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Astronomy & Astrophysics manuscript no. iras16362 July 19, 2007

c ESO 2007

A close look at the hear t of RCW 108
´ F. Comeron and N. Schneider
1

1

2

¨ ESO, Karl-Schwarzschild-Str. 2, D-85748 Garching bei Munchen, Germany e-mail:

fcomeron@eso.org nschneid@cea.fr

2

SAp/CEA Saclay, F-91191 Gif-sur-Yvette, France e-mail:

Received; accepted
ABSTRACT Context. The IRAS 16362-4845 star-forming site in the RCW 108 complex contains an embedded compact cluster that includes

some massive O-type stars. Star formation in the complex, and in particular in IRAS 16362-4845, has been proposed to be externally triggered by the action of NGC 6193.
Aims. We present a photometric study of the IRAS 16362-4845 cluster sensitive enough to probe the massive brown dwarf regime. In

´ particular, we try to verify an apparent scarcity of solar-type and low-mass stars reported in a previous paper (Comer on et al. 2005, A&A, 433, 955).
Methods. Using NACO at the VLT we have carried out adaptive optics-assisted imaging in the J H KS L bands, as well as through

narrow-band

lters centered on the Br
S



and the H

2

S (1) v

=

1



0 lines. We estimate individual line-of-sight extinctions and, for
S

stars detected in the three J H K

lters, we estimate the contribution to the K

ux caused by light reprocessed in the circumstellar

environment. We also resolve close binary and multiple systems. We use the K luminosity function as a diagnostic tool for the characteristics of the underlying mass function.
Results. IRAS 16362-4845 does contain young low-mass stars. Nevertheless, they are far less than those expected from the extrapo-

lation of the bright end of the K luminosity function towards fainter magnitudes. We estimate a total stellar mass of 370 M . Nearly all the cluster members display L excesses, whereas K an extremely red object with ( K
S

-

L)

>

S

excesses are in general either absent or moderate (

<

1 mag). We also detect

9, likely to be a Class I source.

Conclusions. The fact that solar-type and low-mass stars are present in numbers much smaller than those expected from the number

of more massive members hints at an initial mass function de cient in low mass stars as compared to that of other young clusters such as the Trapezium. The origin of this difference is unclear, and we speculate that it might be due to external triggering having started star formation in the cluster, perhaps producing a top-heavy initial mass function. We also note that there are no detectable systematic differences between the spatial distributions of bright and faint cluster members. Such absence of mass segregation in the spatial distribution of stars may also support external triggering having played an important role in the history of the RCW 108 region.
Key words. ISM: HII regions - ISM: individual object: RCW 108 - Stars: luminosity function, mass function - Open clusters and

associations: IRAS 16362-4845

1. Introduction
The RCW 108 HII region in the Ara OB1 association has long been regarded as one of the best case studies illustrating the eroding action of newly formed clusters containing massive stars on the molecular gas in their environments (Shaver & Goss 1970, ´ Straw et al. 1987, Comeron et al. 2005). In visible-light images of the region (see e.g. Petersen 2001) RCW 108 appears as a bright rim nebula (SFO 79 in the catalog of Sugitani et al. 1991) on a size scale of several 10 arcmin, de ning a sharp boundary between an extended molecular cloud in the west and a region predominantely lled with ionized gas in the east. The rim represents the ionization front, produced when the ultraviolet radiation from O-type stars of the neighboring cluster NGC 6193 hit and progressively destroy the molecular cloud. Infrared sources indicating active star forming sites are common among externally ionized molecular clouds (Sugitani et al. 1989, 1991, 1995; Sugitani & Ogura 1994), which has been
´ Send offprint requests to: F. Comeron Based on observations obtained at the European Southern Observatory using the Very Large Telescope (VLT) (programme 077.C0660(A)), on Cerro Paranal, Chile.

interpreted as evidence for star formation triggered by radiationdriven implosion of dense cores (e.g. Bertoldi 1989, Miao et al. 2006). This mechanism may well be at work in RCW 108, as noted by the indications of triggered star formation in the re´ gion recently discussed by Urquhart et al. (2004) and Comeron et al. (2005, hereafter CSR05). The most conspicuous star forming site in RCW 108 is the compact HII region IRAS 16362-4845, rst noted by Shaver & Goss (1970). It is deeply embedded in the dense molecular cloud and associated with a cluster of IR-sources discovered by Straw et al. (1987). Its stellar contents has been more recently studied by Urquhart et al. (2004), based on the 2MASS catalog (Skrutskie et al. 2006); by CSR05 by means of dedicated deeper, higher resolution J H K
S

imaging; and by Wolk et al. (2007) us-

ing X-ray emission as a tracer of young stellar populations. The study of CSR05 indicates that the IRAS 16362-4845 embedded aggregate is a Trapezium-like cluster containing at least one late O-type star, in consistency with the visible spectrum of the heavily obscured HII region. The mass of the cluster was estimated by CSR05 to be



210 M , with a rather large uncertainty.

Furthermore, that study noted as an intriguing feature of the color-magnitude diagram an apparent lack of stars fainter than


2

´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

M

K

+1.

0 and with amounts of foreground reddening in the

Table 1. Log of observations Filter J H K
S

range covered by the brightest stars in the cluster, which might be indicative of a peculiar mass function. While the result hinted at by CSR05 may be potentially relevant to understand the build-up of the initial mass function (IMF) in clusters dominated by massive stars, some practical limitations of the observations presented in that work advised a further analysis based on material of higher quality. First and foremost, though the observations were deep enough to penetrate well into the area of the color-magnitude diagram where the lack of the stars was noted, the completeness of the census in that range was di
L 2.17 2.12 0.175 s 65 65 180 1 1 28 16 16 DI T 13 s 20 s 20 s N DI T 5 4 4 NEXP 16 16 16 Date 29/30 June 2006 29/30 June 2006 27/28 July 2006 22/23 July 2006 24/25 July 2006 29/30 June 2006 29/30 June 2006 22/23 July 2006 22/23 July 2006 17/18 July 2006 28/29 July 2006

ffi

cult to assess. This was due to the presence of bright

nebulosity pervading the cluster with large brightness variations over small angular scales. Secondly, although the observed stellar images have a full-width at half maximum (FWHM) below one arcsecond, the combination of crowdedness and nebulosity still hampered the detection of faint members relatively close to brighter stars. For these reasons, higher quality observations were needed to place our tentative conclusion of a de cit of faint members in IRAS 16362-4845 on a In this paper, we present rm standing. new observations of the

J , H , and K

S

bands and the narrow-band. The resulting

eld

size, 54 , matches well the angular size of the IRAS 163624845 cluster (



30 ), whereas the location of CSR-012A near

the northern border of the cluster allows us to sample a portion of the dark cloud located to the north of IRAS 16362-4845. For the L band observations the maximum pixel. Our observations through that eld size of the eld attainable with NACO is 27 , at a scale of 27 milliarcseconds per lter thus sample only the densest part of the cluster centered on CSR-012A. The observations were obtained on six different nights between 29/30 June and and 28/29 July 2006. The images were performed through two series of exposures in each of the J , KS , L , 2.12 the H

IRAS 16362-4845 cluster carried out using adaptive optics nearinfrared imaging at the Very Large Telescope (VLT), which provide a far deeper and sharper view of the cluster than previously available. This new material allows us to reassess the stellar contents of the cluster, and address questions related to possible peculiarities of its IMF, the frequency of infrared excesses among its members, the abundance of massive binary stars, or the spatial distribution of high- and low-mass members, also providing some further insights on the structure of the associated nebula.

µ

m, and 2.17

µ

m

lters, and one series of exposures in

lter. We used the common technique of stacking a num-

ber N D I T of frames, each with a detector integration time D I T , centered on a number N E X P of closely spaced telescope point-

2. Observations
2.1. Data acquisition and reduction

ings on a random dither pattern, for which we used an amplitude of 10 . The image resulting from each exposure was constructed by combining the sky-subtracted individual stacks after correcting for the telescope offset between pointings. This procedure was also used for the L band observations, rather than the alternative chop-and-nod technique. Table 1 gives the log of the observations, including the individual exposure parameters. Sky subtraction from our on-target images was performed by stacking intercalated images obtained around a separate sky position located 3 away from IRAS 16362-4845, as the bright nebulosity pervading the cluster prevented us from using for that purpose the median- ltered on-target frames uncorrected for telescope offsets, as is normally done on uncrowded, nebulosity-free The detector readout mode was selected for each range available. For the elds. lter as the best

Our observations were carried out in service mode using NACO, the adaptive optics near-infrared camera and spectrograph at the VLT (Rousset et al. 2002, Lenzen et al. 2003), in imaging mode. Broad-band images were obtained through the J (1.26

µ

m), H (1.66

µ

m), KS (2.18

µ

m), and L (3.80

µ

m)

lters,

as well as through two narrow-band and 2.17 H2 S (1)

lters centered on 2.12 brightest star in the

µ

m

µm respectively sampling the v = 1 0 and Br lines. The
1 S

nebular emission in the eld

in KS -band images of the cluster, star #12 from CSR05 (hereafter CSR-012A ) has K

= 8.

36 and was used for wavefront

sensing, using the near-infrared wavefront sensor in NAOS, the adaptive optics module of NACO. Different available dichroics were chosen, depending on the band of the observations: the K dichroic for the J and H observations (90% transmitted light in those bands), the J H K dichroic for the L the KS , 2.12 observations (also 90% transmission in that band), and the N20C80 dichroic for

compromise between the necessary sensitivity and the dynamic lters in which two images were obtained on separate observations, one combined frame was produced and used for further analysis. Figure 1 shows a composite of the nal frames obtained through the J H KS lters.

µ

m, and 2.17

µ

m observations (80% transmission).
2.2. Photometr y

These choices provide the best compromises between the signal from the wavefront reference star needed for a good adaptive optics correction and the light transmitted to CONICA, the camera of NACO. We selected the wide- eld camera optics yielding a pixel scale of 54 milliarcseconds per pixel for the images in the
1

Our observations in the J H KS bands were calibrated using the infrared photometry of non-saturated sources in common with those listed in CSR05. The L photometry was calibrated using as a reference the stars in the Straw et al. (1987). The intrinsic di eld observed in that same band by

We use the nomenclature CSR-nnn to refer to the stars listed in

Table 2 of CSR05. Stars detected in the present work and not listed in CSR05 are named CS-nnn. Members of multiple systems are referred to by adding A, B,... to their numbers. In this way, CSR-012A is the brightest member of the system collectively identi ed as Star #12 in CSR05, which the current observations show to be composed of members. ve

ffi

culty of performing stellar photometry in eld, is

observations using adaptive optics, due to the noticeable variation of the point-spread function (PSF) across the compounded in our case with the additional complication of a variable nebular background pervading the area of the clus-


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

3

Fig. 1. A color composite of the frames of IRAS 16362-4845 obtained through the J (blue), H (green) and KS (red) lters. The frames are centered h m s on the brightest source at near-infrared wavelengths, CSR-012A ((2000) = 16 40 00 2, (2000) = -48 51 40 ). The eld covered is 52 â 57 , with North at the top and East to the left. A comparison with Figure 12 of CSR05 gives an appreciation of the increase in quality between the observations discussed in that paper and those presented here. The small red ring just outside the upper right edge of the nebula is actually an artifact due to out-of-focus ghosts of the brightest star in the cluster. The dashed line corresponds to



(2000)

=

16 40 01 78. As explained in

h

m

s

Section 3.2 contamination by background sources is higher to the East of this line, and we exclude that region from our analysis.

ter, which makes it di

fficult

to accurately estimate the contribu-

The radius of the aperture for photometry was chosen as

ve

tion to the measurement of the extended wings of the PSF. The strong and complicated variation of the PSF, particularly in the J band where the adaptive optics correction is poorer and degrades fastest with the distance to the wavefront reference star, led us not to consider PSF tting as a suitable method. Instead, we obtained better results by performing aperture photometry at the position of each detected object. Source detection was carried out automatically using the DAOFIND task in the DAOPHOT package layered on IRAF (Stetson 1987). The choice of parameters of the task, in particular the sharpness and roundness cutoff parameters, was made by successive trials by comparing the results of the automatic detection with a careful visual inspection of the frames. We identi ed the sources that were left undetected by DAOFIND as well as the false detections, relying on the fact that the human eye ultimately provides the best discriminant between real stellar sources and detection artifacts.
2

times the FWHM of the PSF of a bright, non-saturated star located 15 away from the wavefront reference star, which provided an average PSF of the eld. We veri ed that the relatively large number by which the FWHM of the PSF of that star was multiplied ensured a negligible aperture correction elsewehere in the eld. The counts within this aperture were computed by dividing the aperture into concentric rings each with a width of one pixel. Pixel values deviating by more than 3



from the average

2

value within each ring normally denoted the existence of a close companion to the star being measured, and were thus replaced by the average of the other pixels of the ring. For some faint stars located in regions of a strongly variable background this procedure prevented us from obtaining reliable magnitude measurements. Table 2 gives the photometric measurements obtained for all the stars in our images, and is available in full electronically.

IRAF is distributed by NOAO, which is operated by the Association

of Universities for Research in Astronomy, Inc., under contract to the National Science Foundation.


4

´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

Table 3. Completeness and detection limits Filter J H K
S

inant in the L band, and other bright members appear elsewhere in the cluster. The close correspondence between the cluster and the bright nebulosity is well apparent from Fig. 1, although some cluster members appear projected beyond the boundaries of the nebula. The approximate diameter of the cluster is 40 , corresponding to 0.25 pc at the distance of 1.3 kpc (distance modulus 10.6) that we adopt in this paper (see Arnal et al. 2003). The appearance of the nebulosity is reminiscent of a cocoon surrounding most of the stars, except towards the East where the nebulosity fades away without a border as well de ned as towards the other directions. This is well seen in Fig. 2, where the H
2

completeness 18.6 19.1 18.0 13.5

detection 21.2 21.5 20.6 14.1

L

2.3. Completeness limits

The magnitude completeness limits of our observations are not straightforward to assess due to the same factors that complicated the photometric measurements. On average the better adaptive optics correction obtained near the wavefront reference star, which leads in principle to deeper detection limits, is partly offset by the fact that this star lies near the peak of the surface brightness of the nebula. In the H and KS band images, where the adaptive optics correction is best and improves image quality over a larger eld, the deepest detection limits are actually obtained near the edges of the frames, where stellar images are not as sharp as near the centers but the background nebulosity is nearly absent. Since knowledge of the completeness of the stellar census of the cluster is of great importance for the ensuing discussions on the luminosity function of the cluster, we performed numerical simulations on the reduced frames using arti cial stars. To this end, stars at different positions of the eld were used as local PSF references, and arti cial stars of various magnitudes and the same PSF were added within a radius of 10 from each local PSF reference. The same detection procedure and parameters used for the detection of the real sources were then applied to the image containing the arti cial stars. The completeness limit was then adopted as the magnitude of the arti cial stars for which at least 95% of them were recovered by DAOFIND in the surroundings of any of the local PSF reference stars used to de ne them. The location of such reference stars in all representative regions in terms of nebular brightness and adaptive optics corrections ensures that any star brighter than the completeness limit thus dened (except those in the close proximity of the brightest stars) are indeed detected and incorporated to our census of members of the IRAS16362-4845 cluster. Nevertheless, it was found, as expected, that many regions of the cluster provided considerably deeper detection limits, and our census thus contains numerous stars below the completeness threshold. Table 3 lists both the completeness and 5

nebulosity wisps tracing the edges of photodissociation

regions at the interface between the molecular cloud and the HII region are absent towards the East, and where the Br



also

brightens near the Western edge. The edge is actually the brightest feature of the HII region at longer wavelengths, as shown by the image obtained with the Spitzer Space Observatory using the IRAC camera at 8

µm,

available from the Spitzer archive, where

emission is dominated by PAH molecules (see Fig. 4). The overall morphology of the nebula may be due to the overpressure of the cluster gas originated by its ionization having been released by an out ow towards the East. Both the molecular-line and the H observations presented in CSR05 support this interpretation. As noted in that work, the cluster and HII region are not positionally coincident with the peak intensity of the molecular emission, which lies to the West, indicating that the expansion of the HII region nds less resistance in the Eastern direction. Furthermore, the CO intensity contours presented in CSR05 show a prominent pinching just East from the cluster suggestive of a sharp decrease in column density. Finally, H



radial velocity maps show a rather

sharp change at the same point (see Fig. 8 of CSR05), superimposed on a shallower, larger-scale gradient in the Southwest to Northeast direction seen in both H in our direction. Wider- eld images including the surroundings of IRAS 16362-4845 show extended nebular emission towards the East (see Fig. 2 of CSR05), clearly related to the cluster and probably tracing the less dense regions of the ionized out ow, but disconnected from the main body of the compact HII region by a dark patch due to a dense, foreground cloud. The Western border of this cloud appears near the left edge of our NACO images. Interestingly, our H and KS images show a strip roughly 10



and molecular line observa-

tions, as would be expected from an out ow with a component

â

15

across crowded with numerous faint, lightly reddened

stars located between the cluster and the dark foreground cloud (see Section 3.2). The very con ned location of those stars leads us to discard the possibility that they may be cluster members; if they were, most of them would be brown dwarfs based on their faintness and their relatively blue colors. Instead, we consider far more likely that they are background stars seen through a low column density gap just East of the cluster. Although the



detection limits, taking into account that

the latter apply only to limited areas of the region imaged and are generally brighter elsewhere where the nebulosity is more intense.

3. Results
3.1. Overall structure of the cluster and its nebula

dust column density in their direction is far lower than that estimated from the molecular gas column density obtained from
13

CO J

=

1



0 observations of CSR05 ( A

V

reaching 70 mag

Figure 1 presents an overall view of the structure of the cluster and the associated nebulosity, and Figure 2 gives a more detailed view of the nebula through the narrow-band lters centered on the line emission of Br and H2 . At near-infrared wavelengths the cluster is dominated by the likely O-type star CSR-012A ( IRS29 in Straw et al. 1987), near the Northern edge of the HII region. Our new observations reveal it to be a tight concentration of ve stars within 1 1 of the primary (Fig. 3). As already noted by Straw et al. two other stars to the East of the cluster, CSR-18 and CSR-20 (

near IRAS 16362-4845, as compared to AV as low as size given above is not ruled out by the molecular-line observations at 45
13



10 mag 1

from the colors of those sources), a low extinction hole of the CO J

=



0

angular resolution.

=
3.2. Color-magnitude and color-color diagrams

The color-magnitude diagram of the cluster is presented in Figure 5. In it we have made a distinction between the sources located in the Western region of high extinction on the back-

=

IRS19 and IRS20 in Straw et al.), become dom-


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

5

Fig. 3. L -band image of the region around CSR-012A, the brightest star in the K -band, showing the rich concentration of cluster members around it.

Fig. 2. A comparison between the images obtained through the narrowband lters centered on the Br



(top) and the H

2

S (1) v

=

1



0 Fig. 4. A wider- eld view of the IRAS 16362-4845 region obtained with the IRAC camera on board of the Spitzer space observatory at a central wavelength of 8

(bottom) lines. Both images are on the same intensity scale. More stars are visible in the H 48
2

image due to the lower nebular intensity and the eld covered by each image is

better adaptive optics correction. The

â

48 .

µm.

The western edge of the nebula is marked

by a bright, crescent-shaped rim probably dominated by PAH emis-

ground and the Eastern region where, as discussed in the previous Section, the extinction appears to be much lower allowing the detection of numerous background sources. For simplicity we have separated the sources in these two categories according to their position with respect to the line
h m s

sion. Some of the brightest cluster members have strong mid-infrared excesses and can be identi ed in the Spitzer image. The image corresponds to AOR key number 12486912 (Program # 112, IRAC and IRS observations of RCW 108, Principal Investigator G. Fazio), and has been obtained from the Spitzer archive.

(2000) =

16 40 01 78 and we base our discussion on the cluster contents on the Western region alone. The exclusion of the Eastern region, which occupies 20% of our images, is likely to exclude as well some cluster members, but simple visual inspection of Figure 1 indicates that the vast majority of the cluster members lie in the Western region where contamination by background sources is very low. The overall appearance of the cluster ( H tion along the line of sight towards each star, and to the varying amounts of infrared excess produced by the circumstellar material associated to each object. Given the deeply embedded nature of the cluster and its youth we expect both causes to signi cantly contribute. The existence of sources with large amounts of circumstellar dust reprocessing the light of the central object is conrmed by the ( J by the ( H

-

KS ), K

presented in Figure 5 shows a broad distribution in ( H

-

S

diagram

-

H ), ( H

-

KS ) color

-

KS ), ( KS

-

KS ) diagram (Fig. 6) and, most clearly,

L ) diagram (Fig. 7), where many stars (in

which can be due both to variations of the foreground extinc-

the latter diagram, most of them) lie to the right of the reddening


6

´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

Fig. 6. ( H Fig. 5. ( H

-

KS ), ( J

-

H ) color-color diagram of the IRAS 16362-4845

cluster. Symbols are the same as in Figure 5. The solid line represents

-

KS ) , K

S

Color-magnitude diagram of the IRAS 16362-4845

the position of 1 Myr old unreddened stars with no circumstellar infrared excess from the models of Palla & Stahler (1999). The dashed lines are reddening vector whose lengths corresponds to an extinction AV

cluster. Full large dots are stars located in the region dominated by the cluster, whereas the small open circles are objects located in the lowextinction region to the East of



(2000)

=

16 40 01 7 and dominated

h

m

s

=

40 mag. The top vector has its origin at the position of a 1 M

by faint background sources. The dotted line marks the position of the completeness limits in H and K
S

star, and the bottom has its origin at the position of an early O-type star. This latter vector delimits the area accessible to a normal photosphere reddened by extinction, and the position of objects to the right of this line denotes the existence of infrared excess.

given in Table 3. The solid line is

the 1 Myr isochrone of the models of Palla & Stahler (1999) using the infrared colors derived by Testi et al. (1999). The dashed line marks the position of a 1 M star obscured by varying amounts of extinction.
V

The lower right end of the line corresponds to A

=

40 mag. Finally,

the dot-dashed line delimits the area where CSR05 noted an absence of faint, lightly reddened stars, which is further discussed in Sect. 3.2. The bottom of the area corresponds to the approximate detection limit of the observations of CSR05.

vector delimiting the region of the diagrams accessible to normal photospheres obscured by different amounts of extinction; see Sect. 3.4. For reference, the solid line in those diagrams is the unreddened 1 Myr isochrone using the evolutionary models of Palla and Stahler (1999), complemented with the J H KS colors computed for them by Testi et al. (1999) and shifted to our adopted distance modulus D M

=

10.6.

The isochrone plotted in Fig. 5 demonstrates the sensitivity of our observations to the entire range of stellar masses and even to massive brown dwarfs. The upper end of the isochrone shown in this Figure corresponds to a mass of 30 M , whereas the lowest end marks the position of a 0.01 M brown dwarf.
Fig. 7. ( K

Because of the narrow range of infrared colors covered by normal stars, the main effect of assuming a different cluster age is to shift the isochrone vertically in the color-magnitude diagram, thus changing the mass corresponding to an object of a given absolute magnitude in the sense of the mass being lower for a younger assumed age. Although the precise masses of the intrinsically faintest objects detected in our images critically depend on the age of the cluster, we can con dently state that the completeness limit of our observations probes the brown dwarf regime for any plausible age of the cluster. Indeed, a star at the hydrogen-burning limit ( M magnitude M K (Baraffe et al. 2003). The age of the IRAS 16362-4845 cluster is likely to be much less than that (see Sects. 9 and 4). It is interesting to compare the results obtained here with the claim made in CSR05 about the apparent lack of cluster members in the ( H
S

-

L ), ( H

-

KS ) Color-color diagram of the IRAS 16362-

4845 cluster. Symbols and lines have the same meaning as in Fig. 6, but now we have not plotted the reddening vector with its origin at the position of a O-type star (bottom left end of the solid line) since it does not separate anymore the regions accessible and inaccessible to normal photospheres reddened by extinction.

=

0.078 M ) reaches an absolute

=

18.1 minus the adopted D M
K

K

= 6.

5 (corresponding to our completeness limit

by a foreground extinction of A

= 10.6 of the = 1 mag) at

cluster, obscured an age of 6 Myr

-

KS ), KS diagram lying below an ex-


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

7

tinction vector having its origin near the position of a mainsequence A0 star. Given the severe incompleteness of the observations presented in CSR05 above K
S

eral related. Extensive grids of models such as those computed by Lada & Adams (1992) for Herbig Ae/Be stars and classical T Tauri star disks indicate that families of models produce shifts in the ( J

15 in the area of

the cluster due to the bright nebulosity, the region of the colormagnitude diagram that was claimed to be devoid of stars in that work is effectively delimited by the polygon joining the points [( H

-

H ), ( H

-

KS ) diagram that can be approxi-

mately characterized by an infrared excess vector with a slope

=

(J

-

H )/( H

-

-

KS ), KS ]

=

(0, 11.2), (0, 15.0), (1.3, 15.0), and (1.3, 13.6),

KS ). A good approximation for classical

T Tauri stars with surface temperatures T

=

3000 K is



0.75,

where only the likely foreground star CSR-013 is detected in their observations. Our present work reveals four stars in that area, whose non-detection in CSR05 is readily explained by examining their location: one of them, CSR-012E, is a faint companion just 11 from the bright star CSR-012A. The second is CSR-008B, very close (022) to the blue, likely foreground star CSR-008A. The third is CSR-008A itself, to which CSR05 allocated an infrared excess that our new observations show to be actually due to the presence of the much redder component B. The true ( H

similar to the values derived from the disk models of Meyer et al. (1997). For Herbig Ae/Be stars with disks having central holes and inner edge temperatures T ally provide the best

=

2000 K, which gener-

ts to the observed colors of these objects,


to

0.56. If we assume a negligible infrared excess at J , where ux is lowest, the approximate condition E

the ratio between luminosity reprocessed at the disk and photospheric
J

=

0 leads

-

KS ) color is consistent with it being a foreground, implying E 0.36 E
K
S

E
H

H

= /( +
S

1) E

K

lightly star when component B is excluded. Finally, the fourth star CS-104, is faint and near CSR-012A, only 15. The depopulation of that area of the color-magnitude diagram in CSR05 is thus real, as only close companions to other stars populate it in the present study.

S

(2)
H

=

0.43 E

K

for Herbig Ae/Be stars and E
H

=
K
S

for classical T Tauri stars. We thus adopt E

= 0.

39 E

for our analysis as an average value characterizing the typical value of the circumstellar excess vector. Finally, since the model isochrones provide values of ( J0 , H0 , KS 0 ) at n discrete masses rather than a continuous curve, we must interpolate between the given points. We do this by means of a simple linear interpolation, such that the line joining two points i and i by the continuous variable becomes J H

3.3. Correcting for extinction and infrared excess

Near-infrared observations, most notably in the KS band, present well known advantages for the observational characterization of the population of a cluster: they allow one to probe down to very low masses given that most of the luminous output of the least massive stars and brown dwarfs lies in the near-infrared, and they reduce the effects of extinction with respect to shorter wavelength bands. Nevertheless, the K
S

+

1 of the

isochrone is approximated by a straight segment parametrized



. Taking all this into account, Eq. (1)

=
i

J0

i

+

( J0i+1

-

J0i )

+ 2.

518 A
K

K

(3a) 39 E
K K

band also probes a spec-

tral region where circumstellar emission can signi cantly contribute to the luminosity of star and even dominate over the photosphere, and the greater transparency of dust at that wavelength increases the level of contamination of background sources for clusters residing in clouds of low or moderate column density. Although the large dynamic range in luminosity covered by our observations is in principle a very useful resource to study the stellar population of IRAS16362-4845 over a wide range of masses, the particular conditions of the cluster require special care in both deriving and interpreting the properties of its members. To approximately correct the observed magnitudes for the effects of extinction and circumstellar emission, let us write the observed magnitude of a star in each of the J , H , and KS bands as m
3

=

H0

+
KS
0i

( H0i+1

-

H0i )
1

+ 1.
KS

563 A
0i

- 0.
K

(3b) (3c)

KS

=

0i

+

( KS 0i+

-

)

+

A

-

E

where J0i , H0i , KS

are the magnitudes given by the models at

the isochrone point i. The solution to system (3) yields the values of the three unknowns



, AK , and E K . It may be noticed

that a similar procedure has been recently used by Figueredo et al. (2005) in their analysis of the cluster associated to the giant HII region G333.1-0.4. They correct for infrared excess assuming that it is noticeable only in KS , and then correct of extinction by dereddening along the limiting reddening vector in the (J

-

H ), ( H

-

KS ) diagram marking the boundary of the region

accessible by objects free of infrared excess. As those authors point out, this procedure yields a lower limit to E K , as well as an overestimate of AK , and is equivalent to setting

=

0 in Eq. (2).

However, both models and observations suggest that the excess

=

m0

+

A

+

E
0

(1) is the

at H cannot be neglected in general. Furthermore, the method that we have used does not use the limiting reddening vector, but an actual reddening vector having its origin at the photospheric colors of the best KS tting theoretical model instead. For the stars having measured magnitudes in the J , H , and lters we have solved the system (3) for 1 ing as the best solution the one for which 0 which the best

where m is the magnitude in one of those bands, m

magnitude free of extinction and infrared excess given by the model isochrone used, A is the extinction, and E is the infrared excess. The values of A are related by the extinction law; adopting the standard values from Rieke & Lebofsky (1985), A
J

=

2.518 A

K

S

,A

H

=

1.563 A

K

S

. On the other hand, cir-

i

n

-

1, tak-

1, i.e., for

cumstellar disk models show that the values of E are in gen3

tting model is an interpolation between consec-

utive points along the isochrone. As may be noticed in Fig. 5, isochrones within a certain range of young ages present a kink giving rise to a small mass interval in which stars of a given mass are brighter than those slightly more massive, due to their slower evolutionary rate. In those narrow intervals it is possible to more than one solution to the system (2) with 0 of little practical concern, as the values of A H0 , K
S0 K

L measurements are also available for many stars near the center

of the cluster, and the procedure described here can in principle be easily extended towards that band as well. Nevertheless, since L strongly dominates over the photosphere for many objects in our sample, the t would become greatly sensitive to the approximation used to represent the excess in that band, which is necessarily rough. We have thus preferred to base our analysis on the J H K
S

nd



1. This is

and E K , and of J0 ,

bands alone.

are very similar for all solutions. The derivation of the


8

´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

K luminosity function (see Sect. 3.5) is thus hardly affected by the choice of the solution in those kinks. The A or E tting procedure described above yields unphysical sothat may be due to photometric errors, to the simplilutions for some objects, with slightly negative values of either
K K

over time, and the choice of theoretical pre-main sequence evolutionary tracks. To derive the KLF intrinsic K
S 4

of IRAS 16362-4845 we have derived

magnitudes using the method described in the previ-

ous Section for all the stars detected in at least two of the J H K bands. We have excluded a few sources with H

cations adopted for the reddening and infrared excess corrections, or to deviations between the colors predicted by the models and those of the actual stars. We have dealt with the cases where the solution to Eqs. (3) yields E tem again by least squares, setting E
K K

-

KS

<

0.5,

S

as they are most likely foreground stars. Also, we have considered components of close binary systems as separate stars only when reliable photometry could be obtained for each of them (see Sect. 3.6). Since we are excluding the Easternmost area of the imaged eld due to the apparent extinction hole in that region, as described in Sect. 3.2, we expect the population analyzed here to be strongly dominated by cluster members, as the thick dust column provides an efficient screen against the background population even at the KS band. A few sources have very red colors and may be background to the cluster, but as discussed in Sect. 3.7 this is unlikely to be the case even for the reddest sources. The KLF that we derive is potentially affected by several biases that must be carefully taken into account. In addition to the varying detection thresholds across the eld discussed in Sect. 2.3, the varying level of foreground extinction implies that the sampled volume of the cluster becomes progressive smaller for intrinsically fainter sources, as the most obscured objects of a given intrinsic luminosity fall below the detection limit. On the other hand, some faint members that would have remained undetected if their emission were purely photospheric may be brought above the detection threshold by their infrared excess. However, based on the moderate or altogether absent KS -band excesses that we generally obtain for the stars having J H KS photometry, we estimate the in uence of infrared excesses to be small in the derivation of the KLF. To minimize the bias due to the variable sampling of the cluster volume due to extinction, we have obtained a KLF by considering only stars having A
V

<

0 by solving the sys-

to zero. On the other hand,
K

the only case in which the solution to Eqs. (3) yields A dening vector in the ( J

<

0

corresponds to the object farthest to the right of the limiting red-

-

H ), ( H

-

KS ) diagram. For this star
S

we have assumed that all the excess lies in the K

band and we

have solved Eqs. (3) again, now replacing the coefficient 0.39 in Eq. (3b) that accounts for the H -band excess by zero, thus obtaining solutions in which both A
K

and E

K

are positive. This

might be an object possessing a large central hole devoid of the hot dust component that provides the main contribution to the infrared excess at the shorter wavelengths. Finally, many of the fainter objects in our sample have measurements in only two bands, generally H and KS . We have assumed that they do not have infrared excess and have solved the relevant subset of Eqs. (3) for



and AK . Since our results for
S

the objects with measurements in J , H , and K those that do the appropriate value of E

show that some-

what more than half do not require infrared excess, and that for
K

is generally small (see
K

Sect. 3.4) we do not expect that the assumption of E

for objects

with only one measured color introduces a signi cant bias.

3.4. Infrared excesses

The list of stars for which an infrared excess in the K band is needed to obtain a good t to the colors predicted by the modmeasurements in our t their J H K colors. els using Eqs. (3) is given in Table 4. Taking into account that the total number of sources having J H K
S

<

23.8 and M

K

< 4.

0. Both

constraints are related by the condition that the intrinsically faintest and most obscured stars in such sample are at the overall completeness limit of our observations (see Table 3). Indeed, a star with M have H

sample is 44, our results imply that infrared excesses need to be assumed for 45% of the stars in order to spurious in some cases, as the values of E
S

This fraction is uncertain and the derived infrared excess may be
K

often are of the order

=

19.1, KS

K

=

4.0 obscured by AV

=

23.8 mag would

=

17.3 and would thus be detected any-

of the combined uncertainties arising from the photometry and the modeling of the circumstellar emission. Indeed, even in the cases where the nd E
K

where in the cluster, as would be any star with a brighter absolute magnitude and reddened by a smaller amount. Our choice of the limiting absolute magnitude and extinction is a compromise between the need to include a statistically representative sample of the cluster population (which requires a limiting extinction as high as possible) and to probe the KLF down to low enough masses. Concerning the latter, M a mass M 0.15 M
K

t suggests the presence of infrared excess we

< 0.

75 (implying that most of the emission arises from

the photosphere rather than from reprocessing by the circumstellar environment) for 15 out of 21 stars, and only 4 objects require E
K

> 1.

0. It must be noted that the limitation of this analysis to

= 4.

0 corresponds to

stars detected at J , H , and KS leads us to consider objects that have cleared most of their circumstellar envelopes and are thus detectable at those wavelengths. Observations at longer wavelengths reveal few additional objects that have not reached that stage, and are discussed in Sect. 3.7.

at an age of 1 Myr, the precise mass be-

ing only approximate due to the large uncertainties in low-mass evolutionary tracks at such early ages (Baraffe et al. 2002). Our magnitude-limited KLF, presented in Fig. 8, thus covers nearly the whole range of stellar masses. The most intriguing feature of the KLF of the IRAS 163624845 cluster is its atness over virtually the entire absolute magnitude range covered by our observations, which supports the hints of a top heavy luminosity function reported in CSR05. Similar indications are found in the KLF independently de4

3.5. The K luminosity function

The luminosity function in the K band is frequently used in studies of young clusters as a diagnostic tool of the mass function and the star formation history of their stellar populations. Pioneering work on the interpretation of the K luminosity function was presented by Zinnecker et al. (1993). More recently, Muench et al. (2000) have carried out extensive modeling of the K luminosity function (hereafter KLF) showing its dependence on factors such as the cluster age, the spread of star formation

More rigorously we should instead refer to the K

tion here. However, the difference between the magnitude of any of the stars in our sample in the K and K
S

S

luminosity func-

bands is virtually irrelevant for our

purposes, as it is much smaller than the size of the bins used to build the KLF and the size of the errors introduced by our rather schematic corrections for extinction and infrared excess. We thus use the name K luminosity function as is normally done in the literature.


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108 Table 4. Objects with infrared excess Star CSR-002 CS-096 CSR-014 CSR-009A CSR-018 CS-097 CSR-015 CSR-019A CSR-001 CSR-010A CSR-010B CSR-010C CSR-012B CSR-023 CS-103 CSR-017 CSR-005B CSR-005A CS-099 CS-064
1

9

(2000)
16:39:58.6 16:40:00.3 16:40:00.3 16:39:59.9 16:40:01.2 16:40:00.3 16:40:00.3 16:40:01.3 16:39:58.3 16:40:00.0 16:39:59.9 16:40:00.0 16:40:00.1 16:40:01.7 16:40:00.0 16:40:00.8 16:39:59.5 16:39:59.5 16:40:00.2 16:40:01.4



(2000)

KS 10.568

J 2.098 0.952 1.896 0.938 1.320 2.198 2.355 2.245 2.013 2.247 3.146 0.840 1.795 2.786 2.814 2.595 2.926 3.055 2.266 1.485

-48:52:05 -48:51:56 -48:51:55 -48:51:53 -48:51:52 -48:51:51 -48:51:46 -48:51:45 -48:51:45 -48:51:42 -48:51:42 -48:51:42 -48:51:40 -48:51:38 -48:51:38 -48:51:38 -48:51:37 -48:51:37 -48:51:37 -48:51:33

± 0.010 15.777 ± 0.114 13.537 ± 0.008 11.242 ± 0.006 9.074 ± 0.010 14.388 ± 0.018 13.008 ± 0.008 11.415 ± 0.008 13.020 ± 0.005 10.662 ± 0.006 13.985 ± 0.191 16.315 ± 0.040 10.897 ± 0.025 13.362 ± 0.011 13.959 ± 0.041 11.769 ± 0.009 11.811 ± 0.016 10.679 ± 0.008 16.990 ± 0.019 16.752 ± 0.027

- ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±

H 0.014 0.063 0.030 0.007 0.014 0.056 0.042 0.007 0.028 0.021 0.117 0.199 0.099 0.025 0.161 0.042 0.016 0.011 0.078 0.062

H 1.282 0.611 1.259 0.798 1.181 1.358 1.316 1.430 1.104 1.569 1.743 1.681 1.495 1.494 1.836 1.655 1.970 2.157 2.362 0.969

- ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±

K

S

K

S

0.014 0.127 0.022 0.008 0.014 0.035 0.034 0.010 0.027 0.012 0.207 0.063 0.047 0.024 0.100 0.030 0.019 0.009 0.033 0.052 0.879 0.435 3.564 1.790 2.171 1.302 1.148 2.328 1.366 3.052 1.882 1.658 3.109 2.972 1.989 3.056 2.872

- - - ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± -

L

M

K

A

V

E

K

0.052 0.018 0.021 0.059 0.023 0.026 0.075 0.022 0.206 0.069 0.034 0.045 0.044 0.026 0.086 0.020 0.089

-2.3 5.0 2.9 0.5 -1.6 3.1 0.9 -1.4 1.5 -1.6 0.9 6.4 0.0 0.2 2.6 -1.3 -1.2 -3.0 7.7 5.4

20.6 3.0 8.2 6.6 9.8 12.0 16.1 21.3 11.9 20.0 23.9 3.8 12.8 22.7 16.3 24.4 26.3 30.8 9.2 8.3

0.04 0.16 0.91 0.64 0.98 0.66 0.35 0.20 0.41 0.56 0.25 1.11
1

1.15 0.01 1.05 0.25 0.56 0.40 2.31 0.20

: Fit obtained with

=

0 in Eqs. (3) (i.e., no H -band excess), as A

V

<

0 would be obtained otherwise.

of embedded clusters in the literature, which obtain KLFs rising towards fainter absolute magnitudes, with a particularly sharp increase in number counts in the 0

<

M

K

<

2 interval. For refer-

ence, we show in Fig. 8 the shape of the KLF of the Trapezium cluster (Muench et al 2002). Similar results are obtained on the clusters embedded in galactic giant HII regions (Figueredo et al. 2005 and references therein), as well as for other embedded clusters containing massive stars (e.g. Leistra et al. 2006, Massi et al. 2006, Fujiyoshi et al. 2005, Muench et al. 2003 for recent work). Although very low mass stars and even brown dwarfs appear to be present in the cluster coexisting with its most massive members, they do so in numbers far smaller than expected from a normal IMF (e.g. Kroupa 2001). It is hard to explain the at shape of the extinction-limited KLF as a result of incompleteness or incorrect assumptions solely. The steepness of the KLFs of other clusters is well detected at M
Fig. 8. The K luminosity function (KLF) of the IRAS16362-4845 embedded cluster, expressed as number of stars per 1 mag bin. The solid line corresponds to the extinction-limited sample that considers members with line-of-sight extinctions A down to M
K

K

<

2, which is two full magnitudes brighter than the at-

limit of our extinction- and magnitude-limited sample. The

ness of our KLF that we obtain thus cannot be attributed to the possible inaccuracy of the limiting magnitudes listed in Table 3. On the other hand, an underestimated infrared excesses would lead to a derived absolute magnitude that is brighter than the actual one, as both the extinction and the photospheric ux are then overestimated. Therefore, the possibility exists in principle that objects populating the brighter bins of Fig. 8 may actually be intrinsically fainter objects with strong infrared excess that has not been properly taken into account. The required size of the discrepancy between the actual and the derived infrared excesses should be signi cantly larger than the bin size of our KLF in order to have a noticeable impact on its shape. However, such large infrared excesses would imply negative extinctions for many of the objects of the cluster, and we can thus rule out this as a signi cant contributing effect. A systematic underestimate of the infrared excess may also have taken place among the objects undetected at J , as we have assumed E
K

= 4.

V

<

23.8, which is is completed

0. An unrestricted KLF including all the cluster mem-

bers (as well as possibly a few very reddened background stars) is represented by the dashed line. This latter KLF is virtually complete for the brightest bins, but becomes progressively incomplete as fainter stars become undetectable due to extinction. Finally, the KLF derived by Muench et al. (2002) for the Trapezium cluster is shown for reference, scaled by an arbitrary factor for convenience in the representation. The absolute magnitude scale for the Trapezium has been set by assuming its distance modulus to be 8.0, as used by those authors. No correction for extinction has been applied to the Trapezium KLF. Given that the vast majority of the Trapezium members have extinctions below A shift the Trapezium KLF by less than one bin towards the left.
K

= 0.5

(see Fig. 4a of Muench et al. 2002), a mean extinction correction would

=

0 for them.

rived by Wolk et al. (2007) for the X-ray selected population of RCW 108, where the central region containing IRAS 163624845 has a KLF remarkably atter than that of the surrounding area. This is in contrast with the results of most other studies

However, those tend to be faint objects at H and KS contributing to the faintest magnitude bins of the KLF, which may thus contain objects that would have been excluded had the infrared excess been properly taken into account. Since the lack of J


10

´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

measurements is more common towards fainter magnitudes, the systematically neglected infrared excesses among such objects should lead to our KLF being steeper than the actual one, and thus cannot account for the derived may also result if the value of atness. Finally, a systematic underestimate of the luminosity among the brightest objects

suggests that close binarity arises very early in their evolution (Bosch et al. 2001, Apai et al. 2007). The diffraction-limited quality of our observations in the L band allows us to search for visual pairs among the bright members of the IRAS 16362-4845 cluster with angular separations below 01, corresponding to a projected distance of 130 AU, and J H K
S



that we have adopted in Eqs. (2) 0. We

were smaller than the actual one, leading to the underestimate of the excess at H and the invalidity of our assumption E note nevertheless that a larger value of

photometry of the individual components can be ob02



J

tained for pairs separated by more than

-

03. We constrain

would imply that the

ourselves to the central area of the cluster covered by the L images (Sect. 2.1), as the image quality in that band (and therefore the shortest separations that can be measured) is nearly uniform across the entire eld. We also consider exclusively the brighter pairs in which both components are detected at L , since crowding at fainter magnitudes in the H or KS bands makes it impossible to separate true wide binaries from chance alignments or unrelated sources. We have set the upper limit of the separation for stars to be considered as close pairs to 1 1 (1430 AU projected distance), as this is the distance of the farthest star of the distinct subcluster of ve members around CSR 012A, where the stellar density in our images peaks (see Fig. 3). Unfortunately our results are of limited value in order to determine the physical characteristics of the detected systems, given the frequent existence of excess emission in the L band (Sect. 3.2) and the lack of a precise age determination of the cluster, which prevent us from deriving masses and mass ratios. However, our tions in the 01 ndings should include all the massive stars in the cluster with separa-

dust responsible for the infrared excess should be signi cantly hotter than the dust sublimation temperature in order to produce colors similar to those of the underlying objects, which we consider highly implausible. We are thus inclined to consider the remarkably Section 4. The mass of the IRAS 16362-4845 cluster was crudely estimated by CSR05 to be about 210 M . The new observations allow us to re ne that estimate by providing a deeper and more accurate census of its members, based on the more detailed estimate of individual stellar parameters described in Sect. 3.3. The cluster mass derived from stars detected in at least two of the J H KS bands amounts to 370 M . The difference is due to the fact that CSR05 based their estimate on the number of stars above the absolute magnitude of a main sequence A0V star and assuming that the cluster population follows a log-normal Miller & Scalo (1979) IMF, rather than on the individual masses. Applying the same method to the new observations presented here taking into account the stars more massive than 1 M (the dashed line in Fig. 5) we obtain a cluster mass of only 120 M . The lower masses obtained with this method are easily explained by the fact that the mass function underlying the KLF is atter than the log-normal Miller & Scalo (1979) IMF, and that the average mass of the stars above a certain threshold is therefore greater than the average mass expected from the assumed IMF. We note by passing that our new determination, 370 M , changes little due to the incompleteness of the census at very low masses, since their contribution to the cluster total mass is very small given the atness of the IMF. It is also fairly insensitive to the assumed age, as the massive stars dominating the mass of the cluster evolve very rapidly towards the main sequence. at shape of the KLF of the IRAS 16362-4845 as a real feature of the cluster, and we discuss its implications in

-

11 range, and can thus form a basis for future

statistical studies of its massive binary star content. The multiple systems that we detect are listed in Table 5. These systems include the CSR 012 cluster of ve members, the rather loose possibly triple system CSR 010, and the likely casual arrangement of unrelated objects CSR 008. The bluest component of the latter, CSR 008A, is the brightest star in the nebula at visible wavelengths and the very light reddening indicated by its visible and infrared colors suggests that it is a foreground star. However, it was noted in CSR05 that its infrared photometry shows a clear KS -band excess hinting at true cluster membership. The superior resolution of the observations presented here offer the solution to the puzzle, showing that CSR 008A has a redder companion at only 0 22. Although such close chance alignment is highly unlikely, we believe that the widely different infrared colors of both components convincingly argues for true membership in IRAS 16362-4845 of component B only, leaving component A as an ordinary foreground star. Finally, we include in our list star CS-093, which appears elongated in the KS -band images and shows a faint, short tail to the South in the L -band images. We consider this to be most likely due to a fainter, marginally resolved companion at jected distance) from the primary star. Excluding the likely foreground source CSR 008A and considering CS-093 as two separate stars, the results listed in Table 5 indicate that 18 of the 42 sources detected in the L band reside in 7 multiple systems with projected separations from the primary between 80 and 1430 AU, out of which 5 are binary systems, corresponding to a multiplicity fraction of 0.43 . As noted above,
5

3.6. Bright, close multiple systems

It is a well known observational fact that binaries with mass ratios close to unity are very common among massive stars (e.g. Garmany et al. 1982, Mason et al. 1998, Preibisch et al. 1999, Garc´ & Mermilliod 2001). Close binarity among a massive stars actually places important constraints on their proposed formation mechanisms (e.g. Bonnell et al. 1998, Yorke & Sonnthaler 2002, Bonnell 2005, Beuther et al. 2007) and has been usually linked to the fact that virtually all massive stars are formed in dense cluster environments, where dynamical interactions among cluster members and with the accreting gas can dominate their early evolution. Such interactions favor the formation of bound systems, while subsequent accretion causes high mass ratios and orbital evolution towards shorter separations between their members (Bate et al. 2003; it may be noted however that a high fraction of wide binaries with low mass ratios has been observed in the young but dynamically evolved cluster NGC 6611 by Duchene et al. (2001). Observations of dense clusters containing very young massive stars, or even massive stars still deeply embedded in their parental material,



0 06 (80 AU pro-

5

We de ne the multiplicity fraction as

nS n=2 nS n=1
where S
n

n n

,

is the number of systems containing n stars; this is, the num-

ber of stars in multiple systems divided by the total number of stars. An alternative de nition used by other authors (e.g. Duchene et al. 2001) is S n / n=1 S n , i.e., the number of multiple systems over the total n=2


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108 Table 5. Close pairs in the L band Star CSR-009A CSR-009B CSR-019A CSR-019B CSR-011A CSR-011B CSR-010A CSR-010B CSR-010C CS-093 CSR-012A CSR-012B CSR-012D CSR-012E CSR-012C CSR-008A CSR-008B
3

11

(2000)

1



(2000)

1

J 12.978

H 004 12.040

K 006 11.242

S

L' 006 10.807 11.147 008 10.113 10.969 10.293 11.633 8.334 13

separation 017 044 025 045 015 030 062 018 013

position angle

16:39:59.9

-48:51:53

± 0. -

± 0. -

± 0. -

± 0. ± 0. ± 0. ± 0. ± 0. ± 0.

85 6

16:40:01.3

-48:51:45

15.090

± 0. -

004

12.845

± 0. -

006

11.415

± 0. -

220 2

16:40:00.1

-48:51:45

17.502 15.598

± 0.075 ± 0.014 ± 0.018 ± 0.084 ± 0.193 ± 0.
187

14.286 13.596 12.231 15.728 17.996 16.014 9.198 12

± 0.034 ± 0.029 ± 0.010 ± 0.081 ± 0.049 ± 0.
176

12.894 12.607 10.662 13.985 16.315 14.341 8.356 11 13 10

± 0.016 ± 0.013 ± 0.006 ± 0.191 ± 0.040 ± 0.
088

245 5

16:40:00.0

-48:51:42

14.478 18.874 18.836

12.619

± 0.021 ± 0.076 .263 ± 0.056 ± 0.
030

107 089
2

300 9 34 0

16:40:00.4 16:40:00.1

-48:51:42 -48:51:40

19.382 10.872

12.003 7.105

14.187 17 14 16:39:59.7 -48:51:39

± ± - .441 ± .651 ±

0.010 0.091 0.113 0.099 017

12.392 14 12

± 0.010 ± 0.040 .981 ± 0.064 .571 ± 0.094 .195 ± 0.053 ± 0. ± 0.
022 150

10.897

± 0.010 ± 0.025 .856 ± 0.050 .581 ± 0.099 .953 ± 0.050 ± 0.012 ± 0.150 ± 0.008 ± 0.016

± 0.025 ± 0.023 9.981 ± 0.014 12.628 ± 0.046 9.603 ± 0.020
9.015 11.677 11.017 7.878 9.822

065 087 110 109

358 2 202 4 119 4 23 3

13.127

± 0. -

12.625 16.135 12.823 13.781

12.324 15.028 10.679 11.811

± 0. ± 0.

039 025 022

307 4

CSR-005A CSR-005B
1 2 3

16:39:59.5

-48:51:37

16.146 16.707

± 0.010 ± 0.011

± 0.005 ± 0.011

± 0.018 ± 0.085

027

138 0

: Coordinates of the primary component.

: Extension approximately 006 long towards position angle 170 , probably caused by a faint, marginally resolved companion.
: Probably not a true binary but a chance alignment: CSR-008A is a likely unrelated foreground star, and CSR-008B a cluster member.

our criterion of detection of both components in the L band is rather loose in terms of the lowest masses represented in Table 5, as some components may be detected thanks to their infrared excess rather than to the intrinsic brightness (and mass) of the central star. Assuming that the frequency and amount of L excesses are the same among components of multiple systems and single stars, the multiplicity fraction that we have derived above should nevertheless remain unaffected by this caveat. In principle such assumption may not be taken for granted, as multiplicity can affect the timescale for dissipation of the inner disks responsible for the L excess. However, we nd no evidence for any systematic differences between members of multiple systems and single stars when considering their positions in the ( H diagram (Fig. 7). As illustrated by the discussion of the observations of NGC 6611 by Duchene et al. (2001), completeness corrections in binarity studies require that one carefully takes into consideration instrumental effects, contamination by unrelated sources, extrapolations beyond the range of separations probed by the observations, and assumptions on the statistics of orbital parameters. Comparisons among different studies are furthermore hampered by the different observing techniques, instruments, and wavelength regions used. Leaving aside the di

3.7. Ver y red objects

The bulk of the stars projected on the area of the nebula and its surroundings have near infrared colors indicating extinctions in the 10

<

A

V

<

30 magnitude range. However, a few objects

listed in Table 6 display H

-

KS colors in excess of 2.5 mag, 40. Of these,

indicating extinctions near or in excess of AV

three (CS-010, 014, and 042) lie in the low extinction zone to the East of the cluster noted in Sect. 3.1 and excluded from our analysis of the cluster, and are probably distant luminous stars for which most of the foreground extincion is unrelated to the RCW 108 complex. Of the remaining four, CSR-006 has a derived extinction of A
V

-

KS ), ( K

S

-

L)

=

72 mag which is in good agreement with the extinction
13

produced by the molecular cloud on the foreground in that direction, as derived from the CO maps presented in CSR05. However, Fig. 1 shows that the star lies precisely behind a narrow lane of dust that also obscures the nebula in that direction. It may thus be that the star is actually a cluster member, as supported by its also red ( KS

-

L ) color suggestive of infrared ex-

cess. CS-070 and CS-141 are both faint objects with rather uncertain H -band measurements that may have led to an overestimate of the extinction in their direction, especially in the case of CS-070 which is close to our threshold for the identi cation of possible background objects. Finally, CS-094 is also very red in ( K
S

ffi

culties of an un-

biased comparison, we can say that our results indicate a high multiplicity fraction among the brightest members, and the lower limit that we nd (at least 43% of the massive stars in the cluster residing in binaries) suggests that, like in other clusters, most of the massive stars in IRAS 16362-4845 may be part of multiple systems.
number of systems, including single stars. Using this second de nition we obtain a multiplicity fraction of 0.21, which is close to the value 0.18 found by Duchene for NGC 6611 also using adaptive optics observations.

-

L ) and may actually be a low-mass object with sub-

stantial infrared excess, possibly also in the KS -band. Since it is not detected at J we derived the extinction by assuming that the emission at K
S

is photospheric (Sect. 3.3), which leads to an

overestimate of the extinction if that assumption is incorrect. We can summarize our results regarding very red sources by stating that we nd no strong evidence for any of these four objects being a background star, thus con rming that the dense column


12

´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

Table 6. Very red objects detected at K Star CSR-006 CS-010 CS-014 CS-042 CS-070 CS-094 CS-141

S



(2000)



(2000)

K 14.019 18.009 15.494 15.938 17.238 16.161 16.984

S

H 0.013 0.050 0.010 0.026 0.088 0.025 0.037 4.507 3.058 3.543 2.511 2.676 2.993 3.900

16:39:59.5 16:40:02.8 16:40:02.7 16:40:02.1 16:40:01.2 16:40:00.3 16:39:58.5

-48 51 56 -48 51 48 -48 51 28 -48 52 00 -48 51 49 -48 51 36 -48 51 54

± ± ± ± ± ± ±

- ± ± ± ± ± ± ±

K

S

KS 3.728

0.079 0.204 0.061 0.063 0.168 0.050 0.291

-L ± 0.
-

A 024

V

72.1 46.3 56.7 38.4 40.3

2.569

± 0.
-

100

47.7 62.3

Table 7. The candidate embedded protostar CS-109

(2000) (2000)
f2.2µm f3.8mum f4.5µ f5.8µ f8.0µ
m m m

16:39:59.8 -48:51:30 (6.0

< 3 µJy ± 0.2) mJy (18 ± 9) mJy (53 ± 32) mJy < 123 mJy

of molecular gas located behind the cluster provides an effective screen against such sources. With only one exception, all the sources detected at L are also seen in the KS images. The exception is CS-109, a source well detected with L

=

11.53 projected outside the boundaries
Fig. 9. Spectral energy distribution of the protostellar candidate CS-109.

of the bright nebulosity to the North of CSR-012 subcluster. The darkness of the background in that region and the good adaptive optics correction that is attained so close to the wavefront sensing star, which is only 10 away, allow us to set a very stringent limit of K is difficult to account for such extremely red color by assuming that CS-109 is a background star with normal colors obscured by the full column of dust in the molecular cloud. The necessary extinction is AV from the
13 S

>

20.5 at its position, implying ( KS

-

L)

>

9. It

Fig. 9 shows the spectral energy distribution of this source. Despite the large uncertainty in the IRAC uxes, it seems clear that the distribution rises towards longer wavelengths over the entire interval covered by our observations, consistent with it being a Class I or at spectrum source, commonly interpreted as an embedded protostar (e.g. Adams et al. 1987). We must note however that it is unusual for Class I sources to be unresolved down to the scale probed by the L observations (Whitney et al. 1997, Kenyon et al. 1998, Haisch et al. 2006). Deep diffraction-limited VLT observations in the 4-10 bution of CS-109.

>

150 mag, which is over twice that derived

CO maps but is not ruled out on small areas given the

large beam size of the millimeter observations. However, the absolute magnitude of an object reddened by such amount should be brighter than M
L

= -7.

9. The limit assumes that the object is images. While all

placed just behind the cloud at the distance of RCW 108, and that it is just below the detection limit in our K
S

µ

m range should would be very

valuable in con rming the true nature and spectral energy distri-

this is in principle possible (such an absolute magnitude could be attained by a bright Mira variable or a red supergiant), we deem the con uence of enumerated factors leading to its detection extremely unlikely. An alternative possibility is that CS-109 is a protostellar object still deeply embedded in its parental core. To investigate it we have examined archive Spitzer images of the region obtained with the IRAC camera between 3.6 and 8.0

4. Discussion
IRAS 16362-4845 is a relatively nearby example of the earliest evolutionary stages of a massive star forming cluster. Its youth is supported by its location embedded in the RCW 108 molecular cloud, the evidence for warm circumstellar material around a sizeable fraction of its members, the existence of two likely embedded protostars, and the fact that the HII region produced by the hottest stars of the cluster still has not excavated the surrounding molecular cloud beyond the boundaries of the cluster. However, perhaps the most important characteristic of the cluster is the possibility that it is an example of externally triggered star formation, for which additional tentative evidence in the region has been reported by Urquhart et al. (2004) and CSR05. It is thus interesting to discuss our possibility. The most intriguing feature of the IRAS 16362-4845 cluster found in the present study, which we have discussed at length ndings in the light of this

µ

m, which should

help in extending or constraining the spectral energy distribution beyond the L band. The object is faintly detected at 4.5 marginally detected at 5.8 consistency with the

µ

m. The non-detection at 3.6

µm µm

and is in

ux measured in our L -band images, and

an upper limit is obtained at 8

µ

m. The measurement of

uxes

in the two bands where the object is detected is made difficult by the steep brightness gradient of the nebulosity at its position, which prevents us from accurately determining the background emission level. Our measurements are listed in Table 7. CS-109 appears as a point source in our images despite their excellent resolution, implying the absence of extended structure with a linear size larger than about 70 AU.


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108

13

in Sect. 3.5, is the unusually

at KLF. Translating the shape of

ported here may provide the strongest constraints on future theoretical work.

the KLF to a shape of the IMF requires a number of assumptions (Muench et al. 2000), mainly about the age and star formation history of the cluster that are at present largely unknown. However, the atness of the KLF strongly suggests a top heavy IMF with a de cit of low-mass stars. A second, and possibly related, unusual feature of the cluster is that, contrarily to what is observed in most clusters containing massive stars, the most massive members are not at the center. At near-infrared wavelengths the cluster is dominated by CSR-012A and its subgroup of companions, but several other bright stars for which we nd similar or even brighter absolute magnitudes are scattered over the whole area of the cluster: most notably, CSR-006 and CSR020 are candidate O-type stars based on their absolute magnitudes, in addition to CSR-012A, and are far removed from it. The central location of the most massive stars in most clusters is thought to be of primordial nature due to their preferential formation at the bottom of the cluster potential well, rather than a result of dynamical evolution (Bonnell et al. 2007). This may play an important role in determining the shape of the IMF as a function of the distance to the cluster center in the competitive accretion scenario, leading to the preferential formation of low mass stars in the peripheral regions of the cluster. Lowmass stars might also be depleted at the centers due to coalescence to form more massive stars. We observe no hints of such mass segregation among the detected members of IRAS 163624845. Low-mass stars appear all across the cluster, but they do so in small numbers as implied by the shape of the K luminosity function (Fig. 8). Furthermore, it is obvious from Fig. 1 that the CSR-012 subcluster does not occupy a central position, but is rather located near the northern edge. The overall lack of lowmass stars might be interpreted in terms of competitive accretion as a result of the relatively widespread presence of massive stars, which would hamper the formation of low mass stars everywhere in the volume of the cluster rather than in the central regions only, contrarily to the more common case of clusters that have their most massive members at their centers. Modern theoretical studies of the build-up of the IMF, based on the turbulent fragmentation and gravitational collapse of an isolated molecular cloud (see e.g. Bonnell et al. 2007 and references therein), have proven to be quite successful in reproducing both the shape of the IMF and the mass segregation evidences observed in many clusters. However, the mass spectrum resulting from externally triggered star formation might be markedly different from the one resulting from those studies, as noted by Zinnecker (1989), and remains largely unexplored. Indirect evidence for signi cant differences in the IMF resulting from triggered star formation, leading to the preferential formation of intermediate- and high-mass stars, has been reported (Sugitani et al. 1991, Dobashi et al. 2001, Getman et al. 2007). A recent study by Negueruela et al. (2007) of the young cluster NGC 1893, which contains both both early-type O stars and evidence for triggered star formation, also reports indications of a top-heavy IMF. Clearly, more modeling work is needed in order to obtain more quantitative predictions on the IMF resulting from triggered star formation (Elmegreen 2007). However, it may be noted that already the early modeling of the sequential star formation scenario by Elmegreen & Lada (1977) predicted the preferential formation of massive stars. The reasons for this are related to the warm temperature of the gas in the shocked layer located ahead of the ionization front produced by the triggering massive stars, to turbulence induced by Rayleigh-Taylor instabilities in this shocked layer, and to possible coalescence of unstable fragments in this layer. Observations like the ones re-

5. Conclusions
The new observations of IRAS 16362-4845 presented in this paper represent a great improvement over previously existing ones, in terms of both depth and resolution, and give access to a more detailed study of its stellar population. The main conclusions of our work can be summarized as follows: The IRAS 16362-4845 cluster is embedded in a compact HII region that is ionization-bounded on its western side, where the density of the molecular cloud that harbors it is highest. The HII region appears to be open towards the East, in the direction towards the interface between the molecular cloud and the ionizing radiation of NGC 6193 that gives rise to the RCW 108 rim nebula. The stellar population of IRAS 16362-4845 is dominated by the small cluster around the star CSR-012, which is probably ´ a late O-type star based on the results obtained by Comeron et al. (2005). Unlike in other clusters containing massive stars, CSR-012 is not centrally located but is actually displaced near the northern edge of the HII region. Other stars of similar brightness are seen throughout the cluster, but no concentration of bright stars with respect to the distribution of fainter ones is apparent. Earlier suspicions of an absence of low-mass stars in the IRAS 16362-4845 cluster are not con rmed by these new observations, which do detect a faint stellar component. Our observations sample the entire stellar mass range of IRAS 16362-4845 and reveal some likely brown dwarfs. Nevertheless, the numbers of low mass stars are well below the expectations drawn from the abundance of massive stars when a normal initial mass function is extrapolated to lower masses. We estimate that the cluster has a total stellar mass of



370 M .

We estimate from the J H KS photometry that most of the members of the IRAS 16362-4845 cluster have their KS band ux dominated by photospheric emission, with a contribution by circumstellar emission in that band that is either absent or moderate. Moreover, the vast majority of members display excess emission in the L band. The K luminosity function of the cluster strongly suggests an atypical underlying initial mass function characterized by an overabundance of high mass stars. We are unable to identify any bias in our observations or our analysis that might possibly account for the shape of the K luminosity function if the initial mass function is normal. The diffraction-limited resolution of our L -band observations allows us to identify 7 close (separation tions in other massive star forming regions. IRAS 16362-4845 is still an active star forming site, as hinted by the existence of at least one extremely red object that is undetected below 3.8 at least 8

<

1 1) binary

or multiple systems, consistent with the multiplicity frac-

µ

m. The rising spectrum up to

µ

m suggests that this is a Class I source.

At a distance of 1.3 kpc from the Sun, the stellar and even substellar content of IRAS 16362-4845 is both nearby and accessible to current instrumentation. Its location embedded in the core of a molecular cloud that is being photoevaporated by the stars of a nearby cluster raises the possibility that its formation may have been externally triggered. This makes it an attractive


14

´ F. Comeron and N. Schneider: A close look at the heart of RCW 108
Negueruela, I., Marco, A., Israel, G., Bernabeu, G., 2007, A&A, in press. Palla, F., Stahler, S.W., 1999, ApJ, 525, 772. Petersen, C.C., 2001, Sky & Telescope, vol. 102, no. 11, 54. Preibisch, T., Balega, Y., Hofman, K.-H., Weigelt, G., Zinnecker, H., 1999, New Astr., 4, 531. Rieke, G.H., Lebofsky, M.J., 1985, ApJ, 288, 618. Rousset, G., et al., 2002, SPIE, 4839, 140. Shaver, P.A., Goss, W.M., 1970, Austr. J. Phys. Suppl., 14, 77. Skrutskie, M., et al., 2006, AJ, 131, 1163. Stetson, P.B., 1987, PASP, 99, 191. Straw, S., Hyland, A.R., Jones, T.J., Harvey, P.M., Wilking, B.A., Joy, M., 1987, ApJ, 314, 283. Sugitani, K., Ogura, K., 1994, ApJS, 92, 163. Sugitani, K., Fukui, Y., Mizuni, A., Ohashi, N., 1989, ApJ, 342, L87. Sugitani, K., Fukui, Y., Ogura, K., 1991, ApJS, 77, 59. Sugitani, K., Tamura, M., Ogura, K., 1995, ApJ, 455, L39. Testi, L., Palla, F., Natta, A., 1999, A&AS, 133, 81. Urquhart, J.S., Thompson, M.A., Morgan, L.K., White, G.J., 2004, A&A, 428, 723. ´ Whitney, B.A., Kenyon, S.J., Gomez, M., 1997, ApJ, 485, 703. ´ Wolk, S.J, Spitzbart, B.D., Bourke, T.L., Gutermuth, R.A., Vigil, M., Comeron, F., 2007, ApJ, submitted. Yorke, H.W., Sonnhalter, C., 2002, ApJ, 569, 846. Zinnecker, H., 1989, in Evolutionary Phenomena in Galaxies, eds. J.E. Beckman and B.E.J. Pagels, Cambridge Univ. Press. Zinnecker, H., McCaughrean, M.J., Wilking, B.A., 1993, in Protostars and Planets III, Univ. of Arizona Press.

target to search for observational evidences of such a formation process. The main result of the present study is the derivation of a K luminosity function that suggests an overabundance of massive stars with respect to less massive ones. Together with the lack of a central concentration of massive stars, this might be such a signature of externally triggered star formation, as we speculate in this paper. Further theoretical modeling of the initial mass function produced in such scenario will be needed to assess whether or not this interpretation, which is tentative at the moment, is likely to be correct. Regardless of this, the new observations of IRAS 16362-4845 that we have presented here con rm it as an excellent example for the study of the properties of the youngest clusters containing massive stars.
Acknowledgements. We are pleased to acknowledge the advise of Dr. Lowell Tacconi-Garman at the ESO User Support Department in the preparation of our Service Mode observations. The Paranal Science Operations staff is warmly thanked for the careful execution of this program. We also thank Dr. Francesco Palla for making available to us the evolutionary tracks used in this paper, and Dr. Hans Zinnecker for useful comments on an early draft of this paper. The constructive comments of the referee, Dr. August Muench, helped improve the paper and are greatly appreciated.

References
Adams, F.C., Lada, C.J., Shu, F.H., 1987, ApJ, 312, 788. Apai, D., Bik, A., Kaper, L., Henning, Th., Zinnecker, H., 2007, ApJ, 655, 484. Arnal, E.M., May, J., Romero, G.A., 2003, A&A, 412, 431. Bate, M.R., Bonnell, I.A., Bromm, V., 2003, MNRAS, 336, 705. Baraffe, I., Chabrier, G., Allard, F., Hauschildt, P.H., 2002, A&A, 382, 653. Baraffe, I., Chabrier, G., Barman, T.S., Al lard, F., Hauschildt, P.H., 2003, A&A, 402, 701. Bertoldi, F., 1989, ApJ, 346, 735. Beuther, H., Churchwell, E.B., McKee, C.F., Tan, J.C., 2007, in Protostars and Planets V, eds. B. Reipurth, D. Jewitt, K. Keil, Univ. of Arizona Press. Bonnell, I.A., Bate, M.R., Zinnecker, H., 1998, MNRAS, 298, 93. Bonnell, I.A., 2005, astro-ph/0501260. Bonnell, I.A., Larson, R.B., Zinnecker, H., 2007, in Protostars and Planets V, eds. B. Reipurth, D. Jewitt, K. Keil, Univ. of Arizona Press. Bosch, G., Selman, F., Melnick, J., Terlevich, R., 2001, A&A, 380, 137. ´ Comeron, F., Schneider, N., Russeil, D., 2005, A&A, 433, 955 (CSR05). Dobashi, K., Yonekura, Y., Matsumoto, T., Momose, M., Sato, F., Bernard, J.-P., Ogawa, H., 2001, PASJ, 53, 85. ¨ Duchene, G., Simon, T., Eisloffel, J., Bouvier, J., 2001, A&A, 379, 147. Elmegreen, B.G., 2007, in Massive Stars: from Pop III and GRBs to the Milky Way, eds. M. Livio and E. Villaver, Cambridge Univ. Press. Elmegreen, B.G., Lada, C.J., 1977, ApJ, 214, 725. Figueredo, E., Blum, R.D., Damineli, A., Conti, P.S., 2005, AJ, 129, 1523. Fujiyoshi, T., Smith, C., Moore, T.J.T., Lumsden, S.L., Aitken, D.K., Roche, P.F., 2005, MNRAS, 356, 801. Garc´ B., Mermilliod, J.-C., 2001, A&A, 368, 122. a, Garmany, C.D., Conti, P.S., Massey, P., 1982, ApJ, 242, 1063. Getman, K.V., Feigelson, E.D., Garmire, G., Broos, P., Wang, J., 2007, ApJ, 654, 316. Haisch, K.E., Barsony, M., Ressler, M.E., Greene, T.P., 2006, AJ, 132, 2675. Kenyon, S.J., Brown, D.I., Tout, C.A., Berlind, P., 1998, AJ, 115, 2491. Kroupa, P., 2001, MNRAS, 332, 231. Lada, C.J., Adams, F.C., 1992, ApJ, 393, 278. Leistra, A., Cotera, A.S., Liebert, J., 2006, AJ, 131, 2571. Lenzen, R., Hartung, M., Brandner, W., Finger, G., Hubin, N.N., Lacombe, F., Lagrange, A.-M., Lehnert, M.D., Moorwood, A.F.M., Mouillet, D., 2003, SPIE, 4841, 944. Mason, B.D., Gies, D.R., Hartkopf, W.I., Bagnuolo, W.G., ten Brummelaar, T., McAlister, H.A., 1998, AJ, 115, 821. Massi, F., Testi, L., Vanzi, L., 2006, A&A, 448, 1007. Meyer, M.R., Calvet, N., Hillenbrand, L.A., 1997, AJ, 114, 288. Miao, J., White, G.J., Nelson, R., Thompson, M., Morgan, L., 2006, MNRAS, 369, 143. Miller, G.E., Scalo, J.M., 1979, ApJS, 41, 513. Muench, A.A., Lada, E.A., Lada, C.J., 2000, ApJ, 533, 358. Muench, A.A., Lada, E.A., Lada, C.J., Alves, J., 2002, ApJ, 573, 366. Muench, A.A., Lada, E.A., Lada, C.J., Elston, R.J., Alves, J.F., Horrobin, M., ´ ´ Huard, T.H., Levine, J.L., Raines, S.N., Roman-Zuniga, C., 2003, AJ, 125, 2029.


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108, Online Material p 1

Online Material


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108, Online Material p 2

List of Objects
`RCW 108' on page 1 `SFO 79' on page 1 `NGC 6193' on page 1 `IRAS 16362-4845' on page 1 `G333.1-0.4' on page 7 `NGC 6611' on page 10 `NGC 1893' on page 13


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108, Online Material p 3 Table 2. Photometry of sources detected in the IRAS 16362-4845 cluster

Name CSR-001 CSR-002 CSR-003 CSR-004 CSR-005A CSR-005B CSR-006 CSR-007 CSR-008A CSR-008B CSR-009A CSR-009B CSR-010A CSR-010B CSR-010C CSR-011A CSR-011B CSR-012A CSR-012B CSR-012C CSR-012D CSR-012E CSR-013 CSR-014 CSR-015 CSR-016 CSR-017 CSR-018 CSR-019A CSR-019B CSR-020 CSR-021 CSR-022 CSR-023 CSR-024 CSR-025 CS-001 CS-002 CS-003 CS-004 CS-005 CS-006 CS-007 CS-008 CS-009 CS-010 CS-011 CS-012 CS-013 CS-014 CS-015 CS-016 CS-017 CS-018 CS-019 CS-020 CS-021 CS-022 CS-023 CS-024

RA(2000) 16:39:58.3 16:39:58.6 16:39:58.7 16:39:59.0 16:39:59.5 16:39:59.5 16:39:59.5 16:39:59.8 16:39:59.7 16:39:59.7 16:39:59.9 16:39:59.9 16:40:00.0 16:39:59.9 16:40:00.0 16:40:00.1 16:40:00.1 16:40:00.1 16:40:00.1 16:40:00.2 16:40:00.1 16:40:00.2 16:40:00.3 16:40:00.3 16:40:00.3 16:40:00.3 16:40:00.8 16:40:01.2 16:40:01.3 16:40:01.2 16:40:01.6 16:40:01.6 16:40:01.7 16:40:01.7 16:40:01.9 16:40:02.1 16:40:03.0 16:40:03.0 16:40:02.9 16:40:02.9 16:40:02.9 16:40:02.8 16:40:02.9 16:40:02.8 16:40:02.8 16:40:02.8 16:40:02.8 16:40:02.8 16:40:02.7 16:40:02.7 16:40:02.7 16:40:02.6 16:40:02.6 16:40:02.6 16:40:02.6 16:40:02.5 16:40:02.5 16:40:02.5 16:40:02.4 16:40:02.4

Dec(2000) -48:51:45 -48:52:05 -48:51:59 -48:51:33 -48:51:37 -48:51:37 -48:51:56 -48:51:59 -48:51:39 -48:51:39 -48:51:53 -48:51:53 -48:51:42 -48:51:42 -48:51:42 -48:51:45 -48:51:45 -48:51:40 -48:51:40 -48:51:39 -48:51:41 -48:51:41 -48:51:59 -48:51:55 -48:51:46 -48:51:43 -48:51:38 -48:51:52 -48:51:45 -48:51:45 -48:51:48 -48:51:41 -48:51:52 -48:51:38 -48:51:48 -48:51:55 -48:51:39 -48:51:40 -48:51:47 -48:51:49 -48:51:47 -48:51:16 -48:51:42 -48:51:16 -48:51:39 -48:51:48 -48:51:50 -48:51:53 -48:51:52 -48:51:28 -48:51:47 -48:51:57 -48:51:56 -48:51:54 -48:51:59 -48:51:50 -48:51:53 -48:51:45 -48:51:30 -48:51:15 14.176 20.054 14.909 17.642 17.879 17.441 13.101 16.692 16.679 13.180 16.019 11.575 15.090 14.478 18.874 18.836 17.502 15.598 10.872 14.187 14.651 12.978 15.867 13.127 17.880 16.146 16.707 16.137 13.948

J

H 14.124 11.850 16.419 14.815 12.823 13.781 18.526 13.534 12.625 16.135 12.040 12.231

K 0.027 0.010 0.019 0.004 0.005 0.011 0.078 0.015 0.022 0.150 0.006 13.020 10.568 14.471 13.183 10.679 11.811 14.019 12.250 12.324 15.028 11.242 10.662

± 0.008 ± 0.010 ± 0.007 ± 0.010 ± 0.011 ± 0.010 ± 0.017 ± 0. ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±
004

± ± ± ± ± ± ± ± ± ± ±

± ± ± ± ± ± ± ± ± ± ±

S

L 0.005 0.010 0.011 0.008 0.008 0.016 0.013 0.008 0.012 0.150 0.006 10.611 7.878 11.872

± 0.

075

± 0.028 ± 0.018 9.822 ± 0.085 10.291 ± 0.020
11.677 11.017

0.018 0.084 0.193 0.075 0.014 0.010 0.091 0.099 0.113 0.006 0.022 0.026 0.013 0.031 0.010 0.004 0.010 0.043 0.009 0.014 0.020

15.728

± 0.010 ± 0.081 17.996 ± 0.049 14.286 ± 0.034 13.596 ± 0.029 9.198 ± 0.010 12.392 ± 0.040 12.195 ± 0.053 12.981 ± 0.064 14.571 ± 0.094 12.532 ± 0.024 14.796 ± 0.020 14.324 ± 0.033 11.148 ± 0.010 13.424 ± 0.029 10.255 ± 0.010 12.845 ± 0.006
11.294 16.294 13.105 14.856 15.030 15.289 21.019 20.566 20.925

13.985

± 0.006 ± 0.191 16.315 ± 0.040 12.894 ± 0.016 12.607 ± 0.013 8.356 ± 0.010 10.897 ± 0.025 10.953 ± 0.050 11.856 ± 0.050 13.581 ± 0.099 12.335 ± 0.009 13.537 ± 0.008 13.008 ± 0.008 10.261 ± 0.005 11.769 ± 0.009 9.074 ± 0.010 11.415 ± 0.008
14.377

± 0.039 ± 0.025 10.807 ± 0.017 11.147 ± 0.044 8.334 ± 0.021 12.619 ± 0.076 13.263 ± 0.056 10.293 ± 0.015 11.633 ± 0.030 7.105 ± 0.025 9.015 ± 0.023 9.603 ± 0.020 9.981 ± 0.014 12.628 ± 0.046 ± 0.051 ± 0.022 8.924 ± 0.017 8.797 ± 0.024 5.510 ± 0.019 10.113 ± 0.025 10.969 ± 0.045 7.142 ± 0.021 12.089 ± 0.046 11.204 ± 0.026 11.704 ± 0.044 12.814 ± 0.069
10.837 12.658

± ± ± ± ± ± ± ±

0.010 0.024 0.027 0.021 0.031 0.015 0.235 0.194 382

± 0.007 ± 0.014 12.252 ± 0.007 13.362 ± 0.011 13.558 ± 0.010 13.967 ± 0.009
19.496 18.038 19.390 19.814 20.101 20.404 18.888

9.846

± 0.

20.779 21.067 21.172 20.263 19.037 19.945

± 0.188 ± 0.198 ± 0.195 ± 0.171 ± 0.060 ± 0.142

18.705 18.009 19.101 19.335 19.111 15.494 19.041 19.730 20.063 18.172

21.471 20.298 21.104 20.835

± ± ± ±

0.300 0.122 0.131 0.137

20.420 19.481 19.491 19.618 20.611

± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±

0.162 0.041 0.077 0.267 0.283 0.227 0.156 0.102 0.050 0.134 0.067 0.103 0.010 0.082 0.167 0.122 0.066 0.199 0.080

± 0.071 ± 0.080 ± 0.284


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108, Online Material p 4 Table 2. (Continued)

Name CS-025 CS-026 CS-027 CS-028 CS-029 CS-030 CS-031 CS-032 CS-033 CS-034 CS-035 CS-036 CS-037 CS-038 CS-039 CS-040 CS-041 CS-042 CS-043 CS-044 CS-045 CS-046 CS-047 CS-048 CS-049 CS-050 CS-051 CS-052 CS-053 CS-054 CS-055 CS-056 CS-057 CS-058 CS-059 CS-060 CS-061 CS-062 CS-063 CS-064 CS-065 CS-066 CS-067 CS-068 CS-069 CS-070 CS-071 CS-072 CS-073 CS-074 CS-075 CS-076 CS-077 CS-078 CS-079 CS-080 CS-081 CS-082 CS-083 CS-084 CS-085

RA(2000) 16:40:02.4 16:40:02.4 16:40:02.4 16:40:02.4 16:40:02.3 16:40:02.3 16:40:02.3 16:40:02.3 16:40:02.3 16:40:02.3 16:40:02.2 16:40:02.2 16:40:02.2 16:40:02.2 16:40:02.2 16:40:02.1 16:40:02.1 16:40:02.1 16:40:02.1 16:40:01.9 16:40:01.9 16:40:01.9 16:40:01.9 16:40:01.9 16:40:01.8 16:40:01.8 16:40:01.7 16:40:01.8 16:40:01.6 16:40:01.7 16:40:01.7 16:40:01.7 16:40:01.5 16:40:01.6 16:40:01.5 16:40:01.4 16:40:01.5 16:40:01.5 16:40:01.5 16:40:01.4 16:40:01.5 16:40:01.4 16:40:01.3 16:40:01.3 16:40:01.3 16:40:01.2 16:40:01.1 16:40:01.1 16:40:01.1 16:40:01.1 16:40:01.0 16:40:01.0 16:40:00.9 16:40:00.9 16:40:00.9 16:40:00.9 16:40:00.9 16:40:00.8 16:40:00.8 16:40:00.7 16:40:00.7

Dec(2000) -48:51:43 -48:51:43 -48:51:49 -48:51:45 -48:51:36 -48:51:42 -48:51:57 -48:51:48 -48:51:49 -48:51:52 -48:51:51 -48:51:39 -48:51:51 -48:51:49 -48:52:02 -48:51:43 -48:51:54 -48:52:00 -48:51:49 -48:51:48 -48:51:53 -48:51:46 -48:51:43 -48:51:56 -48:51:37 -48:52:01 -48:51:45 -48:52:08 -48:51:15 -48:51:30 -48:51:44 -48:51:58 -48:51:18 -48:51:57 -48:51:23 -48:51:12 -48:51:53 -48:51:32 -48:52:06 -48:51:33 -48:52:09 -48:51:54 -48:51:37 -48:51:51 -48:51:47 -48:51:49 -48:51:28 -48:51:38 -48:51:50 -48:51:57 -48:51:48 -48:51:39 -48:51:32 -48:51:49 -48:51:39 -48:51:42 -48:51:48 -48:51:46 -48:51:50 -48:51:39 -48:51:49 19.891 20.131 23.125 17.326 19.206 19.660 18.134 18.800 23.685

J

H

K 20.531 19.624

20.485 20.258 20.805 17.943 20.546 19.876 21.880 17.962 21.097 19.981 19.419 21.136 17.484 18.449 20.472 16.855 17.677 19.772

± 0.187 ± 0.120 ± ± ± ± ± ± ± ± ± ± ± ± ±
0.143 0.032 0.139 0.090 0.290 0.033 0.207 0.131 0.090 0.202 0.097 0.057 0.156

19.685 18.676 17.236 18.673 15.511 18.389 21.407 17.399 19.327 19.210 17.698 18.601 19.744 15.787 15.938 19.602 17.681 14.863 16.543 18.562 19.443 13.047 18.944 19.586 20.582 20.078 21.050

± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±

S

L 0.197 0.091 0.143 0.074 0.025 0.058 0.014 0.087 0.353 0.035 0.074 0.122 0.051 0.114 0.126 0.041 0.026 0.136 0.155 0.013 0.043 0.060 0.093 0.009 0.087 0.343 0.270 0.159 0.289 0.160 0.074 0.038 0.021 0.143 0.262 0.150 0.155 0.143 0.027 0.211 0.020 0.162 0.060 0.016 0.088 0.286 0.071 0.111 0.287 0.045 0.116 10.912

± 0.021 ± 0.043 ± 0.098 ± 0.
032 233

± 0.028 ± 0.734

14.950 21.594

± 0.

023

± 0.

21.747 19.437

± 0.032 ± 0.020

18.460 16.503

± ± ± ±

0.294 0.096 0.041 0.025

19.757 17.916 18.180 15.656 20.243 20.608

20.612

± 0.

138

18.976 20.467 18.418

± 0.

043

17.721 16.900 16.032

± 0.

045 019

16.752 19.990 15.262 19.680 14.317 15.134 17.238 20.077 17.195 15.877 20.877 17.172 16.320 19.971 18.101 19.137

± 0.

± 0.

014

15.759 19.914 19.091 17.333 18.482 18.328

± 0.070 ± 0.033 ± 0.143 ± 0.099 ± 0.105 ± ± ± ± ± ± ± ± ±
0.055 0.070 0.257 0.220 0.131

± 0.

269

21.380 20.756 17.909

± 0.162 ± 0.194 ± 0.120 ± 0.398 ± 0.036 ± 0.
146

± 0.

279 21.096 18.671 0.290 0.084 0.165 0.193 19.718 19.482 16.370

± 0.

211

19.148 20.950


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108, Online Material p 5 Table 2. (Continued)

Name CS-086 CS-087 CS-088 CS-089 CS-090 CS-091 CS-092 CS-093 CS-094 CS-095 CS-096 CS-097 CS-098 CS-099 CS-100 CS-101 CS-102 CS-103 CS-104 CS-105 CS-106 CS-107 CS-108 CS-109 CS-110 CS-111 CS-112 CS-113 CS-114 CS-115 CS-116 CS-117 CS-118 CS-119 CS-120 CS-121 CS-122 CS-123 CS-124 CS-125 CS-126 CS-127 CS-128 CS-129 CS-130 CS-131 CS-132 CS-133 CS-134 CS-135 CS-136 CS-137 CS-138 CS-139 CS-140 CS-141 CS-142 CS-143 CS-144 CS-145 CS-146

RA(2000) 16:40:00.6 16:40:00.7 16:40:00.6 16:40:00.6 16:40:00.5 16:40:00.4 16:40:00.4 16:40:00.4 16:40:00.3 16:40:00.3 16:40:00.3 16:40:00.3 16:40:00.3 16:40:00.2 16:40:00.0 16:40:00.1 16:40:00.0 16:40:00.0 16:40:00.0 16:39:59.9 16:39:59.9 16:40:00.0 16:39:59.9 16:39:59.8 16:39:59.8 16:39:59.8 16:39:59.7 16:39:59.8 16:39:59.7 16:39:59.7 16:39:59.7 16:39:59.6 16:39:59.5 16:39:59.5 16:39:59.4 16:39:59.4 16:39:59.4 16:39:59.4 16:39:59.4 16:39:59.3 16:39:59.3 16:39:59.3 16:39:59.3 16:39:59.3 16:39:59.2 16:39:59.2 16:39:59.1 16:39:59.1 16:39:59.0 16:39:58.9 16:39:58.9 16:39:58.9 16:39:58.8 16:39:58.8 16:39:58.8 16:39:58.5 16:39:58.4 16:39:58.2 16:39:58.1 16:39:58.0 16:39:58.0

Dec(2000) -48:51:38 -48:52:05 -48:51:53 -48:51:44 -48:51:40 -48:51:40 -48:51:38 -48:51:42 -48:51:36 -48:51:29 -48:51:56 -48:51:51 -48:52:03 -48:51:37 -48:51:37 -48:52:05 -48:52:02 -48:51:38 -48:51:40 -48:51:34 -48:51:48 -48:52:01 -48:52:10 -48:51:30 -48:51:29 -48:51:41 -48:51:34 -48:51:58 -48:52:01 -48:51:40 -48:52:07 -48:51:52 -48:51:53 -48:51:45 -48:51:39 -48:51:29 -48:51:48 -48:52:07 -48:51:49 -48:51:38 -48:51:28 -48:51:42 -48:51:37 -48:51:52 -48:51:37 -48:51:42 -48:51:58 -48:52:03 -48:51:55 -48:51:40 -48:51:38 -48:51:59 -48:51:17 -48:51:45 -48:51:30 -48:51:54 -48:51:55 -48:51:43 -48:51:28 -48:51:47 -48:52:09 19.548 20.997 15.889 20.904 20.961 20.867 17.791 21.226 19.980 18.609 21.618 20.749 17.340 17.944 19.382 20.451 22.368 21.058

J

H

K 19.321

± 0.

S

L 176

± 0.282 ± 0.196
17.756

± 0.

108 037

17.210 14.594 19.199 14.341 16.161 19.196 15.777 14.388 18.704 16.990

± 0. ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ±

099 11.597 12.003 13.592

± 0.

193 16.088

± 0.

0.013 0.244 0.088 0.025 0.197 0.114 0.018 0.172 0.019

± 0.

028

± 0.

187

16.014 19.154

± 0.176 ± 0.043 ± 0.057 ± 0.030 ± 0. ± ± ± ±
027

± 0.030 ± 0.097 ± 0.
056 087

± 0.026 ± 0.047 ± 0.073 ± 0.275 ± 0.042 ± 0.133

16.388 15.746 19.352 20.458 16.720 15.795 15.241 19.068

12.598 14.118

± 0.

0.148 0.027 0.091 0.059 15.294 13.959 14.438 18.564 18.449 15.632 20.758 20.136 18.994 18.761 0.022 0.041 0.050 0.232 0.174 0.012 0.264 11.532 0.214 0.269 0.254 0.009 0.211 0.299 0.211 0.168 0.295 0.197 0.144 0.259 0.111 0.387 0.093 0.269 0.015 0.097 0.162 0.218 0.015 0.191 0.138 0.160 0.090 0.025 0.019 12.059 11.207 10.850 11.934

± 0.017 ± 0.040

± 0.

103

17.405

± 0.149 ± 0.036

± 0.

026

± 0.

022 580 145

14.694 18.804

± 0.

027 227

13.105 19.229

± 0.

031

± 0.

± 0.

20.042

± 0.

20.035 21.138

± 0.

480

18.750 19.662 19.513

21.678

± 0.

203 19.261 20.298

20.071 21.810

± 0.144 ± 0.296 ± 0.
019 110

18.473 19.480 18.968 18.955

16.544

14.886 19.307

± 0.

054

± 0.

133

18.634

± 0.

18.623 20.445 17.353

± 0.

029

15.573 19.664 18.340 18.304

20.466 19.997

± 0.

011

15.074 22.202 20.884

± ± ± ± ±

0.200 0.126 0.024 0.281 0.289

18.377 17.362 15.207 16.984 20.810

± 0.037 ± 0.264 ± 0.015 ± 0.045

± 0.

167 17.934 047 18.165 21.321

± 0.

± 0.041 ± 0.051 ± 0.251

15.933 17.328


´ F. Comeron and N. Schneider: A close look at the heart of RCW 108, Online Material p 6 Table 2. (Continued)

Name CS-147

RA(2000) 16:39:57.9

Dec(2000) -48:51:43

J

H 15.976

K 026 15.159

± 0.

± 0.

S

L 009