Rcw36: characterizing the outcome of massive star formation††thanks: Based on observations performed with the ESO New Technology Telescope at La Silla Observatory, as part of program 074.C-0728, and with the ESO Very Large Telescope on Cerro Paranal, Chile, as part of programs 078.C-0780, 084.C-0604 and 087.C-0442.
Key Words.:Stars: formation – Stars: massive – Stars: pre-main-sequence – Stars: variables: T Tauri, Herbig Ae/Be
Context:Massive stars play a dominant role in the process of clustered star formation, with their feedback into the molecular cloud through ionizing radiation, stellar winds and outflows. The formation process of massive stars is poorly constrained because of their scarcity, the short formation timescale and obscuration. By obtaining a census of the newly formed stellar population, the star formation history of the young cluster and the role of the massive stars within it can be unraveled.
Aims:We aim to reconstruct the formation history of the young stellar population of the massive star-forming region RCW 36. We study several dozens of individual objects, both photometrically and spectroscopically, look for signs of multiple generations of young stars and investigate the role of the massive stars in this process.
Methods:We obtain a census of the physical parameters and evolutionary status of the young stellar population. Using a combination of near-infrared photometry and spectroscopy we estimate ages and masses of individual objects. We identify the population of embedded young stellar objects (YSO) by their infrared colors and emission line spectra.
Results:RCW 36 harbors a stellar population of massive and intermediate-mass stars located around the center of the cluster. Class 0/I and II sources are found throughout the cluster. The central population has a median age of Myr. Of the stars which could be classified, the most massive ones are situated in the center of the cluster. The central cluster is surrounded by filamentary cloud structures; within these, some embedded and accreting YSOs are found.
Conclusions:Our age determination is consistent with the filamentary structures having been shaped by the ionizing radiation and stellar winds of the central massive stars. The formation of a new generation of stars is ongoing, as demonstrated by the presence of embedded protostellar clumps, and two exposed protostellar jets.
Most massive stars (M) form in clusters. As they evolve fast and are sources of ionizing radiation and stellar winds, they impact the evolution of their surrounding young stellar population and birth cloud. While massive stars are usually detected while on or already off the main sequence, their surrounding lower mass population is often still forming and provides a window, or “clock”, on the star formation history. Therefore, a strategy towards understanding the complex process of clustered massive star formation is to study the outcome of star formation; to obtain a sample of star-forming regions, quantify the physical properties of the embedded young stellar population, and reconstruct the star formation history. With the combination of intermediate-resolution spectroscopy () and multi-band photometry, it is possible to derive the stellar parameters of low- and high-mass stars and compare these to stellar evolution models to derive masses and ages. It also allows to characterize circumstellar material (disks and outflows) surrounding YSOs and, by virtue of the line-of-sight extinction, the local abundance of dust. Combining all these findings, a complete census of the young stellar population is obtained. Scenarios can then be conceived to describe the progression of star formation throughout the cluster, and the causes and effects of the formation of the massive stars within it.
Over the past two decades, near-infrared imaging and spectroscopic surveys within the Galactic plane have revealed many young embedded stellar clusters showing a rich diversity in stellar content and evolutionary history (Hanson et al. 1997; Walborn & Blades 1997; Blum et al. 2000; Feigelson & Townsley 2008). In a few of these studies, a massive star is identified as the source of ionizing radiation that formed the ultra-compact H ii region (UCHII, e.g. Watson & Hanson 1997; Alvarez et al. 2004; Bik et al. 2005). In some cases, evidence of age spread and hence sequential star formation is found. In other cases, star formation seems to be triggered by the expansion of an H ii region (e.g. Zavagno et al. 2006, 2007). However, it is generally difficult to derive a causal connection between the different generations of young stars. This is because evidence for triggered star formation is at best indirect, and because the uncertainties on stellar age estimates are large (e.g. Preibisch 2012). This is particularly the case for massive (proto)stars, as their position on the Hertzsprung-Russell diagram (HRD) strongly depends on their accretion history (Davies et al. 2011).
As part of the Formation and Early evolution of Massive Stars (FEMS) collaboration (Bik et al. 2010), we have obtained near-infrared images and spectra of several young embedded massive clusters, following up on a near-infrared survey of 45 southern star-forming regions centered on IRAS point sources exhibiting colors characteristic of UCHII regions (Bik 2004; Kaper et al., in prep.). Bik et al. (2010) presented a spectroscopic census of RCW 34 in the Vela Molecular Ridge (VMR). They detected three distinct regions of star formation, suggesting that star formation progressed from south to north. Maaskant et al. (2011) studied the high-mass star-forming region GGD 12-15 centered on IRAS 06084-0611. They showed that the youngest generation of stars is centrally located, while somewhat more evolved objects are spread out over a larger area, suggesting sequential star formation along the line of sight. Wang et al. (2011) detected different evolutionary stages of star formation in the S255 complex. They conclude that their observations are best explained by the so-called triggered outside-in collapse star formation scenario, in which the filaments on the outskirts of the cluster collapse first, enhancing the instability of the massive star-forming cluster core.
The aim of this paper is to study the massive star-forming region RCW 36 (Gum 20, BBW 217) using a combination of photometry and spectroscopy that covers a broad range in wavelength (m). Fig. 1 contains an overview of the observations. RCW 36 is located in cloud C of the VMR, along a high-column density cloud filament which extends north to south (NS, Fig. 1, left; Hill et al. 2011). It includes a young star cluster (Massi et al. 2003) associated with the H ii region G265.151+1.454 of Caswell & Haynes (1987). The region comprises the IRAS point source 08576-4334 with UCHII colors, also known as IRS 34 (Liseau et al. 1992), co-located with an UCHII region (Walsh et al. 1998). Hunt-Cunningham et al. (2002) suggest that star formation in RCW 36 is induced by a collision of two molecular gas clumps. These clumps are detected in several molecular emission lines at different velocities north and south of the star-forming region. Minier et al. (2013, hereafter MTHM13) detect a tenuous, high-temperature bipolar nebula extending up to at least 10 (2 pc) both east and west (EW) from the cluster. Around the origin of the nebula, a ring-like structure with high column density is detected; this structure is also visible in Fig. 1. These authors propose a scenario where both the ring and the bipolar nebula are shaped (or “blown out”) by the radiation pressure of the central O-star(s) in the cluster.
RCW 36 is very well suited for our scientific purposes, due to its relative proximity ( kpc, Liseau et al. 1992; Yamaguchi et al. 1999), the presence of several O and B stars (Bik et al. 2005) and YSOs (Bik et al. 2006; Ellerbroek et al. 2011). Based on observations in the far-infrared, Verma et al. (1994) estimate a total luminosity of L, corresponding to the luminosity of two O9 V stars. This indicates that the emerging far-infrared luminosity is dominated by thermal reprocessing of radiation of the massive stars by dust. This confirms that the energy output from other massive (proto)stars embedded in the UCHII does not reach the observer, consistent with the high column density in this region ( cm, mag, Hill et al. 2011).
Apart from the peaking far-infrared emission possibly indicating the presence of at least two (massive) protostellar cores (Hill et al. 2011; Giannini et al. 2012), a more evolved young stellar population is detected in the near-infrared. Baba et al. (2004) have performed near-infrared imaging in the -bands and have detected more than 350 cluster members within the central ( pc). The same authors derive an age of 2–3 Myr and an average extinction of mag. Bik et al. (2005) classify two objects (objects 1 and 3 in this study, see Table 4) as O9 V – B1 V stars based on their K-band spectra, which is consistent with the result from Verma et al. (1994) and a distance of 0.7 kpc.
Bik et al. (2006) report two near-infrared bright YSOs exhibiting Br and CO emission (08576nr408 and 08576nr292, our objects 2 and 4, respectively). Bik & Thi (2004) show that the CO emission from object 4 likely originates in a circumstellar Keplerian rotating disk. Ellerbroek et al. (2011) report the discovery of bipolar jets around two sources (HH 1042 and HH 1043, associated with 08576nr292 and 08576nr480; our objects 4 and 97, respectively; see also Ellerbroek et al. 2013), adding to the evidence for ongoing star formation in this region.
In this paper we perform a detailed analysis of the stellar content of RCW 36 using optical and near-infrared spectroscopy as well as near- and mid-infrared imaging. In Sect. 2 we describe the observations. Sect. 3 identifies the stellar content of the region by the different photometric datasets; sources are assigned Lada classes (Lada 1987) according to their near- to mid-infrared spectral energy distributions (SED). We present detailed optical and near-infrared photospheric spectral classification of the pre-main sequence (PMS) population in Sect. 4. In Sect. 5 we combine all the results and refine the age estimate of the stellar population of RCW 36. We present a possible scenario for its star formation history, in which the massive and intermediate-mass stars have preceded a new generation of embedded protostars. Sect. 6 summarizes the conclusions of this work.
2 Observations and data reduction
To obtain a complete picture of the stellar content of RCW 36 we use archival near- to mid-infrared photometry to complement our near-infrared integral field VLT/SINFONI spectroscopy as well as optical to near-infrared VLT/X-shooter spectra of selected sources. The Herschel images (observed as part of the HOBYS program, observation ID 1342196658, P.I. Motte; Hill et al. 2011), which are used to show the large-scale structures with respect to the stellar population, were retrieved from the Herschel Science Archive.
2.1 Near-infrared imaging and photometry: NTT/SOFI
We have retrieved broad-band near-infrared data from the ESO archive. The observations (P.I. Bialetsky) were carried out with SOFI (Moorwood et al. 1998) on the New Technology Telescope (NTT) at La Silla Observatory in Chile. The observations were performed on 18 May 2005 under decent seeing conditions (0.9 in , 0.8 in and 1.2 in the -band). The Detector Integration Time (DIT) was 10, 8, and 6 seconds for , , and , NDIT = 12 and a total number of 6 frames were taken on source. This results in a total exposure time of 12 min (), 9.6 min (), and 7.2 min () for the three bands. Offset sky positions were taken to ensure a good sky subtraction.
The data were reduced using the ESO pipeline (version 1.5.2) for SOFI. We corrected the frames with darks and flat fields which were obtained on the same morning as the science observations. After that the data were sky subtracted and the final mosaic was created. We obtained an astrometrical solution by matching the positions of the stars with those of 2MASS (Skrutskie et al. 2006).
Photometry was performed using daophot (Stetson 1987) under the IRAF environment. First, stellar sources were detected using the daofind task and aperture photometry was performed using the task phot with an aperture equal to the Full With Half Maximum (FWHM) of the stellar sources. Using the tasks pstselect and pst a reference point spread function (PSF) was constructed by using over 25 bright, isolated stars in the SOFI images. Finally, PSF-fitting photometry was performed using the task allstar on all the sources detected with a 3 threshold. The absolute calibration of the photometry was performed by comparing the photometry of bright, isolated stars with their 2MASS values. No significant color terms were found in the photometric calibration.
The stars with 10.8 mag (20 stars) are saturated in the 2005 SOFI images. For those stars the 2MASS photometry is used instead. For two objects (3 and 7 in Table 4), the 2MASS photometry is contaminated by a neighboring bright star and the SOFI and band magnitudes of Bik (2004) are used. The magnitudes of these objects in the -band (which was not covered by Bik 2004) are calculated by performing spectrophotometry on their SINFONI spectra. The seeing conditions during these observations were sufficient to avoid contamination by neighboring stars. The SINFONI -band values of these stars agree well with the Bik (2004) observations.
The limiting magnitudes at 10 are approximately 19.1, 19.3, and 18.4 mag in , , and , respectively. A total of 745 sources are detected in the -band with a photometric error of mag. Of these, 395 are detected in all three bands with a photometric error of mag and a positional agreement of . The point sources are numbered according to their magnitude.
2.2 Mid-infrared imaging and photometry: Spitzer/IRAC
Imaging data taken with IRAC (Fazio et al. 2004) on board of Spitzer have been retrieved from the Spitzer archive (program ID 20819, P.I. Tsujimoto). The data of RCW 36 have been taken in the high-dynamic range mode, consisting of a set of deep images with a frame time of 10.4 s and a set of images taken with a frame time of 0.4 s to ensure that the brightest sources were not saturated. The raw data have been processed with the standard IRAC pipeline (version 18.18.0) to create the basic calibration data (BCD). These BCD were downloaded from the Spitzer archive and processed by custom IDL routines as described in Balog et al. (2007).
As the point spread function of the IRAC images is undersampled, obtaining aperture photometry is preferred. Aperture photometry on the reduced mosaic was performed using daophot inside IRAF. The sources (618 in band 1) were detected using daofind and photometry was performed using phot with an aperture equal to the FWHM of the stellar PSF. The background is measured in an annulus between 2 and 6 pixels around the star. The photometry is corrected with the aperture corrections taken from the IRAC handbook. Sources with a positional agreement of 2 pixels () between the different IRAC bands were matched. Some matches were discarded because of the faulty detection of some of the filaments in the 8 m IRAC-band as point sources.
The stellar population under study is located on the forefront of a dense dust cloud with up to 100 mag (Hill et al. 2011). Contamination from extragalactic background sources is thus expected to be negligible. Moreover, only two sources dimmer than the brightness limit formulated by Gutermuth et al. (2009) for extragalactic sources (i.e. [3.6 m] mag) were detected.
IRAC detections were matched with SOFI sources within . Increasing this radius with a factor 2 did not lead to a different result. Some spurious matches were discarded upon careful examination of the images. Only the sources with a photometric error of less than 0.1 mag (0.2 mag for IRAC) are used for the analysis. This results in 250 matches between IRAC band 1 and SOFI and 23 matches between all IRAC bands and SOFI . The limiting magnitudes at 10 were 15.3, 12.7, 11.4 and 7.5 mag in bands 1, 2, 3 and 4 respectively.
2.3 Near-infrared integral field spectroscopy: VLT/SINFONI
Near-infrared - and -band spectra have been taken with the integral field spectrograph SINFONI (Eisenhauer et al. 2003; Bonnet et al. 2004), mounted on UT4 of the ESO Very Large Telescope (VLT) on Cerro Paranal in Chile. The data were obtained in service mode between February 28 and March 23, 2007, with a typical seeing of . RCW 36 was observed using the non-adaptive optics mode with the pixel scale, resulting in an field of view. To obtain an - and -band spectrum the grating was selected resulting in a spectral resolution of and a wavelength coverage from m and m.
To cover the area of the cluster as shown in Fig. 1, a mapping pattern was applied with offsets of and . The offsets were designed such that every pixel in the field of view is covered as least twice. A detector integration time (DIT) of 30 seconds per integration was chosen. Sky frames were taken every 3 minutes on carefully selected offset positions with the same integration times to ensure an accurate sky subtraction. After every science observation a telluric standard star was observed at the same airmass to enable correction for the telluric absorption lines.
The SINFONI data are reduced using the SPRED software package version 1.37 (Schreiber et al. 2004; Abuter et al. 2006). The data reduction procedure is described in detail in Bik et al. (2010) and consists of dark and flat field correction of the raw data. After a distortion correction, the merged 3D datacubes were created. Telluric standard stars were used to correct for the telluric absorption lines, as is described in Ellerbroek et al. (2011).
2.4 Optical to near-infrared spectroscopy: VLT/X-SHOOTER
Eight objects have been observed with X-shooter, mounted on UT2 on the VLT, resulting in optical to near-infrared ( nm) spectra, see Table 1. The slits used were (UVB, nm), (VIS, nm) and (NIR, nm). This resulted in a spectral resolution of 5000, 9000 and 11,000 in the three arms, respectively. Directly before or after these observations the A0V star HD80055 was observed in order to remove telluric absorption lines in the near-infrared. A spectrophotometric standard was observed each night for flux-calibration. The spectra were reduced using the X-shooter pipeline (version 1.3.7 Modigliani et al. 2010).
The X-shooter spectra of the early-type stars (1, 2, 3 and 10) ensure a more precise spectral classification than that obtained with SINFONI, while the spectrum of object 2 also contains information on its circumstellar material. This is also the case for the young stellar objects (4, 9 and 97). Finally, the late-type PMS star 26 was observed with X-shooter in order to check the consistency of the spectral classification of the SINFONI spectra of late-type stars.
|Object||HJD||Exp. time||continuum S/N|
|#||B05111See Bik et al. (2005, 2006); Ellerbroek et al. (2013).||(s)||460 nm||800 nm|
3 Results from photometry
In this section, we use the near-infrared SOFI photometry to identify the stellar population (Sect. 3.1) and the IRAC photometry to reveal the circumstellar material surrounding it. The stellar density increases towards the location of the massive stars, objects 1 and 3 (see Fig. 1), which we define as the center of the cluster.
3.1 Near-infrared imaging
The () color-magnitude diagram (Fig. 2a) shows a reddened stellar population consistent with a distance of 0.7 kpc. The population blueward of mag probably consists of low-mass foreground stars and is not considered to be part of the cluster.
Fig. 2b shows the (, ) color-color diagram. Also here the foreground population is clearly visible. To correct for reddening due to interstellar extinction. We adopt the extinction law from Cardelli et al. (1989) with the total-to-selective extinction parameter set to the average Galactic value of (which may be an underestimate, see also Sec. 4.4). We conclude that the majority of the sources are found along the reddened location of the main sequence. However, several objects are located below this reddening line. These objects possess a near-infrared excess, indicative of a circumstellar disk and hence of their young age.
Using the above mentioned extinction law, we calculated the average extinction towards the stellar population of RCW 36. Excluding the foreground population, the average reddening detected towards the main sequence is mag with a 1 spread of 5.5 mag. This large spread on the average value of the extinction suggests that differential extinction is strongly affecting the appearance of the stellar population.
An alternative way to estimate the average interstellar extinction is obtained by dereddening to the locus of the classical T Tauri stars (cTTS, Meyer et al. 1997); this assumes for every source an intrinsic infrared excess due to a circumstellar disk. This likely leads to a more accurate estimate of . The mean extinction obtained by dereddening all sources to the cTTS locus is mag, about 30% less than the aforementioned value, but consistent within the uncertainty. This estimate agrees with the average value of mag found by Baba et al. (2004) who use the same method.
The photometric data of RCW 36 do not show distinct subgroups within the stellar population that have a different extinction or infrared colors, apart from the “foreground population” defined above. For a more thorough treatment of the extinction properties, see Sect. 4.4.
3.2 Mid-infrared imaging
Fig. 3a displays the IRAC colors of the 210 sources which were detected (with a photometric error mag) in the first two IRAC bands and with a -band counterpart. Using the classification scheme of Gutermuth et al. (2009), which makes use of the and first two IRAC bands, we classified 18 of these sources as class 0/I and 70 sources as class II. This classification is consistent with the classification scheme of Megeath et al. (2004) and Allen et al. (2004) for the objects that are detected in all four IRAC bands (Fig. 3b). We find that about half of the sources detected in , and at least in the first two 2 IRAC bands also have an intrinsic infrared excess. This suggests a high disk fraction in this cluster and hence a young age (Lada & Lada 2003; Hernández et al. 2008). A more accurate estimate of the disk fraction would be obtained by a completeness analysis, which is beyond the scope of this paper.
The spatial distribution of the Lada classified sources (Fig. 4a, c) shows that many of the class 0/I sources are associated with the filamentary structures. Although they trace point sources in all IRAC bands, their flux in band 4 is possibly contaminated by emission from the filaments, resulting in a classification as class 0/I.
4 Results from spectroscopy
We have obtained SINFONI and X-shooter spectroscopic observations of the brightest objects. In this section we analyze the early-type spectra (X-shooter), the late-type spectra (SINFONI) and the YSOs (X-shooter). Fig. 4e displays the coverage of these observations. The extinction map of the cluster is displayed in Figs. 4b and d; the SINFONI map of the nebular lines is shown in Fig. 4f.
4.1 Early-type stars: objects 1, 2, 3 and 10
The results of the spectral analysis of the four brightest photospheric spectra found in RCW 36 are summarized in Table 2. In order to constrain the spectral types of the O- and B-type stars (Objects 1, 3 and 10), non-LTE atmosphere models (FASTWIND; Puls et al. 2005) were fitted to the spectra using a genetic algorithm approach. This allows for the simultaneous determination of the main stellar and wind parameters using a selection of 11 helium and hydrogen lines. We refer to Mokiem et al. (2005) for a description of the algorithm, the parameters and the fitted lines. The resulting best-fit line profiles for both stars are overplotted on the spectra in Fig. 5.
The spectra of objects 1, 3 and 10 have no spectral signatures of circumstellar material, nor anomalously strong mass loss ( M yr) or rotation. We determine optical spectral types of the O stars by quantitative EW measurements following Conti & Alschuler (1971) and Mathys (1988, 1989), as revised by Sana et al. (in prep.). The spectral types of objects 1 and 3 are O8.5 – 9.5 V and O9.5 – B0 V, respectively. Object 10 is a B2 – 3.5 IV star based on comparison with synthetic spectra from Munari et al. (2005). Its surface gravity is not well constrained by the atmosphere fitting and may indeed be lower than the ZAMS value. This could indicate that the star is still in a PMS contraction phase, although it is already on the ZAMS. Note that the narrow, deep central absorption in the H i lines is due to oversubtraction of the nebular spectrum.
Object 2 is classified by comparing its spectrum in the region nm to synthetic spectra from Munari et al. (2005) to obtain a spectral classification. Using the Ca K line as the main temperature indicator in the optical, a spectral type A2 – A4 is determined. Based on the shape of the Balmer line wings, the luminosity class is IV. Object 2 exhibits a significant excess emission starting at 1 m and a flat (class II disk) SED. For its placement in the HRD, the absolute magnitude is corrected for the intrinsic infrared excess (amounting to 1.21 mag in ). Some of the emission lines in its spectrum (H i, He i, Fe ii) have asymmetric or P-Cygni profiles, indicating the presence of a stellar wind (See Fig. 6). The emission profiles of higher H i transitions are more symmetric and double-peaked, pointing to their origin in a high-density medium (i.e. an inner gas disk). Based on the above, we classify object 2 as a Herbig Ae star. Its radius is “bloated” with respect to its main sequence size, consistent with what is expected (Palla & Stahler 1993) and observed (Ochsendorf et al. 2011) in intermediate-mass PMS stars.
The extinction towards objects 1, 2, 3 and 10 is determined within 0.5 mag uncertainty by fitting the slope of the SED to a Kurucz atmospheric model (Kurucz 1993) in the entire X-shooter wavelength range (for object 2 up to 1 m, see Fig. 7). Finally, the stellar radius is determined by scaling the observed flux to the flux at the stellar surface given by the model; this scaling is degenerate with the distance estimate. With the spectral types and confirmed ZAMS nature of the OB stars, the spectroscopic parallax method is used to estimate the distance towards the cluster. For this we use the -band magnitudes and extinction of objects 1, 3 and 10, and adopt the absolute magnitude calibration defined by the “observational” scale defined in Martins & Plez (2006). The results (Table 2) are in the range of 0.6 – 0.8 kpc, consistent with the distance of 0.7 kpc obtained from literature and with the spectroscopic distance estimates by Bik et al. (2005). The error in the spectroscopic distance determination is dominated by the error in .
A radial velocity of km s could be determined from the photospheric lines of object 3. This is in concurrence with the nebular velocity km s, measured by Bronfman et al. (1996) from the CS(2-1) emission line. Based on its colors and magnitude, object 14 (located outside the SINFONI field) is probably also a main-sequence OB star, but due to its high extinction ( mag) and location outside the cluster center it is either an unrelated background star, or an ejected “runaway” star. This would have to be confirmed by spectroscopy.
4.2 Late-type stars
A total of 138 SOFI point sources with magnitudes between mag and mag have an associated point source in the SINFONI continuum images. The observations provide an - and -band spectrum of every source with a spectral resolution of . We have classified the spectra with a signal-to-noise ratio S/N in the -band (corresponding to mag); for these a spectral type could be determined within two subtypes. As a result, 47 sources are assigned a spectral type and luminosity class. See Fig. 4e for their location; part of the central cluster is not covered by the SINFONI observations. Table 4 summarizes the spectral types and characteristics of the classified sources; their SINFONI spectra are displayed in Fig. 11.
Most of these stars have photometric spectral types later than F. The late-type spectra are compared with reference spectra from Cushing et al. (2005) and Rayner et al. (2009). The spectral classification is based on atomic and molecular absorption lines (e.g., Mg i, Na i, Ca ii, and CO). The depth of the CO lines at 2.3 m serves as an indicator for the luminosity class. The photometric spectral type of object 26 (K4 V) agrees with the spectral type determined from its X-shooter spectrum, which was obtained using the region around the Ca ii triplet at 850 nm. Many spectra exhibit H i emission, particularly in the Br line; the spectra for which this emission appears as a point source in the line map are indicated in Table 4.
The uncertainty in the spectral classification (typically one subtype) is reflected in the errors on the effective temperature, , and on the extinction, . The effective temperature is determined by the calibration in Kenyon & Hartmann (1995). Since this calibration overestimates (by K for G to K for mid-K, Cohen & Kuhi 1979) for PMS stars, a correction was made which contributes significantly to the error budget. Using the intrinsic colors listed by Kenyon & Hartmann (1995) and the -colors from SOFI, we determine for every classified source. Where available, we based the determination of on the - and -band fluxes so as to avoid contamination of continuum excess emission caused by circumstellar material. The thus derived estimates agree reasonably well with the values derived from dereddening to the cTTS locus (Sect. 3.1); their mean ratio is . The spread in these values can in large part be explained by the spread in the intrinsic colors of cTTS stars found by Meyer et al. (1997).
The above procedure results in an absolute magnitude for every classified star, from which the stellar radius can also be determined. Also, the stellar mass and age were obtained from the absolute magnitudes and effective temperatures, by interpolating between evolutionary tracks and isochrones; see Sect. 5.1.
4.3 Young Stellar Objects
In this section, we discuss the spectra of three YSOs in RCW 36 which stand out by their infrared brightness and/or associated outflows: objects 4, 9 and 97. The spectra of 4 and 97 are also discussed in detail in Ellerbroek et al. (2013); a selection of emission lines from object 9 is displayed in Fig. 6.
The spectrum of object 4 (also known as 08576nr292) exhibits no photospheric features, and is dominated by continuum emission from a circumstellar disk. It contains many emission lines that originate in a disk-jet system. This intermediate-mass YSO has a high accretion rate ( M yr, Ellerbroek et al. 2013), and is associated with the Herbig-Haro jet HH~1042, demonstrating its current accretion activity. For an extensive study of object 4 (08576nr292) and its disk-jet system, we refer the reader to Ellerbroek et al. (2011, 2013). We adopt the mag based on SED fitting in the former paper. Object 4 is classified as a class 0/I YSO, but it could also be class II as its red bands may be contaminated by emission from the surrounding cloud.
Object 97 (08576nr480) is another class 0/I YSO associated with a jet (HH~1043), but it is much more embedded than object 4. Objects 4 and 97 are both superposed on the filamentary structures west of the central cluster, demonstrating star formation is ongoing in this region. Their jets contain emission lines from high ionization species like [S iii] and [O ii], which are also detected in the ambient nebular spectrum. This medium is probably ionized by the central O stars (Ellerbroek et al. 2013).
Object 9 is a class 0/I YSO whose severely reddened spectrum exhibits many emission lines (predominantly H i) with asymmetric line profiles with a strong blue-shifted “wing” with velocities up to km s indicative of an outflow (Fig. 6). However, the blue wing is only detected in lines associated with high densities and not in forbidden lines, indicating the wind is dense and optically thick, shrouding the stellar photosphere. A dust shell or disk might further obscure the central object, although no CO bandhead emission is seen in this object, unlike in objects 4 and 97.
|Region||Offset555R.A. and Dec. offsets from object 1, in .||Size||(mag)||He i 1.70 m/Br|
|(Fig. 4f)||()||()||(Br10 / Br)||( corr.)|
|I||14 3||13.1 0.4||0.070 0.007|
|II||3 2||19.6 0.8||0.059 0.015|
|IV||7 14||7.3 1.6|
4.4 Extinction and nebular spectrum
An extinction map (Fig. 4b, d) was produced based on the colors of the SOFI detections. The method for constructing the extinction map was based on that described by Lombardi & Alves (2001), which is in turn based on the color excess method first presented by Lada et al. (1994). We consider sources for which -photometry exists and exclude foreground sources (Sect. 3.1) and sources for which an intrinsic near-infrared excess is detected (Sect. 3.2). For every source, is determined by dereddening to the ZAMS (see Fig. 2a). These values are represented by the colored dots in Fig. 4b. Then a spatial grid is defined with a spatial resolution of 4. For every grid point, an value was calculated by taking the weighted average of of the 20 nearest neighboring stars. The weight used is the inverse squared distance to the nearest neighbors, with a minimum of (a smoothing parameter).
The spatial resolution is variable across the extinction map, depending on the local stellar surface density. The effective resolution element (defined as the mean distance to the 20 nearest neighbors) decreases radially from in the central arcminute of Fig. 4b to at the edges of the map. The error in the extinction measurement of a gridpoint can be expressed as the square root of the variance in of the 20 nearest neighbors, weighted with their inverse squared distance. We find that the relative error in varies between across the map, increasing up to in areas west and southeast of the center, where not many sources are found. However, the uncertainty is probably dominated by the systematic error in the adopted extinction law ( may be higher, and vary across the cluster) and the fact that some sources may possess an intrinsic infrared excess. Both cases would result in an overestimate of .
A gradient is visible from high extinction ( mag) in the eastern part to lower extinction ( mag) westward of the filamentary structures. As this western part is also where the 500 m flux peaks (see §5), the apparent low extinction value is due to the fact that we observe only those sources which are at the forefront of the molecular cloud. In this region, the SOFI sample is thus biased toward sources which are on average less embedded than the rest of the cluster population. Therefore low-extinction regions in Fig. 4b may in fact have a large column density, as the completeness of the dataset is affected by the presence of cold dust. Also, isolated sources that have a very different extinction compared to their nearest neighbors cause a local maximum or minimum in the extinction map, which should be attributed to circumstellar rather than interstellar material. No trend is found between the extinction and the evolutionary status of the sources.
A more reliable estimate of the extinction of individual sources can be obtained by spectral classification, as in that case the intrinsic colors are known. The 47 sources for which this is the case are overplotted on the extinction map in Fig. 4d. Most of the extinction values derived by spectroscopy are similar to the local value of the extinction map. A few sources have a lower individual than the ambient value. This may be because the sources of which a photospheric spectrum could be classified are typically located on the forefront of the cluster. Alternatively, the ambient extinction may be overestimated for the reasons mentioned above.
An independent estimate of may be obtained by using the observed correlation between the equivalent width (EW) of diffuse interstellar bands (DIBs) and extinction (e.g. Herbig 1993; Vos et al. 2011). Fig. 8 shows the correlation between the EW of the two strongest DIBs and (see Sect. 4.1) of the spectra in which DIBs were detected. The DIB strength to extinction ratios of these spectra are comparable, even though the extinction towards objects 2 and 4 is enhanced by circumstellar material. However, these ratios deviate from the correlation found by Vos et al. (2011). This discrepancy may reflect an underestimate of , which is seen to be higher than the average Galactic value of 3.1 in lines of sight towards star-forming regions (see e.g. Cardelli et al. 1989; Hoffmeister et al. 2008; Dahlstrom et al. 2013). Alternatively, the DIB carrier(s) may be less abundant in star forming regions, which would provide a constraint on its nature. The limited spread in extinction values in our sample prevent us from testing these possibilities.
The cluster extinction can also be determined by tracers in the nebular spectrum. Nebular spectra were extracted at five subregions of the H ii region in the SINFONI field of view (Fig. 4f). The locations and sizes of these regions are listed in Table 3. We have measured the fluxes of selected lines to calculate the extinction and the temperature of the radiation field. In an H ii region, a deviation of the predicted Br10/Br flux ratio (0.33) is mainly dependent on (Storey & Hummer 1995). The third column in Table 3 lists the derived extinction values. Regions I–III have an increasing amount of extinction, with the highest extinction at the location where the Spitzer mid-infrared flux also peaks (Region III). This is consistent with the high column density at this location (Hill et al. 2011). The region north of the O stars (IV) has a low amount of extinction, while region V has too low S/N to determine the extinction. The values derived by the Br10/Br ratio (Table 3) are higher than those derived in the extinction map, possibly because these nebular lines originate in the medium behind the stellar population. The extinction values in regions I, II and IV agree within error with those calculated with the color excess method (Fig. 4b).
The spectrum of the ionized nebula also provides an estimate on the temperature of its ionizing source. Lumsden et al. (2003) predict the (extinction-corrected) value of the He i 1.70 m/Br line flux ratio as a function of the temperature of the ionizing star. Only in regions I and II the He i 1.70 m flux is bright enough; the derived ratios ( and ) are consistent with a temperature of K. This coincides with the temperature determined from the optical spectra of objects 1 and 2 (Sect. 4.1). This corroborates the finding of Verma et al. (1994) that the emergent flux of the nebula is dominated by the contribution from the central O stars (see also Bik et al. 2005).
5 The star formation history of RCW 36
In this section, we summarize and interpret the results of the analysis of the stellar population, extinction and nebular properties. We first present the HRD and discuss the age of the stellar population (Sect. 5.1). We then compare our findings with the far-infrared observations of the large-scale molecular cloud in which RCW 36 is embedded and propose a scenario for the star formation history (Sect. 5.2).
5.1 Stellar population: age and spatial distribution
In the previous section we have derived and discussed the stellar parameters of the sources in RCW 36, summarized in Table 4. With this information we can construct a HRD (Fig. 9a–c). The stellar population is scattered along the 1 and 2 Myr PMS isochrones, with the O and B stars (objects 1, 3 and 10) already having arrived on the main sequence. Object 10 is possibly a naked and bloated PMS B star based on its apparently low surface gravity, although it is already (nearly) on the main sequence. Object 2 is the brightest class II object that was classified and will likely be the first PMS star arriving on the main sequence after object 10, within the next 0.5 Myr. Along the evolutionary tracks, these stars are followed by the population of low-mass stars, the brightest ones of which are confirmed as class II objects and possess a disk SED similar to object 2.
The PMS age distribution of the classified stars is displayed in a generalized histogram (Fig. 9d). For every classified star, a gaussian profile centered on its age (derived from the isochrones in Fig. 9b) with the age error as its width, was added to this histogram. These profiles were normalized so that every star contributes the same area to the histogram. The age distribution peaks at Myr, which is the weighted mean age of the stellar population. Many sources have ages within this range, while a number of faint, low-mass ( M) sources are located around older isochrones; this is reflected by the “tail” towards high ages in Fig. 9d. However, this “tail” disappears when only considering sources with age errors Myr (which make up about 75% of the sample). Thus, the age distribution is consistent with being the result of the large uncertainties in the ages of the fainter sources, which arise because of the dense spacing of isochrones in this region of the HRD. This does not necessarily imply the presence of an older generation of stars.
No significant age gradient can be found in the spatial distribution of classified stars. Based on the position of the bright, centrally located sources in the HRD we infer that the most massive stars in the cluster are in the center of the cluster, as was also concluded by Baba et al. (2004).
5.2 Comparison with molecular cloud structure
RCW 36 is embedded in an elongated dense filamentary structure of pc length (Hill et al. 2011). The far-infrared observations with Herschel reveal a bipolar nebula extending EW, with a ring-like structure with a major axis of 2 pc at its center. MTHM13 propose that this morphology is the result of a “blowout”: the ionizing radiation of massive stars cleared out the central region of the ring, ionized its inner edges, and heated the nebula on either side of the ring resulting in the bipolar nebula. Such parsec-scale bipolar molecular outflows are commonly associated with massive star-forming regions (Arce et al. 2007). According to the modeling by MTHM13, the timescale for the currently observed ring and bipolar nebula to form out of an initial filament is less than 1 Myr.
Our study of the stellar population shows that the massive and intermediate-mass stars in the center of RCW 36 have the right age and characteristics for having shaped the ring-like structure and bipolar nebula. The ring is likely blown out by the stellar winds of the O stars. The stellar parameters of objects 1 and 3 yield a combined wind luminosity of 10 erg s (Vink et al. 2000). Assuming an initial ISM density of 100 cm and a lifetime of 1 Myr, such a wind would have blown out a cavity with radius pc (Koo & McKee 1992, equation 4.4), which is the observed size of the ring. Radiation escapes into the regions east and west of the ring, which have a lower density, resulting in the bipolar nebula.
Fig. 10 shows the location of the stellar population with respect to the large-scale cloud structure. The stellar surface number density, the number of SOFI detections per pc, is calculated as the inverse of the squared mean distance to the 20 nearest neighbors on a grid with a spatial resolution of 5. It peaks at the location of the central O stars (at pc) and decreases outward (down to pc at the edge of the SOFI field of view). The ring structure (with its major axis extending NS) found by MTHM13 is clearly visible in this image, as well as the bipolar nebula extending EW. The O stars are located along the major axis of the ring, although south of the center. Within the ring, around the O stars, a co-evolving retinue of intermediate-mass (1-5 M) stars is found. Class 0/I objects, including the two jet sources HH 1042 and HH 1043, are found mainly along the ring structure, although some (e.g. object 9) also occupy the central region. Fainter class II objects are spread over the cluster; most are found within or along the ring structure.
Star formation in RCW 36 is currently ongoing. Some protostellar cores (Giannini et al. 2012, MTHM13) and a UCHII region (Walsh et al. 1998) are found in the region where the far-infrared flux peaks (region II/III, Fig. 4f). The column density in this region is very high (Hill et al. 2011), consistent with the fact that radiation only escapes from its surface and not from the embedded sources. It is entirely possible that new massive stars are being formed within this thick column of dust. Star formation is actively taking place at the forefront of this region, as demonstrated by the presence of two HH objects (Ellerbroek et al. 2013). Their positioning with respect to the cloud is reminiscent of the HH 901 and HH 902 jets emerging from the pillars north of the cluster Trumpler 14 in the Carina Nebula (Smith et al. 2010).
We have presented a multi-wavelength photometric and spectroscopic study of the massive star-forming region RCW 36. The detailed characterization of the stellar population has resulted in a formation scenario which is in line with (and can partly explain) the current large-scale structure of the molecular cloud. Our most important findings are summarized below:
RCW 36 contains a young PMS stellar population; about half of the classified objects have an intrinsic infrared excess. This suggests that many sources in the cluster have disks, which would be consistent with its young age.
The radiative energy output is dominated by two O stars (objects 1 and 3) in the cluster center, which illuminate the filamentary edges of the dissipated molecular cloud. These objects, along with a B star (object 10), are located on or close to the ZAMS and are thus probably co-evolving with the lower-mass PMS population.
A Herbig Ae star (object 2) is the most massive object on its way to the main sequence; its bloated radius and circumstellar disk confirm its PMS nature.
The classified stellar population has a mean PMS age of Myr. A number of low-mass ( M) objects appear to be older, although large uncertainties exist in their age determination. Our findings are consistent with a single-age population; the uncertainties are too large to make quantitative statements about an age spread in this cluster.
A likely scenario for the star formation history of the cluster is obtained by combining the studies of the molecular cloud and the stellar population. The brightest (most massive) objects that formed 1 Myr or more ago, are found along the “sheet”, with the O stars in the middle of a ring-shaped structure, which may be shaped by the central stars’ radiation field. The timescale for the formation of the ring agrees with the age of the central stellar population.
While an embedded forming massive stellar population may exist beyond a large column of dust (the purple circles in Fig. 4a, and the UCHII), a new generation of lower-mass stars is observed at the forefront of this region (objects 4 and 97). Other class 0, I and II YSOs are found in the cluster center as well as the periphery.
RCW 36 is revealed as a very rich cluster which can serve as a excellent “laboratory” to study sequential star formation. The impact of newly formed massive stars on the ambient molecular cloud is clearly exposed. Our study suggests that the surrounding intermediate-mass stellar population has an age compatible with the dynamical age of the cluster. The presence of an even younger generation of protostellar objects (massive protostellar cores, a UCHII region, YSOs and protostellar jets) shows that RCW 36 is an active site of star formation, which will considerably reshape the molecular cloud within the next millions of years.
Acknowledgements.The authors thank the anonymous referee for providing useful comments and suggestions that helped improve the paper. A.B. acknowledges the hospitality of the Aspen Center for Physics, which is supported by the National Science Foundation Grant No. PHY-1066293. This research project is financially supported by a grant from the Netherlands Research School for Astronomy (NOVA).
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Appendix A Classified sources in RCW 36
Table 4 contains an overview of the stellar properties of a selected sample of sources which are assigned a photometric spectral type and/or are individually described in the text. Fig. 11 displays the normalized SINFONI spectra of the same sources. The locations of some of the absorption (and emission) lines used for spectral classification are indicated with vertical lines.
|# (B05)||RA (h m s)||Dec ( )||Data||Sp. Type||Lada Class||Br em.||T (K)||(mag)||(mag)||(mag)||(R)||(M)|
|1 (462)||8 59 27.34||-43 45 25.84||JHK, 1234, S, X||O9 V||34200||7.02||-3.32||10.4||8.2||22.0|
|2 (408)||8 59 28.51||-43 46 02.94||JHK, 1234, S, X||A3 IV||II||+||8720||7.40||-1.54||8.5||5.1||4.9|
|3 (413)||8 59 27.55||-43 45 28.41||JHK, 1234, S, X||O9.7 V||34900||7.55||-2.71||9.6||6.0||20.0|
|4 (292)||8 59 21.67||-43 45 31.05||JHK, 1234, S, X||…||0/I||+||…||9.32||…||8.0||…||…|
|6||8 59 26.24||-43 45 27.81||JHK, 123, S||G6 III||5450||9.55||-1.30||15.1||8.7||4.0|
|7||8 59 27.01||-43 45 28.36||JHK, 12, S||K1 IV||4980||9.59||-0.61||9.0||6.7||3.5|
|9||8 59 27.40||-43 45 03.74||JHK, 1234, S, X||…||0/I||+||…||9.63||…||…||…||…|
|10 (179)||8 59 27.74||-43 45 38.23||JHK, 123, S, X||B2.5 IV||19750||9.69||-0.59||9.8||2.9||6.0|
|13||8 59 28.02||-43 45 19.40||JHK, 1234, S||K0 V||II||+||5150||10.08||-0.50||12.5||6.2||3.3|
|15||8 59 27.86||-43 45 59.75||JHK, 1234, S||K0 V||II||5150||10.21||-0.30||12.0||5.7||3.2|
|18||8 59 29.90||-43 46 26.04||JHK, 1234, S||G8 V||II||+||5270||10.27||-0.43||13.7||5.9||3.2|
|26||8 59 24.18||-43 45 25.91||JHK, 12, S, X||K4 V||+||4490||10.77||0.51||9.6||4.3||1.5|
|27||8 59 26.23||-43 45 44.43||JHK, 1234, S||G0 V||II||+||5780||10.79||-0.23||16.7||5.2||2.7|
|32||8 59 26.74||-43 45 30.07||JHK, 1234, S||K1 IV-V||II||+||4980||11.05||0.17||15.3||4.8||2.8|
|34||8 59 28.40||-43 44 42.84||JHK, 1234, S||G3 V||II||5580||11.17||0.20||16.1||4.3||2.5|
|40||8 59 28.56||-43 46 30.43||JHK, 123, S||G4 IV-V||+||5550||11.39||0.24||17.9||4.3||2.5|
|42||8 59 29.73||-43 46 05.44||JHK, 123, S||K6 V||II||4105||11.47||1.13||10.4||3.4||0.8|
|44||8 59 26.52||-43 45 38.02||JHK, 1234, S||K1 V||II||+||4980||11.57||0.50||17.1||4.1||2.5|
|47||8 59 27.58||-43 45 25.00||HK, S||K1 V||+||4980||11.72||1.29||11.1||2.8||2.0|
|51||8 59 28.03||-43 45 16.57||JHK, S||K4 IV-V||4490||11.79||0.86||15.8||3.7||1.4|
|56||8 59 22.63||-43 45 15.49||JHK, 12, S||K0 IV||II||5150||11.87||1.67||9.0||2.3||1.7|
|67||8 59 28.06||-43 45 00.58||JHK, 12, S||K3 IV||4630||12.16||1.30||15.1||2.9||1.6|
|71||8 59 24.59||-43 45 13.87||JHK, S||M0 V||3800||12.23||2.15||7.9||2.4||0.5|
|72||8 59 26.20||-43 45 23.83||JHK, 12, S||G5 V||+||5520||12.25||1.32||15.8||2.6||1.8|
|82||8 59 22.61||-43 45 49.67||JHK, 12, S||K2 IV||4800||12.36||2.02||10.4||2.0||1.6|
|83||8 59 25.83||-43 45 35.83||JHK, 12, S||M0 IV-V||3800||12.39||1.82||12.5||2.7||0.5|
|85||8 59 29.04||-43 45 04.50||JHK, 123, S||K3 V||II||4630||12.45||1.77||13.6||2.3||1.5|
|88||8 59 26.87||-43 44 47.40||JHK, 123, S||K3 V||II||+||4630||12.49||1.78||13.8||2.3||1.5|
|93||8 59 29.99||-43 45 51.88||JHK, 12, S||K1 V||II||4980||12.55||2.10||11.4||1.9||1.5|
|94||8 59 31.21||-43 45 52.41||JHK, 12, S||K3 V||4630||12.56||1.90||13.3||2.2||1.5|
|95||8 59 30.61||-43 46 10.10||JHK, 12, S||K6 V||4105||12.60||1.99||12.9||2.3||0.8|
|97 (480)||8 59 23.66||-43 45 30.51||HK, S, X||…||0/I||…||12.61||…||…||…||…|
|100||8 59 24.71||-43 45 03.32||JHK, 12, S||M1 V||3670||12.64||2.40||9.5||2.1||0.4|
|101||8 59 25.29||-43 45 13.99||JHK, 12, S||K1 V||+||4980||12.68||2.12||12.4||1.9||1.5|
|102||8 59 26.46||-43 45 53.99||JHK, 12, S||M0 V||3800||12.69||1.82||15.2||2.9||0.5|
|105||8 59 25.35||-43 45 50.67||JHK, 123, S||M2 V||+||3530||12.72||2.09||13.1||2.7||0.4|
|108||8 59 27.25||-43 46 14.62||JHK, 12, S||K4 V||4490||12.73||2.75||7.0||1.5||1.2|
|110||8 59 31.25||-43 46 25.86||JHK, 12, S||K3 III||4630||12.77||2.20||12.5||1.9||1.4|
|116||8 59 25.08||-43 45 45.90||JHK, 12, S||M1 V||3670||12.81||1.94||15.4||2.7||0.4|
|120||8 59 28.39||-43 46 28.46||JHK, S||K6 IV-V||4105||12.85||2.39||11.4||1.9||0.8|
|124||8 59 31.06||-43 45 55.54||JHK, 12, S||K6 V||II||4105||12.91||2.42||11.8||1.9||0.8|
|132||8 59 26.99||-43 45 52.10||JHK, 123, S||K3 V||II||4630||12.98||1.78||18.4||2.4||1.5|
|134||8 59 28.32||-43 45 20.85||JHK, S||K3 V||+||4630||13.00||2.43||12.5||1.7||1.3|
|144||8 59 28.12||-43 45 55.49||JHK, 1, S||K4 V||4490||13.14||2.76||10.7||1.5||1.2|
|148||8 59 27.05||-43 45 32.34||JHK, 123, S||K3 IV-V||4630||13.19||2.47||13.8||1.7||1.3|
|161||8 59 28.64||-43 45 13.54||JHK, 12, S||M1 IV||II||3670||13.26||2.55||13.7||1.9||0.4|
|172||8 59 27.91||-43 45 04.93||JHK, S||K5 IV||4250||13.30||2.36||16.0||1.9||1.0|
|180||8 59 28.88||-43 45 51.73||JHK, 12, S||K6 V||4105||13.39||3.26||8.4||1.3||0.9|
|186||8 59 25.30||-43 45 31.20||JHK, 12, S||K3 III||4630||13.46||2.74||13.8||1.5||1.2|
|191||8 59 29.67||-43 45 02.16||JHK, 12, S||G7 V||II||5380||13.48||3.15||10.3||1.1||1.0|
|: Id from Bik et al. (2005, 2006); Ellerbroek et al. (2013). : Available data: SOFI (JHK), IRAC (bands 1234), SINFONI (S), X-shooter (X).|
|: Based on comparison with reference spectra from Rayner et al. (2009). For objects 1, 3 and 10, see Sec. 4.1.|
|: Evolutionary class (Lada 1987) derived from IRAC colors as defined by Megeath et al. (2004); Gutermuth et al. (2009).|
|: From Kenyon & Hartmann (1995), averaged between the spectral type range. For objects 1, 3 and 10, see Sec. 4.1.|
|: Assuming a distance of 0.7 kpc. : Derived using intrinsic colors of Kenyon & Hartmann (1995) and Martins & Plez (2006). For objects 1, 2, 3 and 10, SED-fitting was used (Sect. 4.1).|
|: Result of fitting with Kurucz model (Kurucz 1993) : Result of interpolating evolutionary tracks and isochrones (Siess et al. 2000; Da Rio et al. 2009; Ekström et al. 2012).|