A&A 451, 195-206 (2006)
DOI: 10.1051/0004-6361:20054380
G. A. Wade1 - A. W. Fullerton2,3 - J.-F. Donati4 - J. D. Landstreet5 - P. Petit6 - S. Strasser7
1 - Department of Physics,
Royal Military College of Canada,
PO Box 17000, Station "Forces'',
Kingston, Ontario, K7K 4B4, Canada
2 -
Dept. of Physics and Astronomy,
University of Victoria,
PO Box 3055,
Victoria, BC V8W 3P6, USA
3 -
Dept. of Physics and Astronomy,
The Johns Hopkins University,
3400 North Charles Street,
Baltimore, MD 21218, USA
4 -
Observatoire Midi-Pyrénées, 14 avenue Edouard Belin,
31400 Toulouse, France
5 -
Physics & Astronomy Department,
The University of Western Ontario,
London, ON, N6A 3K7, Canada
6 -
Max-Planck Institut für Aeronomie
Max-Planck-Str. 2,
37191 Katlenburg-Lindau, Germany
7 -
Dept. of Astronomy,
University of Minnesota,
116 Church St. S.E.,
Minneapolis, MN 55455, USA
Received 20 October 2005 / Accepted 20 December 2005
Abstract
Aims. In this paper we confirm the presence of a globally-ordered, kG-strength magnetic field in the photosphere of the young O star Orionis C, and examine the properties of its optical line profile variations.
Methods. A new series of high-resolution MuSiCoS Stokes V and I spectra has been acquired which samples approximately uniformly the rotational cycle of Orionis C. Using the Least-Squares Deconvolution (LSD) multiline technique, we have succeeded in detecting variable Stokes V Zeeman signatures associated with the LSD mean line profile. These signatures have been modeled to determine the magnetic field geometry. We have furthermore examined the profile variations of lines formed in both the wind and photosphere using dynamic spectra.
Results. Based on spectrum synthesis fitting of the LSD profiles, we determine that the polar strength of the magnetic dipole component is
G and that the magnetic obliquity is
,
assuming
.
The best-fit values for
are
and
.
Our data confirm the previous detection of a magnetic field in this star, and furthermore demonstrate the sinusoidal variability of the longitudinal field and accurately determine the phases and intensities of the magnetic extrema. The analysis of "photospheric'' and "wind'' line profile variations supports previous reports of the optical spectroscopic characteristics, and provides evidence for infall of material within the magnetic equatorial plane.
Key words: stars: early-type - stars: magnetic fields - stars: mass-loss
The detection of magnetic fields in O-type stars has proven to be
a remarkably challenging observational problem (e.g., Donati et al. 2001).
The apparent absence of magnetic signatures has often been interpreted as a
selection effect, since about 5% of B- and A-type stars
(see, e.g., Johnson 2005) do display organized magnetic fields with disk-averaged
strengths ranging from a few hundred G to several tens of kG (Mathys 2001).
These fields are believed to be the fossil remnants of either interstellar fields
swept up during the star formation process or fields produced by a pre-main sequence
envelope dynamo that has since turned off.
As there is no particular reason to suspect that similar processes should not
occur during the formation of more massive stars, it seems reasonable a priori to
expect that magnetic fields of similar structure and intensity should also exist
in O-type stars.
In the absence of direct magnetic measurements, this view has been supported by the
wide-spread occurrence of variability in the winds of O-type stars, which may provide
indirect evidence for the presence of dynamically important magnetic fields in their
atmospheres.
During the past 20 years, sustained optical and UV spectroscopic observations have shown
that the winds of O-type stars are highly structured, exhibiting both coherent and stochastic behaviour, as well as cyclical variability on timescales ranging from
hours to days.
A key result of this work has been the conclusion that a major component of this
variability results from rotational modulation of structures imposed on the
wind by some deep-seated process.
Magnetic fields are presently considered a likely source
of such structures (e.g. Cranmer & Owocki 1996)
.
Although by no means prototypical, Orionis C (HD 37022; HR 1895) is perhaps
the best example of an O-type star with distinctive, periodic variability of its
spectroscopic stellar wind features.
It is a very young, peculiar
O7 star, and the brightest and hottest member
of the Orion Nebula Cluster.
It exhibits strictly periodic spectroscopic variability which is strongly
suggestive of a magnetic rotator: H
and He II
4686
emission, peculiar ultraviolet C IV
1548, 1550 wind lines, photospheric absorption lines, and ROSAT X-ray emission all appear
to vary with a single well-defined
period of
d (Stahl et al. 1996,1993; Gagné et al. 1997; Walborn & Nichols 1994).
The similarity of this behaviour to that of some magnetic A- and B-type stars
(see, e.g., Shore & Brown 1990) has motivated the suggestion first made by
Stahl et al. (1996) that
Ori C also hosts a fossil magnetic field
which confines the stellar wind and produces the observed variability by
rotational modulation. Such a phenomenon appears to have been first suggested (for the O star
Pup) by Moffat & Michaud (1981) (although no field has ever been detected in that star).
In order to test this suggestion, Donati & Wade (1999) obtained longitudinal
magnetic field measurements of Ori C, but failed to detect the
presence of a photospheric field.
However, in a seminal paper, Donati et al. (2002) reported the detection of a fossil
magnetic field in the photosphere of
Ori C based on 5 measurements
of the Stokes V profiles of selected photospheric absorption lines.
This detection provided the first empirical support that dynamically important
magnetic fields exist in O-type stars, and that their behaviour is directly linked
to spectroscopic variability
.
It also represents the youngest main sequence star in which a fossil-type field has
been detected (comparable in age to NGC 2244-334; Bagnulo et al. 2004).
Since
Ori C is clearly a pivotal object, it is important that
this detection be corroborated and the determination of its magnetic properties
be refined.
The primary goal of the present paper is to report confirmation of
the detection of a magnetic field in the photospheric lines of Ori C.
Our new spectropolarimetric observations, which provide a substantially larger data set
that samples the
rotational period more fully, are described in Sect. 2.
The analysis of the Least-Squares Deconvolved circular polarisation spectra,
both in terms of the mean longitudinal magnetic field,
,
and the LSD mean Stokes V profiles, is described in Sect. 3.
In Sect. 4 we discuss the variability characteristics of features in the Stokes
I spectrum, and comment on consistency with both the magnetically-confined
wind shock model of Donati et al. (2002) as well as very recent MHD simulations of
magnetically-channelled winds by ud-Doula & Owocki (2002) and Gagné et al. (2005).
Finally, in Sect. 5 we summarise our results, and discuss implications for
our understanding of the origin and evolution of magnetic fields in
intermediate- and high-mass stars, and for our understanding of wind
modulation in massive stars.
Circular polarisation (Stokes V) spectra of Ori C were obtained
during the period 1997-2000 using the MuSiCoS spectropolarimeter
mounted on the 2 metre Bernard Lyot telescope at Pic du Midi
observatory.
The spectropolarimeter consists of a dedicated polarimetric module mounted at
the Cassegrain focus of the telescope and connected by a double optical fibre
(one for each orthogonal polarisation state) to the table-mounted
cross-dispersed échelle spectrograph.
The spectrograph and polarimeter module are described in detail by
Baudrand & Bohm (1992) and by Donati et al. (1999), respectively.
The standard instrumental configuration allows for the acquisition of circular
or linear polarisation spectra with a resolving power of about 35 000
throughout the range 4500-6600 Å.
Table 1:
Journal of observations.
Phases are calculated according to the ephemeris of
Stahl et al. (1996),
.
The peak S/N per pixel in the continuum is quoted in the final column.
A complete circular polarisation observation consists of a series of
4 sub-exposures between which the polarimeter quarter-wave plate is
rotated back and forth between position angles (of the plate fast
axis with respect to the polarising beamsplitter fast axis) of
and
.
This procedure results in exchanging the orthogonally polarised beams
throughout the entire instrument, which makes it possible to reduce systematic errors
(due to interference and other effects in the telescope and polarisation optics, instrumental drifts, astrophysical variability, etc.) in spectral line polarisation measurements of sharp-lined stars to below a level of about
10-4 (Wade et al. 2000).
In total, 45 Stokes V spectra of Ori C were obtained over 4
observing runs in 1997 February, 1998 February/March, 2000 February/March,
and 2000/2001 December/January, with peak signal-to-noise ratios (S/N) of
typically 250 per pixel in the continuum.
The 6 spectra obtained during 1997 and 1998 have already been discussed by
Donati & Wade (1999).
Some fringing is visible in the Stokes I spectra (maximum amplitude
1% peak-to-peak in the red), although no fringing is evident in
Stokes V. Such effects are present in most polarimeters, and
generally result from internal reflections producing secondary beams
coherent with the incident beam, but with important (wavelength
dependent) phase differences (Semel 2003). This fringing limits the effective S/N of the Stokes I spectra to about 150:1 at H
,
and to about 350:1 at He II
.
In addition to exposures of Ori C, during each run observations of various
magnetic and non-magnetic standard stars were obtained which confirm the
nominal operation of the instrument (see, e.g., Shorlin et al. 2002).
The observing log is shown in Table 1.
As demonstrated by, e.g., Donati & Wade (1999), the Least-Squares Deconvolution
(LSD) multi-line procedure can provide enormous improvement in the precision of
spectral line polarisation measurements as compared with individual line
measurements.
We began using the 13 lines employed by Donati & Wade (1999) for the LSD
analysis.
After exploring the effect of removing lines from this list, it became clear
based on the shape of the Stokes I profile that only 3 of the lines in the
mask were contributing relatively uncontaminated photospheric profiles.
In the end, an LSD mask including only O III 5592,
C IV
5801 and C IV
5811 was used, with
a mean wavelength
Å and a mean Landé factor
.
We found that the resultant LSD S/N was dependent on the velocity bin size of the extracted LSD profiles. Ultimately, experiment yielded a best S/N for an bin size of 13.5 km s-1. Therefore, each pixel in the extracted Stokes I and V LSD profiles corresponds to a velocity interval of 13.5 km s-1.
Each extracted LSD profile was determined from 4 individual photospheric
features (C IV
5801 appears in two separate orders in each
spectrum).
To further reduce the noise, LSD profiles were binned in phase, according to the 15.422 day period of Stahl et al. (1996), weighting each pixel according to the inverse of its squared error bar. The final binned profiles, which will be used for all further LSD analysis, are summarised in Table 2.
Table 2: Binned LSD Stokes V profiles and longitudinal magnetic field measurements.
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Figure 1:
Longitudinal magnetic field variation (in G) of ![]() |
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Following the approach of Donati et al. (2001), the best-fit geometry of the dipolar magnetic field was determined using two complementary procedures.
First we modeled the variation of the mean longitudinal magnetic field.
The longitudinal field
and its associated error
were inferred from each set of LSD profiles, in the manner
described by Donati et al. (1997) and modified by Wade et al. (2000).
Measurements were made in the range
km s-1 around line centre, in both the Stokes V and diagnostic N LSD profiles.
The N, or null profiles are calculated from analysis of Stokes V CCD frames obtained at identical waveplate angles, and are nominally
consistent with zero (see, e.g., Donati et al. 1997). The N profiles provide a powerful diagnosis of systematic errors
in the polarisation data.
The measurements are summarised in Table 2 and are shown, phased
according to the 15.422 d period, in Fig. 1.
The reduced
of the (phased) Stokes V measurements of
is 1.0 for
a first-order sine fit (fitting zero-point, amplitude and phase), and 10.2 for
the null-field hypothesis (i.e. a straight line through
.
The reduced
of the N
measurements is 0.8 for a
first-order sine fit (fitting zero-point, amplitude and phase), and 1.2
for the null-field hypothesis.
Based on these results, for the 10 longitudinal field data discussed here,
the null-field hypothesis can be confidently ruled out for the Stokes
V
variation (false-alarm probability
0.1%),
whereas a null magnetic field is consistent with the N measurements within 2
.
The 1-order sine fit to the Stokes V
data
is characterised
by a maximum of
G, a minimum of
G and a phase of
maximum
(all uncertainties
).
Then, using a modified version of the programme F LDCURV, we calculated
model longitudinal field variations corresponding to a large number of dipolar
magnetic field configurations, varying the dipole intensity
as
well as the obliquity angle
.
F LDCURV computes the surface distribution of magnetic field corresponding
to a particular magnetic geometry, then weights and integrates the longitudinal
component over the hemisphere of the star visible at a given phase (for this
procedure we have assumed a limb-darkening coefficient of 0.4).
In order to uniquely identify the magnetic geometry, we must specify four parameters: the magnetic dipole polar strength ,
the inclination of the stellar rotational axis to the observer's line-of-sight i, the obliquity of the magnetic axis with respect to the rotational axis
,
and the phase of closest passage of the positive magnetic pole to the line-of-sight,
.
The phase variation of the longitudinal field allows us to constrain 3 of these parameters; typically
and
are determined by fitting the magnetic curve, whereas i is constrained using other data. Donati et al. (2002) derive
based on the rotational period and inferred
of
Ori C, as well as the characteristics of the optical and UV spectroscopic variability. Their conclusions are independently supported by the optical spectroscopic study of Simón-Díaz et al. (2006), who find
,
and qualitatively by the original Magnetically-Confined Wind Shock modelling of the X-ray variation of this star by Babel & Montmerle (1997b). We therefore adopt
for the modelling performed in this paper
.
For each computed model (over 40 000 in total), the reduced
of
the model fit to the observations was calculated.
The "map'' of these
s, shown in the upper frame of Fig. 2
(where black regions represent acceptable models, whereas white regions represent
models rejectable at more than the 95% [approximately 2
]
confidence
level
), indicates that for
a rotational axis inclination
,
acceptable magnetic models
are characterised by
G and
(where associated uncertainties are quoted at 1
).
If we allow a
uncertainty in our assumed value of
i, we obtain
G and
(
)
and
G and
(
).
Therefore our modeling of the longitudinal field variation constrains the
dipole magnetic field geometry of
Ori C to
G and
.
![]() |
Figure 2:
Map of reduced ![]() ![]() ![]() ![]() ![]() ![]() |
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![]() |
Figure 3:
Map of reduced ![]() ![]() ![]() ![]() ![]() ![]() |
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Our second procedure involves direct modeling of the LSD Stokes Iand V profiles, in a manner similar to that described first by
Donati et al. (2001).
In particular, all line profile modeling was accomplished using the LTE
polarised synthesis code Z EEMAN2 (Landstreet 1988; Wade et al. 2001)
and assuming an A TLAS9 model atmosphere with
K
and
.
We began by finding a phase-independent model fit to the Stokes Iprofiles, varying the line equivalent width, projected rotational
velocity ,
and microturbulent broadening
as free
parameters.
Assuming
km s-1 (Donati et al. 2002), we found a microturbulent
broadening
km s-1 (corresponding to a total thermal+turbulent
broadening of about 31 km s-1) provided a best-fit to the profiles.
As reported by Donati et al. (2002), the LSD Stokes I profiles vary systematically
according to rotational phase, with a maximum depth of about 12% of the
continuum at phase 0.0, and a minimum depth of about 9% of the continuum
at phase 0.5.
This variation is quite probably related to variable contamination of the
photospheric profiles by the wind spectrum; see Sect. 5.
The next step was to calculate Stokes V profiles corresponding to a
large number of dipole surface magnetic field configurations, varying
the dipole parameters
and
for selected values of i.
Finally, we compared each of these calculated profiles (about 28 000 in all)
with the observed Stokes V profile, calculating the reduced
and building up a "map'' of their agreement as well.
![]() |
Figure 4:
Averaged Stokes I and V and diagnostic null (N) LSD profiles of ![]() |
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Figure 5:
Dynamic spectra illustrating the line profile variations
of circumstellar features as a function of phase of the
15.422-day period.
Quotient spectra computed by dividing each spectrum
in the time series by the mean spectrum are illustrated.
The original spectral time series is also overplotted in the lower panel, where
the mean spectrum is shown in black.
Note that different lines have different dynamic ranges
("stretches'').
These are indicated by qmax and qmin, which are the
maximum and minimum values of the quotient that are plotted,
respectively.
White (qmax) corresponds to wavelengths or times when the
local flux is greater than its mean value.
Strong nebular emission components
have been excised from the central region of H![]() ![]() ![]() ![]() ![]() ![]() |
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![]() |
Figure 6: Same as Fig. 5, with a small dynamic range ("stretch'') that is fixed for all features. This presentation emphasizes weak spectral features like the red-shifted emission near phase 0.5. |
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Figure 7:
Same as Fig. 5, only for a selection of features
predominantly formed in the photosphere of ![]() |
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Figure 8:
Same as Fig. 5, only for a selected He I lines.
A strong nebular emission component has been excised from
the central region of He I ![]() |
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The results are illustrated in the lower frame of Fig. 2
(for
)
and in Fig. 3 (for
and
)
for 95% (2
)
confidence.
For
,
we obtain a somewhat smaller value for the intensity of the
best-fit dipole,
G, and
(again, all quoted uncertainties are
).
For
we obtain
G and
,
while for
we obtain
G and
.
Therefore our modeling of the LSD mean Stokes profile
variation constrains the dipole magnetic field geometry of
Ori C to
G and
.
The reduced
for the null-field hypothesis for the V measurements
is 1.11, versus 0.76 for the best-fit dipole model, indicating that the
null-field hypothesis can be confidently ruled out for the Stokes V spectrum
(false-alarm probability
0.1%).
At the same time, a similar analysis of the N diagnostic null LSD profiles
provides a best-fit model of
G with reduced
of 0.69, consistent with a null field and indicating that no significant
fringing contamination of the N profiles exists.
The low reduced
statistics obtained for both Stokes V and N suggest that the error bars are slightly overestimated (by about 10%). Based on the results of Wade et al. (2000), this is not unexpected. Moreover, the difference in best-fit reduced
obtained for the
V versus the N profiles (0.76 versus 0.69) is marginally significant,
indicating that (not surprisingly) the Stokes V signatures are not quite
fit to within their error bars, and suggesting that other unmodelled effects
may be at play. The computed Stokes profiles corresponding to the best-fit model are compared with the mean LSD profiles in Fig. 4.
The general solutions for the magnetic field geometry obtained using these
two methods are in good agreement.
The
solutions for
differ at the 1.5
level.
Although not especially significant, this difference could be attributable
to the more approximate weighting of the local field in the F LDCURV
model, or possibly contributions to the amplitude and variability of the
Stokes V signatures by emission contamination of the photospheric profiles
by the circumstellar material.
Due to the more sophisticated nature of the direct Stokes V fitting
procedure, we adopt these results as our formal constraints on the surface
magnetic field of Ori C.
These results are perfectly consistent with those derived by Donati et al. (2002),
but are to be preferred since they are derived from a larger sample of spectra.
We therefore confirm the detection and the geometrical characteristics of
the magnetic field discovered by Donati et al. (2002).
Moreover, we can affirm the sinusoidal nature of the longitudinal field
variation and confidently define the phases of maximum and minimum
(
and
,
respectively)
according to the ephemeris of Stahl et al. (1996).
The geometry of the magnetic field and circumstellar material of Ori C is sketched in Fig. 1 of Smith & Fullerton (2005) or Gagné et al. (2005), and is illustrated as a series of animations at www.astro.udel.edu/t1oc.
Phase-resolved variations of the Stokes I profiles from the 2000 time series are displayed in an image format for selected transitions in Figs. 5-8. In these images, which are commonly referred to as "dynamic spectra'', individual spectra occupy horizontal strips, which are stacked vertically to show systematic variations with time. The transitions are grouped to illustrate variations that are predominantly circumstellar (Figs. 5 and 6) and photospheric (Fig. 7), with selected He I lines serving as hybrid cases (Fig. 8).
The line-profile variations of Ori C have been previously
illustrated as dynamic spectra by
Stahl et al. (1996, for H
,
He II
4686,
C IV
1548, 1550,
Si IV
1393, 1402, and
O III
5592)
and by
Reiners et al. (2000, for He I
4471 and
4713,
C IV
5811, and
O III
5592).
Stahl et al. (1996) represented the variations as differences with respect to
either
(a) line profiles from a similar star [15 Monocerotis;
spectral type O7 V((f))] that are at most modestly contaminated by
stellar-wind emission; or
(b) minimum absorption profiles from the time series.
Reiners et al. (2000) illustrated the variations in the spectra themselves, i.e.,
without renormalizing by a template spectrum.
In contrast, the dynamic spectra illustrated here represent variations
with respect to the mean spectrum, which is reasonably well defined by the
approximately uniform sampling of this time series.
The advantage of displaying quotient spectra with respect to the mean
profile is that the relative amplitude of variations can be compared in a
meaningful way between (unsaturated) lines of different shape and strength.
Whatever template is used to normalize a time series, it is worth remembering
that the emergent spectrum is a very nonlinear function of, e.g., hydrodynamic
variables such as density.
Consequently, quantitative interpretation of the line profile variations in
terms of physical parameters generally requires comparison with a model.
Figures 5 and 6 illustrate the line-profile variations for
H,
H
,
and He II
4686, all of which
are dominated by stellar wind material in the magnetosphere.
All three lines exhibit qualitatively similar patterns of variability.
In Fig. 5, the contrast is adjusted separately for each line in order that
the entire variation can be seen, while in Fig. 6 a fixed high-contrast
"stretch'' is used to enhance the visibility of weaker variations.
Figure 5 confirms that the basic modulation consists of a broad emission
excess that attains maximum strength at phase 0.0 (when the magnetic equator is
viewed approximately face-on), decreases to a minimum near phase 0.5 (when the magnetic
equator is viewed approximately edge-on), and increases again starting near phase 0.75.
Although the overall symmetry of the emission feature is difficult
to determine due to contamination from the nebular component in Hand H
,
it is blue-shifted near phase 0 by perhaps as much as
100 km s-1.
As the emission fades, the primary variations are more nearly centered on
the rest velocity of the star.
However, Fig. 6 also indicates the presence of a second
component, which is particularly prominent in He II 4686.
It consists of an approximately stationary excess with
respect to the template at
km s-1 between
phases
0.4 and 0.6, which achieves maximum strength near phase 0.5
(i.e., when the magnetic equator is viewed edge-on).
Although it is often difficult to determine whether an increase in the
quotient spectrum represents excess emission or reduced absorption with
respect to the template spectrum, it is clear from examination of the
line profiles themselves that this component is due to emission.
This component has been noted previously in H
by Stahl et al. (1993),
who concluded that it was responsible for the smaller minimum near phase
0.5 in the "M-shaped'' equivalent width curve; see, e.g., Fig. 4 of
Stahl et al. (1996).
However, the origin of this feature was not discussed.
Figure 7 shows the line-profile variations of "photospheric''
lines, i.e., lines that are not obviously contaminated by emission from
the stellar wind or material trapped in the magnetosphere.
The lines included in this category are
C IV
5801, 5811,
O III
5592, and
He II
5411.
The first three of these are combined to obtain the LSD measurements
of the magnetic field of
Ori C.
The variations in all these lines are qualitatively similar,
including C IV 5801, which is not illustrated.
They are deeper with respect to the mean profile between phases 0.0 and
0.15; less deep between phases 0.2 and 0.6 (possibly later); and
deeper again between phases
0.8 and 1.0.
As previously noted by Stahl et al. (1996), the variations appear to move from
blue-to-red (i.e., in the sense of rotation), particularly near phase 0,
but are confined to the central region of the line profile.
The behaviour illustrated here is very similar to that exhibited in
other dynamic spectra, despite differences in the templates used;
see, e.g., Stahl et al. (1996) and Reiners et al. (2000).
Although not illustrated here, several other high-excitation lines were also
investigated.
No significant variations were visible in the He II 4542
and Si IV
4654 absorption lines or the
C III
5696 selective emission feature.
A weak modulation might be present in Si IV
4631.
The line profile variations of the He I lines illustrated
in Fig. 8 span the behaviour of "wind'' and "photospheric'' lines.
In particular, the strong He I 5876 triplet exhibits
"circumstellar'' variations.
This is not surprising, since the sensitivity of He I
5876
to low-density environments is a well-known consequence of non-LTE physics.
However, the "circumstellar'' variations in He I
5876 occur
at smaller velocities than indicated in Fig. 5.
In contrast to He I 5876, the weaker
He I
4713 triplet and
4921 singlet show
"photospheric'' variations that are very similar in phasing and character
to those illustrated in Fig. 7.
The variations in He I
4713 are substantially weaker
than those in He I
4921, even though the lines themselves
have similar strength.
Although not illustrated here, the He I
5015 singlet
also shows weak "photospheric'' variations.
The primary modulation of the circumstellar line profiles has been
modelled by Donati et al. (2002) in terms of the "magnetically confined wind
shock model'' of Babel & Montmerle (1997a,b).
The key feature of this model is that a sufficiently strong, dipolar
magnetic field channels wind material to the magnetic equator, where it
collides with material similarly channeled from the opposite hemisphere.
The consequences include substantial shock heating of the gas, the development
of extended cooling zones, and the creation of a zone of enhanced density at the
magnetic equator.
With this model and their revised mapping between rotational and magnetic phase
(which we confirm), Donati et al. (2002) were able to reproduce the fundamental
features of the primary modulation; see, e.g., their Fig. 12.
As a result, the maximum emission from circumstellar material is now understood
to come from the cooling region above and below the magnetic equator, which is
viewed face-on at phase 0.0.
In contrast, the observer views the magnetic equatorial plane edge-on at phase 0.5, at which
time the cooling disk is viewed edge-on against the stellar photosphere and its
contributions to the H
profile are minimized.
Despite the generally good fit, Donati et al. (2002) noted some difficulties with
their model profiles, which also did not reproduce the red-shifted emission feature near
phase 0.5.
Recently, Gagné et al. (2005) presented a magnetohydrodynamic (MHD) model of the
wind of
Ori C that confirms and extends the basic picture provided
by the "magnetically confined wind shock'' model.
The new "magnetically channeled wind shock'' model follows the same
computational approach described by ud-Doula & Owocki (2002) by allowing for the dynamic
competition between the outward forces of radiation pressure and the channeling and
confinement of the magnetic field.
However, in addition to allowing for the feedback of the outflow on the geometry of
the magnetic field, the newer model also incorporates a detailed treatment of the
energy balance in the wind, including the effects of compressive heating due to the
strong shocks in the vicinity of the magnetic equator.
Gagné et al. (2005) show that this model successfully explains both the periodic
modulation of X-ray emission exhibited by
Ori C and the basic features
of the X-ray emission lines.
The new MHD model also suggests that the circumstellar environment of Ori C
is much more dynamic than previously supposed.
In particular, the models show that the amount of material confined to the region of
the cooling disk is limited by outflow through the disk beyond the Alfvén radius
(i.e., the point at which the magnetic field can contain material channeled to the
region of the magnetic equator) as well as infall in the inner regions.
This infall occurs sporadically when cool, compressed material in the magnetic
equator becomes too dense to be supported by radiative driving, and consequently
falls back onto the stellar surface along distorted field lines.
Although much of the infalling material is too cool to emit X-rays, Gagné et al. (2005)
found evidence that the radial velocities of X-ray emission lines were systematically
red-shifted by as much as 93 km s-1 when the disk was viewed edge on
(i.e., near phase 0.5).
Smith & Fullerton (2005) also recognized that the presence of infalling material near the
magnetic equator provides a clue to the origin of the red-shifted emission feature
near phase 0.5 in line profiles dominated by circumstellar material
(Fig. 6).
They attributed this feature to a column of optically thick material flowing inward
at the magnetic equator, and used this interpretation to explain previously unnoticed
modulations of the C IV and N V resonance doublets at modest
red-shifts of 250 km s-1 or less.
We similarly attribute the red-shifted emission features near phase 0.5 in
Fig. 6 (and also for He I 5876 in Fig. 8)
to the infalling material seen in the MHD models.
In this context, the fact that the emission feature is seen at smaller red-shifted
velocities in He I
5876 (
50 km s-1; possibly over a
smaller range of phases) than its counterpart in
He II
4686 (
200 km s-1) suggests that
He+ is formed closer to the radius where infall is initiated, and that
the material heats up as it accelerates on its return to the star.
Although the presence of infalling material appears to be a robust consequence of
magnetic channeling for a star like Ori C, the current generation
of MHD simulations indicates that it is an occasional occurrence, which is not
tied to any specific rotational or magnetic phase.
Consequently, the consistent presence of the red-shifted features near phase 0.5 in
time series of optical, UV, and X-ray lines obtained over many rotational cycles is
surprising.
It remains to be seen whether refined MHD models can reproduce the approximate
steady-state of infall implied by these observations.
Since the mass-loss rate of
Ori C is poorly constrained, one
straightforward (though arbitrary) solution might be to increase the value assumed
in the MHD simulations.
An increase in the basal mass flux may slightly alter the magnetic geometry above the stellar surface,
but it will certainly increase the rate at which material stagnates near the magnetic
equator, perhaps to the point where the frequency of infall in the simulations
matches the high duty-cycle implied by the observations. However, a similar steady-state accretion also seems to be required for the magnetic B star
Cep, for which the mass-loss rate is well constrained. Such a scenario therefore does not appear to be applicable to
Cep, but it is not obvious that this conclusion can be extended to the case of
Ori C.
The red circumstellar emission component at phase 0.5 may also play a role in the variations of "photospheric'' lines. Stahl et al. (1996) and Reiners et al. (2000) interpreted these variations in terms of an absorption excess at phase 0.0, because the alternative - an emission excess at phase 0.5 - contradicted the behaviour of the primary modulation observed in emission lines, which achieves maximum at phase 0.0. Reiners et al. (2000) modelled the photospheric variations in terms of asymmetrical distributions of spots characterized by (a) reduced abundance near the magnetic poles and (b) enhanced abundance around the magnetic equator. They found reasonable agreement in both cases, but particularly for (b). Unfortunately, the mapping between rotational phase and magnetic geometry used by Stahl et al. (1996) and Reiners et al. (2000) turned out to be incorrect (see Donati et al. 2002), so their interpretations require revision.
More recently, Simón-Díaz et al. (2006) attempted to explain the variations of the photospheric lines as a consequence of additional continuum light from the disk-like structure at the magnetic equator. However, they also assumed that rotational phase 0.0 corresponds to the configuration when the magnetic equator is viewed edge-on, in order that the maximum contamination (hence minimum line depth and equivalent width) will be seen at rotational phase 0.5. This mapping between rotational phase and magnetic geometry is not supported by the magnetic field measurements. Furthermore, continuum variations on the rotational period of the size they predict (0.16 magnitudes, peak-to-peak) are excluded by the photometry of van Genderen et al. (1985). Consequently, the explanation for the photospheric line-profile variations proposed by Simón-Díaz et al. (2006) is not viable.
Instead, we note that the photospheric line profiles are less deep at
the same time that the circumstellar lines exhibit the red emission component.
This behaviour is seen particularly well in Fig. 8
which shows that the red-wing emission of the "circumstellar''
He I 5876 line is in phase with the mid-cycle absorption
minima in the other, "photospheric'' He I lines.
Consequently, if the explanation for the red-shifted component at phase 0.5 in
circumstellar lines in terms of infall is correct, then the synchronized
variations of the "photospheric'' lines might also be attributable
to the infalling material.
From this perspective, the change in photospheric line depth is due to partial
emission filling near phase 0.5, rather than excess absorption.
Quantitative modelling based on magnetohydrodynamic simulations is required
to determine whether the geometry and properties of the infalling material
can affect high-excitation lines like C IV
5801, 5811,
e.g., by producing an optically thick screen that shadows a small fraction
of the stellar disk (thereby minimizing continuum photometric variations) but is
sufficient to change the appearance of the red half of a photospheric line profile.
In particular, it remains to be seen whether sufficient density accumulates near the
stagnation point to explain the presence of high-excitation emission at small
infall velocities.
Using a new series of 45 Stokes I and V spectra obtained with the MuSiCoS
spectropolarimeter at Pic du Midi observatory, we have detected the
photospheric magnetic field of the young O7 Trapezium member Ori C.
We confirm and extend the conclusions of Donati et al. (2002): that the Stokes V
variations are consistent with a dipolar magnetic field with a polar strength
between 1150 and 1800 G, and an obliquity
if
.
Moreover, we demonstrate that the variation of the longitudinal magnetic
field is sinusoidal to within the errors; the phase of maximum
longitudinal field is
according to the ephemeris of
Stahl et al. (1996); and the longitudinal field varies between
G and
G.
We have also exploited our high-resolution Stokes I spectra to study the
cyclical variations of spectral absorption and emission lines formed in the
photosphere and wind of
Ori C.
We confirm the variability properties reported previously
by, e.g., Stahl et al. (1996), and highlight evidence that suggests the presence
of infalling material in the magnetic equatorial plane, which is consistent
with earlier suggestions of such phenomena by Donati et al. (2001) and the
theoretical predictions of Gagné et al. (2005).
Ori C holds a special place in our understanding of magnetism in
intermediate- and high-mass stars, because it is one of only two O-type stars in which a magnetic field has been detected and characterized at multiple epochs.
It is also one of the youngest stars (with an age of about 106 years) in which a
magnetic field has been detected, demonstrating once again (e.g. Bagnulo et al. 2004) that
magnetic fields are apparent at the surfaces of some intermediate and high
mass stars at very early main sequence evolutionary stages.
The confirmation of a field in
Ori C extends the range of known stars
hosting fossil magnetic fields by a factor of about 2 in effective
temperature, and by about 4 in stellar mass.
Acknowledgements
G.A.W. warmly acknowledges David Bohlender (Herzberg Institute of Astrophysics, Canada) for first bringing the intriguing case ofOri C to his attention. We also extend thanks to Otmar Stahl for helpful advice and fruitful discussions. G.A.W. and J.D.L. acknowledge Discovery Grant support from Natural Sciences and Engineering Research Council of Canada.