A&A 370, 100-111 (2001)
DOI: 10.1051/0004-6361:20010227
Th. Henning1 - M. Feldt2 - B. Stecklum3 - R. Klein1
1 - Astrophysikalisches Institut und Universitäts-Sternwarte Jena, Schillergäßchen 3,
07745 Jena, Germany
2 -
Max-Planck-Institut für Astronomie, Königsstuhl 17, 69117 Heidelberg, Germany
3 -
Thüringer Landessternwarte Tautenburg, Sternwarte 5, 07778 Tautenburg, Germany
Received 2 October 2000 / Accepted 7 February 2001
Abstract
We present the first near-infrared images of the two massive southern
star-forming regions G11.11-0.40 and G341.21-0.21, obtained with adaptive
optics at sub-arcsecond resolution. Together with conventional narrow-band
near-infrared images and polarization maps, mid-infrared data, millimetre
and radio data, we were able to determine the morphological structure
and to constrain the stellar content of the regions. In both cases the
molecular cloud cores have masses of about 500
and contain clusters
of OB stars. This supports the view that massive stars do not form in
isolation. In the case of G11.11-0.40 we could identify the central
O star responsible for the ionization of the ultracompact H II region.
Our results for G341.21-0.21 indicate that this source is a
precursor of an ultracompact H II region.
Key words: stars: formation - stars: early type - ISM: H II regions - methods: observational
To continue our study, we selected the two UCH IIs G11.11-0.40 and G341.21-0.21 which were observed with the ALFA and ADONIS AO systems, respectively. Both objects are associated with IRAS sources (G11.11-0.40 with IRAS18085-1931 and G341.21-0.21 with IRAS16487-4423) which have colours typical of ultracompact H II regions (Wood & Churchwell 1989). The CS(2-1) survey performed by Bronfman et al. (1996) demonstrated that the sources are located in dense molecular cloud cores.
The object G11.11-0.40 was identified as an ultracompact radio source
(Kurtz et al. 1994; Becker et al. 1994). The
source velocity fits the kinematics of the 3 kpc arm. Kurtz et al.
adopted a distance of 5.2 kpc which gives an IRAS luminosity of
(Walsh et al. 1997). G11.11-0.40 is also
associated with high-velocity gas seen in the 12CO (J=1-0) line
(Shepherd & Churchwell 1996).
The other object, G341.21-0.21, is associated with an H2O maser
detected at the position of the IRAS source (Filho & Escalise
1990), which is generally taken as evidence for massive
star formation. In addition, an OH maser was detected by
Caswell (1998). In this paper, we adopt a kinematic
distance towards this source of 3.7 kpc. Note that this distance
corresponds to the near solution of the kinematic model,
whereas Forster & Caswell (1999) prefer the far solution
of 15.3 kpc. That far distance would, however, yield a total
luminosity of more than 106
,
a value that seemed unrealistically
large to us. For our chosen distance, the luminosity resulting from
the IRAS fluxes is
.
Date | ![]() |
Tel./Instr. | FOV![]() |
PSF | Ref. Star![]() |
Limiting |
FWHM | Mag.![]() |
|||||
1999 Sep. | J/H/K' | CA3.5(ALFA/OMEGA) | 40
![]() |
0
![]() ![]() ![]() ![]() |
FS28 (10.744/10.644/10.602) | 19.5/19.8/19.1 |
1998 Jul. | Br![]() |
ESO2.2m/IRAC2b | 130
![]() |
1
![]() |
HD 210427 (8.852) | ![]() |
1998 Mar. | 1.3mm | ESOSEST(Bolo) | 4![]() |
23
![]() |
Uranus | 50mJy/beam |
![]() ![]() ![]() ![]() ![]() ![]() ![]() |
Date | ![]() |
Tel./Instr. | FOV![]() |
PSF | Ref. Star![]() |
Limiting |
FWHM | Mag.![]() |
|||||
1999 Mar. | K' Pol | ESONTT / SOFI | 4.9![]() |
0
![]() |
- | - |
1998 Jul. | Br![]() |
ESO2.2m/IRAC2b | 130
![]() |
1
![]() |
HD 210427 (8.852) | ![]() |
1996 Oct. | N | ESO2.2(MANIAC) | 43
![]() |
1
![]() |
![]() |
6.3 |
1998 Mar. | 1.3mm | ESOSEST(Bolo) | 4![]() |
23
![]() |
Uranus | 50mJy/beam |
1995 Aug. | H/K' | ESO3.6(ADONIS/SHARP2) | 13
![]() |
0
![]() ![]() ![]() |
HD 97048/HD 2811 (6.19/7.067) |
![]() ![]() ![]() ![]() ![]() ![]() |
Both targets were observed in June 1998 using the IRAC2b NIR-camera at
ESO's 2.2m telescope on La Silla/Chile. Four filters were used with
central wavelengths 2.105
m (BP4), 2.121
m (BP5, used as
H2(1-0)S1 filter), 2.148
m (BP7), and 2.164
m (BP8, used as
Br
filter). The effective wavelengths were taken from the IRAC2b
manual. The use of lens C resulted in a pixel scale of 0
507 per
pixel. A mosaic of five frames was taken with the object once in the
centre of the detector and once in the centre of each quadrant. The
total integration time on source added up to 20 min. The same
procedure was performed for the standard star HD 210427 with 10 min
of total integration time. Sky subtraction, flat-fielding, bad pixel
removal and photometric calibration were achieved in the same way as
for the broad-band images. The procedure of subtracting the continuum
for photometry of the Br
emission is described in Feldt et
al. (1999). Astrometric calibration was achieved in the
same way as for the broad-band images via a comparison with the DSS.
The continuum radiation of both targets was mapped in March 1998 at
the 15m SEST telescope at La Silla/Chile. The detector system was
the He-cooled single-channel bolometer system (Kreysa 1990).
The equivalent bandwidth of the bolometer is
50GHz centred
on a frequency of
GHz (
).
The effective beam size at this wavelength is
.
The source was mapped four times with the "double beam'' technique
described first by Emerson et al. (1979). To generate
the dual beams, a focal plane chopper with a chopping frequency of
6 Hz was used. Chopping was done in azimuth. The chopper throw was
67
.
The map rows were generated by moving the telescope
continuously along the direction of the beam separation (i.e. in
azimuth) with a scanning velocity of 8
/s and an elevation
separation between adjacent scans of 8
.
Calibration maps of the planet Uranus (adopted brightness temperature
96K; Griffin & Orton 1993) were obtained with the same
technique and parameters as used for the two science targets. The
atmospheric transmission was measured by sky dips. The telescope
pointing was found to be repeatable within 5
.
Data reduction was performed with the SEST standard software and with the software package MOPSI (written by R. Zylka) which use the NOD2 and GAG libraries.
First, we present the results of the 1.3mm continuum mapping of the
two targets to give an overview of the regions where the
UCH IIs are located. The resulting maps are shown in Figs.
1 and 2 together with contours of the
8.8
m emission based on MSX images
.
The coordinate systems
in all figures relating to G11 are systems relative to the position of the
UCH II given by Kurtz et al. (1994). Coordinates for G341
are centred on the OH maser position of Caswell (1998).
Note that the 1.3mm emission consists of the sum of free-free- and dust emission at that wavelength. The free-free contribution has been subtracted from the G11 map in Fig. 1. This was achieved by multiplying the 2cm map from KCW94 with a factor of (2/0.13)0.1 (assuming optically thin free-free emission) and convolving the resulting map with the SEST beam before subtracting it from the 1.3mm map. We note that the free-free contribution to the 1.3mm flux derived by this method is not really significant. At the maximum of the 2cm map we compute a relative contribution of 5% from free-free emission to the 1.3mm emission.
As for G341, no high-resolution continuum maps currently exist at centimetre wavelengths. The non-detection of Br
-radiation from this source implies,
however, that very little ionisation might be present in that source
and thus no correction has to be made to the 1.3mm flux.
Both targets exhibit a rather irregular structure at 1.3mm. The diameters are
of the order of 2
5. However, G11 appears to be more concentrated
towards a single core, while G341 comprises two major peaks with a
clear gap in between, an indication for the presence of two dense cores.
There seems to be an east-west elongation of the core
of G11 while the 8.8
m emission of this source is extended in north-south
direction. This suggests that the heating sources are
embedded in a filament or surrounded by a dust torus, similar to NGC6334A
(Kraemer et al. 1999; Sandell 1999).
The strongest 1.3mm peak of G341 lies immediately to the
northeast of the reference position. The weaker secondary peak,
more to the southeast, coincides roughly with the position of the
infrared source IRAS 16487-4423 which Bronfman
et al. (1996) listed as a source with colour characteristics
typical of UCH IIs. We note that the 8.8
m fluxes
from the two cores of G341 behave in the opposite way to the behaviour of 1.3mm emission.
Thus, it seems reasonable that the source associated with the maximum of
the 1.3mm emission is less evolved compared to the UCH II-type
object to the southeast. The offsets between the peaks of 8.8
m and 1.3mm
emission are presumably indicating that star formation started at the outskirts
of these cores and is progressing towards their centres.
Region | Ass. temp. [K] | mm-flux [Jy] | Mass [
![]() |
G11 Core | 40 | 2.2 | 620 |
G11 Halo | 15 | 8.2 | 1.2
![]() |
G341 Core 1 | 40 | 2.8 | 400 |
G341 Core 2 | 40 | 2.6 | 370 |
G341 Halo | 15 | 8.1 | 6.0
![]() |
The flux density emitted from an optically thin dust cloud with mass
at frequency
is given by:
G11 can be divided into a central region of
diameter
exhibiting an integrated flux density of 2.2 Jy and a halo region
with
diameter and a total flux density of 8.2 Jy.
Assuming both regions to be concentric spheres and correcting for the
halo contribution in the core aperture, we derive a cloud mass of
620
in the core and 1.2
in the halo. The
temperatures used for this estimate were 40K in the core and 15K
in the halo. When using such values, one should keep in mind that by a
slight error of, e.g., 5K in the temperatures assumed above, the
mass estimates can vary by a factor of up to 2 for the halo (assumed
temperature of 15K) or 1.3 for the core (assumed temperature of
40K).
For a classification of the embedded stellar sources, it is crucial to de-redden their colours and thus to know the extinction towards these sources. This extinction can be measured in two ways, via the dust mass and via a comparison of radio free-free and NIR recombination emission. We will now derive K-band extinction values from the measured dust masses; a comparison to values derived from the second method will be given in the next section.
The complete dust mass seen in the core aperture (core plus
contributing halo) corresponds to a mass column density of
6.7
gcm-2. With a dust mass absorption coefficient of
2.15
cm-1g-1, again taken from Ossenkopf &
Henning (1994) - this time for 2.2
m, we derive a
K-band extinction of 5.5mag. Assuming the UCH II to be embedded
halfway into this cloud, we thus expect an extinction of
AK=2.8mag. Note that when using Ryter's (1996)
expression for the extinction,
![]() |
(2) |
For G341, we find two core regions (see Fig. 2).
The main peak has an integrated flux density of 2.8Jy, the secondary
peak to the southeast of 2.6Jy. A flux density of 8.1Jy remains in
the extended emission of the halo. Using the same temperatures as for
G11, we find cloud masses of 400
and 370
in the primary and
in the secondary cores, respectively, and 6000
in the halo.
Applying the appropriate extinction estimation, we derive a K-band extinction of AK=2.1 mag (AV =19 mag) in the coagulation case or AK=1.7 mag (AV =16 mag) in the interstellar grain case. Both estimates assume the source to be embedded halfway into the halo.
![]() |
Figure 1:
1.3mm continuum map of G11 with 8.8
![]() ![]() |
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![]() |
Figure 2:
Same as Fig. 1 but for G341. The contour levels for the overlayed
8.8
![]() ![]() |
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![]() |
Figure 3:
Narrow-band emission from G11 (Strong contours).
Underlayed is the narrow-band continuum image at 2.148
![]() ![]() ![]() ![]() ![]() ![]() ![]() |
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Figures 3 and 4 show the
continuum-subtracted Br
emission from G341 and G11, respectively.
All printed images taken in the infrared spectral range were filtered with the multi-scale maximum-entropy method described by Pantin & Starck (1995). This is to enhance the visibility of details in the printed version only.
No line emission is detected from G341. G11 does show Br
emission,
the integrated flux density is 0.019Jy. A comparison of the Br
flux density distribution and the 2 and 3.6cm maps were used to
derive the NIR extinction towards the source. For a thorough
description of the technique used for this comparison, we refer the
reader to an earlier paper (Feldt et al. 1998). As the
spectral index between the 2cm and the 3.6cm emission measured by
KCW94 is not exactly 0.1, as expected for optically thin free-free
emission, we use only the 2cm map for the comparison with our Br
data and assume this radiation to be optically thin. Assuming an
electron temperature of 104K we derive a mean emission measure
across G11 of 7.2
pccm-6 with a peak value of
2.8
pccm-6. Assuming the peak size to be of the order
of the VLA resolution element, i.e. 0
5 or 0.01 pc, this implies a
peak electron density of 1.6
.
The mean electron density
in G11 is about 4.3
.
When predicting the Br
flux
density from these values and comparing them to those measured, we derive a mean extinction of 3.5mag at the wavelength of
Br
.
The peak value is 4.3mag. Given the uncertainty of how deep
the source actually is embedded in the cloud core, these values agree
reasonably well with the extinction of 2.8 mag derived in Sect. 3.1.1.
![]() |
Figure 4:
Narrow-band continuum image of the area around G341 at 2.148
![]() ![]() ![]() ![]() ![]() ![]() |
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The size of the 10
m emitting region is about 1
or 3700AU
in radius. This gives a lower limit of 6.4
g
of total
(gas + dust) density or a hydrogen number density of 2.8
.
Band ID | ![]() ![]() |
Flux [Jy] | ![]() |
A | 8.8 | 9.13 | 5 |
B1 | 4.29 | <16.3 | |
B2 | 4.25 | < 8.8 | |
C | 12.13 | 11.2 | 3 |
D | 14.65 | 16.9 | 4 |
E | 21.41 | 82.3 | 6 |
![]() |
Figure 5:
Polarization map of G341. The greyscale image is the sum of our narrow-band
observations made at the NTT. Polarization vectors are plotted in
regions where the signal-to-noise ratio exceeds 3![]() ![]() |
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Imaging polarimetry at near-infrared wavelengths is a helpful tool to fix the position of the illuminating source when it his hidden from direct view and to find evidence for disk-like structures (see, e.g., Ageorges et al. 1996; Burkert et al. 2000).
Figure 5 shows the resulting polarization map of G341.
It can be seen that the degree of
polarization across the extended emission that marks G341 itself
varies between a minimum of 2% and a peak value of 35%. The peak
polarization occurs about 3
south of the intensity peak. At
this location, the polarization vectors are aligned almost perfectly
in an east-west direction.
Scattering theory implies that for single scattering, the
polarization vector is perpendicular to the line connecting the light
source and the scatterer. Thus, the location of the illuminator can be estimated
by computing the centre of gravity of all intersection points of the normals to the
polarization vectors.
In the case of G341 this was done using vectors with
polarization degrees exceeding 10%, i.e. arising from single scattering.
The formal error of the position was derived from the scatter of the
intersection points.
The result of this
procedure is shown in Fig. 5, where the ellipse
represents the 2-error of the light source
location derived from the polarization pattern. The primary illuminator
is presumably situated slightly southeast of the peak of the K' emission.
The shift may indicate an increase of the column density towards
this direction, which is also evident from the comparison of the H and K' images
(see Fig. 8). The near-infrared source geometry as inferred
from the polarization map might be as follows. Light from a luminous source which
is hidden from direct view by dust in the foreground emerges to the northeast and
is scattered by dust lanes towards the observer. This causes the arc-like pattern of
northern polarization vectors. A fraction of the light is also scattered towards the
south, where it is reflected from dust in the foreground. Since the solid angle
of the northern dust lanes as seen from the southern foreground cloud is much larger
then the solid angle of the star as seen from the northern dust lanes, the alignment
of the polarization vectors is no longer centrosymmetric but almost parallel.
Figures 6 and 7 show the results of
our adaptive optics imaging of G11. To identify the ionizing source of
the UCH II, we performed photometry on the 5 presumably stellar
sources seen inside or close to the ionized region in Fig.
6. For photometry, an IDL adaption of DAOPHOT
(Stetson 1987) was used. All five stars were used to
derive a PSF. The results of the photometry are presented in Table 5. Offset positions in this table are given from the
reference position in the figures. Columns 4 and 5 give the results of
the photometric measurements. Columns 6 and 7 summarize the absolute
magnitudes of the sources which have been derived by correcting for
the distance modulus of 13.6mag and the mean extinction of 3.5mag
in K determined in Sect. 3.2. It should be noted
that all sources except the source No. 1, are outside the region where
the extinction could be determined from the comparison of Br
and
radio emission, but at the location of No. 1 the extinction is indeed
3.5mag. Converting the absolute magnitudes into spectral types using
a zero age main sequence (ZAMS) from Straizys (1995), we
find an O5ZAMS star (source No. 1) and two stars around spectral
type B1ZAMS. For two other sources, de-reddening with the standard
extinction produces values far from the ZAMS.
ID | Offset RA [
![]() |
Offset Dec. [
![]() |
K [mag] | H-K [mag] | ![]() |
(
![]() |
Spectral Type (ZAMS) |
(1) | (2) | (3) | (4) | (5) | (6) | (7) | (8) |
1 | 0.8 | 1.3 | 13.3 | 2.1 | -3.8 | -0.1 | O5 |
2 | 2.8 | -0.9 | 15.5 | 1.9 | -1.6 | -0.2 | B0.5 |
3 | 4.1 | -2.6 | 15.5 | 3.0 | -1.6 | 0.8 | - |
4 | -3.8 | -1.4 | 14.9 | 1.1 | -2.1 | -1.1 | - |
5 | -1.1 | -1.8 | 16.6 | 2.3 | -0.5 | 0.1 | B1.5 |
![]() |
Figure 6:
K' image of G11. The logarithmic gray scale ranges
from 0.22mJy/![]() ![]() ![]() ![]() ![]() |
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Figure 7:
H-band image of G11. The logarithmic gray scale ranges
from 0.25mJy/![]() ![]() ![]() ![]() |
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The results of the AO imaging of G341 are presented in
Figs. 8 and 9. Similar to G11, the
IDL adaption of DAOPHOT was used to perform photometry on the 9 point
sources identifiable in both bands. Due to the lack of detectable Br
and free-free emission, we cannot derive exact spectral types for G341.
The results of the photometric measurements are presented as a
colour-magnitude diagram in Fig. 10. The source
magnitudes are corrected for the distance module of 12.8mag. For
the brightest K'-band source, a de-reddening vector is given that
represents the core extinction of AK = 2.5mag derived in
Sect. 3.1.1. Regions of different grey-shades represent
the reddening areas of several ZAMS spectral types. From these areas,
we conclude that we have detected two late O-type stars and several
stars of spectral type BZAMS. The bright K'-band source at the
reference position with the extended feature attached to it is not
shown in the diagram. Its distance, corrected K'-magnitude is
-0.6 mag while it has a colour index of H-K' of 5.5mag. This
source does not fall into the reddening area of the ZAMS. Such colour
characteristics may arise from excess emission from large
amounts of heated dust close to the star. The extended appearance of
this object might favour such an explanation.
As well as thermal emission of NIR radiation, such dust could
scatter photons into the line of sight that would have otherwise
escaped unnoticed by the observer. The moderate polarization level
across the source of roughly 10% supports this hypothesis.
From Fig. 1 and the estimates regarding
dust masses and resulting extinctions made in Sect. 3.1.1,
it becomes clear that G11 is located close to the centre of a large
molecular cloud. The K-band extinction, estimated in
Sect. 3.1.1 under the assumption of G11 being located
halfway into the observed cloud, amounts to 3.8mag. This matches
fairly well with the mean extinction of 3.5mag measured by the comparison
from free-free and Br
radiation in Sect. 3.2. The peak
value of 4.3mag derived by that comparison cannot be compared
directly to the dust-based estimate because it appears on much
smaller spatial scales than the resolution of the 1.3mm map can
reproduce.
We find the density of the ionized gas of 4.3
derived in
Sect. 3.2 to be slightly lower than the general hydrogen
density of 3.6
based on the dust mass. For the outer
regions, the millimetre-continuum measurements yield a lower value of
5.8
.
This means that the electron density measured from
the radio continuum at 2cm is of the same order as the measured halo
density of the G11 molecular cloud, but falls short of the measured
core density.
The coarse spatial resolution of the 1.3mm map would not reveal
the presence of a potential cavity blown
free of dust by the central O star. The density inside this cavity
might be significantly lower than in the rest of the cloud core.
Indeed, when judging from Figs. 6 and
7, the ionized region seems to be located at the rim
of a region of much higher extinction immediately to the south (no
fainter background sources south of G11 in
Fig. 6). When judging the position from
Fig. 1, the UCH II is located on the southern rim of
the cloud core. This means that a small central cavity might indeed
exist.
![]() |
Figure 8:
K' image of G341. The logarithmic gray scale ranges
from 0.14mJy/![]() ![]() ![]() ![]() ![]() ![]() ![]() ![]() |
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![]() |
Figure 9:
H-band image of G341. The logarithmic gray scale ranges
from 0.23mJy/![]() ![]() ![]() ![]() ![]() ![]() |
Open with DEXTER |
![]() |
Figure 10: Colour-magnitude diagram of the point sources in G341 |
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Figure 11 shows the spectral energy distribution of G11.
The data are taken from the IRAS point source catalogue (including the
LRS spectrum), the MSX catalogue (see Sect. 3.4) and from our
observations. Apart from the steeply rising shape of the SED up to
100
m, which is typical for UCH IIs, we notice a strong emission
feature around 7
m. This feature is probably due to polycyclic aromatic
hydrocarbons (PAHs). When we integrate the flux using a simple
trapezoidal integration, we end up with a total luminosity of
7.2
.
This corresponds to a difference of 12.14mag from
the bolometric absolute magnitude of a G2V star. According to
Straizys (1995), the resulting bolometric absolute
magnitude of -7.53mag points to a star of spectral type O7ZAMS.
We note that although the error of the rough integration may be
considerable, we expect it to be no larger than about 30%. This
estimate comes from a comparison with an artificial SED for G11 that
peaks at 200
m at 400Jy and then stays constant out to 1.3mm.
![]() |
Figure 11: Spectral energy distribution of G11 |
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Since both the SED and the ionization point to an O7ZAMS star, it
appears that the determination of the spectral type from the colour
magnitude diagram is contaminated by some additional flux in K and
H. Since the de-reddening procedure described in
Sect. 3.6.1 places the central star exactly on the ZAMS
within the photometric errors of about 0.1mag, it seems unlikely
that large errors arise from this step. Another error source might be
the slope of NIR reddening vector, but the existing
reddening laws differ little in this near-infrared region.
Contamination of the Kmagnitude by Br
-flux might be a reasonable
cause, but our measurements show that the Br
-flux contributes less
than 0.1mag to the total K-magnitude, whereas the difference in
brightness between an O7 and an O5ZAMS star is about 0.7mag. This
brightness difference means that we measure about twice the flux in
the NIR than expected from a star of spectral type O7.
Two possible explanations remain: first, the surrounding dust might scatter light into the line of sight that would not have reached the observer without this dust. The resulting anisotropic radiation distribution might resemble a source brighter in the NIR than is actually present. Future polarization measurements might clarify whether this is the case or not. Second, an unresolved binary system or even a very compact cluster might be contained in what we identify as the ionizing source. The most massive star allowed to contribute to the total luminosity of the system, besides an O7 star, would be of spectral type B0. An unresolved binary consisting of an O7 and a B0 star would have a combined K-band magnitude of 13.4mag, close to the measured 13.3mag and within the 0.1mag error limit. The colour index of the system would not be affected by the binary. The total luminosity of the system would be 30% above the result of the trapezoidal integration, only just within the limit of the above error estimate. A somewhat fainter star of spectral type B would be a better fit here. Several other stars of spectral type B have been identified inside G11 (see Table 5), so this might be asuggestive of an embedded cluster. However, our detection limit does not allow us to probe for more deeply embedded or lower mass members.
The total mass of the G11 region is 1
.
As the central
heating object appears to be of spectral type O7, its mass should be
about 25
.
Two additional stars have been identified inside G11,
both of spectral type around B1 and thus contributing a total mass of
about 20
.
Assuming that these stars contain the most
significant part of the total (young) stellar mass inside G11, we
derive an absolute lower limit for the star formation efficiency of
about 0.6% for G11. Together with the total luminosity of
7.2
,
we derive an L/M ratio of about 7 for the
complete region while that of the central source is 2500
/
.
The L/M value is often used as a measure for the star formation
efficiency. The values obtained here compare well with the results
discussed by Henning et al. (2000b) for other regions of
massive star formation.
The millimetre continuum map (Fig. 2) shows the image of a filamentary and fragmented cloud. The region of G341 contains two clearly distinct cores. These cores are separated by a projected distance of about 1.6 pc. This distance points to two independently collapsing cores of an extended cloud rather than to a binary system of two massive young stars. The object G341 studied in this paper by AO and other observations is close to the northwestern core as can be seen from Figs. 8 and 4. The southeastern core is the known IRAS source 16487-4423, probably another region of massive star formation.
In the following discussion, we will concentrate on the
northwestern core, which is also the location of the OH maser (Caswell
1998). The projected distances between the individual
infrared sources visible in our AO image in Fig. 8
range from 2
to 4
(7000-14 000AU). These
distances are comparable to those between the massive stars in the
central area of the Orion nebula near the Trapezium region. A similar
range of distances was found for the "pseudo Trapezium systems'' by
Abt & Corbally (2000).
From Fig. 10 we learn that most of the sources inside the region covered by our AO image are probably of spectral types O and B. The brightest source in the K'-band image (see Fig. 8) close to the OH maser position of Caswell (1998) is not plotted in the colour-magnitude diagram. Its H-K colour is 5.5mag. This might indicate that we are seeing extended emission from heated dust or scattered light. Scattering certainly does occur, as can be seen from the polarization map in Fig. 5.
The rather steep spectral index
between 2.2 and
10.6
m would qualify this object as Lada class I, following
Wilking et al. (1989). Dereddening can obviously not be
applied to determine the spectral type of the embedded star. Using our
four data points at 1.6, 2.2, 10.6, and 1300.0
m for a simple
trapezoidal integration, we derive a total luminosity of 195
.
Due to the lack of data points between 10.6
m and the millimetre
measurement, this should be seen as an absolute lower limit because
we miss the peak emission at mid-infrared/far-infrared wavelengths.
We note that the spectral properties of the central source are very
similar to those of the luminous young stellar object MIR1 found by
Feldt et al. (1998) slightly north of the cometary UCH II G45.45+0.06.
Evidence for an embedded massive star comes from the association of the source with an OH maser(Caswell 1998). OH masers in star-forming regions are generally associated with (compact) UCH IIs. It is mostly accepted that the OH masers are located in the compressed shell between the ionization and shock front (see, e.g., Elitzur 1987). The H II/OH masers are probably excited by a combination of collisions and far-infrared continuum radiation. Luminous sources with a low geometrical dilution factor are required to explain the high brightness temperatures of the OH masers (Pavlakis & Kylafis 1996).
Another hint towards a region of high-mass star formation is the mass of
the cloud core of 400
.
This mass is comparable to the mass of G11's
cloud core (620
)
and the core to the southeast, which is known
to be associated with a luminous IRAS source and has a mass of 370
.
Starting from our hypothesis that there is indeed a young high-mass star embedded in G341, we need to find a reason why no detectable ionization occurs in its immediate surroundings.
Due to the probable combination of reddening and excess emission it is
virtually impossible to determine the spectral type of the embedded
star from photometry. If we rather arbitrarily assume a spectral type of O8, we end up with a
factor of 40 fewer UV photons than for the central source of G11. Such
a star would emit
1048.35 Lyman continuum photons per second
(Panagia 1973). If we assume that this is also the
number of recombinations in the nebula and that the ratio of
recombinations to Br
photons is about
(Hummer & Storey 1987), we can derive the expected Br
flux in Wm-2 from
The above estimate holds for the assumption that the free-free
emission from the ionized gas is optically thin. However, this may not
be the case. The strong, extended K'-band emission indicates very
dusty surroundings very close to the central source. UV photons might
be absorbed before they can cause ionization of gas atoms. In such a
scenario, the central star might be of earlier spectral type than
derived above and still no detectable ionization would be expected.
In Sect. 3.3, we derived a lower limit of 2.8
for the hydrogen density around the embedded source from the N-band
flux. Computing the UV optical depth using Ryter (1996) and
Mathis (1990), we find that such a density would lead to
an optical depth of 1 for energetic UV photons after a line of sight
of about 100AU in length. Should this density be present very close
to the star, all of the ionizing radiation might be absorbed within a
few hundred AU. Another possible cause for the lack of visible
ionization might be the infall of material to be ionized
(see also below).
Another complication is the second brightest star in
Fig. 8 at position (+5
,
-1.5
). This
star appears in the colour-magnitude diagram about 0.7mag short of
the ZAMS after de-reddening, as can be seen from
Fig. 10. The rough determination of the extinction in
Sect. 3.1.1, the uncertainty of the embedding depth plus
additional circumstellar contributions can easily account for this
shortfall. Hence, this second brightest source is almost certainly a
star of spectral type O. Yet it doesn't produce a visible
HII region. The source might be located at a much smaller
distance and thus be not as bright in absolute magnitudes as it
appears in Fig. 10, but this is unlikely because of
the extinction which fits the one measured towards G341 very well. On
the other hand, this source might be a fully evolved O star which has
no detectable HII of its own any more. The projected distance
towards G341 is about 22200AU. The linear size of G341 in
Fig. 8 is about 11000AU. Thus the solid angle
under which G341 is seen from the O star is about 0.16sr. Assuming
the O star to be of spectral type O7, this means that
1046.72Lyman continuum photons per second should hit the dust shell of G341.
Accordingly, gas atoms at the surface of this dust shell should be
ionized. Again, we can convert this number into an expected Br
flux
using Eq. (4), and we end up with an expected flux
density of
mJy. This is still above our detection
limit but only by a factor of about three. Intermediate absorption by
dust or a slight error in the assumption of a spectral type can
explain why we would not detect Br
radiation from this type of
ionization.
So far, we have been unable to determine exactly the type of the central source embedded in G341. However, we have argued that it might be a massive star and if it is, it must be so heavily enshrouded by dust that all of its ionizing radiation is absorbed before it can produce a detectable HII region or the HII region is suppressed by the accretion of matter (see, e.g., Yorke 1986; Henning 1990; Testi et al. 1997). The OH maser and the total mass of the cloud core are arguments in favour of a massive star-forming region. For a discussion of similar objects using radio continuum and OH maser data, we refer to Forster & Caswell (1999).
The strong excess emission in K' shows that large
amounts of hot dust exist very close to the star. Its
morphology would be governed both by local density variations and
the temperature gradient. A very crude comparison to the model of
Testi et al. (1997) shows that we should expect a colour
index of K-H of about 5 mag on the central peak close to the star
and about 1.5mag in outlying regions. We derived a colour index of
5.5mag for the central peak in K and measure a gradually
decreasing colour index towards north/northeast which reaches 2.5mag
at the northern end of the extended emission. Differences in
foreground extinction and the density profiles can easily account for
the differences. We should note however, that in contradiction to the
model of Testi et al., in G341, intensity variations do not happen on
scales of 1
pc but on scales of 1
pc.
Our data demonstrate that both massive molecular clouds investigated by us contain clusters of OB stars, supporting the view that massive stars do not form in isolation. The object G11.11-0.40 is one of the rare cases where the central engine of an ultracompact H II region could be identified. This embedded O star is an excellent target for future near-infrared spectroscopy with very large telescopes in order to derive the exact spectral type of the object. There is some evidence that the ionizing source may even be a binary. The features of the object G341.21-0.21 suggest that it is in an early stage of massive star formation, presumably before the UCH II phase. Sensitive radio maps and molecular line data taken with interferometers are required to check this statement.
Acknowledgements
We would like to thank our referee, John Mathis, whose comments helped to improve the paper. B.St. acknowledges support from the DFG grant Ste 605/17-1.