may appear that this implies that the line forming re-
gion enters in to the regime where wind turbulence develops.
However, for this statement no compelling evidence is avail-
able, as we do not find any correlation between vturb and the
distance between the line forming region and (for instance) the
sonic point. Moreover, for most objects at R(τline = 2/3) > 1.04
values of vturb close to the maximum allowed value are found.
Consequently, these values should be interpreted as lower lim-
its. Therefore, we can only conclude that our analysis points
towards a gradient in the turbulent velocity, possibly connected
to a link between microturbulence and wind instabilities, and
suggest further investigation in this direction.
4.6. Wind parameters
The wind parameters and their uncertainties determined using
the automated method are listed in Tabs. 2 and 3. Compared
to our SMC analysis we find that we were able to accurately
determine these parameters for a significantly larger number
of objects. This is mainly the result of an on average higher
signal-to-noise ratio of the spectra as well as of denser winds
for the LMC objects compared to their SMC counterparts. In
total we determined 22 mass loss rates and 6 upper limits. The
upper limits are defined (and can be identified in Tab. 3) by an
error bar − log Ṁ > 1.0 dex.
To provide a meaningful comparison of the mass loss rates
we place the LMC objects in the modified wind momentum
luminosity diagram. This diagram shows as function of stellar
luminosity the distribution of the so-called modified wind mo-
mentum, which is defined as Dmom ≡ Ṁv∞R
⋆ . Not only does
this allow for an assessment of the behaviour of Ṁ within our
12 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
4.5 5.0 5.5 6.0 6.5
log(L/Lsun)
Fig. 8. Modified wind momentum (Dmom) in units of
g cm s−2 R⊙
1/2 vs. luminosity. Symbols have the same mean-
ing as in Fig. 2. Upper limits are shown as inverted triangles.
Dashed lines correspond to the predicted wind-momentum lu-
minosity relations (WLR) from Vink et al. (2000, 2001). The
upper, middle and lower relation, respectively, correspond to
predictions for Galactic, LMC and SMC metallicity. The ob-
served modified wind momenta show a strong correlation with
luminosity with an average relation that lies in between the pre-
dicted Galactic and SMC WLR. This is quantified by the em-
pirical WLR that was constructed for the LMC objects and that
is shown by the dotted line. The open square corresponds to
the wind momentum corrected for clumping of the supergiant
Sk −67 166. Shown as a dashed-dotted is the empirical WLR
obtained using this corrected Dmom.
sample, it also provides a convenient method to compare the
observed wind strengths to the predictions of line driven wind
theory. According to this theory Dmom is predicted to behave as
log Dmom = x log
L⋆/L⊙
+ log D◦ , (1)
where L⋆ is the stellar luminosity (Kudritzki et al. 1995;
Puls et al. 1996). In this equation x is the inverse of the slope of
the line-strength distribution function corrected for ionisation
effects (Puls et al. 2000). The vertical offset D◦ is a measure
for the effective number of lines contributing to the accelera-
tion of the outflow.
In Fig. 8 the distribution of the modified wind-momenta
for our programme stars are presented. Indicated using circles,
triangles and squares are objects of, respectively, luminosity
class V, III and I-II. Upper limits are shown as grey inverted
triangles. A clear correlation between L⋆ and Dmom can be ob-
served in this figure. Over an order of magnitude in L⋆ the av-
erage modified wind-momentum decreases by approximately
1.5 dex. A comparison of the behaviour of Dmom with pre-
dictions is facilitated by the theoretical WLRs calculated by
Vink et al. (2000, 2001) that are shown as a set of dashed lines.
The upper, middle and lower of these predicted power laws
were calculated for, respectively, Solar, LMC and SMC metal-
licity. Compared to these predictions we find that the LMC
wind-momenta are approximately confined between the theo-
retical WLR for Galactic and SMC metallicity. In other words,
compared to the Galactic and SMC case, stars in the LMC have
intermediate wind strengths.
To further quantify the behaviour of LMC winds relative to
that of Galactic and SMC outflows we have fitted a power law
to the observed modified wind-momentum distribution, consis-
tently accounting for both the symmetric errors in L⋆ and the
asymmetric errors in Dmom, This yielded the following empiri-
cal WLR
log Dmom = (1.81 ± 0.18) log
L⋆/L⊙
+ (18.67 ± 1.01) . (2)
The theoretical LMC relation from Vink et al. is given by
x = 1.83 and Dmom = 18.43. The error bars of the theoretical
and empirical relations are in agreement. More importantly, in
Fig. 8 the empirical relation, shown as a dotted line, is found to
lie between the predicted Galactic and SMC relations. These
predictions have been found to be in good agreement with
the observed Galactic WLR (Repolust et al. 2004, Paper I) and
observed SMC WLR (Paper II). Consequently, our empirical
LMC WLR is quantitative evidence for the fact that massive
stars in this system have mass loss rates intermediate between
those of massive stars in the Galaxy and SMC.
The differences between the empirical and theoretical LMC
WLR at the start and end of the observed luminosity range
are, respectively, 0.17 and 0.16 dex. In Fig. 8 this seems to
imply a systematic offset between the two relations. However,
we note that these differences are still smaller than the typi-
cal uncertainty in Dmom of 0.2 dex. More importantly, no cor-
rection was applied for the possibility that the winds of our
sample stars are, in contrast to our assumption, not smooth but
structured. In recent years evidence has been mounting for this
so-called clumping in the winds of O- and early B-type stars.
In particular spectroscopic modelling of UV (resonance) lines
(Crowther et al. 2002; Hillier et al. 2003; Bouret et al. 2003;
Massa et al. 2003; Martins et al. 2004, 2005b; Bouret et al.
2005; Fullerton et al. 2006) seems to suggest the existence of
clumping factors in the range of 10-100, implying correspond-
ing reductions of Ṁ by factors 3 up to 10. More recently,
Puls et al. (2006) also found clumping factors of the order 5
to 10 (normalized to the unknown clumping properties in the
outermost, radio-emitting wind) from the analysis of Hα, in-
frared, millimetre and radio fluxes. In the present study we try
to account for possible wind clumping effects by correcting the
mass loss rates of stars with Hα in emission. Markova et al.
(2004) and Repolust et al. (2004) argue that the mass loss rates
of these stars could be overestimated as a result of the fact that
Hα emission lines are formed over a relatively large volume
where clumping might have set in. In contrast, for stars with
Hα in absorption the line is formed relatively close to the stel-
lar surface, where clumping effects are negligible. Based on the
comparison of dwarfs and supergiants in their Galactic sample
Repolust et al. derived a numerical correction factor of 0.44 for
the mass loss rates of supergiants with Hα in emission.
We have applied the clumping correction to the super giant
Sk −67 166, which has a Hα emission profile. In Fig. 8 its new
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 13
wind momentum is indicated using an open symbol. Using this
value the following empirical WLR is obtained1
log Dmom = (1.43 ± 0.17) log
L⋆/L⊙
+ (20.77± 0.97) . (3)
In Fig. 8 we see that for log L⋆/L⊙ & 5.3 the new WLR com-
pares better to the Vink et al. relation. For lower luminosities
the situation is less clear. Due to the large uncertainties, this
range has a relatively low weight in the fit. Consequently, a
discrepancy for low L⋆ is less significant than the good agree-
ment obtained for the higher luminosities. For this reason it
is difficult to investigate the existence of a “weak wind prob-
lem” for stars at log L⋆/L⊙ . 5.3, first reported by Bouret et al.
(2003). This study, as well as later studies (Hillier et al. 2003;
Evans et al. 2004a; Martins et al. 2004, 2005b) report a steep-
ening in the WLR relation relative to predictions starting at
about the above mentioned L⋆, leading to an over prediction of
the wind strength by up to a factor 100 at log L⋆/L⊙ ∼ 4.5.
Our LMC results do not appear to confirm this break be-
tween observations and seem to follow the predictions down
to log L⋆/L⊙ ≈ 5.0. In a forthcoming paper we will present a
comprehensive overview of the observed WLR relations in our
Galaxy and the Magellanic Clouds as well as a thorough dis-
cussion of the successes and failures of the theory of radiation
driven winds in predicting the WLR, including possible causes
for the weak wind problem (Mokiem et al. in preparation).
5. The mass discrepancy
In Fig. 9 we investigate the mass discrepancy of our LMC tar-
gets as a function of the helium surface abundance. On the ver-
tical axis a measure for this discrepancy is shown, which is
scaled to the mean of the evolutionary and spectroscopic mass.
This ensures that positive and negative discrepancies follow
the same linear scale. For non-enriched stars, i.e. YHe . 0.10,
approximately three times as many objects lie above the one-
to-one correlation than below it. Therefore, in contrast to our
finding for the equivalent SMC case (Paper II) a significant
mass discrepancy is found for our sample of non-enriched LMC
stars. The reason for this is unclear. As we stated before our
analysis employs state-of-the-art atmosphere models, and is in
principle not hampered, as were previous studies, by poten-
tially unoptimised fits. As stated, in our SMC dataset, analysed
in an identical manner, no evidence was found for a mass dis-
crepancy for YHe < 0.10.
Massey et al. (2005) also study the Ms vs. Mev problem in
a set of LMC stars. For objects hotter than 45 kK they find
a mass discrepancy that is even stronger than what we find.
This behaviour could be the result of an underestimate of the
photospheric line pressure in this high temperature regime. In
Fig. 9 we have highlighted the objects with Teff ≥ 45 kK using
open symbols. Although their mass discrepancy is considerable
1 The small errors in the relevant parameters of Sk −67 166, which
dominates the WLR at very high luminosity, cause the significant dif-
ference between Eqs. 2 and 3. Note that our least square fitting does
not account for the correlation of both quantities (due to R⋆). In view
of the well known distance these effects are likely small compared to
the Galactic case (see Markova et al. 2004; Repolust et al. 2004).
-1.00
-0.50
0.05 0.10 0.15 0.20 0.25 0.30
Fig. 9. Mass discrepancy as a function of helium abundance
for the LMC sample. Evolutionary masses used to calculate
the discrepancy were derived from the non-rotating tracks of
Schaerer et al. (1993). Symbols have the same meaning as in
Fig. 2. The open symbols correspond to stars with Teff ≥ 45 kK.
they do not stand out as a separate group.To further illustrate
this, note that the hottest supergiant in our sample Sk −67 166
(Teff = 40 kK) at YHe = 0.28 has a Ms of 75 M⊙ which is
in very good agreement with its Mev of 70 M⊙. Consequently,
though we can not explain the mass discrepancy we do not an-
ticipate that it is connected to a flawed treatment of the photo-
spheric radiation pressure that manifests itself at metallicities
as high as that of the LMC environment (but not yet at values
typical for the SMC).
If we accept the analysis of stellar and photospheric pa-
rameters, the presence of a mass discrepancy may point to an
oversimplified picture of the evolution of massive stars used
to determine Mev and/or, possibly, to a breakdown of the as-
sumption of a spherically symmetric atmosphere. One obvious
simplification may be that we have used tracks for non-rotating
stars. One of the effects introduced by rotation is a wide bi-
furcation in the evolutionary tracks (Maeder 1987). Stars rotat-
ing faster than roughly half the surface break-up velocity will
essentially follow tracks representative of homogeneous evo-
lution. Langer (1992) showed that as a result of this the M/L-
ratio will be a monotonically decreasing function for increasing
helium enrichment. Consequently, stars evolving along homo-
geneous tracks are expected to be increasingly under-massive
for increasing helium abundance. One would therefore derive a
positive mass discrepancy if the evolution of a rapidly rotating
star was incorrectly described using non-rotating or modestly
rotating model tracks.
So, can rapid rotation be used to explain the observed mass
discrepancy? Let us first focus on the supergiants. For these
stars the mass discrepancy problem appears absent, though we
note two exceptions at YHe = 0.07 and YHe = 0.09. This ab-
sence of signs of a distinct mass discrepancy was also found
in our study of SMC bright giants and supergiants (Paper II).
The conclusion seems to be that the supergiants follow non-
rotating or modestly rotating evolutionary tracks. This implies
that, apparently, all supergiant targets that we have selected in
14 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
both the SMC (Paper II) and LMC happen to have started out
their evolution with low or modest initial rotational velocities.
As the initial rotational velocity distribution derived in Paper II
for the NGC 346 cluster implies that only some 5 − 15 percent
of stars are expected to evolve along homogeneous tracks this
should not be alarming. One supergiant, however, might not fit
this picture: Sk -67 166. This source does not feature a mass
discrepancy, but does show an increased helium abundance.
Given its position in the HRD this is quite difficult to explain in
term of evolution without rotation. A scenario accounting for a
more delicate interplay between mass loss and rotation may be
required to explain its properties (see Herrero & Lennon 2004).
Now let us turn to the dwarfs and giants. For our SMC
sample we found a correlation between the helium surface en-
richment and the mass discrepancy for class V and III stars,
with Ms being systematically smaller than Mev for YHe & 0.11.
Interestingly, this dependence was only found for dwarfs and
giants (see above). It was suggested that this behaviour was
the result of rotationally enhanced mixing, enriching the at-
mospheres with primary helium. The dwarfs and giants at
YHe > 0.11 displayed in Fig. 9 also systematically suffer from
a positive mass discrepancy. Consequently, this scenario also
seems to be a good explanation for the situation in the LMC
sample. However, we realise that these seven objects are a rel-
atively small sample and that the typical uncertainties in YHe
are 0.03. This statement should therefore be seen as a working
hypothesis.
For helium enriched dwarfs and giants the above scenario
seems a logical one, as helium enrichment early on in the evo-
lution must imply efficient mixing. For the dwarfs and giants
that do not show excess helium in their surface layers such
a clue or indication is not present (unless they evolve left of
the main sequence as do chemically homogeneous stars, see
Sect. 5.1). In principle these non YHe-enriched stars could ro-
tate sufficiently rapid to cause a mass discrepancy, but not to
the extent of chemically homogeneous evolution. However, if
so, a very significant fraction of the stars should have started
their evolution at super-critical rotation. At least for the stars
studied in the SMC cluster NGC 346 (Paper II) this is not the
case.
5.1. Chemically homogeneously evolving O2 stars
Three out of the four O2 stars are found to lie to the left of the
ZAMS in Fig. 5 and exhibit relatively large mass discrepan-
cies (see open symbols in Fig. 9). These are N11-026, BI 237
and BI 253. For a discussion of the effective temperature of the
fourth O2 object, N11-031 (Teff ∼ 45 kK), we refer to Sect. 4.1.
Even though the hottest three (Teff ∼ 53 − 54 kK) are not sig-
nificantly enriched in helium it is tempting to speculate on a
possible (near) homogeneous evolution of these objects. This
would not only provide an explanation for their mass discrep-
ancy, it would also explain their peculiar location in the HRD.
The possibility that these are true ZAMS stars is very exciting.
However, due to the short evolutionary time scales in this part
of the HR-diagram, it would also be unlikely. Based on the non-
rotating model tracks of Schaerer et al. (1993) we estimate that
a 60 M⊙ star will within 2 Myr evolve away from the ZAMS
to a location in the HRD at Teff ≈ 45 kK. This is well beyond
the error bars on the parameters of these objects. Given the fact
that N11-026 and BI 253 are thought to be associated with the
LH 10 and 30 Doradus clusters, respectively, which have ages
of approximately 3 (see Sect. 6.2) and 2 Myr (de Koter et al.
1998), a normal evolutionary scenario appears unlikely. Based
on its large radial velocity the field star BI 237 probably is a
runaway star (Massey et al. 2005), therefore it is also likely to
be relatively old.
The problems with a reconciliation of the hottest three O2
stars with fully homogeneous evolution and an age of at least
2 Myr, are i) that their surface helium abundances are not sig-
nificantly enriched, and ii) that their rotational velocities are
not extreme. The latter issue need not be “a smoking gun” con-
sidering the possibility that we may see them relatively pole-
on and the fact that their vr sin i values, ranging between 110-
190 km s−1, are above the sample average. Concerning the first
point, it appears that we must concede to the possibility that
stars may evolve along tracks similar to those for homogeneous
evolution while in fact they are not fully homogeneous – i.e.
the near surface layers are not yet strongly affected by the mix-
ing that must occur deeper in. Having said this, we do point to
the fact that the error bars on YHe do allow for relatively large
ages even within the hypothesis of fully homogeneous evolu-
tion. Within the error bars the maximum YHe is ∼ 0.13 for all
three stars. It takes fully homogeneously evolving stars of 60
and 40 M⊙ about ∼ 1.5 and ∼2.0 Myr, respectively, to build up
this amount of helium enrichment (S. -C. Yoon, private com-
munication). This is close to the derived cluster ages of LH 10
and 30 Doradus. For reference, if the star would evolve along
near-homogeneous tracks they would be older.
We tentatively conclude that the HRD position, helium
abundance, and rotational velocity of the three hottest O2 stars
in our sample could be consistent with (near-)homogeneous
evolution. Evidence for efficient rotation-induced mixing dur-
ing the main sequence phase of O stars has also been presented
by Lamers et al. (2001), based on the chemical abundance pat-
tern of ejected circumstellar nebulae.
5.2. Large sample trends
To firmly establish our findings with respect to the mass dis-
crepancy problems, we compare in Fig. 10 spectroscopic and
evolutionary masses for all stars that have been analysed using
our automated fitting method so far. Presented in this figure are
the combined Galactic (Paper I), SMC (Paper II) and the cur-
rent LMC samples, corresponding to a total of 71 O- and early
B-type stars.
The figure confirms our two main findings. First, the good
agreement between Ms and Mev for the supergiants is clearly
visible. This is an encouraging finding and we believe that this
means that the improvements in the stellar atmosphere mod-
els, evolutionary calculations and spectral analysis techniques
have finally resolved the long standing mass discrepancy as
found by Herrero et al. (1992). The correspondence between
the spectroscopic and evolutionary mass scale, in particular for
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 15
Fig. 11. The FLAMES field for the star forming region N11 with the associations LH9 (south of centre) and LH10 (north of
centre). The FLAMES targets are identified using circles and their star identification is given. From Evans et al. (2006). [Included
as separate JPG file in astro-ph submission.]
-1.00
-0.50
0.05 0.10 0.15 0.20 0.25 0.30
Fig. 10. Mass discrepancy as a function of helium abundance
for the combined samples from Paper I (Galactic), Paper II
(SMC) and the current LMC sample. Symbols have the same
meaning as in Fig. 2. Note the good correspondence between
Ms and Mev for the bright giants and supergiants (square sym-
bols).
increasing helium enrichment is striking. The reason for this is
probably related to the fact that the class I-II objects are found
in the region of the HR-diagram in which stars are not expected
to undergo extreme evolutionary phases, such as homogeneous
evolution. Consequently, the enriched bright giants and super-
giants have evolved along relatively simple evolutionary tracks,
which (in this respect) appear well understood.
Our second finding, i.e. that of a correlation of mass dis-
crepancy with the helium abundance, is also corroborated by
the large sample in Fig. 10. For YHe > 0.10 all dwarfs and gi-
ants, except for one dwarf, are found above the Ms = Mev line.
In all fairness, given the typical uncertainty of 0.03 in YHe, from
a statistical point of view the region 0.09 < YHe < 0.12 should
be regarded with care. For larger helium abundance, however,
the correlation can be regarded to be statistically significant. A
total of 11 objects show a positive mass discrepancy, with only
a single counter example. Moreover, the magnitude of the dis-
crepancy seems to be related to the amount of enrichment. This
strongly points to efficient mixing in the main sequence phase,
leading to (near-)chemically homogeneous evolution.
6. The evolutionary status of N11
In this section we explore the evolutionary status of the N11
field. We will first briefly outline the current understanding of
N11 with respect to the OB associations LH9 and LH10 and
in particular the sequential evolutionary link between the two.
Based on the analysis of the 22 stars in our sample associated
with these clusters we will then estimate their ages and discuss
whether they are compatible with a sequential star formation
scenario.
6.1. LH9 and LH10 in N11
N11 is an intricate giant H ii region containing several massive
star forming complexes. The largest of these are the OB associ-
ations LH9 and LH10 (Lucke & Hodge 1970). An image of the
region, including the adopted star identifications, is presented
in Fig. 11 (Evans et al. 2006). Several studies have suggested
that the star forming activity in these associations are linked.
In particular the study by Parker et al. (1992) has provided a
key to understanding of the structure and formation of LH9
and LH10. Their analysis of the stellar content of N11 revealed
the presence of several O3-O5 stars and possible ZAMS stars
in LH10. In contrast the earliest spectral type associated with
LH9 was found to be O6. They also found the slope of the ini-
tial mass function of LH10 to be significantly flatter than that
of LH9, indicating that the former contains a higher ratio of
high mass to low mass stars. Combined with the fact that the
reddening of LH10 is larger than that of LH9, it was concluded
that LH10 is the younger of the two clusters. Based on these
findings the authors propose an evolutionary link between the
two, where star formation in LH10 could possibly be triggered
by the stellar winds and supernovae of massive stars in LH9.
Further evidence for a sequential star formation sce-
nario was given by Walborn & Parker (1992), who found a
dual structural morphology of N11 analogous to that of the
30 Doradus H ii region. In the latter substantial evidence sug-
gests the presence of current star formation in regions sur-
rounding the central star cluster (e.g. Walborn & Blades 1987;
Hyland et al. 1992, also see Walborn & Blades 1997). This sec-
ondary burst seems to be set off approximately 2 Myr after the
initial star formation took place, possibly initiated by the en-
ergetic activity of the evolving cluster core. Walborn & Parker
argue that this is very similar to N11 where the stellar content of
LH9 and LH10 also suggests an age difference of ∼2 Myr (also
see Walborn et al. 1999). The process in N11, however, would
be advanced by∼2 Myr, classifying it as an evolved 30 Doradus
analogue, though less massive.
More recently, several bright IR sources showing charac-
teristics of young stellar objects were discovered in the N11B
nebula surrounding LH10 by Barbá et al. (2003). These objects
are probably intermediate mass (pre-)main-sequence Herbig
Ae/Be stars belonging to the same generation as do the LH10
objects (the pre-main-sequence evolutionary timescales of in-
termediate mass stars being longer than that of massive OB
stars). Barbá et al. also found that the massive stars in LH10
have blown away their ambient molecular material and are cur-
rently disrupting the surface of the parental molecular cloud
material surrounding LH10 (i.e. the material in N11B).
16 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
4.44.54.64.7
log Teff
65 66
Fig. 12. Comparison of the programme stars located in N11 with isochrones derived from the evolutionary tracks of Schaerer et al.
(1993). The symbols used to denote dwarf, giants and bright- and supergiants are, respectively, circles, triangles and squares.
Filled and open symbols are used to distinguish stars associated with LH9 and LH10, respectively. Black symbols refer to objects
located within a radius of two arc minutes of the LH9 and LH10 core, whereas the grey symbols correspond to stars found outside
this radius. Labels correspond to the N11 identification numbers given in Tab. 1. Isochrones (dashed lines) are shown from one
up to ten million years with 1 Myr intervals. For reference evolutionary tracks from Schaerer et al. (1993) are drawn as grey lines,
where for clarity purpose only the blueward part of the evolution for the most massive tracks is shown. The ZAMS corresponding
to these tracks is given by the black solid line.
6.2. The ages of LH9 and LH10
To estimate the ages of our programme stars we compare their
location in the HR-diagram to theoretical isochrones in Fig. 12.
The luminosity classes V, III and I-II are represented using, re-
spectively, circles, triangles and squares. To differentiate be-
tween stars associated with LH9 and LH10, respectively, filled
and open symbols are used. Membership is defined on the ba-
sis of minimum distance to either the LH9 or LH10 cluster
core. This, therefore, should be taken with some care: the cores
are about 4 arcmin or 60 parsec apart, which can be traversed
in 2 Myr if the proper motion of the star is some 30 km s−1.
Runaway O- and B-type stars have typical velocities of 50–
100 km s−1, therefore, it is entirely possibly that (a few) indi-
vidual objects are assigned to the wrong association. For posi-
tion reference see Fig. 11. Isochrones in Fig. 12 are shown as
dashed lines for one up to ten million years with 1 Myr intervals
and were derived from the evolutionary tracks of Schaerer et al.
(1993). Note that these tracks do not account for the effects of
rotation.
The distribution of the stars in Fig. 12 is such that they can
be separated in a group of objects younger than three million
years and a group of objects older than this age. We also see that
the oldest objects are predominantly found in LH9, whereas
LH10 contains the largest fraction of young objects, suggest-
ing that the clusters can indeed be separated in terms of age. At
first sight it seems that they can be characterised by an age of
∼7 Myr and ∼2 Myr, respectively. However, a significant num-
ber of objects in both clusters are found to be several million
years younger and older than these preliminary ages. A possi-
ble explanation for this large age scatter could be “contamina-
tion” by field stars. To assess this possibility we have assigned
grey symbols to all objects outside a radius of two arc minutes
from the cluster cores. Disregarding these objects reduces the
age scatter significantly; still a number of stars appear to con-
tradict with the notion of two coeval populations.
To investigate the age distributions in more detail the indi-
vidual age estimates are shown in Fig. 13. LH9 and LH10 ob-
jects, which are shown using the identical symbols as in Fig. 12,
are placed in, respectively, the left and right part of the diagram
and are separated using a dashed line. First concentrating on the
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 17
8 29 33 36 42 45 58 65 66 123 4 26 31 32 38 48 51 60 61 68 72 87
N11-#
Fig. 13. Individual age estimates for stars in N11 based on the
non-rotating tracks of Schaerer et al. (1993). Symbols have the
same meaning as in Fig. 12. The left and right part of the
diagram, respectively, contain stars associated with LH9 and
LH10 and are separated using a dashed line. The upward point-
ing arrow denotes a lower age limit for N11-065 determined
from its surface helium abundance using chemically homoge-
neous evolutionary models (Yoon et al. 2006; Yoon & Langer
2005).
LH9 objects we see that they are near-coeval with exception of
four dwarfs. Of these the two non-core members are located
at a distance of approximately six arc minutes from the central
concentration. Therefore, it is probable that they are either spa-
tially not related to N11 or that possibly their formation was
triggered more recently by the stellar activity in LH9. The third
dwarf N11-065 occupies a location very close to the ZAMS and
in principle would be the youngest member of LH9. However,
we also find a considerable helium enrichment for this object
(YHe= 0.17). Combined with the fact that N11-065 has a large
mass discrepancy, this suggests that it might be evolving chem-
ically homogeneously. Consequently, a more appropriate age
estimate should be derived from tracks appropriate for this kind
of evolution. Adopting such tracks from Yoon et al. (2006, also
see Yoon & Langer 2005) we derive a lower limit of 6 Myr for
the age of N11-065 based on the surface helium abundance.
In Fig. 13 this estimate is indicated using an upward pointing
arrow, and is in good agreement with the bulk of the LH9 stars.
The dwarf N11-123 exhibits no chemical peculiarities and
its location on the sky places it in the central concentration of
LH9, suggesting that it should have been formed in the burst
of star formation that formed the cluster. However, given the
fact that N11-123 is in principle the only star deviating from
the coeval nature of LH9 we suspect that it is like N11-058
and N11-066 part of the periphery of the cluster and that its
relatively close position to the core is due to a projection ef-
fect. Consequently, we conclude that the central concentration
of LH9 is coeval with an age of ∼7.0 ±1.0 Myr. The error bar
is of the same order of magnitude as the 0.5–1.0 Myr intro-
duced by uncertainties in the evolutionary tracks due to effects
of (relatively modest) rotation (see Paper II).
The right part of Fig. 13 shows that the stars in the cen-
tral concentration of LH10 have ages ranging from one up to
approximately six million years. Despite this large scatter it is
clear that the majority of the stars are younger than 4.5 Myr
and that only one object (N11-087) in the cluster core is older
than this age. An explanation for the large age of the latter ob-
ject might be that it formed in LH9 and, over time, migrated
towards LH10. As explained at the start of this section, this
is a possibility given that the distance from the centre of LH9
to the current position of N11-087 is only three arc minutes.
Considering the error bars on the age determinations and the
additional uncertainty of 0.5–1.0 Myr introduced by rotation
(see Paper II), we finally estimate an age best describing LH10
of ∼3.0 ±1.0 Myr.
6.3. Sequential star formation?
If the starbursts in N11 are sequential the formation of LH10
should be induced by supernova explosions and/or stellar winds
in LH9. Given the age difference and distance between the
two clusters we find that both means of triggering are possi-
ble. In the first case the age difference of approximately four
million years is compatible with the time of ∼3 Myr it takes
before the most massive stars end their life in a supernova (e.g.
Schaller et al. 1992). The time needed for the supernova shock
to cross the distance of∼3 arc minutes between LH9 and LH10,
corresponding to approximately 40 parsec at the distance of
the LMC, is only ∼105 years (see e.g. Falle 1981); hence it
is a possible scenario. If we consider triggering through stellar
winds, the time scales are also compatible. Garcia-Segura et al.
(1996), for instance, using hydrodynamical simulations have
shown that a 60 M⊙ star can create a wind-driven bubble in the
interstellar medium of ∼50 pc during its main sequence lifetime
of 3 Myr. Consequently, this scenario seems to be appropriate
as well and in agreement with the determined time scales.
In view of the above we conclude that a sequential scenario
for LH9 and LH10 seems very likely. A combination of super-
novae and stellar winds from stars in LH9 may have initiated
star formation in LH10.
7. Summary and conclusions
We have analysed a sample of 28 massive OB-type stars lo-
cated in the LMC. For a homogeneous and consistent treatment
of the data we employed the automated fitting method devel-
oped by Mokiem et al. (2006), which combines the genetic al-
gorithm based optimisation routine pikaia (Charbonneau 1995)
with the fast non-LTE unified stellar atmosphere code fastwind
(Puls et al. 2005). The sample is mostly drawn from the targets
observed within the context of the VLT-FLAMES survey of
massive stars (Evans et al. 2005). In total 22 of these stars are
located in the LH9 and LH10 clusters within the giant H ii re-
gion N11. This region is believed to have been the scene of
sequential star formation, with the stellar activity in LH9 ignit-
ing secondary starbursts in different associations in N11. Our
main findings are summarised below.
18 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
i) The effective temperature per spectral sub-type of the LMC
stars is found to be intermediate between that of Galactic
and SMC O- and early B-type stars, with the LMC objects
being, respectively, cooler and hotter by typically ∼2 kK.
ii) Based on the helium and hydrogen lines it was possible to
determine the effective temperatures, though with relatively
large error bars, of the four O2 stars in our sample. Three
of these are found to be hotter by more than 3 to 7 kK (see
Sect. 4.2) compared to the O3 star in our sample, suggesting
that O2 stars indeed represent a hotter subgroup within the
O-type class. However, we note that for the one ON2 star
a relatively low Teff of 45 kK was obtained, indicating that
the N iv and N iii classification lines (Walborn et al. 2002b)
are not fully compatible with the helium lines traditionally
used for classification.
iii) The spectroscopically determined masses of the dwarf and
giant stars in our set of programme stars are found to be
systematically smaller than those derived from non-rotating
evolutionary tracks. For helium enriched dwarfs and giants,
i.e. those having YHe > 0.11, we find that all show this mass
discrepancy. The same was found in an analysis of SMC
stars using the same methods (Mokiem et al. 2006). We in-
terpret this as evidence for efficient rotationally enhanced
mixing leading to the surfacing of primary helium and to
an increase of the stellar luminosity.
iv) The bright giants and supergiants do not show any mass
discrepancy, regardless of the surface helium abundance.
This also is consistent with the finding for Galactic and
SMC class I-II objects studied with the same methodology
(Mokiem et al. 2005, 2006). This suggests that shortly after
birth all these stars must have rotated at less than about 30
to 40 percent of the surface break-up velocity.
v) A weak correlation is found between microturbulent veloc-
ity and surface gravity. More extended atmospheres (i.e.
lower gravity stars) require a relatively large vturb to fit the
lines. The reason for this relation is unclear, however, it
does not seem to be connected to the lines being formed
closer to the sonic point of the wind flow in low gravity
stars.
vi) From a comparison of modified wind momenta Dmom we
find that the wind strengths of LMC stars are weaker com-
pared to Galactic stars, and stronger compared to SMC
stars. Comparing the derived Dmom as a function of lumi-
nosity with predictions for LMC metallicities by Vink et al.
(2001) yields good agreement in the entire luminosity range
that was investigated (5.0 < L/L⊙ < 6.1).
vii) We have determined an age of, respectively, ∼7.0±1.0 Myr
and ∼3.0±1.0 Myr for the clusters LH9 and LH10. The age
difference and relative distances are in good agreement with
a sequential star formation scenario, in which stellar activ-
ity in LH9 triggered the formation of LH10.
Acknowledgements. We would like to thank Sung-Chul Yoon and
Alexander van der Horst for constructive discussions, and George
Meynet for providing us with the evolutionary models of the Geneva
group. M.R.M. acknowledges financial support from the NWO
Council for Physical Sciences. S.J.S. acknowledges the European
Heads of Research Councils and European Science Foundation
EURYI (European Young Investigator) Awards scheme, supported by
funds from the Participating Organisations of EURYI and the EC Sixth
Framework Programme. A.H. and F.N. acknowledge support from the
Spanish MEC through project AYA2004-08271-C02. JSV acknowl-
edges RCUK for his Fellowship. Spectral fits were calculated us-
ing the LISA computer cluster at SARA Computing & Networking
Services.
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20 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
Appendix A: Fits and comments on individual
objects
The observed spectra shown in this section were corrected for
radial velocities. If not noted differently the lines that were fit-
ted are the hydrogen Balmer lines Hα, Hγ and Hδ; the He i
singlet line at 4387 Å; the He i triplet lines at 4026, 4471 and
4713 Å; and the He ii lines at 4200, 4541 and 4686 Å. Over
plotted are the best fit spectra, unless noted differently. For a
discussion of the line weighting scheme adopted in our fitting
procedure we refer to Paper I (especially to their Table 3.).
N11-004 (Fig. A.1) With exception of He i λ4471 good fits
for all lines of this OC9.7 supergiant were obtained. The under
prediction of the He i λ4471 line is possibly connected to the
so-called generalised dilution effect (Voels et al. 1989, see also
Repolust et al. 2004).
N11-008 (Fig. A.1) All lines are reproduced well. Note that
given the weakness of He ii λ4686 and lower weight assigned
to this line, the reproduction of its profile can be considered
good.
N11-026 (Fig. A.1) A high effective temperature of 53 kK
was needed to fit the spectrum of this O2 giant. Note that even
though on its relative plot scale the very weak He i λ4471 line
is not fitted perfectly, the quantitative fit quality is good and
is comparible to that of the He ii lines. Still, trying to assess
whether the fit could be improved we ran test fits with increas-
ing relative weight of the He i λ4471 line. It turned out that an
increase of a factor of five was necessary to obtain a solution
with an improved He i λ4471 line fit, which is shown in Fig. A.1
as a set of dotted lines. This best fit has a Teff lower by 3.7 kK
and all other fit parameters approximately equal compared to
the regular fit. The overal fit quality of the other diagnostic lines
is somewhat less, though this is difficult to discern by eye be-
cause these lines are relatively strong (one has to zoom in to the
0.5 to 1 percent level, as we do for He i λ4471). Note, however,
that the error estimates for Teff are relative large and within the
lower error estimate of 3.9 kK (cf. Tab. 3) the solutions agree.
Also note that even for the lower Teff solution this O2 star is
still considerably hotter than an O3 star (see Sect. 4.1).
N11-029 (Fig. A.1) A good fit was obtained. The slight un-
derestimation of the core strength of the He ii λ4200 and
He ii λ4541 lines is, given the relative weakness of these lines,
not significant.
N11-031 (Fig. A.2) To fit the spectrum of this O2 star a
relatively low effective temperature of 45 kK was necessary.
Walborn et al. (2004) derived a Teff of 55 kK for this object.
However, their determination of this parameter was based on
the analysis of the N iv-v ionisation adopting a fixed gravity of
log g = 4.0. Based on the analysis of the helium ionisation we
exclude a Teff higher than ∼47 kK. This is illustrated by the dot-
ted lines, which correspond to a model calculated for an effec-
tive temperature higher by the upper error estimate we derived
for Teff. For even higher temperatures the He i λ4471 would, in
contradiction to the observations, disappear completely (also
see Sect. 4.1).
N11-032 (Fig. A.2) The He i and He ii blend at 4026 Å was
not observed for this object. In the final fit, therefore, also the
Hδ Balmer line is shown.
N11-033 (Fig. A.2) This B0 giant is fast rotator. To repro-
duce the observed line profiles a projected rotational velocity
of 256 km s−1 was necessary. In the presented fit the strength
of He ii λ4541 seems to be under predicted. However, given the
relative weakness of this line this small discrepancy is negligi-
ble. Note that the sharp change in the central part of the Hα
profile is the result of an over subtraction of the core nebular
feature and was not taken into account in the fitting procedure.
N11-036 (Fig. A.2) The line profiles of N11-036 could be re-
produced quite accurately because of its relatively slow rota-
tion.
N11-038 (Fig. A.3) The relatively poor reproduction of
He i λ4387 in this O5 bright giant is the result of the low weight
assigned to this line based on the spectral type of N11-038. A
good fit to the He i λ4471 could not be obtained due to its pecu-
liar line profile shape. Possibly this triangular profile is related
to macroturbulence.
N11-042 (Fig. A.3) A relatively low projected rotational ve-
locity was found for this B0 giant allowing for a nearly perfect
N11-045 (Fig. A.3) A good fit was obtained for this O9 giant.
Despite the good fit quality a problem is apparent when Ms and
Mev are compared in Tab. 2. The spectroscopic mass is found
to be less than half its evolutionary equivalent.
N11-048 (Fig. A.3) Based on the fact that the He ii λ4686 is
the strongest line in its spectrum Parker et al. (1992) classified
this star as a Vz star. The fit parameters, however, place this
object at a considerable distance from the theoretical ZAMS in
Fig. 5. We also find that the spectrum can not be reproduced ac-
curately. The width of the neutral helium lines is systematically
unpredicted, whereas the width of the He ii is over predicted.
This and the fact that the helium abundance that is recovered
from the spectrum is rather low (YHe = 0.06) are indicative of
a possible binary nature.
N11-051 (Fig. A.4) The spectrum of the O5 dwarf N11-051
shows very broad lines, indicating fast stellar rotation. To ob-
tain the final fit a projected rotational velocity of 333 km s−1
was required.
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 21
N11-004   OC9.7 Ib
 4330  4340  4350
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4192  4200  4208
HeII 4200
 4537  4541  4545
HeII 4541
 4683  4686  4689
HeII 4686
N11-008   B0.7 Ia
 4335  4340  4345
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4192  4200  4208
HeII 4200
 4537  4541  4545
HeII 4541
 4684  4686  4688
HeII 4686
N11-026   O2 III(f*)
 4335  4340  4345
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4681  4686  4691
HeII 4686
N11-029   O9.7 Ib
 4335  4340  4345
 6552  6562  6572
 4023  4026  4029
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4194  4200  4206
HeII 4200
 4537  4541  4545
HeII 4541
 4684  4686  4688
HeII 4686
Fig. A.1. Comparison of the observed line profiles of N11-004, -008, -026 and -028 with best fitting synthetic line profiles
obtained using the automated fitting method (grey lines). Wavelengths are given on the horizontal axis in Å. The vertical axis
gives the normalised flux. Note that this axis is scaled differently for each line. The dotted line profiles for N11-026 correspond
to the best fit obtained adopting a five times larger relative weight for the He i λ4471 line. This best fit has a Teff lower by 3.7 kK
and all other fit parameters approximately equal compared to the fit for the nominal He i λ4471 weight.
N11-058 (Fig. A.4) A good fit was obtained for all lines of
this O5.5 dwarf. The small under prediction of the He i4387
line strength is the result of the relatively low weight assigned
to this line for the spectral type of this star.
N11-060 (Fig. A.4) Note that the He i λ4471 line of this O3
dwarf has a strength that is comparable to the strength of this
line in the spectrum of the O2 giant N11-031. As the effective
temperature we find for these two objects are also compara-
ble, the helium spectrum suggest a spectral type of O3 for both
objects.
N11-061 (Fig. A.4) All lines of this O9 dwarf are reproduced
correctly. The mass loss could be reliably determined using the
automated method at a rate of 2.1 × 10−7 M⊙. Note that the
22 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
N11-031   ON2 III(f*)
 4333  4340  4347
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4465  4471  4477
HeI 4471
 4710  4713  4716
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4680  4686  4692
HeII 4686
N11-032   O7 II(f)
 4096  4102  4108
 4335  4340  4345
 6547  6562  6577
 4386  4388  4390
HeI 4387
 4469  4471  4473
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-033   B0 IIIn
 4330  4340  4350
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4384  4388  4392
HeI 4387
 4467  4471  4475
HeI 4471
 4709  4713  4717
HeI 4713
 4192  4200  4208
HeII 4200
 4537  4541  4545
HeII 4541
 4681  4686  4691
HeII 4686
N11-036   B0.5 Ib
 4330  4340  4350
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4192  4200  4208
HeII 4200
 4537  4541  4545
HeII 4541
 4684  4686  4688
HeII 4686
Fig. A.2. Same as Fig. A.1, however, for N11-031, -032, -033 and -036. Shown as dotted profiles for N11-031 is the effect of and
a 2.2 kK increase in Teff .
value of 1.8 obtained for the wind acceleration parameter β is
rather high for a dwarf type object.
N11-065 (Fig. A.5) A relatively high helium abundance of
0.17 was recovered from the spectrum. The good fit quality
obtained for all lines, however, does not indicate an overesti-
mation of this parameter. A strong mass discrepancy is found
with Ms being smaller than Mev by approximately a factor of
N11-066 (Fig. A.5) Apart from a slight under prediction of the
width of He ii λ4686 all lines of this O7 dwarf could be fitted
with good accuracy. Like the previously discussed object, again
a large mass discrepancy is found with a ratio of Mev to Ms of
N11-068 (Fig. A.5) The final fit reproduces the observed line
profiles very accurately. No further comments are required.
N11-072 (Fig. A.5) The star has a very sharp lined spectrum.
The best fit required a vr sin i of 14 km s
−1. Note that given the
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 23
N11-038   O5 II(f+)
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4466  4471  4476
HeI 4471
 4709  4713  4717
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-042   B0 III
 4330  4340  4350
 6547  6562  6577
 4023  4026  4029
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4712  4713  4714
HeI 4713
 4192  4200  4208
HeII 4200
 4536  4541  4546
HeII 4541
 4684  4686  4688
HeII 4686
N11-045   O9 III
 4335  4340  4345
 6552  6562  6572
 4022  4026  4030
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4683  4686  4689
HeII 4686
N11-048   O6.5 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4465  4471  4477
HeI 4471
 4710  4713  4716
HeI 4713
 4194  4200  4206
HeII 4200
 4535  4541  4547
HeII 4541
 4682  4686  4690
HeII 4686
Fig. A.3. Same as Fig. A.1, however, for N11-038, -042, -045 and -048.
lower error estimate, given in Tab. 3, and the spectral resolution
of the data this value can also be interpreted as an upper limit.
N11-087 (Fig. A.6) To fit the spectrum of N11-087 a vr sin i
of 276 km s−1 was needed. The final fit shows a good repro-
duction of the line profiles with a slight under prediction of the
He ii λ4541 line. Given the signal-to-noise ratio of the spectrum
and the relative strength of this line, we do not believe this to
be significant.
N11-123 (Fig. A.6) Like the previously discussed object N11-
123 is also classified as an O9.5 dwarf. The effective tempera-
ture needed to fit the spectrum is, however, higher by more than
2 kK. This is the result of a log g that is higher by ∼0.2 dex.
BI 237 (Fig. A.7) A good fit was obtained for this O2 dwarf
star. The He i λ4471 line is hardly visible. Fortunately, the hy-
drogen and He ii lines provide the automated fitting method
with enough information to determine an effective temperature.
The error estimates on this parameter, however, are consider-
able (see Tab. 3). The effect of lowering Teff to the lower error
estimate of 49 kK is illustrated by the dotted profiles.
BI 237 was recently also analysed by Massey et al. (2005).
Compared to this study we find large differences for Teff
24 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
N11-051   O5 Vn((f))
 4330  4340  4350
 6547  6562  6577
 4020  4026  4032
HeI + HeII 4026
 4382  4388  4394
HeI 4387
 4466  4471  4476
HeI 4471
 4706  4713  4720
HeI 4713
 4192  4200  4208
HeII 4200
 4533  4541  4549
HeII 4541
 4678  4686  4694
HeII 4686
N11-058   O5.5 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4468  4471  4474
HeI 4471
 4709  4713  4717
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-060   O3 V((f*))
 4333  4340  4347
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4466  4471  4476
HeI 4471
 4710  4713  4716
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4680  4686  4692
HeII 4686
N11-061   O9 V
 4335  4340  4345
 6547  6562  6577
 4023  4026  4029
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4194  4200  4206
HeII 4200
 4536  4541  4546
HeII 4541
 4683  4686  4689
HeII 4686
Fig. A.4. Same as Fig. A.1, however, for N11-051, -058, -060 and -061.
(+5 kK), log g (+0.2 dex) and Ṁ (−0.4 dex). Our solution
for higher effective temperature and surface gravity can be ex-
plained by the fact that the ionisation structure of the atmo-
sphere is set by both these parameters (also see Paper I). The
reduced mass loss rate we find is the result of the increased
value for β (+0.4) obtained by the automated method.
BI 253 (Fig. A.7) The star has as an identical spectral type
as the previously discussed object. The photospheric parame-
ters determined from the spectrum are also very similar to the
parameters of BI 237. Note that in contrast to this agreement
BI 253 is found to have a much denser wind. Its mass loss rate
is larger by a factor 2.5 and reflects the fact that compared to
BI 237 the wind lines Hα and He ii λ4686 are more filled in.
In their analysis Massey et al. (2005) could only determine a
lower estimate for Teff of 48 kK, which is consistent with our
lower error estimate. The other parameters determined by these
authors are in fair agreement with our analysis. An exception
to this is the surface gravity for which Massey et al. estimate
a value lower by 0.3 dex, which is the result of the lower Teff
used for their fit.
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 25
N11-065   O6.5 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-066   O7 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-068   O7 V((f))
 4330  4340  4350
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
N11-072   B0.2 III
 4335  4340  4345
 6547  6562  6577
 4024  4026  4028
HeI + HeII 4026
 4386  4388  4390
HeI 4387
 4468  4471  4474
HeI 4471
 4712  4713  4714
HeI 4713
 4192  4200  4208
HeII 4200
 4536  4541  4546
HeII 4541
 4684  4686  4688
HeII 4686
Fig. A.5. Same as Fig. A.1, however, for N11-065, -066, -068 and -072.
Sk −66 18 (Fig. A.7) With exception of the He ii lines at
4541 Å and 4686 Å we obtained a near perfect fit. The helium
abundance YHe = 0.14 indicates an enriched atmosphere.
Sk −66 100 (Fig. A.7) was previously analysed by Puls et al.
(1996). Compared to this study we find similar parameters with
exception of a lower Teff which is the result of the inclusion of
blanketing and an increased helium abundance of YHe = 0.19.
The final fit shows that with this large value all lines could
be reproduced quite accurately. The slight mismatch of the
He ii λ4686 line can be improved by increasing Ṁ to its up-
per error estimate. The effect of this increase of 0.12 dex on the
line profiles is shown using dotted lines.
Sk −67 166 (Fig. A.8) This supergiant has both Hα and
He ii λ4686 in emission indicating the presence of a dense stel-
lar outflow. To correctly reproduce the line profiles a very high
mass loss rate of 9.3×10−6 M⊙yr
−1 was required. This value is
in good agreement with the findings of Crowther et al. (2002)
who studied this objects using the model atmosphere code cm-
fgen (Hillier & Miller 1998) in the optical, UV and far-UV.
Compared to these authors we also find that the other fit pa-
rameters are in good agreement. An exception to this is the he-
26 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
N11-087   O9.5 Vn
 4325  4340  4355
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4465  4471  4477
HeI 4471
 4707  4713  4719
HeI 4713
 4192  4200  4208
HeII 4200
 4534  4541  4548
HeII 4541
 4680  4686  4692
HeII 4686
N11-123   O9.5 V
 4330  4340  4350
 6542  6562  6582
 4022  4026  4030
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4192  4200  4208
HeII 4200
 4536  4541  4546
HeII 4541
 4683  4686  4689
HeII 4686
Fig. A.6. Same as Fig. A.1, however, for N11-087 and -123.
lium abundance. Crowther et al. (2002) adopted a fixed value
of YHe = 0.2 whereas our automated fitting method was able to
self consistently determine a value of YHe = 0.28.
Sk −70 69 (Fig. A.8) To fit the spectrum of this O5 dwarf
Sk −70 69 a helium abundance of YHe = 0.17 was required.
Note that, similar to N11-065, which is also strongly enriched
with helium, the spectroscopic mass is found to be much
smaller than the evolutionary mass.
M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC 27
BI237   O2 V((f*))
 4333  4340  4347
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4465  4471  4477
HeI 4471
 4710  4713  4716
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4680  4686  4692
HeII 4686
BI253   O2 V((f*))
 4333  4340  4347
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4465  4471  4477
HeI 4471
 4710  4713  4716
HeI 4713
 4195  4200  4205
HeII 4200
 4536  4541  4546
HeII 4541
 4680  4686  4692
HeII 4686
Sk -66 18   O6 V((f))
 4335  4340  4345
 6552  6562  6572
 4021  4026  4031
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4468  4471  4474
HeI 4471
 4711  4713  4715
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
Sk -66 100   O6 II(f)
 4335  4340  4345
 6552  6562  6572
 4021  4025  4029
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4469  4471  4473
HeI 4471
 4711  4713  4715
HeI 4713
 4196  4200  4204
HeII 4200
 4537  4541  4545
HeII 4541
 4676  4686  4696
HeII 4686
Fig. A.7. Same as Fig. A.1, however, for BI 237, BI 253, Sk −66 18 and Sk −66 100. The dotted profiles for BI 237 correspond
to a model calculated with a Teff reduced by 3.8 kK compared to the best fit value. For Sk −66 100 the dotted profiles show the
effect of an increase in Ṁ by 0.12.
28 M. R. Mokiem et al.: Wind properties and evolution of hot massive stars in the LMC
Sk -67 166   O4 If+
 4335  4340  4345
 6547  6562  6577
 4021  4025  4029
HeI + HeII 4026
 4385  4388  4391
HeI 4387
 4469  4471  4473
HeI 4471
 4711  4713  4715
HeI 4713
 4196  4200  4204
HeII 4200
 4537  4541  4545
HeII 4541
 4677  4686  4695
HeII 4686
Sk -70 69   O5 V
 4335  4340  4345
 6547  6562  6577
 4021  4026  4031
HeI + HeII 4026
 4383  4388  4393
HeI 4387
 4468  4471  4474
HeI 4471
 4709  4713  4717
HeI 4713
 4195  4200  4205
HeII 4200
 4537  4541  4545
HeII 4541
 4682  4686  4690
HeII 4686
Fig. A.8. Same as Fig. A.1, however, for Sk −67 166 and Sk −70 69.
This figure "6489fig11.jpg" is available in "jpg"
 format from:
http://arxiv.org/ps/0704.1113v1
http://arxiv.org/ps/0704.1113v1
	Introduction
	Data description
	Analysis method
	The assumption of spherical symmetry
	Fundamental parameters
	Effective temperatures
	The Teff scale of O2 stars
	Gravities
	Helium abundances
	Microturbulence
	Wind parameters
	The mass discrepancy
	Chemically homogeneously evolving O2 stars
	Large sample trends
	The evolutionary status of N11
	LH9 and LH10 in N11
	The ages of LH9 and LH10
	Sequential star formation?
	Summary and conclusions
	Fits and comments on individual objects
