Issue |
A&A
Volume 515, June 2010
|
|
---|---|---|
Article Number | A45 | |
Number of page(s) | 13 | |
Section | Interstellar and circumstellar matter | |
DOI | https://doi.org/10.1051/0004-6361/200913209 | |
Published online | 08 June 2010 |
The origin of mid-infrared emission in
massive young stellar
objects: multi-baseline VLTI observations of W33A![[*]](/icons/foot_motif.png)
W. J. de Wit - M. G. Hoare - R. D. Oudmaijer - S. L. Lumsden
School of Physics & Astronomy, University of Leeds, Woodhouse Lane, Leeds LS2 9JT, UK
Received 30 August 2009 / Accepted 9 December 2009
Abstract
Aims. In this paper we aim to determine the
structure on 100 AU scales of the massive young stellar
object W33A, using interferometric observations in the
mid-infrared. This emission could be caused by a variety of elements,
for example, the inner protostellar envelope, outflow cavity walls, or
a dusty or gaseous accretion disk.
Methods. We used the Unit Telescopes of the VLT
Interferometer in conjunction with the MIDI instrument to obtain
spectrally dispersed visibilities in the N-band on
4 baselines with an angular resolution between 25 and
60 milli-arcsec (equivalent to 95 and 228 AU
at 3.8 kpc). The visibility spectra and spectral energy
distribution were compared to 2D-axi-symmetric dust radiative transfer
models with a geometry that includes a rotationally flattened envelope
and outflow cavities. We assumed an O 7.5 ZAMS star
as the central source, consistent with the observed bolometric
luminosity. The observations were compared to models with and without
(dusty and gaseous) accretion disks.
Results. The visibilities are between 5% and 15%,
and the non-spherically symmetric emitting structure has a typical size
of 100 AU. A satisfactory model is constructed to reproduce
the visibility spectra for each (u,v)
point. It fits the N-band flux spectrum, the
mid-infrared slope, the far-infrared peak, and the (sub)mm regime of
the SED. It produces a 350 m morphology consistent with the observations.
Conclusions. The mid-infrared emission of W33A on
100 AU scales is dominated by the irradiated walls of the
cavity sculpted by the outflow. The protostellar envelope has an
equivalent mass infall rate of ,
and an outflow opening angle of
.
The visibilities rule out the presence of any dust disk with total (gas
and dust) mass more than
.
Within the model, this implies a disk
of less than
,
where
is the viscosity of the Shakura-Sunyaev prescription. However,
optically thick accretion disks, which are inside the dust sublimation
radius, are allowed to accrete at rates equalling the envelope's mass
infall rate (up to
)
without substantially affecting the visibilities due to the extinction
by the extremely massive envelope of W33A.
Key words: stars: formation - stars: early-type - ISM: jets and outflows - accretion, accretion disks - techniques: interferometric
1 Introduction
Circumstellar accretion disks are an essential element in the formation of stars. The relatively long time scales involved in low-mass star formation (SF) and the proximity of the regions where they form allow the disks to be imaged at millimetre (Dutrey et al. 1996), near-infrared (near-IR; Padgett et al. 1999) or even optical wavelengths (Grady et al. 1999) and studied in great detail. The short time scales on which high-mass stars are formed and the distance of massive SF regions amongst others render the study of the accretion process particularly difficult. Whether or not a high-mass star grows in mass through the accretion of matter delivered to the stellar surface by means of a circumstellar disk is therefore still left unanswered. Recent 3-D radiation hydrodynamic simulations demonstrate that accretion through a disk can continue in the presence of strong radiation pressure (Krumholz et al. 2009). The pressure is greatly reduced by radiation escaping through the outflow cavities (Krumholz et al. 2005). Observationally a growing number of direct and indirect pieces of evidence point to large-scale (10 000 AU) rotating toroids (e.g. Beltrán et al. 2004; Beltrán et al. 2005; Furuya et al. 2008), but also to intriguingly disk-like structures on scales of several 100 AU in a small number of comparatively nearby examples (e.g. Hoare 2006; Jiménez-Serra et al. 2007, for a review see Cesaroni et al. 2007).
The reality of a disk accretion scenario in massive star
formation can be
assessed by observations of massive stars during the main
mass-gathering phase
at significant angular resolution. Prime examples of such young massive
stars
constitutes the class of object referred to as massive YSO (MYSO), but
also as
high-mass protostellar object, and BN-object. The class is
characterised by
luminous (>
)
infrared objects with spectral energy
distributions that peak around 100
m. The luminosity indicates a single
ZAMS star that would be able to ionise its surroundings, although too
little or
no radio emission is observed. This absence has been used as an
argument to
claim ongoing mass infall onto the central object. As a result the star
is
extended hence relatively cool (Hosokawa
& Omukai 2009) and therefore unable to ionise its
surroundings (Hoare
& Franco 2007). Alternatively, high-mass infall rates
(Walmsley 1995)
or gravitational entrapment (Keto
2002) could quench the development of an H II
that dominates the radio continuum emission. These alternative
scenarios invoke, however, very high emission-measure (EM) H II
regions to make them optically thick in the radio, e.g.
(Osterbrock 1989).
However, it is unlikely in this picture
that the region would still be optically thick in the near-IR H I
recombination lines Br
and Br
,
as these have line-centre optical depths that are 105-6
times less than
(Hummer & Storey
1987). This would imply that these objects should have strong
near-IR H I lines, and yet
they only show weak, broad lines indicating a stellar wind origin (Bunn et al. 1995).
Nonetheless, the MYSO class has the potential to test the
disk-accretion scenario. The actual accretion region is probably found
on AU to tens of AU scales near the central stellar object.
This region is, however, buried by hundreds of
,
the dust that makes up the extended (
0.1 pc) protostellar
envelope (van der Tak
et al. 2000). Only
long-wavelength radiation is capable of carrying away information on
the
physical processes in play, limiting the attainable spatial resolution
with single-dish apertures. Interferometers can overcome the problem of
limited resolution
and the near and mid-infrared wavelength regions currently deliver the
highest
angular resolution by means of e.g. the Very Large Telescope
Interferometer (VLTI).
The inner parts of the dusty MYSO protostellar envelope is
heated to a few 100 K and emits strongly at mid-IR
wavelengths. The 10 m
emission is usually unresolved at the 0.3
resolution of single 8 m dish telescopes, but it is partially
resolved at
24.5
m (de Wit et al. 2009).
Here, we present new observations
made of the MYSO W33A with the MID-infrared Interferometric
(MIDI) instrument
at the VLTI. MIDI is a two-beam combiner that operates in the N-band
(8-13
m)
and that delivers spectrally dispersed interferometric observables
(see Leinert
et al. 2003). W33A is one of the
most intensely studied MYSO thanks to its IR spectrum rich in ice
features (Gibb
et al. 2000). Its physical
nature has received much less attention. The object has an IRAS
luminosity of
(Faúndez et al.
2004) for a
kinematical distance of 3.8 kpc (Bronfman et al. 1996).
The object shows
weak, compact, and optically thick radio continuum emission (Rengarajan
& Ho
1996; van der Tak & Menten
2005) and broad (
100 km s-1),
single-peaked H I recombination emission
consistent with an ionised
stellar wind origin (Bunn
et al. 1995). Modelling
of the protostellar envelope with spherical dust radiative transfer
(RT) model
at constant density by Gürtler et al. (1991) required the
presence of an
inner dust-free cavity of 135 AU (35 milli-arcsec at
3.8 kpc), a size
that would be well resolved by MIDI. In de Wit et al.
(2007, hereafter Paper I) we presented a single MIDI baseline
for W33A in order to probe this hundred AU scale
region. Via 1-D modelling we found an inner dust-free radius at
7 milli-arcsec (26.5 AU), corresponding to the dust
sublimation
radius. We concluded that the dispersed MIDI visibilities and SED can
be
reproduced by a spherical dust model provided that the density law is
relatively flat. In this paper we present three additional VLTI/MIDI
observations of W33A with baseline position angles close to
perpendicular and
with different projected lengths. This allows us to constrain the
distribution
of the emitting material further using 2-D axi-symmetric RT models.
We describe and discuss the new and archived MIDI observations in Sects. 2 and 3. The basic features of the RT model, our modelling approach, and the final preferred model are explained in Sect. 4. The implications for W33A and their effect for MYSOs in general are presented in Sect. 5. We conclude our work in Sect. 6
2 Observations and data reduction
Table 1: List of MIDI observations of W33A.
2.1 MIDI observations
W33A was observed with the VLTI and the MIDI instrument in service mode on four different occasions using the UTs as aperture stations (see Table 1). We aimed at covering different position angles on the shortest possible UT baselines. The baselines are presented in Fig. 1 projected onto the near-IR reflection nebula of W33A as imaged by UKIDSS (Casali et al. 2007; Warren et al. 2007; Lucas et al. 2008). The shortest VLTI baselines are attained with the ATs, but W33A cannot be observed in this mode because of its relatively low brightness. The obtained angular resolution for the data set presented here stretches from 25 to 60 milli-arcsec. Detailed descriptions of the MIDI observation procedure are given in Przygodda et al. (2003), Chesneau et al. (2005) and Leinert et al. (2004).
![]() |
Figure 1: UKIDSS K-band observations of W33A. Projected VLTI/MIDI baselines are indicated by white bars. |
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The W33A observations were executed in
the so-called High-Sens MIDI mode, which uses all
the incoming light for
beam combination and fringe tracking. The incoming beams were combined
producing two complementary interferometric channels that have by
definition a
phase difference of
radians. The uncorrelated flux
spectrum is measured immediately after the
interferometric observation in order to obtain final visibilities. The
high-sens mode is advantageous when observing
faint targets or low visibilities, however the accuracy of the final
visibility spectrum is limited by the sky brightness variation between
the
interferometric and photometric measurement. This can amount to 10-15%
uncertainty, depending on the atmospheric stability. During the nights,
the
unresolved interferometric calibrator star W33A
was observed in order to determine the instrumental visibility and
correct for it. We also inspected
all other observed calibrators from different programmes for each night
(a total of 18 objects), which allowed us to estimate the
stability and accuracy of the
instrumental visibility. We found it to be uncertain by not more than
5%, much smaller than the
uncertainty in the determined photometric measurement (20%, see next
section). The calibrator star also provides an approximate flux
calibration (Cohen
et al. 1999) in addition to the visibility
calibration. A prism with a spectral resolution of 30 was
employed to
disperse the incoming beams.
The correlated flux spectra were estimated by the MIA+EWS software package (version 1.6; see Jaffe 2004; Koehler 2005). Interferograms were coherently added to maximise the signal-to-noise. First, fringe spectra had to be corrected, because they are affected by the instantaneous atmospheric and instrumental piston. The fringe scan was used to estimate the group delay due to the atmosphere. Removing the atmospheric and the (known) instrumental group delays constitutes a correction to the dispersed fringe signal, and straightens the dispersed fringe spectra; i.e., the phase is independent of wavelength. Next, the phase offset due to varying water refraction index between the time of recording of the fringe spectra had to be accounted for. In principle, all spectra can then be added to a final fringe spectrum and the amplitude of the power at each frequency estimated, which will give the correlated flux.
The total flux spectra were taken in a chopping mode and
each observation resulted in two spectra as MIDI splits the beams. The
spectra were extracted from the raw data in the EWS mode. An estimate
of the sky background was made and subtracted. The area corresponding
to the mask with which the correlated spectra were extracted is used
here.
The counts were then simply summed in the y-direction.
The final
spectrum is the sum of the two geometric means of the four individual
spectra:
.
This quantity is also what was obtained for the
correlated flux after beam combination and thus ensures consistency in
deriving visibilities.
2.2 JCMT/HARP-B observations
W33A was observed on the night of 5 June 2007 in 12CO, 13CO, and C18O using the HARP-B instrument on the James Clerk Maxwell Telescope. The 12CO data have a resolution of 488 kHz, and the other two lines were rebinned to the same resolution, giving a velocity resolution of approximately 0.4 km s-1. The maps had a linear baseline removed. The core emission was presumed to be traced by the C18O emission (which has line profile very close to a Gaussian). We used this as a model to trace the emission in the 12CO and 13CO maps that arises from the actual outflow. Full details of this work, which is part of a larger project studying outflows from the Red MSX Source Survey (see e.g. Urquhart et al. 2008), will be published separately in due course.
3 Description of observational data of W33A
3.1 N-band from MIDI
The measured MIDI observables (flux, correlated flux) and derived visibilities of W33A are presented in Fig. 3. Each column represents a VLTI configuration, and we have included the observations already presented in Paper I (configuration D). The top panels show the flux-calibrated N-band spectra. They show a variation of around 20% as expected for data taken in the observational high-sens mode. The errorbars represent two uncertainties that are added in quadrature: (1) the rms estimates of the flux; and (2) the systematic error determined from a total of five W33A MIDI spectra. (More than the standard one flux spectrum per interferometric measurement were taken). The distinct feature of W33A in the N-band is the extremely deep silicate absorption, the actual depth is not detected at the central wavelength where it drops to less than 1% times the level of the pseudo-continuum.
The second row of panels shows the correlated flux spectra.
The detector noise level was estimated by processing
that part of the detector directly adjacent to the W33A fringes,
indicative of
the background noise. The noise produces a signal a few times
0.1 Jy around 10 m, despite the absence of incoming flux at the
base
of the silicate absorption. Similar noise levels have been found
previously
(e.g. Jaffe 2004;
Matsuura et al.
2006). The errorbars of the W33A correlated
flux spectra reflect the uncertainties due to spurious correlated noise
level.
We discard any signal below this noise level in further analysis.
The third row of figures presents the resulting visibilities
for the wavelengths
not affected by low flux levels. The errorbars are the uncertainties
from
correlated and total flux added in quadrature. This implies that the
systematic uncertainty inherent to the high-sens mode are taken into
account
in the visibility uncertainties. We find that W33A has visibility
values between 5% and 15% on the used baselines. The spectral
shape of the
visibilities reveal a clear trend toward decreasing visibilities on the
blue wing of the silicate absorption feature. The lower visibilities
imply a larger emitting region, hence that the increasing optical
depth due to the silicate absorption comes from structures on larger
size scales than the pseudo-continuum. The visibilities corresponding
to the two
similar configurations (B and C) are practically the same, as
are their
correlated flux levels. It gives an independent indication of the
stability of the set-up and of the relatively good weather.
Configuration A
shows a correlated flux level twice higher than the ones of
configuration B and C, despite the relatively little
difference in total flux. Given that the weather during the
observations was
stable, we believe that the difference in visibilities between the
various
configurations is real. If we translate the visibilities at 8 m to
Gaussian emission regions, then
the FWHM decreases from 115 AU to about
95 AU rotating in PA
from 15
(configuration A) to
(B and C) and changing the
baseline length accordingly.
![]() |
Figure 2:
|
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3.2 The molecular outflow and near-IR reflection nebula
The star formation activity in W33A produces an archetypical bipolar
molecular outflow (Fig. 2).
The southeastern lobe constitutes the approaching,
blue-shifted flow. Detailed millimetre line maps that would provide
constraints on the outflow opening angle and the inclination of the
system are
currently unavailable. However, outflow activity produces reflection
nebulae
particularly well traced by near-IR scattered light (see e.g. Tamura et al. 1991).
The near-IR image of W33A already presented in
Fig. 1
is consistent with the CO map in position angle and orientation.
We used this image therefore in order to constrain these system
parameters, which will further help in modelling the MIDI visibilities
and SED later
on. The full extent of the
long K-band nebula is shown in
Fig. 1.
The nebula has a position angle of
,
but patchy extinction distorts the morphology on scales of 1
.
A fainter extension can also be seen in the northwestern
direction. This lends support to
position angle to be the
actual orientation of the system. The high-angular resolution IFU
images
presented in Davies
et al. (2010) confirms the reflected nature of the
K-band emission at the base of the outflow and
position angle. The outer regions of the nebula may have contributions
from shocked
emission as traced by extended
4.5
m
emission in Spitzer IRAC images (Cyganowski et al. 2008).
In Davies
et al. (2010), spectro-astrometry of the
Br
transition is also used to show that the near-IR source is at the base
of the bipolar outflow on scales of 200
-arcseconds. The near-IR source most likely
corresponds
to the central object in W33A, as was assumed during the
execution of the MIDI observations. This source shows a K-band
profile in
the UKIDSS image that is broader than the unresolved stars in the
nearby
field. In the H-band, the scattering nebula is much
weaker (see
Fig. 4).
Given that photons are more efficiently scattered at
shorter wavelengths the relatively dim nature of the H-band
nebula argues for
strong extinction in the line-of-sight towards W33A. Moreover, the
central
source is absent in the H-band.
![]() |
Figure 3: MIDI observations of W33A on four different occasions (see Table 1). Top: flux spectra. Middle: correlated flux spectra (points with errorbars), and the detector noise level (dashed curve). Bottom: derived visibilities. |
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Most of the emission seen at K-band is dust
scattered. Contribution
to the emission by molecular hydrogen gas is small as proven by
intermediate-resolution spectra, obtained as part of the rms
survey. In the H-band
a foreground
star is visible projected onto the nebula, located 2.5
south and
1.0
east of the main K-band
source. We re-measured the flux
densities coming from the full extended nebula as seen on the UKIDSS
images, after first
removing all the interloping stars from the images and replacing the
pixels values using the adjacent regions. We find
that in the K-band the nebula is 0.39 Jy (
)
and in the H-band
0.025 Jy (
).
4 Modelling
In this section we present the simultaneous modelling of the visibility
and
spectral energy distribution of W33A. Historically, the SEDs of MYSOs
have
been interpreted as hot stars embedded in a spherically symmetric dusty
envelope (e.g. Guertler
& Henning 1986). In
Paper I we adopted the approach of modelling the observables
with a 1-D
spherical symmetric radiative transfer code. We concluded that models
with shallow
density laws are preferred, provided that a relatively low effective
temperature for the central star is adopted. Here we construct a
fiducial model of a hot star surrounded by a protostellar envelope
with outflow cavities. Our initial model aims to be as simple as
possible, and we
begin by investigating the imprint on N-band
visibilities by a protostellar envelope that includes outflow cavities.
The walls of the cavities are
suggested to be the source of mid-IR emission seen on 1
scales in the
nearly edge-on MYSO W33A (De Buizer 2006)
and in W33A (de
Wit et al. 2009).
4.1 Description of the radiative transfer code and input
We adopt the 2D-axi-symmetric dust RT model by Whitney et al.
(see for details Whitney
et al. 2003a,b). The model
calculates radiation transfer (absorption, re-emission, and scattering)
through
a dusty structure. The inner rim of the structure is the dust
sublimation
radius that follows from a fit to models with different stellar
temperatures
(see Whitney
et al. 2004: );
space is empty within the inner rim except for the star
itself. The structure
can include three distinct geometrical elements: an accretion disk, a
protostellar envelope, and low-density polar outflow cavities. The
protostellar envelope is described by the analytical TSC solution (Ulrich 1976; Terebey et al. 1984)
of a simultaneously rotating and collapsing spherical
structure. This solution is governed by the infall parameter (
)
and the centrifugal radius
,
i.e. the radius where rotational motion
dominates infall.
The envelope mass infall rate is a parameterization of the
density distribution and does not necessarily represent an actual
determination of
infalling material. The presence of an accretion disk is
optional in the code. When present, the disk structure follows the one
for a flared
-type disk
prescription, and the released accretion luminosity within
the dust inner rim is added to the central luminosity source. The
outflow cavity geometry can have two possible shapes, either a
streamline or polynomial geometry. The streamline is conical on large
scales and the polynomial has a
opening angle at the stellar surface. Each of these geometrical
elements are potential mid-IR emitters and could contribute
to the visibilities. For most model runs, the
is comparatively small (tens of AU) and unresolved by the
employed VLTI baselines in the N-band.
![]() |
Figure 4:
Left: UKIDSS H-band (
upper panel) and K-band ( middle
panel) observations of the base of the W33A outflow
scattering nebula. Images are logarithmically scaled, in units of
MJy/sr and rotated by |
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A model run is necessarily characterised by a large number of free parameters. A fraction of these can be constrained by the known properties of W33A. The spatial information at milli-arcsecond angular resolution provided by MIDI allows the RT model to be seriously tested on its correctness in describing an MYSO.
4.1.1 Fixed input parameters
The central star parameters are chosen close to an
O7.5 ZAMS according
to the scale of Martins
et al. (2005). The bolometric luminosity is
(Faundez et al. 2004), and it is
located at the kinematical distance of 3.8 kpc (Bronfman et al. 1996).
The luminosity is based on IRAS flux density, and could be
overestimated. We take the opportunity here to correct our 70
m Spitzer
MIPS photometry presented in Paper I, which was not corrected
for detector nonlinearity. Taking this correction into account we find
a value of
Jy,
quite similar to the IRAS measurement at 60
m. This
indicates that the luminosity based on IRAS is probably reasonable. We
think it is unlikely that a significant fraction of the IR luminosity
is because of the secondary millimetre source noted in van der Tak et al.
(2000) at 6
from the main
source. The W33A MIPS 24
m and 70
m data show no sign of a secondary
object, any more than a new 24.5
m image taken with Subaru/COMICS at an
angular resolution of 0.6
(in prep.). This
suggests that the IR luminosity is dominated by a single object.
A number of the important parameters that determine the
properties of the TSC
envelope are fixed for the following reasons. We have chosen a single
type of
dust that consists of an interstellar mixture of ``warm silicates'' (Ossenkopf et al. 1992)
and of amorphous carbon with a size
distribution according to Mathis et al. (MRN; 1977). The type
of silicates provides a better fit to the observed shape of the
absoprtion feature than MRN DL grain-mixture (Draine & Lee 1984),
in particular to the blue wing (see Paper I, and see also Capps et al. 1978;
Guertler & Henning
1986). The centrifugal radius is set at 33 AU which
is comparable
to the dust sublimation radius. The TSC envelope makes a transition
from a pure infall, -1.5 radial density law to a shallower rotating,
-0.5 law at the
centrifugal radius. A TSC envelope with a centrifugal radius smaller
than the sublimation radius is dominated by the infalling part of the
envelope's
structure solution; i.e., the envelope is quite spherically symmetric
apart
from the outflow cavities. The shape of the outflow cavities is
polynomial (see Whitney
et al. 2003a),
which in the presence of a stellar wind could be a more appropriate
geometry
than a conical streamline. The density in the outflow cavity is chosen
to
be typical of MYSO outflows, i.e. between
and a few times 104 cm-3
(see Beuther
et al. 2002). The cavity material consists of gas
and dust. We initially base the outer edge of the envelope on
350
m
observations by van der
Tak et al. (2000). The radial intensity profile
shows a downturn in submm emission at a radial distance of
,
or
AU. We require a
posteriori that the model densities at the outer edge are similar to
that of the large-scale molecular cloud material in which W33A is
embedded, viz. 104 cm-3.
4.1.2 Cavity opening angle
One can exploit the observed shape of the near-IR scattering
nebula to
constrain the opening angle of the outflow. To this end we use the
less computationally intensive scattering-only version of the dust RT
code (Whitney &
Hartmann 1992;
Whitney & Hartmann
1993). The scattering code provides images
at ten inclination angles
simultaneously, which one can readily
compare to the observations. We model the near-IR nebula under the
assumption that the emission is free of direct thermal emission. We
put particular emphasis on matching the outflow opening angle with the
observed one and constraining the inclination angle. Model input
parameters conform to the fixed parameters discussed in the previous
subsection, except for the dust properties. The detailed dust
properties are not important for the goal to constrain the geometrical
parameters
of the W33A system from the near-IR nebula. For completeness,
the scattering dust used here is coated with a thin layer of water ice.
The dust sizes follow a distribution that fits the extinction curve
from (Cardelli
et al. 1989) for a prescription of
.
This distribution is not a single power law, and it is
presented in (Kim
et al. 1994). For more
details on the dust properties see Fig. 6 in Whitney et al. (2003a).
![]() |
Figure 5:
Left: SED of the preferred model overplot
W33A flux
densities. Small filled circles trace the ISO SWS spectrum, small open
circles are the MIDI spectrum, the filled square is the
Spitzer MIPS point, open square is the IRAS measurement at
60 |
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![]() |
Figure 6: Model (full line and filled circles) fits to the observed visibilities. Model parameters are listed in Table 2. Corresponding images are presented in Fig. 8. Note the discrepancy with the spherically symmetric model (dotted lines) used to fit configuration D in Paper I. |
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Figure 4
compares the H and K-band
images of a model run with the W33A scattering nebula. The images of
the observed nebula
have been rotated by .
The observed and model images are in
the same units (MJy sr-1). Model images
are convolved with a
Gaussian point spread function that simulates the one derived from the
UKIDSS images, and Gaussian noise derived from the observed images is
added to the model ones. The outflow angle measured from the images
directly would be
;
however, the presented model has an
envelope geometry with a full cavity opening angle of less than half
that, viz.
.
In our fitting procedure, attention is
devoted to the absence of an H-band counterpart of
the central
source. Whether or not a central source is detectable depends on the
system inclination. The model that is presented in Fig. 4
has an inclination of
.
An inclination less than
or
more than
are excluded. In the former case, the central star
is visible in the H-band, whereas the redshifted
outflow lobe becomes prominent in the latter. These features are not
present in the UKIDSS near-IR images.
Although the shape on the sky can be reproduced reasonably
well, not a single
model run reproduces the observed H-K
colour of 3.5.
Photometry done on the model images shows that they produce colours
that are much bluer
instead. The model images of Fig. 4 match the colour
of the observed images, but additional foreground extinction of
is required. The model H and K-band
fluxes need to be scaled
up accordingly. The central source visible in the K-band
only constitutes a minor contribution to the total observed flux
density at this wavelength. We note that MYSO models invariably have
problems matching
the short wavelength fluxes of MYSO scattering nebulae (e.g. Alvarez et al. 2004).
The deviant model colours might be
rectified by including a clumpy structure within the outflow cavity
(Indebetouw et al.
2006) and/or
very small dust particles that are transiently heated to high
temperatures. We
find that introducing an outflow cavity with a smooth density
distribution into an infalling envelope cannot be the only physical
element
determining the observed near-IR emission of MYSOs. Regardless of the
solution, a deeper investigation of this problem is beyond the purpose
of this paper.
Nonetheless, a comparison of near-IR images with scattering
RT models has led
to relatively stringent constraints on the outflow opening angle and
comparatively less stringent limits on the inclination.
4.1.3 Methodology
The prime model parameters that remain to be determined are
the density as
parametrized by the envelope mass infall rate, the presence of an
accretion
disk (and its properties), and the outer radius. They determine the
peak of
the SED and, depending on inclination angle, the depth of the silicate
feature, and the dominant emission regions on scales of
100 AU. The RT code
is computationally demanding which excludes a grid search for the
optimum
value for the decisive model parameters. Building on a simple fiducial
model,
we try to find a model that fits the SED and MIDI visibilities
simultaneously. One run delivers the SEDs at 10 inclinations
simultaneously,
evenly spaced in .
Model images are produced for a single
inclination at a time but for any number of desired
wavelengths/filters. A
total of
``photons'' are used to produce images, for SED-only-runs
the code uses four times less. We divided the N-band
up in 6 square filters of 1
m width, running from 7.5
m to
12.5
m.
The resulting
narrow-band model images are multiplied with a Gaussian model for the
VLT-UT
Airy disk. This represents the aperture field of view with an FWHM
appropriate
to the particular wavelength. The resulting image is then
Fourier-transformed
and the visibilities extracted for the projected interferometer
baseline and
position angle. The model images have a pixel size of
,
which is
less than half the angular resolution (0.024
)
at the longest baseline
for the shortest wavelength. The SED is matched from
8
m
to
1 mm. The input data are the MIDI flux spectrum, the slope of
the SED
around 40
m determined
from the ISO-SWS spectrum (see Paper I) and the
MIPS 70
m flux
density. At the longer wavelengths, the integrated
fluxes from the following observations are used: 350
m from
with the
CSO/SHARC instrument (van
der Tak et al. 2000), 850
m
with JCMT/SCUBA (Di Francesco et al. 2008), and
1.2 mmm with SEST/SIMBA (Faúndez et al. 2004). The
far-IR and (sub)mm observations are compared to model SEDs representing
all the emitted energy at each wavelength, whilst the MIDI flux
spectrum is compared to the model SED with an aperture of 0.6
corresponding to the
MIDI slit width. Matching of the SED and visibilities is performed by
eye.
Table 2: Parameters for the preferred model discussed in Sect. 4.2.
![]() |
Figure 7: Logarithmically scaled slices through the spatial distribution of the density ( left) and temperature ( right) of the preferred model perpendicular to the outflow axis. The cavity has a very small, but non-zero density consisting of gas and dust (see text). |
Open with DEXTER |
![]() |
Figure 8:
Logarithmically scaled images ( top row) and
corresponding Fourier transforms ( bottom row) of
the preferred model (Table 2)
at wavelengths 7.5 |
Open with DEXTER |
4.2 Preferred model fit
With the parameter values described in the previous section, a model is obtained that produces a reasonable match to both the SED and the visibilities. We ran a total of some 30 models in order to narrow down the parameter values by varying the mass infall rate, inclination angle, and the outer radius. A satisfactory model is presented in Figs. 5 and 6. The parameter values of this particular model can be found in Table 2. The figures show that the model fits both the SED and visibilities reasonably. A discrepancy between model and observed SED is found in the near-IR, as expected (see Sect. 4.1.2). The model reproduces the spectral trend of smaller visibilities in the wings of the silicate feature, especially clear on the blue wing quite well. Figure 6 also displays the visibilities from the spherically symmetric solution that matches the SED and visibilities of configuration D only (from Paper I). Each panel shows the spherically symmetric model visibilities calculated for the appropriate baseline length. It is clear that they provide a considerably worse fit to the visibilities obtained on the additional configurations A, B, and C, illustrating the need for multiple baselines.
The actual model images of the preferred model are shown in
Fig. 8,
along with a slice through the spatial temperature and density
distribution in
Fig. 7.
They make clear that the dominant emission regions in the
N-band at a few 100 AU scale are the
envelope regions close to the surface
of the outflow cavity. These envelope regions are warmed up because of
irradiation by the star. We refer to these regions as ``cavity walls''
for brevity. The density distribution in Fig. 7 illustrates that the
density of the paraboloidal cavity is a few orders of magnitude less
than that
of the envelope. The cavity itself is hotter than the bulk of the
envelope but
contributes very little to the N-band emission due
to the low density.
Figure 8
shows that the emission is relatively extended on the probed scales and
produces consequently very low visibilities. An increase in optical
depth, by either probing a
higher inclination angle or increasing the mass infall rate,
extinguishes the
warmest, inner part of the irradiated walls. As a result visibilities
generally decrease with an optical depth increase. However the
extremely deep
silicate absorption feature of W33A requires a large optical depth.
This is
the trade off for obtaining a decent fit. We find that for an
inclination
angle larger than ,
the mass infall rate needs to be lower than
in order to see the warm cavity wall regions and avoid spatially
resolving the N-band emission out. The increase in
optical depth because of the silicate absorption has a similar effect.
Figure 8
makes clear that the warm inner regions of the cavity walls become
progressively fainter with wavelength causing an increase in the
typical emitting size, hence lower visibilities.
The same effect is found when one increases the centrifugal
radius. The
preferred model with an increased
to a value of 300 AU shows
that the density in the equatorial region increases, but the increase
is
confined to within a fraction of a degree of the equator. The density
actually
decreases away from the equator allowing a larger area of the cavity
walls to
be heated up. Visibilities drop therefore with increasing
.
Further
fine-tuning of the model parameters would require additional
observations at
preferentially shorter baselines.
In summary, MIDI visibilities and SED of W33A can be matched
with a simple
geometry that consists of a protostellar envelope with a -1.5 radial
density
distribution that has polar outflow cavities with a
opening angle.
The modelling of the MIDI data is not very sensitive to the exact
density
power law in the envelope within the framework of the applied RT model.
The
density power law in this case is effectively governed by the
centrifugal
radius. Spherical models as applied to W33A in de Wit et al. (2007)
are
inadequate for modelling structure on 100 AU scales, as
Fig. 5
illustrates. On the other hand, resolved observations of MYSOs at
25
m at
average inclination angles (de
Wit et al. 2009) and at
m (van der
Tak et al. 2000) provide insight into the density law starting
at radii of 1000 AU and indicate a radial density law of
-1.0. Matching images of W33A with model images at these respective
wavelengths will provide much stronger constraints on the centrifugal
radius. For now we show the 350
m morphology of our
model in Fig. 9.
We convolved the model image with a 10
beam appopriate for
comparison with the CSO observations presented in van der Tak et al.
(2000). These observations show a structure that is more
flattened than the
model image; however, the FWHM of the model image
of 16
is bracketed by the
observed FWHM range of 14.4
-19.4
,
showing a good
correspondance. The density at
the outer radius of the model complies with expected densities of
molecular clouds. Observations
in the submm give a lower limit to the size, while in reality the outer
parts of the envelope may become even steeper than -1.5(e.g. Mueller et al. 2002).
The outer radius is
therefore uncertain by a factor of 2. The best-matching RT
model presented does not need contributions to the emission from other
components within the MYSO environment; instead, N-band
emission on a 100 AU scale is completely dominated by the warm
cavity walls.
5 Discussion
In several studies the accretion based formation scenario for massive
stars
has been forwarded partially based on observations with 1
angular
resolution and/or SED studies of MYSOs (e.g. De Buizer et al. 2005).
If MYSOs are to be considered young massive
stars still in their accretion phase, and this accretion occurs in a
similar
fashion to low-mass stars, then circumstellar accretion disks ought to
be
present. The MYSO accretion disk is expected to be luminous
(approaching that
of the star itself), and is estimated to have a size a few times
100 AU. Our
analysis raises the question to what degree an accretion disk can
contribute
to the mid-IR emission.
5.1 The contribution of an accretion disk
The dust RT code allows the inclusion of a dust disk truncated at the
25 AU
dust sublimation radius. As a first step we consider the effect on the
model
fit by adding such a dusty disk to the preferred model presented
previously
(Table 2). We emphasize that the disk is directly irradiated
by the star, as
the space between disk and star is empty in the model. The mass of the
dusty
disk is initially chosen to be 1% of the stellar mass, i.e.
0.25 .
The radial density law of the disk agrees with that expected of an
-type disk
(i.e. a
power law exponent of 1.875, Frank
et al. 1985).
Because of the large inner truncation radius, the actual accretion
luminosity
generated by the dust disk is marginal and about 1% of the
stellar
luminosity (see Eq. (6) in Whitney
et al. 2003a). The formal accretion
rate (see Eq. (5) in Whitney
et al. 2003a) for
is then
.
The accretion luminosity inside the dust
sublimation radius, where no actual disk structure is present in the RT
code,
is emitted with the stellar spectrum (see Whitney et al. 2003a).
The disk has
a constant opening angle; i.e., the scaleheight increases linearly with
radius.
The model disk's outer radius is arbitrary because it does not emit in
the
N-band and is better constrained by millimetre data;
for now, it is
set to 500 AU.
The resulting model visibilities of a protostellar envelope
plus dust disk are
shown in Fig. 10.
It is clear that such a model violates the observed
MIDI visibilities shortward of the silicate absorption feature. At
these
wavelengths the dust disk contributes 50% to the total flux density,
and one can deduce that the disk has a visibility of
approximately 70%. For this model to be consistent with the
observations, we computed the fractional contribution to the total
8
m flux
by the dust disk
as probably at most 1%, if the disk emission at this wavelength is
fully
unresolved. If it is marginally resolved (as it is), the
1% upper limit
becomes less. By running models with increasingly lower disk masses, we
find that this flux fraction is reached for a total mass of the dusty
disk of
0.01
,
or equivalently
.
The
SED of W33A is still fitted well by the dust disk+envelope model. We
thus
find that the disk mass is very low compared to the envelope mass and,
for
reasonable values of
,
that the dust disk accretion rate is more than
three orders of magnitude lower than the envelope mass infall rate.
These
results on the disk masses are in accordance with other studies of
MYSOs (e.g. Murakawa
et al. 2008) and also with some evolved early-type
Herbig Be stars (Alonso-Albi
et al. 2009). Were the disk
outer radius much smaller than the assumed 500 AU and similar
to the small outer radii found for the Herbig Be stars, then the disk
mass should be scaled down accordingly: a 50 AU outer radius
results in a disk mass of
.
A
small disk mass may imply that such a dust disk is not present, or that
it is
simply extremely faint in the mid-IR because of the absence of gas and
dust opacities in the model disk, which would lead to strong shielding
from the stellar radiation field.
![]() |
Figure 9:
Morphology at 350 |
Open with DEXTER |
![]() |
Figure 10:
Model predictions for baseline configuration A. Represented are the
preferred model (full line), with the addition of a dust disk (
|
Open with DEXTER |
Most of the accretion luminosity, however, would be released within the
25 AU
dust sublimation radius by the hot gaseous extension of the dust disk.
Self-shielding means that such a disk probably contains dust at smaller
radii than the 25 AU sublimation radius determined from direct
irradiation.
Inclusion of a self-luminous accretion disk within the dust
sublimation radius is required to fit the SED and the near-IR
interferometric data of the Herbig B6e star W33A (Kraus et al. 2008).
Accretion luminosity generated by an accretion rate of
for this particular object is, however, not essential for modelling the
star's mid-IR interferometry, which is already reproduced reasonably
well by a passive disk. We explored to what extent the mid-IR
interferometry of W33A is affected by including an optically thick
accretion disk.
We approximated the reslulting fluxes and visibilities with an
optically thick,
geometrically thin
disk model (see Malbet
& Bertout 1995; Malbet
et al. 2007) and added these to the preferred
envelope model in Table 2. The disk is heated
only by stationary accretion and for an accretion rate equalling the
mass
infall rate (
)
the
generated accretion luminosity is
for our adopted
central star. The spatial properties of the optically thick disk
emission is such that 99% of the N-band is emitted
within 50 AU. Most of the
8
m flux
is emitted at 20 AU. The disk thus remains largely unresolved
on the employed VLTI baselines with MIDI. Disks accreting at a slower
rate than the one assumed are cooler and the mid-IR emitting region
more compact. We chose the inner radius of the disk to be at 1 R*.
The disk temperature falls below 1500 K at 2.5 AU,
beyond which the disk opacity is given by both gas and dust.
As in the previous paragraph, if
the preferred model visibilities are to be affected by the optically
thick disk, the ratio of disk emission to the emerging envelope
emission should be at
least 1% at 8 m.
We calculated the resulting disk spectrum for the
said disk accretion rate and applied an extinction of
due to the
overlying envelope, as found in the line-of-sight to W33A from the
preferred model. We show the resulting visibilities in
Fig. 10.
The figure shows that an
disk model embedded in a
protostellar envelope is compatible with the observed MIDI
visibilities. The
SED for this particular model does not differ from the SED with only an
envelope contributing
to the emission, provided that the total luminosity of each model is
chosen the same. The
disk contributes less than
1% to the total 8
m
flux, despite its high luminosity. The envelope of W33A is simply
too opaque for any disk emission to be directly observed. MIDI
observations of
MYSOs with less massive envelopes or at smaller inclinations, however,
may well disclose the presence of an
accretion disk in such systems. The expected visibilities would be much
higher
than the ones observed for W33A (depending on distance and
accretion rate), considering the compactness
and brightness of the 8
m emission in accretion disks.
The popular YSO SED model grid developed by Robitaille et al. (2006) allows an estimate of the
extent to which MYSO SEDs are dominated by dust disk emission. This YSO
SED grid makes use of the same 2-D axi-symmetric RT calculations as we
use in this
paper (Whitney
et al. 2003a), and the disks in the SED model grid
are therefore
dust disks. We extract all (135) SEDs corresponding to ZAMS
hot stars with a
luminosity of
and an envelope mass infall rate of
at least
.
Inspection of the fractional disk flux contribution at 8
m (just
short of the extended wings of the
silicate absorption feature) shows that all except three of the
extracted models have a fractional contribution over 15%, and half of
the models have a
contribution over 70%. The geometry of the accretion disk determines
the total visibility at these wavelengths. The SED model grid thus
contains a
preponderance of MYSO models with relatively high visibilities, similar
to the
envelope plus dust disk model presented in Fig. 10. At present, the
MIDI
observations of MYSOs demonstrate the opposite trend, however.
Visibilities
presented for W33A are in line with the low N-band
visibilities observed for
other MYSOs (see Linz
et al. 2008; Vehoff
et al. 2008; de Wit et al. in prep.). The
emission of these objects could equally well be explained partially or
fully in the context of cavity wall emission as proposed here for W33A.
Table 3: Models from the YSO SED model grid (Robitaille et al. 2006), obtained using the web-based fit procedure.
5.2 SED model grid fits
To what extent are models obtained from SED fits consistent
with the spatial
information on milli arcsecond scales? To address this question, we
applied the
web-based SED fitter procedure developed by Robitaille et al. (2007)
to the W33A SED data shown in Fig. 12. The procedure
returns a ranking on the basis of a
minimization method, and we discuss the grid models following this
ranking.
![]() |
Figure 11:
Logarithmically scaled K-band image for grid model
3004217 to be compared with the observations in Fig. 4. The contours are
at 1%, 40%, 50%, and 65% of the peak flux. The image is convolved with
a 2-D Gaussian function with an FWHM of 1.6
|
Open with DEXTER |
The procedure delivers a model consisting of a central hot star and a
dust disk as the best fit (model
grid number 3002520, see Table 3). We see from
Fig. 12
that the grid model fit to the SED is acceptable, as expected. This
particular grid model has a system inclination of ,
consistent
with the one of the reflection nebula. The best-fit model consists of a
protostellar
envelope with a relatively large cavity opening angle compared to the
one
derived from the scattering nebula. Moreover, a fit of the model to the
observed SED requires a distance of
1 kpc, which is a factor
4 too short. The discrepancy
between the modelled and observed SED is mainly at the depth of the
silicate
feature. This depth is better accounted for by dust partially
consisting of ``warm silicates'' (see Ossenkopf et al. 1992)
as discussed previously, while the appropriate set of optical constants
is not used in the SED model grid. When it is broken down in its
various contributing components, we find that the grid model SED is
completely
dominated by disk emission at wavelengths shortward of the silicate
absorption
feature. The emission accounts for the relatively high flux levels
observed shortward of the silicate feature. This is a common property
of MYSO SEDs. Spherical and indeed
2D-axisymmetric RT codes for the protostellar envelope are unable to
reproduce
these high flux levels (see e.g. de
Wit et al. 2009). Contrary to the good SED fit,
Fig. 13
shows that the corresponding visibilities are far too high. The
visibilities are obtained by producing images with the dust RT code
using the
same parameter file as produced the SED. The grid model visibilities
are inconsistent simply because of the dominance of the (compact) disk
emission. Clearly, a good model fit to the SED does not guarantee a
proper
description of the object in consideration, because the best-fit grid
model violates the spatial information delivered by MIDI.
If we use the known properties of W33A and limit the model
distance to
be between 3.0 and 4.6 kpc with a maximum foreground
extinction of
Av=50,
then we find that the returned fitting models have a strong
preference for small inclinations. In the best 10 models, 7
have an
inclination of 18.19 degrees (the smallest present within the
grid), also the best-fit grid model (number 3004217). In
addition, the visibilities fit the MIDI observations reasonably well
(Fig. 13).
The
model is characterized by a central cool star with a large stellar
radius. Such a description of the MYSO M8E-IR is preferred in Linz et al. (2009).
In the case of W33A, the small inclination angles are inconsistent with
the allowed range we derived from the
scattering nebula. The predicted K-band morphology
of this model is shown in
Fig. 11,
and it is strongly dominated by the central source in contradiction
with the observations of W33A (Fig. 4).
The ranked fourth model however has a more reasonable
inclination angle of
(number 3009737). The latter model consists of a hot star and
an accretion disk and encounters the previously discussed problems in
fitting the visibilities (Fig. 13). If we applied
all the W33A
fiducial parameters (
,
cavity opening angle
)
in selecting a model from the grid for a hot star (
)
without a
dust disk, then the model complying with these constraints is ranked
245th
(model number 3003047).
![]() |
Figure 12: Predicted SEDs of the discussed models obtained from the SED web fit procedure listed in Table 3. |
Open with DEXTER |
![]() |
Figure 13:
Predicted visibilities along the major (lower curve for each model) and
minor axis (top curve for each model) at 8 |
Open with DEXTER |
It is clear that the SED model grid can deliver models that fit the SED, but all the SED fitting models presented in this section violate the MIDI visibilities. Regarding W33A, the SED grid does not provide a fit to both SED and visibilities simply because the appropriate models are not present in the grid (see also the discussion in Robitaille 2008; Linz et al. 2009).
6 Conclusions
Observations with MIDI at the VLTI of the MYSO W33A indicate that N-band
emission on scales of 100 AU stems from the warm parts of the
envelope close to the outflow cavity, i.e. the ``cavity walls''. When
described by a TSC type geometry, the overlying envelope is
characterised by an equivalent mass infall rate of order
and is
quite massive. We find that the presence of a dust disk located
beyond the dust sublimation radius violates the observed visibilities.
To make the model
consistent with the observations, the dust disk must be marginal in
mass. In contrast, the addition of an optically thick accretion disk
inside the dust sublimation radius has little effect on the
visibilities,
even when the disk dominates the luminosity. The disk emission is
highly
extincted by the massive overlying protostellar envelope. MYSOs viewed
at smaller inclination and/or with less massive protostellar envelopes
have the
potential of revealing such accretion disks.
We would like to thank O. Chesneau for his prompt advice, and T. Fujiyoshi for the W33A COMICS data. It is a pleasure to thank B. Whitney for discussions of the subject, and to an anonymous referee for the clear remarks that improved this manuscript.
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Footnotes
- ... W33A
- Based on observations with the VLTI, proposal 381.C-0602.
- ...
survey
- http://www.leeds.ac.uk/RMS
- ...Robitaille et al. (2006)
- http://caravan.astro.wisc.edu/protostars/index.php
All Tables
Table 1: List of MIDI observations of W33A.
Table 2: Parameters for the preferred model discussed in Sect. 4.2.
Table 3: Models from the YSO SED model grid (Robitaille et al. 2006), obtained using the web-based fit procedure.
All Figures
![]() |
Figure 1: UKIDSS K-band observations of W33A. Projected VLTI/MIDI baselines are indicated by white bars. |
Open with DEXTER | |
In the text |
![]() |
Figure 2:
|
Open with DEXTER | |
In the text |
![]() |
Figure 3: MIDI observations of W33A on four different occasions (see Table 1). Top: flux spectra. Middle: correlated flux spectra (points with errorbars), and the detector noise level (dashed curve). Bottom: derived visibilities. |
Open with DEXTER | |
In the text |
![]() |
Figure 4:
Left: UKIDSS H-band (
upper panel) and K-band ( middle
panel) observations of the base of the W33A outflow
scattering nebula. Images are logarithmically scaled, in units of
MJy/sr and rotated by |
Open with DEXTER | |
In the text |
![]() |
Figure 5:
Left: SED of the preferred model overplot
W33A flux
densities. Small filled circles trace the ISO SWS spectrum, small open
circles are the MIDI spectrum, the filled square is the
Spitzer MIPS point, open square is the IRAS measurement at
60 |
Open with DEXTER | |
In the text |
![]() |
Figure 6: Model (full line and filled circles) fits to the observed visibilities. Model parameters are listed in Table 2. Corresponding images are presented in Fig. 8. Note the discrepancy with the spherically symmetric model (dotted lines) used to fit configuration D in Paper I. |
Open with DEXTER | |
In the text |
![]() |
Figure 7: Logarithmically scaled slices through the spatial distribution of the density ( left) and temperature ( right) of the preferred model perpendicular to the outflow axis. The cavity has a very small, but non-zero density consisting of gas and dust (see text). |
Open with DEXTER | |
In the text |
![]() |
Figure 8:
Logarithmically scaled images ( top row) and
corresponding Fourier transforms ( bottom row) of
the preferred model (Table 2)
at wavelengths 7.5 |
Open with DEXTER | |
In the text |
![]() |
Figure 9:
Morphology at 350 |
Open with DEXTER | |
In the text |
![]() |
Figure 10:
Model predictions for baseline configuration A. Represented are the
preferred model (full line), with the addition of a dust disk (
|
Open with DEXTER | |
In the text |
![]() |
Figure 11:
Logarithmically scaled K-band image for grid model
3004217 to be compared with the observations in Fig. 4. The contours are
at 1%, 40%, 50%, and 65% of the peak flux. The image is convolved with
a 2-D Gaussian function with an FWHM of 1.6
|
Open with DEXTER | |
In the text |
![]() |
Figure 12: Predicted SEDs of the discussed models obtained from the SED web fit procedure listed in Table 3. |
Open with DEXTER | |
In the text |
![]() |
Figure 13:
Predicted visibilities along the major (lower curve for each model) and
minor axis (top curve for each model) at 8 |
Open with DEXTER | |
In the text |
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