11institutetext: Purple Mountain Observatory, Chinese Academy of Sciences,
210008 Nanjing, China
11email<EMAIL_ADDRESS>22institutetext: Max-Plank-Institute for Astronomy,
Königstuhl 17, 69117 Heidelberg, Germany
33institutetext: Graduate University of the Chinese Academy of Sciences, 19A
Yuquan Road, Shijingshan District, 100049 Beijing, China
44institutetext: Centro de Astrobiología (CSIC-INTA), Ctra. de Torrejón a
Ajalvir km-4, E-28850, Torrejón de Ardoz, Madrid, Spain
55institutetext: Instituut voor Sterrenkunde, Katholieke Universiteit Leuven,
Celestijnenlaan 200D, 3001 Leuven, Belgium
66institutetext: Laboratoire dAstrophysique de Marseille, UMR6110 CNRS, 38 rue
F. Joliot-Curie, 13388 Marseille, France 77institutetext: National
Astronomical Observatory of Japan, and GUAS, National Institutes of Natural
Sciences, Japan

# Different Evolutionary Stages in the Massive Star Forming Region S255
Complex

Y. Wang 112233 H. Beuther 22 A. Bik 22 T. Vasyunina 22 Z. Jiang 11 E. Puga
4455 H. Linz 22 J. A. Rodón 66 Th. Henning 22 and M. Tamura 77

(Received 7 August 2010; accepted 11 November 2010)

###### Abstract

Aims. Massive stars form in clusters, and they are often found in different
evolutionary stages located close to each other. To understand evolutionary
and environmental effects during the formation of high-mass stars, we observed
three regions of massive star formation at different evolutionary stages that
reside in the same natal molecular cloud.

Methods. The three regions S255IR, S255N and S255S were observed at 1.3 mm
with the Submillimeter Array (SMA) and follow-up short spacing information was
obtained with the IRAM 30m telescope. Near infrared (NIR) $H+K$-band spectra
and continuum observations were taken for S255IR with VLT-SINFONI to study the
different stellar populations in this region.

Results. The combination of millimeter (mm) and near infrared data allow us to
characterize different stellar populations within the young forming cluster in
detail. While we find multiple mm continuum sources toward all regions, their
outflow, disk and chemical properties vary considerably. The most evolved
source S255IR exhibits a collimated bipolar outflow visible in CO and H2
emission, the outflows from the youngest region S255S are still small and
rather confined in the regions of the mm continuum peaks. Also the chemistry
toward S255IR is most evolved exhibiting strong emission from complex
molecules, while much fewer molecular lines are detected in S255N, and in
S255S we detect only CO isotopologues and SO lines. Also, rotational
structures are found toward S255N and S255IR. Furthermore, a comparison of the
NIR SINFONI and mm data from S255IR clearly reveal two different (proto)
stellar populations with an estimated age difference of approximately 1 Myr.

Conclusions. A multi-wavelength spectroscopy and mapping study reveals
different evolutionary phases of the star formation regions. We propose the
triggered outside-in collapse star formation scenario for the bigger picture
and the fragmentation scenario for S255IR.

###### Key Words.:

stars: formation – ISM: jets and outflows – ISM: molecules – stars: early-type
– Hertzsprung-Russell (HR) and C-M diagrams – stars: individual: S255IR,
S255N, S255S

##  1 Introduction

Massive stars ($M>8M_{\odot}$) are one of the paramount components in the
evolution of the universe, yet their formation is significantly less well
understood. than that of their low-mass counterparts. S255IR is a famous
massive star formation complex at a distance of $1.59^{+0.07}_{-0.06}$ kpc
(Rygl et al., 2010), embraced by the Sharpless regions S255 and S257 that are
already evolved HII regions. The SCUBA 850 $\mu$m observation (Di Francesco et
al., 2008; Klein et al., 2005) shows three main continuum sources: G192.60-MM2
(Minier et al., 2005) lies in the center region S255IR, and two additional mm
continuum peaks G192.60-MM1 (Minier et al., 2005) and G192.60-MM3 (Minier et
al., 2005) toward the northern region S255N and southern region S255S,
respectively (Figure 1).

The central region S255IR: The central IRAS source with a luminosity of
$5\times 10^{4}$ L⊙ (Minier et al., 2005) harbors three UCHII regions region
(Snell & Bally, 1986) that is associated with Class II CH3OH and H2O maser
emission (Goddi et al., 2007), which indicate the presence of massive young
stellar objects. The luminosity would be $2\times 10^{4}$ L⊙ if we apply the
new distance of 1.59 kpc. The region hosts a cluster of low-mass sources that
is surrounded by a shocked bubble of H2 gas (Ojha et al., 2006; Miralles et
al., 1997). Tamura et al. (1991) resolved the central near infrared source
into two compact sources, NIRS 1 and NIRS 3. Furthermore, Minier et al. (2007)
reported HCO+ infall signatures toward this region. In the mid-infrared,
Longmore et al. (2006) resolved a massive proto-binary which coincide with
NIRS 1 and NIRS 3 (Figure 3). And NIRS 1 has been identified as a massive disk
candidate by near-infrared polarization observations (Jiang et al., 2008, see
also Simpson et al. (2009) for a different interpretation).

The northern region S255N: This region with 105 L⊙ (Minier et al., 2005) also
hosts a UCHII region G192.584-0.041 (Kurtz et al., 1994) as well as Class I
CH3OH and H2O maser emission (Kurtz et al., 2004; Cyganowski et al., 2007),
which indicate the presence of the massive young stellar objects. The
luminosity would be $4\times 10^{4}$ L⊙ if we apply the new distance of 1.59
kpc. The recent Submillimeter Array observations toward that region at a
spatial resolution of 3.6′′ resolved three mm sources and strong molecular
line emission, and none of the sources is associated with NIR point source
(Cyganowski et al., 2007). Outflow activity is evidenced by various tracers
from shocked H2 emission via cm continuum and SiO emission (Miralles et al.,
1997; Cyganowski et al., 2007), and global infall was reported by Minier et
al. (2007).

The southern region S255S: This sub-source is the least studied so far. It
exhibits strong mm continuum emission (Figure 1) but no other sign of active
star formation yet such as near infrared and mid infrared emissions.
Therefore, it is proposed to be in a pre-stellar phase of the evolutionary
sequence (Minier et al., 2007).

While G192.60-MM2 in S255IR still shows signs of active star formation, it is
associated with a NIR cluster and appears to be the most evolved one of the
three region. G192.60-MM1 in S255N is of similar luminosity, and G192.60-MM3
is S255S is a High-Mass Starless Core candidate. Therefore, S255 complex is
the ideal candidate source to simultaneously investigate several sites of
massive star formation at different evolutionary stages within the same
larger-scale environment.

##  2 Observations and Data reductions

###  2.1 Submillimeter Array observations

The S255 complex was observed with three fields with the Submillimeter Array
(SMA) on November 3rd 2008 in the compact configuration with seven antennas
and on February 8th, February 17th 2009 in the extended configuration with
eight antennas, and February 13th in the extended configuration with seven
antennas. The phase centers of the fields, which are known as S255IR, S255N
and S255S, were R.A. 06h12m54.019s Dec.
$+17^{\circ}59^{\prime}23.10^{\prime\prime}$ (J2000.0), R.A. 06h12m53.669s
Dec. $+18^{\circ}00^{\prime}26.90^{\prime\prime}$ (J2000.0) and R.A.
06h12m56.58s Dec. $+17^{\circ}58^{\prime}32.80^{\prime\prime}$ (J2000.0),
respectively. The SMA has two spectral sidebands, both 2 GHz wide and
separated by 10 GHz. The receivers were tuned to 230.538 GHz in the upper
sideband ($v_{lsr}$ = 10 km s-1) with a maximum spectral resolution of 0.53 km
s-1. The weather of February 8th and February 13th was mediocre with zenith
opacities $\tau$(225 GHz) larger than 0.2 measured by the Caltech
Submillimeter Observatory (CSO). But the weather of November 3rd 2008 and
February 17th 2009 was good with zenith opacities $\tau$ (225 GHz) around 0.1.
So we only used the data observed on these two days. For the compact
configuration on November 3rd 2008, bandpass was derived from the quasar
3c454.3 observations. Phase and amplitude were calibrated with regularly
interleaved observations of the quasar 0530+135 (11.4∘ from the source). The
flux calibration was derived from Uranus observations, and the flux scale is
estimated to be accurate within 20$\%$. For the extended configuration on
February 17th 2009, bandpass was derived from the quasar 3c273 observations.
Phase and amplitude were calibrated with regularly interleaved observations of
the quasar 0530+135. Because of the lack of flux calibrator observations, the
flux was estimated by the SMA calibrator database for the gain calibrator, to
be accurate within 20$\%$. We merged the two configuration data set together,
applied different robust parameters for the continuum and line data, and got
the synthesized beam sizes between 1.4′′ $\times$1.1′′ (PA 86∘) and
1.9${}^{\prime\prime}\times 1.6^{\prime\prime}$ (PA 77∘), respectively. The
3$\sigma$ rms of 1.3 mm continuum image is $\sim$ 4.5 mJy and the 3$\sigma$
rms of the line data is 0.14 Jy/beam at 2 km s-1 spectral resolution. The
flagging and calibration was done with the IDL superset MIR (Scoville et al.,
1993) which was originally developed for the Owens Vally Radio Observatory and
adapted for the SMA111The MIR cookbook by Chunhua Qi can be found at
http://cfa-www.harvard.edu/$\sim$cqi/mircook.html.. The imaging and data
analysis were conducted in MIRIAD (Sault et al., 1995).

###  2.2 Short spacing from the IRAM 30 m

To complement the CO$(2-1)$ observations with the missing short spacing
information and to investigate the large-scale general outflow properties to
find the connection of the three regions, we observed them with the HERA array
at the IRAM 30 m telescope on November 10th 2009. The 12CO$(2-1)$ line at
230.5 GHz and 13CO$(2-1)$ line at 220.4 GHz were observed in the on-the-fly
mode. We mapped the whole region with a size of 5${}^{\prime}\times$ 3′
centered at R.A. 06h12m54.02s Dec.
$+17^{\circ}59^{\prime}23.10^{\prime\prime}$ (J2000.0). The sampling interval
was 3.5′′, a bit better than Nyquist-sampling to minimize beam-smearing
effects. The region was scanned two times in the declination and right
ascension direction, respectively, in order to reduce effects caused by the
scanning process. The spectra were calibrated with CLASS which is part of the
GILDAS software package, then the declination and right ascension scans were
combined with the plait algorithm in GREG which is another component of the
GILDAS package. The 12CO data has a beam size of 11′′, and the rms noise level
of the corrected $T_{mb}$ is around 0.23 K at 0.6 km s-1 spectral resolution,
the 13CO has the same beam size but the rms noise is 0.12 K at 0.6 km s-1
spectral resolution.

After reducing the 30m data separately, single dish 12CO and 13CO data were
converted to visibilities and then combined with the SMA data using MIRIAD
package task UVMODEL. With different uv-range selections, the synthesized beam
of the combined data varies from $2.1^{\prime\prime}\times 1.8^{\prime\prime}$
(PA $-84^{\circ}$) to $11.1^{\prime\prime}\times 10.7^{\prime\prime}$ (PA
$-74^{\circ}$).

###  2.3 VLT-SINFONI integral field spectroscopy observations

Near Infrared $H$\- and $K$-band observations were performed using the
Integral Field instrument SINFONI (Eisenhauer et al., 2003; Bonnet et al.,
2004) on UT4 (Yepun) of the VLT at Paranal, Chile. The observations of S255IR
were performed in service mode on February 8th, February 12th, February 13th
and March 29th 2007. The non-AO mode of SINFONI, used in combination with the
setting providing the widest field of view ($8^{\prime\prime}\times
8^{\prime\prime}$) with a spatial resolution of 0.25′′ per slitlet. The
typical seeing during the observations was 0.7′′ in $K$-band. The H+K grating
was used, covering these bands with a resolution of R=1500 in a single
exposure.

S255IR was observed with a detector integration time (DIT) of 30 seconds per
pointing. The observations were centered on the central near infrared source
NIRS 1 at coordinates: R.A. 06h12m53.85s Dec.
$+17^{\circ}59^{\prime}23.71^{\prime\prime}$ (J2000.0). We observed this area
with SINFONI using a raster pattern, covering every location in the cluster at
least twice, resulting in an effective integration time of 60 seconds per
location in the field. The offset in the east-west direction was 4′′ (i.e.
half the FOV of the detector) and the offset in the north-south direction was
6.75′′, results a total observation field of 70${}^{\prime\prime}\times
70^{\prime\prime}$. A sky frame was taken every 3 minutes using the same DIT
as the science observations. The sky positions were chosen based on existing
NIR imaging in order to avoid contamination. Immediately after every science
observation, a telluric standard star was observed, matching as close as
possible the airmass of the object.

The data were reduced using the SPRED (version 1.37) software developed by the
MPE SINFONI consortium (Schreiber et al., 2004; Abuter et al., 2006). The
procedure described by Davies (2007) was applied to remove the OH line
residuals. To calibrate the flux and remove the telluric absorption lines, the
extracted spectrum of one standard star in each OB (OB1: Hip064656, OB2:
Hip048128, OB3: Hip032108, OB4: Hip031899) is used. The flux calibration of
the spectra uses the 2MASS (Cohen et al., 2003) magnitude of the standard
stars (see Bik et al. (2010) for the details of the data reduction).

##  3 Results

###  3.1 Millimeter continuum emission

We averaged the apparently line-free parts of upper and lower sideband spectra
of each region (presented in Figure 2) and got the continuum images of the
three regions shown in Figure 3.

Assuming optical thin dust emission, we estimated the gas mass and column
density of the continuum sources following the equations outlined in
Hildebrand (1983) and Beuther et al. (2005a). We assumed a dust temperature of
40 K, grain size of 0.1 $\mu$m, grain mass density of 3 g cm3, gas-to-dust
ratio of 100 and a grain emissivity index of 2 (corresponding to
$\kappa\approx 0.3$ for comparison of Ossenkopf & Henning (1994)). The results
are shown in Table 1. With the same set-up, we also calculated the gas mass
and the column density of the SCUBA 850 $\mu$m (Di Francesco et al., 2008;
Klein et al., 2005) continuum peaks shown in Table 2. To get the impression of
how much flux we lost in the interferometer observations, we produced the SMA
continuum map of the each region with the same beam size as the SCUBA 850
$\mu$m map (14${}^{\prime\prime}\times 14^{\prime\prime}$, Di Francesco et al.
(2008)). We measured the continuum flux, calculated the gas mass and compare
the results with the ones in Table 2; the mass we obtained from the SMA
observations for S255IR, S255N and S255S is only 7.8$\%$, 8.7$\%$ and 2$\%$ of
the SCUBA 850 $\mu$m measurements, respectively. Since our interferometer
observations are not sensitive to spatial scales $>$ 24′′ (the shortest base
line =8.5 k$\lambda$), we filter out the smoothly distributed large-scale halo
and left only the compact cores. Also, the fact that we filtered out more flux
for the youngest region S255S indicates that at the earliest evolutionary
stages the gas is more smoothly distributed than during later stages where the
collapse produces more centrally condensed structures. Nevertheless, all
derived column densities are of the order or above the proposed threshold for
high-mass star formation of 1 g cm-2 (Krumholz & McKee, 2008).

Figure 1: SPITZER IRAC 4.5 ${\mu}$m image (Chavarría et al., 2008) overlaid
with SCUBA 850 $\mu$m contours (Di Francesco et al., 2008; Klein et al.,
2005). In the upper panel, the grey scale is the SPITZER IRAC 4.5 ${\mu}$m
image, the black contours are the NVSS 1.4 GHz emission and the green contours
are the SCUBA 850 $\mu$m continuum. In the bottom left panel, the grey scale
is the SPITZER MIPS 24 ${\mu}$m image, the contours are the SCUBA 850 ${\mu}$m
continuum, the three dashed circles mark the primary beam of our SMA
observations. In the bottom right panel, the grey scale is the SPITZER MIPS 70
$\mu$m image, the contours and the circles are the same as the ones in the
left panel. The contour levels of the NVSS start at 10$\sigma$ (1$\sigma=0.6$
mJy beam-1) with a step of 5$\sigma$. For the SCUBA 850 $\mu$m, the contour
levels start at 10$\sigma$ (1$\sigma=0.1$ Jy beam-1) with a step of
10$\sigma$. The SPITZER/IRAC post-bcd data processed with pipeline version
S18.7.0 and the MIPS post-bcd data processed with pipeline version S17.2.0
have been downloaded from the SPITZER archive to create these images.

Figure 2: Lower and Upper sideband spectral vector-averaged over all baselines with a resolution of 2 km s-1 per channel. In the Upper sideband image, the 12CO line in S255IR and S255N panels was not fully plotted, which goes to 11.6 Jy and 5.9 Jy respectively. Table 1: Millimeter continuum peaks properties. Source | R.A. | Dec. | $S_{peak}$ | Sint. | Mass | N${}_{H_{2}}$
---|---|---|---|---|---|---
| (J2000.0) | (J2000.0) | mJy/beam | mJy | [$M_{\odot}$] | cm-2
S255IR-SMA1 | 06:12:54.01 | +17:59:23.1 | 95 | 171 | 12 | 5.5 $\times$ 1024
S255IR-SMA2 | 06:12:53.77 | +17:59:26.1 | 52 | 150 | 10 | 3.0 $\times$ 1024
S255IR-SMA3 | 06:12:53.88 | +17:59:23.7 | 30 | 30 | 2 | 1.7 $\times$ 1024
S255N-SMA1 | 06:12:53.71 | +18:00:27.3 | 91 | 327 | 22 | 5.5 $\times$ 1024
S255N-SMA2 | 06:12:52.95 | +18:00:31.7 | 33 | 47 | 3 | 2.0 $\times$ 1024
S255N-SMA3 | 06:12:53.65 | +18:00:18.3 | 18 | 19 | 1 | 1.1 $\times$ 1024
S255S-SMA1 | 06:12:56.65 | +17:58:41.2 | 14 | 17 | 1 | 7.7 $\times$ 1023
S255S-SMA2 | 06:12:56.85 | +17:58:35.0 | 12 | 15 | 1 | 6.6 $\times$ 1023
Table 2: SCUBA sub-millimeter continuum data. Source | Speak | Sint | Mass | N${}_{H_{2}}$
---|---|---|---|---
| Jy/beam | Jy | [$M_{\odot}$] | cm-2
S255IR | 7.81 | 31 | 372 | 6.5 $\times$ 1023
S255N | 7.75 | 30 | 357 | 6.4 $\times$ 1023
S255S | 2.33 | 9 | 110 | 1.9 $\times$ 1023
Figure 3: The SMA 1.3 mm continuum map of S255IR, S255N and S255S. contour
levels start at 5$\sigma$ with a step of 5$\sigma$ in S255IR (1$\sigma=1.7$
mJy/beam) and S255N (1$\sigma=1.6$ mJy/beam) image and 3$\sigma$ in S255S
(1$\sigma=0.9$ mJy/beam) image. The dotted contours are the negative features
due to the missing flux. The synthesis beam is shown at the bottom left of
each figure. In the S255IR panel, the asterisk marks the position of NIRS 3
and the right one marks NIRS 1 (Tamura et al., 1991). The inset is a zoom in
the continuum image of the inner region of S255IR-SMA1 and the contour levels
start at 5$\sigma$ with a step of 10$\sigma$. The triangles are the water
masers (Goddi et al., 2007) and the two open diamond mark the positions of two
6.7 GHz methanol masers detected by Xu et al. (2009). In S255N panel, the
crosses mark the position of the Class I 44 GHz ($7_{0}-6_{1}$) A+ methanol
masers detected by Kurtz et al. (2004) and the triangle is the water maser
(Cyganowski et al., 2007). The (0,0) point in each panel from left to right is
R.A. 06h12m54.019s Dec. $17^{\circ}59^{\prime}23.10^{\prime\prime}$ (J2000.0),
R.A. 06h12m53.669s Dec. $+18^{\circ}00^{\prime}26.90^{\prime\prime}$ (J2000.0)
and R.A. 06h12m56.58s Dec. $+17^{\circ}58^{\prime}32.80^{\prime\prime}$
(J2000.0), respectively.

##### S255IR

The continuum image of S255IR is presented in the left panel of Figure 3. We
resolved two main continuum peaks (i.e. S255IR-SMA1 and S255IR-SMA2) and one
unresolved sub peak (i.e. S255IR-SMA3) in this region. The stronger peak named
as S255IR-SMA1 in the southeast coincides with the near infrared source NIRS 3
(Tamura et al., 1991). S255IR-SMA1 coincides with a UCHII region (Snell &
Bally, 1986) generated by NIRS 3, and the inset in the S255IR panel in Figure
3 shows that it is also associated with Class II CH3OH and H2O maser emissions
(Goddi et al., 2007; Minier et al., 2000, 2005). An unresolved peak S255IR-
SMA3 that coincides with the near infrared source NIRS 1 (Tamura et al.,
1991), which lies 2.4′′ west to S255IR-SMA1 and has been identified as a
massive disk candidate by NIR polarization observation (Jiang et al., 2008).
These two continuum sources are both respectively coincident with the mid-
infrared massive proto-binary source 1 and 2 identified by Longmore et al.
(2006). S255IR-SMA2, which has never been detected before, is not associated
with any infrared source and likely to be in a very young source, it also
shows only a few lines, which will be discussed in detail in Sec. 3.2.

##### S255N

The middle panel of Figure 3 is the continuum image of S255N. The strongest
peak named as S255N-SMA1 which coincides with the UCHII region G192.584-0.041
(Kurtz et al., 1994) is associate with Class I 44 GHz ($7_{0}-6_{1}$) A+
methanol masers which are detected by Kurtz et al. (2004) and water maser
(Cyganowski et al., 2007). All the methanol masers which are marked with
crosses in the S255N panel in Figure 3 are distributed along the direction of
the outflow (Figure 10, see Sec. 3.3). S255N-SMA1 is also associated with
centimeter continuum emission (Cyganowski et al., 2007) which shows an
elongation aligned with the outflow (Figure 10, left panel). However, all
three continuum sources in S255N have no corresponding near infrared point
sources.

##### S255S

The right panel of Figure 3 presents the continuum image of S255S. There are
no near infrared sources or mid infrared peaks (SPITZER MIPS 24 $\mu$m and 70
$\mu$m, Figure 1) that are associated with these two sources. Meanwhile, the
gas mass is only 2$\%$ of the SCUBA 850 $\mu$m measurements, which indicates
that the gas in this region is at an early stage and smoothly distributed. All
these features indicate that S255S is an extremely young region.

###  3.2 Spectral line emission

The observed spectra of lower sideband (LSB) and upper sideband (USB) of three
region are shown in Figure 2.

Table 3: Observed lines in S255IR. $\nu$ | line | E${}_{\rm lower}/h$
---|---|---
$[$GHz$]$ |  |  $[$K$]$
LSB |  |
219.560 | C18O(2$-$1) | 5.3
219.734 | HNCO(102,9$-$92,8) | 219
219.798 | HNCO(100,10$-$90,9) | 48
219.949 | SO(65$-$54) | 24
220.038 | HCOOH(100,10$-$90,9) | 28.2
220.079 | CH3OH(80,8$-$71,6)E | 85
220.339 |  13CO(2$-$1) | 5.3
220.585 | HNCO(101,9$-$91,8) | 91
220.594 | CH3CN(126$-$116) | 315
220.641 | CH3CN(125$-$115) | 237
220.679 | CH3CN(124$-$114) | 173
220.709 | CH3CN(123$-$113) | 123
220.730 | CH3CN(122$-$112) | 87
220.743 | CH3CN(121$-$111) | 65
220.747 | CH3CN(120$-$110) | 58
USB |  |
229.590 | HCOOCH3(193,16$-$184,15)E | 106
229.759 | CH3OH(8-1,8$-$70,7)E | 77
229.858 |  34SO2(42,2$-$31,3) | 7.7
229.864 | CH3OH(195,15$-$204,16)A$+$ | 568
229.939 | CH3OH(195,14$-$204,17)A$-$ | 568
230.027 | CH3OH(32,2$-$41,4)E | 28
230.538 |  12CO(2$-$1) | 5.5
231.061 | OCS(19$-$18) | 100
231.221 |  13CS(50$-$40) | 22
231.281 | CH3OH(102,9$-$93,6)A$-$ | 98.8
Table 4: Observed lines in S255N. $\nu$ | line | E${}_{\rm lower}/h$
---|---|---
$[$GHz$]$ |  |  $[$K$]$
LSB |  |
219.560 | C18O(2$-$1) | 5.3
219.949 | SO(65$-$54) | 24
220.079 | CH3OH(80,8$-$71,6)E | 85
220.339 |  13CO(2$-$1) | 5.3
220.743 | CH3CN(121$-$111) | 65
220.747 | CH3CN(120$-$110) | 58
USB |  |
229.759 | CH3OH(8-1,8$-$70,7)E | 77
230.538 |  12CO(2$-$1) | 5.5
231.061 | OCS(19$-$18) | 100
231.221 |  13CS(50$-$40) | 22
Table 5: Observed lines in S255S. $\nu$ | line | E${}_{\rm lower}/h$
---|---|---
$[$GHz$]$ |  |  $[$K$]$
LSB |  |
219.560 | C18O(2$-$1) | 5.3
219.949 | SO(65$-$54) | 24
220.339 |  13CO(2$-$1) | 5.3
USB |  |
230.538 |  12CO(2$-$1) | 5.5

##### S255IR

In S255IR, we detect 25 lines from 10 species. Besides three normal CO
isotopologues, we also detect some sulfur-bearing species (SO, 34SO2, OCS,
13CS) and some dense gas molecules which are usually used to trace high mass
hot cores, such as CH3OH, CH3CN and HCOOCH3 (Nomura & Millar, 2004; Beuther et
al., 2009; Sutton et al., 1985). All the lines we detect are with lower energy
levels $E_{lower}/k$ between 5.3 to 568 K (Table 3).

Figures 4 and 5 present the integrated line images of all species (except
34SO2 which is blended with CH3OH and too weak for imaging, and 12CO and 13CO
will be discussed in details in 3.3) with different transitions. Most lines
show compact emission peaked at S255IR-SMA1, and only a few show emission at
S255IR-SMA2. Several CH3OH lines (CH3OH(8${}_{0,8}-7_{1,6}$)E,
CH3OH(8${}_{-1,8}-7_{0,7}$)E and CH3OH(3${}_{2,2}-4_{1,4}$)E, in Figure 4)
exhibit extended emission towards S255IR-SMA2. The
CH3OH(8${}_{-1,8}-7_{0,7}$)E line even extends to the northwest of S255IR-SMA1
following the direction of the blue-shifted outflow of S255IR-SMA1 (Figure 8,
see Sec. 3.3). We suggest that particular the (8${}_{-1,8}-7_{0,7}$)E line is
due to the shock heating excited by the outflow, which could also be mixed up
with some Class I methanol maser emission (Sutton et al., 2004; Sobolev et
al., 2005; Kalenskii et al., 2002; Slysh et al., 2002). SO line emission
associated with S255IR-SMA1 is elongated along the outflow and could be
affected by the outflow. 13CS line emission forms a shell around S255IR-SMA1
(Figure 5, similar shell-like C34S emission has also been found by Beuther et
al. (2009)). Only a few lines are associated with S255IR-SMA2, SO, C18O
(Figure 5) and several methanol lines (Figure 4). S255IR-SMA2 appears to be
chemically younger than S2555IRmm1. We extracted C18O spectra from these two
continuum peaks. The spectrum toward S255IR-SMA1 shows two peaks and can not
be fitted by a single Gaussian profile, therefore, we calculated the FWZI
(full width at zero intensity) of these two C18O spectra, which is 12 km s-1
and 6.8 km s-1 for S255IR-SMA1 and S255IR-SMA2, respectively. The difference
of the line width also suggests that S255IR-SMA2 is at a younger evolutionary
stage.

Figure 4: S255IR CH3OH line integrated intensity images with the SMA 1.3 mm
continuum emission in the background. Contour levels start at 3$\sigma$ with
2$\sigma/$level. For the top panels, the $\sigma$ of the contours in each
panel, from left to right, is 12, 15, 11 mJy beam-1, respectively, and for the
bottom panels is 12, 14, 15 mJy beam-1, respectively. The dotted contours are
the negative feature. The integral velocity ranges are shown at the bottom-
left of each panel in km s-1. The synthesized beams are shown at the bottom-
left of each panel. The (0,0) point in each panel is R.A. 06h12m54.019s Dec.
$+17^{\circ}59^{\prime}23.10^{\prime\prime}$ (J2000.0). Figure 5: S255IR
molecular line integrated intensity contour with the SMA 1.3 mm continuum
emission in the background. All the contour levels start at 3 $\sigma$. The
contour step is 3 $\sigma$ in SO (1$\sigma=22$ mJy beam-1), HCOOH
(1$\sigma=10$ mJy beam-1), HCOOCH3 (1$\sigma=12$ mJy beam-1), C18O
(1$\sigma=16$ mJy beam-1), CH3CN (1$\sigma=11$ mJy beam-1) and HNCO
(1$\sigma=13$ mJy beam-1) images, and 1 $\sigma$ in OCS (1$\sigma=13$ mJy
beam-1)and 13CS(1$\sigma=11$ mJy beam-1) image. The dotted contours are the
negative features with the same contours as the positive ones ini each panel.
The integral velocity ranges are shown at the bottom-left of each panel in km
s-1. The synthesized beams are shown at the bottom left of each panel. We only
plotted the $k=0\&1$ line for CH3CN and $(10_{0,10}-9_{0,9})$ line for HNCO.
The (0,0) point in each panel is R.A. 06h12m54.019s Dec.
$+17^{\circ}59^{\prime}23.10^{\prime\prime}$ (J2000.0).

##### S255N

In S255N, we detected 10 lines from 6 species (12CO, SO, CH3OH, CH3CN, OCS,
13CS) and 2 additional CO isotopologues (13CO and C18O) with lower energy
levels $E_{lower}/k$ between 5.3 to 65 K (Table 4). Figure 6 shows the
integrated line images of all species (except three CO isotopologues which
will be discussed in Sec. 3.3 and 3.4). All the lines show a peak associated
with S255N-SMA1, CH3OH and SO lines also show extended emission along the
direction of the outflows (Figure 10, see Sec. 3.3).

Figure 6: S255N molecular line integrated intensity contour with the SMA 1.3
mm continuum emission in the background. All the contour levels start at 3
$\sigma$. The contour step is 1$\sigma$ in CH3OH($8_{0,8}-7_{1,6}$)A-
(1$\sigma=9$ mJy beam-1) and CH3CN($k=0\&1$, 1$\sigma=14$ mJy beam-1) images,
2$\sigma$ in 13CS (1$\sigma=25$ mJy beam-1) image, 3 $\sigma$ in SO
(1$\sigma=13$ mJy beam-1) and CH3OH($8_{1,8}-7_{0,7}$)E images, and 0.5
$\sigma$ in OCS (1$\sigma=10$ mJy beam-1) image. The dotted contours are the
negative features. The integral velocity ranges are shown at the bottom-left
of each panel in km s-1. The synthesized beams are shown at the bottom left of
each panel. The (0,0) point in each panel is R.A. 06h12m53.669s Dec.
$+18^{\circ}00^{\prime}26.90^{\prime\prime}$ (J2000.0).

##### S255S

In S255S, we detected 4 lines from two species (12CO and SO) and 2 additional
CO isotopologues (13CO and C18O) with lower energy levels $E_{lower}/k$
between 5.3 to 24 K (Table 5), which again indicates that S255S is an
extremely young and cold region. We measure the single dish 13CO line width
toward the S255S SCUBA 850 $\mu$m continuum peak, and calculate the virial
mass in this region, which is $\sim$ 80 $M_{\odot}$. The virial mass is
smaller than the mass we got from the SCUBA 850 $\mu$m measurement which
implies that S255S region will likely collapse.

From the single dish 13CO observations, we got the vlsr of the three regions,
which are shown in Table 6. However, the vlsr of the high-mass mm cores in
S255IR and S255N based on the interferometer observations is different from
the one we got from the single dish data. We extracted dense gas spectra, i.e.
CH3OH for S255N-SMA1, CH3CN for S255IR-SMA1 and C18O for S255IR-SMA2, and then
a Gaussian profile was fitted to get the vlsr of the cores. All the velocities
are shown in Table 6. For S255S we can not get a proper vlsr from the SMA
observations due to the missing flux problem, so only the vlsr we got from 30m
observation is listed.

Table 6: The vlsr measured with different telescopes. Sources | vlsr got from 30m | vlsr got from SMA
---|---|---
| km s-1 | km s-1
S255N-SMA1 | 8.6 | 10.4
S255IR-SMA1 | 7.7 | 5.2
S255IR-SMA2 | 7.7 | 9.3
S255S | 6.6 | …

###  3.3 The molecular outflows

After combining the single dish 30m data with the interferometer data, we
integrated the line-wing of 12CO and produced the outflow image of each
region. Figure 7 shows the single dish only outflow map. While the northeast-
southwest outflow which is associated with S255IR is clearly shown in Figure
7, the high velocity outflow is also found in S255N region, however, the
direction of the outflow could not be well defined in single dish outflow map.
In the youngest region S255S, the single dish outflow map does not show much
emission. We combine the SMA data and 30m data together to study the outflow
properties.

Figure 7: The single dish outflow map of the S255 complex. The blue-shifted
outflow is shown in dashed contours with a velocity-integration regime of
[$-$40,0] km s-1, and the red one is shown in full contours with a velocity-
integration regime of [16,56] km s-1. The contours start at 4$\sigma$ with a
step of 5 $\sigma$($1\sigma$=1.2 K km s-1). The three crosses mark the
position of the three SCUBA 850 $\mu$m continuum sources.

##### S255IR

Figure 8 shows the combined SMA and 30m outflow images of S255IR. The
velocity-integration regime of the blue-shifted part is [$-$40,0] km s-1 and
[16,56] km s-1 for red-shifted part. We selected different uv-ranges resulting
in different resolutions. Although at different resolution the structure of
the outflow changes a lot, the northeast-southwest (NE-SW) outflow is
confirmed in all images (also reported by Miralles et al. (1997)). In the
bottom-left panel, the H2 jet-like sources (region (a) and (b) in Figure 19,
see Sec. 3.6.1) also follow the NE-SW outflow direction. The water masers
which are marked with triangles in the inset of S255IR panel in Figure 3 also
follow the direction of the outflow, which confirms that S255IR-SMA1 is the
driving source of this NE-SW outflow. In all panels except the bottom-right
two, the red-shifted part of the outflow bends a little toward northwest,
which might be a signature of the precessing jets. In the S255IR panel in
Figure 3, the water masers also follow the NE-SW direction. We believe that
NIRS 3 is the driving source of this outflow. As we are applying different uv-
range selection, two red-shifted outflow components reveal themselves out to
the north and south of the continuum sources, respectively, and these two
components are both only shown in the lower velocity regime which is offset
from the vlsr 10 to 20 km s-1. Finally, in the 30m-only panel and Figure 7,
the red-shifted outflow component to the north of the continuum sources
becomes connected with the component in S255N.

Figure 9 presents the position-velocity (pv) diagram of the molecular gas
along the outflow. The diagram shows that the red- and blue-shifted outflow
resembles the Hubble-law with increasing velocity at longer distance from the
outflow center, and the blue-shifted outflow also resembles the jet-bow-shock
gas entrainment model (Arce et al., 2007).

##### S255N

Figure 10 shows the combined SMA and 30m outflow images of S255N. The
velocity-integration regime of the blue-shifted part is [$-$44,0] km s-1 and
[16,48] km s-1 for the red-shifted part. The northeast-southwest (NE-SW)
outflow is shown clearly in the SMA and SMA combined with 30m panels but is
hardly seen in the 30m-only panel. In the middle panel, the crosses mark the
position of the Class I 44 GHz ($7_{0}-6_{1}$) A+ methanol masers detected by
Kurtz et al. (2004) and the triangle is the water maser (Cyganowski et al.,
2007). In the left panel, the back ground grey scale is VLA Q-band 3.6 cm
continuum (Cyganowski et al., 2007) which shows a jet-like structure. All
these features follow the NE-SW direction. Based on these features we confirm
the direction of the main outflow as NE-SW. The red-shifted outflow to the
north of S255N-SMA1 on the line (b) top left panel in Figure 10 might be part
of the outflow cavity, but we can not exclude the possibility of multiple
outflows. The red-shifted gas at the bottom part of line (b) seems to be
associated with S255N-SMA3 (the middle panel of Figure 10), however, another
blue-shifted feature at the southern part of the map shows up as the
resolution changes. These two components both can also only be detected at the
relatively lower velocity regime $\sim$10 to 20 km s-1 offset from the vlsr.

Figure 11 presents the position-velocity (pv) diagrams of the molecular gas
along the outflows. In the top panel, the pv-plot cut follows the direction of
the VLA jet-like emission (Fig. 10, left panel, line a). The diagram shows
that the high-velocity gas on the blue-shifted side remains very close to the
outflow center and the red-shifted part does not show the Hubble-law
signature. But in the bottom panel, in which the pv-plot cut follows the
direction of the two elongate red-shifted emission (Fig. 10, left panel, line
b), the red-shifted northern side of the outflow resembles the Hubble-law.
However, the red-shifted southern part of the outflow seems to have nothing to
do with our S255N-SMA1 nor S255N-SMA3. The different blue/red pv-
characteristics may be explained by the blue 12CO emission tracing
predominantly the jet like component also visible in centimeter continuum
emission, whereas the red 12CO emission traces mainly the cavity-like walls of
the outflow or another outflow.

##### S255S

Figure 12 presents the combined SMA and 30m outflow images of S255S. The
velocity-integration regime of the blue-shifted part is [$-9$,1] km s-1 and
[14,32] km s-1 for the red-shifted part. All the panels except the 30m-only
one, show two blue-shifted components and one red-shifted component. Although
the line wing emission is less pronounced than that for the other two regions,
we clearly identify blue- and red-shifted 12CO emission associated with both
mm peaks. The blue-shifted part in the south of the map is only detected in
the 1 km s-1$\geq$v$\geq-6$ km s-1, and the red-shifted part in the north of
the map is only detected in the v$\leq$16 km s-1. Since the line wing emission
does not extend to very high velocities, the outflow should not be oriented
directly along the line of sight. Therefore, the small spatial extent of the
outflow indicates as well the youth of this region.

Figure 8: S255IR 12CO$(2-1)$ outflow images observed with IRAM 30 m telescope
and the SMA. The blue-shifted outflow is shown in dashed contours with a
velocity-integration regime of [$-$40,0] km s-1, and the red one is shown in
full contours with a velocity-integration regime of [16,56] km s-1. The top-
left panel presents the SMA data only, the bottom-right presents the IRAM 30 m
data only. The others are all combined SMA$+$30 m data, which are inverted
with different uv-range selections (shown at the top-right of each panel in
k$\lambda$) resulting in different beam sizes shown at the bottom left of each
panel. All the 12CO emission contour levels start at 5 $\sigma$ with a step of
10 $\sigma$. The dotted contours are the negative features and the stars mark
the position of NIRS 3. In the bottom-left panel, the back ground scale is
SINFONI H2 line emission. The line in the SMA-only panel shows the pv-cut
presented in Fig. 9, and the cross marks the central position of the line. In
panel (0,120), the SMA 1.3 mm continuum map is over plotted (black full
contours), and the contours start at 5 $\sigma$ with a step of 5 $\sigma$. In
panel (0,70), the same CH3OH (8${}_{-1,8}-7_{0,7}$)E integrated intensity map
in Figure 4 is over plotted in black full contours . The $\sigma$ of the
outflow data is shown in Table 7. The (0,0) point in each panel is R.A.
06h12m54.019s Dec. $+17^{\circ}59^{\prime}23.10^{\prime\prime}$ (J2000.0).
Figure 9: Position-velocity diagrams of S255IR for the 12CO$(2-1)$ SMA-only
outflow observations with a velocity resolution of 1 km s-1. The pv-diagram is
in northeast-southwest direction with a PA of 75∘ from the north (the cut is
marked in Figure 8, top-left panel). The contour levels are from 10 to 90%
from the peak emission (4.2 Jy beam-1) with a step of 10%. The $v_{lsr}$ at
7.7 km s-1 and the central position (marked by the cross in Fig. 8 top-left
panel) is marked by vertical and horizontal line.  Figure 10: S255N
12CO$(2-1)$ outflow images observed with IRAM 30 m telescope and the SMA. The
blue-shifted outflow is shown in dashed contours with a velocity-integration
regime of [$-$44,0] km s-1, and the red one is shown in full contours with a
velocity-integration regime of [16,48] km s-1. The left panel presents the SMA
data only, the right presents the IRAM 30 m data only. The middle one is
combined SMA$+$30 m data, the uv-range selection is shown at the top-right of
each panel in k$\lambda$. All the CO emission contour levels start at 5
$\sigma$ with a step of 6 $\sigma$. The dotted contours are the negative
features and the asterisks mark the position of continuum peaks. In the middle
panel, the crosses mark the position of the Class I 44 GHz ($7_{0}-6_{1}$) A+
methanol masers detected by Kurtz et al. (2004) and the triangle is the water
maser (Cyganowski et al., 2007). In the top left panel, the back ground grey
scale is VLA Q-band continuum (Cyganowski et al., 2007) and the line shows the
pv-cut presented in Fig. 11, which follows the masers direction and crosses
the continuum peak. The arrow in the 30m-only panel marks the outflow size we
used to calculate the outflow physical parameters. The $\sigma$ of the outflow
data is shown in Table 7. The (0,0) point in each panel is R.A. 06h12m53.669s
Dec. $+18^{\circ}00^{\prime}26.90^{\prime\prime}$ (J2000.0).

Figure 11: Position-velocity diagrams of S255N for the 12CO$(2-1)$ SMA-only outflow observations with a velocity resolution of 1 km s-1. The top panel is the pv-plot in northeast-southwest direction with a PA of 54∘ from the north (the cut is marked as the line a in Figure 10, left panel), and the bottom panel is the pv-plot in northwest-southeast direction with a PA of 7∘ from the north (the cut is marked as the line b in Figure 10, left panel). The contour levels are from 10 to 90% from the peak emission (4.15 Jy beam-1 in the top panel and 2.82 Jy beam-1 in the bottom panel) with a step of 10%. The $v_{lsr}$ at 8.6 km s-1 and the central position (i.e. the S255N-SMA1 position in the top panel and the cross point of line a and line b in Figure 10 top-left panel in the bottom panel) are marked by horizontal and vertical line. Figure 12: S255S 12CO$(2-1)$ outflow images observed with IRAM 30 m telescope and the SMA. The blue-shifted outflow is shown in dashed contours with a velocity-integration regime of [$-$9,1] km s-1, and the red one is shown in full contours with a velocity-integration regime of [14,32] km s-1. The left panel presents the SMA data only, the right presents the IRAM 30 m data only. The middle one is combined SMA$+$30 m data, and the uv-range selection is shown at the top-right of each panel in k$\lambda$. All the 12CO$(2-1)$ emission contour levels start at 5 $\sigma$ with a step of 4 $\sigma$, except the one in IRAM 30 m only image which starts at 5 $\sigma$ with a step of 2 $\sigma$. The dotted contours are the negative features and the stars mark the positions of continuum peak S255S-SMA1 and S255S-SMA2. The arrows in the 30m-only panel mark the outflow sizes we used to calculate the outflow physical parameters. The $\sigma$ of the outflow data is shown in Table 7. The (0,0) point is each panel corresponds to position R.A. 06h12m56.58s Dec. $+17^{\circ}58^{\prime}32.80^{\prime\prime}$ (J2000.0). Table 7: Outflow data. Region | uv-range | beam size | $\sigma_{\rm red}$ | $\sigma_{\rm blue}$ | Data
---|---|---|---|---|---
| [k$\lambda$] | [′′] | [mJy beam-1] | [mJy beam-1] |
S255IR | (8, 173) | 1.5$\times$1.5 | 15 | 26 | SMA
| (0, 120) | 2.1$\times$1.8 | 15 | 26 | SMA+30m
| (0, 70) | 3.1$\times$2.5 | 18 | 32 | SMA+30m
| (0, 45) | 4.0$\times$3.4 | 22 | 38 | SMA+30m
| (0, 35) | 6.3$\times$5.8 | 26 | 71 | SMA+30m
| (0, 20) | 9.8$\times$7.4 | 38 | 86 | SMA+30m
| (0, 16) | 11.1$\times$10.7 | 40 | 110 | SMA+30m
| … | 11.3$\times$11.3 | 0.8 K km s-1 | 1.1 K km s-1 | 30m
S255N | (8, 173) | 1.8$\times$1.4 | 19 | 12 | SMA
| (0, 40) | 4.2$\times$3.9 | 34 | 18 | SMA+30m
| … | 11.3$\times$11.3 | 0.5 K km s-1 | 1.1 K km s-1 | 30m
S255S | (8, 173) | 1.8$\times$1.4 | 14 | 18 | SMA
| (0, 40) | 4.2$\times$3.9 | 20 | 26 | SMA+30m
| … | 11.3$\times$11.3 | 0.4 K km s-1 | 0.5 K km s-1 | 30m

Applying the method of Cabrit & Bertout (1990, 1992) we derived properties of
the outflows such as outflow mass, outflow energy and dynamical age. Since the
SMA data suffer a lot from the missing flux problem, only the 30m data are
used in the calculations. These calculations assume that the
13CO$(2-1)$/12CO$(2-1)$ line wing ratio is 0.1(Choi et al., 1993; Levreault,
1988; Beuther et al., 2002). The results are shown in Table 8. The dynamical
age calculation depends on the tangent of the outflow’s inclination angle
(with 0∘ being in the plane of the sky). Since we cannot get the inclination
angle information of these sources, we assumed an inclination angle of 45∘ for
all the dynamical age calculations. For S255N and S255S, the outflows are not
clearly defined in the 30m-only maps (Figures 10 and 12). Therefore, we used
with the SMA+30m map to determine the outflow sizes (see right panels of
Figures 10 and 12). For S255S, since we just detected only one outflow lobe
associated with each continuum source, we just extended the only one lobe to
the opposite side of the source to get the proper size and mass of the
outflow. The arrows in the 30m-only panels in Figures 10 and 12 mark the
outflow size we use.

Table 8: Outflow parameters.

Sources | $M_{\rm t}$ | $p$ | $E$ | Size | t | $\dot{M}_{\rm out}$ | $F_{\rm m}$
---|---|---|---|---|---|---|---
| [$M_{\odot}$] | [$M_{\odot}$ km s-1] | [erg] | [pc] | [yr] | [$M_{\odot}$/yr] | [$M_{\odot}$/km/s/yr]
S255IR-SMA1 | 2.9 | 137 | 6.5$\times 10^{46}$ | 0.7 | 7.4$\times 10^{3}$ | 4$\times 10^{-4}$ | 2$\times 10^{-2}$
S255N-SMA1 | 1.0 | 43 | 1.9$\times 10^{46}$ | 0.2 | 2.4$\times 10^{3}$ | 4$\times 10^{-4}$ | 2$\times 10^{-2}$
S255S-SMA1 | 0.02 | 0.4 | 8$\times 10^{43}$ | 0.1 | 2.5$\times 10^{3}$ | 8$\times 10^{-6}$ | 2$\times 10^{-4}$
S255S-SMA2 | 0.03 | 0.8 | 2$\times 10^{44}$ | 0.07 | 1.7$\times 10^{3}$ | 2$\times 10^{-5}$ | 5$\times 10^{-4}$

Entries include total outflow mass $M_{\rm t}$, momentum $p$, energy $E$,
size, outflow dynamical age $t$, outflow rate $\dot{M}_{\rm out}$, mechanical
force $F_{\rm m}$.

###  3.4 Rotational structures

##### S255IR

Figure 13 shows the velocity moment maps of HCOOCH3, CH3CN ($k=2$) and
C18O$(2-1)$. In the velocity maps of HCOOCH3 and CH3CN ($k=2$), we see a clear
velocity gradient perpendicular to the outflow axis, which indicates the
existence of a rotational structure. The C18O$(2-1)$ velocity map shows a
little different picture compared to the other ones. C18O traces a more
diffuse gas compared to the other two molecules. The C18O velocity map also
shows a big velocity difference between S255IR-SMA1 and S255IR-SMA2.

Figure 14 shows the position-velocity diagrams of the HCOOCH3 and C18O$(2-1)$
emission. The cuts, which are shown in Figure 13, go through the peak of the
dust continuum and have position angles perpendicular to the direction of the
outflow presented in Figure 8. The pv-diagram of HCOOCH3 shows that the
rotational structure is not in Keplerian motion, hence maybe it is just a
rotating and infalling core similar to the toroids described by Cesaroni et
al. (2007). The approximately 2.3′′ diameter of the rotational structure
corresponds to 3 700 AU at the given distance of 1.59 kpc. The pv-diagram of
CH3CN ($k=2$) has a similar structure to the one of HCOOCH3, thus we do not
show it.

However, the picture is different for C18O, the pv-diagram shows a Keplerian-
like rotation structure. The full line in the C18O panel shows a Keplerian
rotation curve with a central mass of 28 $M_{\odot}$, which is the mass of the
whole continuum structure in S255IR covering both mm peaks. For a more
detailed discussion, see Sec. 4.3.

Figure 13: HCOOCH3, CH3CN ($k=2$) and C18O$(2-1)$ velocity (1st) moment maps
overlaid with the SMA 1.3 mm dust continuum of S255IR. The contours start at
5$\sigma$ and increase with a step of 10$\sigma$ in all panels ($\sigma$=1.7
mJy beam-1). The lines in each panel show the pv-diagram cut presented in Fig.
14. All moment maps were clipped at the five sigma level of the respective
line channel map. The synthesized beam is shown in the lower left corner of
each plot. The (0,0) point in each panel is R.A. 06h12m54.019s Dec.
$+17^{\circ}59^{\prime}23.10^{\prime\prime}$ (J2000.0). Figure 14: Position-
velocity diagrams derived for the cuts along the observed velocity gradient in
Fig. 13. The offset refers to the distance along the cut from the dust
continuum peak. The contour levels are from 3$\sigma$ with a step of 1$\sigma$
in both panels (1$\sigma$=60 mJy beam-1). The full line in the C18O panel
shows a Keplerian rotation curve with a central mass of 28 $M_{\odot}$, which
is the mass of the whole continuum structure in S255IR. The negative features
are shown in dashed lines.

##### S255N

Figure 15 shows the velocity moment maps of CH3OH(8${}_{-1,8}-7_{0,7}$), CH3OH
($8_{0,8}-7_{1,6}$) and C18O$(2-1)$. The C18O velocity map shows a complicated
velocity structure, the velocity gradient is neither aligned with nor
orthogonal to the outflow orientation. Because C18O is an isotopologue of
12CO, it may be influenced by infall and the outflow. Methanol is well known
as a molecule that traces cores, shocks and masers in star formation regions
(Jørgensen et al., 2004; Beuther et al., 2005b; Sobolev et al., 2007). It has
also been reported as a low mass disk tracer (Goldsmith et al., 1999). In the
velocity maps of the two methanol transitions (left and middle panel of Figure
15), we see a velocity gradient perpendicular to the outflow axis, which
indicates the existence of a rotational structure.

Figure 16 shows the position-velocity diagram of the (8${}_{0,8}-7_{1,6}$)
methanol line emission, and the pv-diagram of the other methanol transition
shows the similar velocity structure. The cuts go through the peak of the dust
continuum and have position angles perpendicular to the direction of the
outflow (Figure 15). The pv-diagram shows that the rotational structure is not
Keplerian. The structure has a size of 4.4′′ corresponding to 7 000 AU at the
given distance of 1.59 kpc. It has a narrow velocity range of 3 km s-1, and
may likely be a large rotating and infalling core similar to the toroids
described by Cesaroni et al. (2007).

There are two low-velocity components at the northeast and southwest of the
continuum peak in the methanol velocity maps which coincide with the outflow.
They may be caused by the shock heating emission, and the line profile is
consist with both thermal emission and maser emission (Sutton et al., 2004;
Sobolev et al., 2005; Kalenskii et al., 2002; Slysh et al., 2002).

Figure 15: CH3OH(8${}_{-1,8}-7_{0,7}$), CH3OH ($8_{0,8}-7_{1,6}$) and
C18O$(2-1)$ velocity (1st) moment maps overlaid with the SMA 1.3 mm dust
continuum of S255N. The contours start at 5$\sigma$ and increase with a step
of 10$\sigma$ (1$\sigma$=1.6 mJy beam-1). All moment maps were clipped at the
five sigma level of the respective line’s channel map. The lines in each panel
show the pv-diagram cut presented in Fig. 14. The synthesized beam is shown in
the lower left corner of each plot. The (0,0) point in each panel is R.A.
06h12m53.669s Dec. $+18^{\circ}00^{\prime}26.90^{\prime\prime}$ (J2000.0).
Figure 16: Position-velocity diagrams derived for the cuts along the observed
velocity gradient in Fig. 15. The offset refers to the distance along the cut
from the dust continuum peak S255N-SMA1. The contour levels are all from
3$\sigma$ and with a step of 1$\sigma$ (1$\sigma$=40 mJy beam-1).

###  3.5 Temperature from CH3CN(12${}_{k}-11_{k}$) in S255IR

Since we detected 7 lines of the CH3CN$(12_{k}-11_{k})$ $k$-ladder with $k$ =
0…6 in S255IR, we can utilize the varying excitation levels of the lines with
lower level energies $E_{lower}/k$ between 58 to 315 K (Table 3) to estimate a
temperature for the central gas core. Figure 17 shows the observed
CH3CN$(12_{k}-11_{k})$ spectrum toward the continuum peak S255IR-SMA1 with
only the compact configuration data. We did a simple Gaussian fitting of the
spectrum and plot the level populations $N_{j,k}$ we calculated from the
fitting result in Figure 18 with the assumption of optically thin emission
(Zhang et al., 1998). The linear fitting result of the lower five levels is
also plotted in Figure 18. It is clear that we can not fit the whole spectrum
with one single temperature, which reveals a temperature gradient of the
source and not optical thin emission. We modeled this spectrum in the local
thermodynamic equilibrium using the XCLASS software developed by Peter Schilke
(private communication). This software package uses the line catalogs from JPL
and CDMS (Poynter & Pickett, 1985; Müller et al., 2001). The model spectrum
with a temperature of 150 K is shown in Figure 17 in dotted line. The main
difference between the model spectrum and the observed spectrum is that the
model one is optically thin whereas the lower line intensity of the observed
$k=3$ line indicates moderate optical depth of the CH3CN lines.

Figure 17: CH3CN(12k\- 11k) spectrum toward the mm continuum peak S255IR-SMA1.
The dotted line shows a model spectrum with $T_{\rm rot}$=150 K and $N_{\rm
CH_{3}CN}=3.5\times 10^{14}$ cm-2. The $k=6$ line is excluded in the model
because it is blurred by the HNCO(10${}_{1,9}-9_{1,8}$) line. Figure 18: From
the Gaussian fitting result, we calculated the $N_{j,k}$ (Zhang et al., 1998),
and plotted in the figure above. $E_{u}$ is the upper level energy of each
transition. The line shows the linear fitting of the first five points.

###  3.6 SINFONI results

####  3.6.1 SINFONI line and maps

The SINFONI observations are centered on the mm peak S255IR-SMA1 covering most
part of the near-infrared cluster. Figure 19 shows the 3 color composite of 3
line maps (Red: Br$\gamma$ line emission, green: H2 (2.12 $\mu$m) line
emission and blue: Fe II (1.64 $\mu$m) line emission). The two massive young
stars (NIRS 1 and NIRS 3) are in the center of the field, surrounded with the
cluster of about 120 stars down to $K=$17 mag. For all the point sources
marked in Figure 24 a SINFONI $H$\- and $K$-band spectrum is available and a
spectral classification is obtained of the brighter members (Sec. 3.6.2). The
PDR and jet like emission is traced by the molecular H2 emission (2.12 $\mu$m,
green). And to the north of the S255IR-SMA1 (the asterisk in the Figure 19),
two point sources (sources #17 and #18 in Figure 24) show strong Br$\gamma$,
which may indicate the existence of accreting signature (Muzerolle et al.,
1998). The two strong Fe II (blue) features around S255IR-SMA1 show
destructive shock (J-type, Hollenbach & McKee (1989)) emission and also follow
the direction of the outflow and the H2 jets, which indicate S255IR-SMA1 is
the energetic driving source of the jets.

To identify the excitation mechanism of the H2 emission (the arc which is
marked in Figure 19 by the red contour and the jet like emission regions (a)
and (b)), we extracted the spectra of these three regions and construct the
excitation diagrams for these regions. Figure 20 shows the excitation diagrams
of them. In these diagrams, the measured column densities of lines are plotted
against the energy of the upper level (see Martín-Hernández et al. (2008) for
the detailed description of this diagram). The total column densities were
calculated using the description of Zhang et al. (1998). Different symbols
represent different vibrational levels (Fig. 20). The arc region (Fig. 19) has
the most lines detected as it covers the largest area. For the jet-like
emission region (a) and (b), the weaker lines seen in the spectrum of the arc
region are not detected and 3$\sigma$ upper limits are given instead. The
solid line in the diagrams is a single temperature fit to all the data points,
while the dashed line in the diagram of the arc region presents a linear fit
to only the $1-0$ S lines. For the arc region it is clear that the column
densities are not represented by a single temperature gas, the line fluxes of
the $2-1$ and $3-2$ vibrational levels are higher than expected from the $1-0$
line fluxes. A likely excitation mechanism of the H2 gas in the arc is
fluorescence by non-ionizing UV photons (Davis et al., 2003).

The excitation diagrams of regions (a) and (b) are well-represented by a
single temperature. However the $3-2$ lines are not detected in either region,
and only one $2-1$ line is detected in region (a) left most of them only upper
limits. We conclude that the excitation diagrams of both regions are
consistent with shock excited emission in an outflow emission. Besides, the
bottom left panel of Figure 8 shows region (a) and (b) follow the direction of
the outflow, also their elongated shapes that all suggest outflow origin. The
linear fits of the excitation diagrams also allow us to determine the column
density and temperature of the emitting gas (Table 9).

Table 9: Physical properties of the H2 gas. Region | Trot | N(H2)
---|---|---
| $[K]$ | [cm-2]
the arc | $2416\pm 64$ | $2.6\pm 0.2\times 10^{17}$
a | $1655\pm 62$ | $5.8\pm 1.7\times 10^{17}$
b | $1091\pm 91$ | $1.9\pm 1.2\times 10^{18}$
Figure 19: Three-color image created from SINFONI line maps. Red: Br$\gamma$
line emission, green: H2 (2.12 $\mu$m) line emission and blue: Fe II (1.64
$\mu$m) line emission. The cross marks the position of S255IR-SMA1. The
contour and the dashed lines mark the regions of which the excitation diagrams
are constructed to identify the nature of the emission (Fig. 20).

Figure 20: The excitation diagram of the whole arc in the south-east of the
Fig. 19 (the top panel) and two jet-like emissions (a) and (b). The solid line
in all diagrams is a single temperature fit to all the lines detected, while
for the dash-dotted line in the top panel a only the 1$-$0 S lines are
included in the fit. The excitation diagram of the arc suggests that UV
fluorescence is the excitation mechanism in this area. Regions (a) and (b) are
most likely outflows as their excitation diagrams suggest thermal excitation.

####  3.6.2 SINFONI stellar spectral type classification

Our SINFONI observations provide an $H$\- and $K$-band spectrum of each source
with a spectral resolution of R=1500. SOFI $J$\- and $K$-band photometry
results are taken from Bik (2004). Objects with $K<$14 have high enough S/N
spectra to obtain a reliable spectral type. This reduces the sample to 39
sources, however, we can only get the spectral type of 16 sources which are
listed in Table 11. For many point sources, we could not get a spectral type,
some of them have very flat spectra and basically no absorption lines and also
no infrared excess, they are likely to be gas clumps. Others have very red
spectra, which are likely dominated by dust emission from the surrounding
environment or circumstellar disks, therefore, the spectrum of the underlying
objects might not be visible, such as sources #17, #18 in Table 11 and our
infrared sources NIRS1 and NIRS3 which will be discussed later.

The SINFONI $H+K$-band spectra of the three brightest cluster members (sources
#1, #2 and #3) show Br$\gamma$ absorption in the $K$-band and Br$10-14$
absorption in the $H$-band (Figure 21). Weaker lines of He I are also visible
in $H$ (1.70 $\mu$m) and K (2.058 $\mu$m). We apply the $K$-band
classification scheme for B stars from Hanson et al. (1996) which links a
$K$-band spectral type to an optical spectral type. For the $H$-band, we used
the classification of Hanson et al. (1998) and Blum et al. (1997). Table 10
shows the measured Equivalent Widths (EW) of the relevant lines in the $H$\-
and $K$-band spectra of the three B/A stars. Star #1 has a strong Br$\gamma$
absorption but does not show He I (2.11 $\mu$m), therefore is classified as
B3V$-$B7V. Star #2 and #3 show much stronger Br$\gamma$ absorption and also do
not have He I (2.11 $\mu$m) are therefore classified as B8V$-$A3V. The EWs of
the $H$-band lines are in agreement with the $K$-band spectral type. Source #1
and source #2 are also associated with UCHII regions (Snell & Bally, 1986),
which suggests that they are high to intermedia mass stars, and this consists
with our spectral type classification results. The results are listed in Table
11.

Table 10: Equivalent width and $K$-band spectral types of the early type stars in S255IR. Star | Br11 (1.68 $\mu$m) | He I (1.70 $\mu$m) | He I (2.11$\mu$m) | Br$\gamma$ (2.166 $\mu$m) |  $K$-Spectral type | Optical Spectral type
---|---|---|---|---|---|---
1 |  $<$0.4 Å | 0.6$\pm 0.4$ Å |  $<$0.4 Å | 5.9$\pm 0.4$ Å | KB4$-$B7 | B3V$-$B7V
2 | 8.3$\pm 0.4$ Å |  $<$0.4 Å |  $<$0.4 Å | 11$\pm 0.4$ Å | KB8$-$A3 | B8V$-$A3V
3 | 8.9$\pm 0.4$ Å |  $<$0.4 Å |  $<$0.4 Å | 8.9$\pm 0.4$ Å | KB8$-$A3 | B8V$-$A3V

Besides the 3 B/A stars we found 13 stars showing absorption lines typical of
later spectral type (Figure 22). The most prominent lines we used are the CO
first overtone absorption bands between 2.29 and 2.45 $\mu$m, and absorption
lines of Mg I and other atomic species in the $H$-band as well as Ca I and Na
I in the $K$-band. To classify the late-type stars we use the reference
spectra of Cushing et al. (2005) and Rayner et al. (2009). The atomic lines,
such as Mg I and Na I, are used for the temperature determination, while the
CO lines are used to determine the luminosity class. See Bik et al. (2010) for
a detailed description.

Compared with the reference spectra, the CO (2.29 $\mu$m) absorption of our
SINFONI spectra is usually deeper than in dwarf reference spectra, but not as
deep as in the giant spectra. In a few cases where a luminosity class IV
reference spectrum was available, the spectrum provided a better match to the
observed spectrum. This suggests that our late type stars are low- and
intermediate PMS stars, and indeed, Luhman (1999) and Winston et al. (2009)
find that PMS spectra have a surface gravity intermediate between giant and
dwarf spectra. If the stars were giants, dust veiling could make the K-band CO
lines weaker. However, the H-band CO and OH lines are also much weaker, while
the atomic lines have the expected EWs. Therefore, this seems to be a surface
gravity effect, suggesting a PMS nature of our late type stars.

The ”double blind” procedure was applied during the classification of the late
type stars to check the accuracy of the classification. In this procedure, Y.
Wang and A. Bik did their own classification separately and independently
without knowing results of each other, and their results were compared
afterwards to get the differences which are the errors of the spectral type in
Table 11. Our classification results showed an error of 1 to 2 subclass in
spectral type.

The relation from Kenyon & Hartmann (1995) was used to convert the spectral
type into effective temperature. However, this relation applies only for main
sequence stars, for PMS stars a different relation may hold (Hillenbrand,
1997; Winston et al., 2009). Cohen & Kuhi (1979) show that the temperature of
PMS stars might be overestimated by values between 500 K (G stars) and 200 K
(mid-K). We took this source of error into account when calculating the errors
in the effective temperature.

With the knowledge of the spectral type, we can estimate the extinction from
the observed color. As the intrinsic $J-K$ color for dwarfs and giants can
differ as much as 0.4 for late K stars (Koornneef, 1983), for the PMS stars,
we used as the intrinsic $J-K$ color the mean of the color for dwarfs and
giants. The difference between the mean value and the dwarf and giant values
is used as the error in the intrinsic color. The derived extinctions are
listed in Table 11.

Table 11: Photometric and spectroscopic properties of the detected stars brighter than $K=14$ mag. Star | R.A. (J2000) | Dec. (J2000) | $K$ | $J-K$ | Teff | Sp. Type | Lum.class | AV
---|---|---|---|---|---|---|---|---
| (h m s) | (∘ ′ ′′) | (mag) | (mag) | (K) |  |  | (mag)
1 | 06:12:54.91 | 17:59:21.05 | 11.26$\pm$0.03 | 1.7$\pm$0.06 | 15620$\pm$2850 | KB4-B7(B3-B7) | V | 10.0$\pm$0.3
2 | 06:12:55.06 | 17:59:28.93 | 10.60$\pm$0.02 | 0.8$\pm$0.03 | 9800$\pm$1590 | KB8-A3(B8-A3) | V | 4.8$\pm$0.2
3 | 06:12:54.68 | 17:59:32.82 | 11.68$\pm$0.04 | 2.2$\pm$0.09 | 9800$\pm$1590 | KB8-A3(B8-A3) | V | 12.8$\pm$0.4
4 | 06:12:54.82 | 17:59:12.98 | 13.39$\pm$0.08 | 2.3$\pm$0.21 | 3650$\pm$271 | M1.5$\pm$1 | V | 8.2$\pm$1.0
5 | 06:12:53.60 | 17:59:28.18 | 13.40$\pm$0.08 | 2.9$\pm$0.26 | 4590$\pm$460 | K4$\pm$1 | V | 12.3$\pm$1.2
6 | 06:12:55.31 | 17:59:15.90 | 13.17$\pm$0.07 | 1.7$\pm$0.20 | 4730$\pm$446 | K3$\pm$1 | PMS | 5.4$\pm$1.1
7 | 06:12:54.59 | 17:59:17.23 | 12.57$\pm$0.05 | 2.2$\pm$0.19 | 4730$\pm$446 | K3$\pm$1 | PMS | 8.6$\pm$1.0
8 | 06:12:54.13 | 17:59:29.18 | 13.37$\pm$0.08 | 2.6$\pm$0.23 | 4900$\pm$492 | K2$\pm$1 | V | 11.1$\pm$1.1
9 | 06:12:56.38 | 17:59:32.75 | 13.05$\pm$0.07 | 4.0$\pm$0.39 | 4900$\pm$492 | K2$\pm$1 | PMS | 19.1$\pm$1.6
10 | 06:12:54.55 | 17:59:02.76 | 12.78$\pm$0.06 | 3.8$\pm$0.33 | 5080$\pm$513 | K1$\pm$1 | PMS | 18.2$\pm$1.4
11 | 06:12:55.10 | 17:59:21.34 | 13.14$\pm$0.07 | 2.3$\pm$0.24 | 5080$\pm$513 | K1$\pm$1 | PMS | 10.0$\pm$1.2
12 | 06:12:53.17 | 17:59:05.93 | 12.50$\pm$0.05 | 3.0$\pm$0.25 | 5385$\pm$630 | G8/K0$\pm$2 | PMS | 14.4$\pm$1.3
13 | 06:12:55.06 | 17:59:31.09 | 12.24$\pm$0.05 | 1.5$\pm$0.20 | 5630$\pm$470 | G7$\pm$1 | PMS | 6.0$\pm$1.2
14 | 06:12:54.90 | 17:59:40.96 | 12.33$\pm$0.05 | 4.7$\pm$0.42 | 5385$\pm$630 | G8/K0$\pm$2 | PMS | 24.0$\pm$1.7
15 | 06:12:54.45 | 17:59:37.10 | 13.97$\pm$0.11 | 2.4$\pm$0.29 | 4730$\pm$446 | K3$\pm$1 | Giant | 9.1$\pm$1.4
16 | 06:12:53.29 | 17:59:21.70 | 13.37$\pm$0.08 | 3.0$\pm$0.27 | 5250$\pm$519 | K0$\pm$1 | S. Giant | 13.9$\pm$1.2
17 | 06:12:54.18 | 17:59:47.58 | 12.65$\pm$0.00 | 4.9$\pm$0.00 | … | … | Br$\gamma$ | …
18 | 06:12:54.15 | 17:59:34.19 | 12.29$\pm$0.00 | 5.3$\pm$0.00 | … | … | Br$\gamma$ | …
NIRS 1 | 06:12:53.83 | 17:59:23.71 | 11.36$\pm$0.00 | 6.2$\pm$0.00 | … | … | … | …
NIRS 3 | 06:12:54.01 | 17:59:23.68 | 12.56$\pm$0.00 | 5.0$\pm$0.00 | … | … | … | …
Figure 21: Normalized SINFONI $H$-band (left panel) and $K$-band (right panel)
spectra of the three B/A stars detected in S255IR. Their spectra are
characterized by absorption lines of the hydrogen Bracket series and He I
(dotted vertical lines). The stars are numbered according to Table 11. Figure
22: Normalized SINFONI $H$-band (left panel) and $K$-band (right panel)
spectra of the 11 objects showing late-type stellar spectral type. And most of
them are PMS stars. The vertical lines show the location of some of the
absorption lines used for the classification of the stars. The stars are
numbered according to Table 11

####  3.6.3 YSOs

Source #17 and #18 have very red spectra and show strong Br$\gamma$ point
emission, but the spectra are featureless due to the strong dust emission,
which indicates that they are very young and with ongoing accretion activity
(Figure 23). Source #17 also shows strong $K$-band CO line emissions (Bik &
Thi, 2004), which indicates the existence of circumstellar disk. The two
infrared sources NIRS3 and NIRS1 which coincide with our mm sources S255IR-
SMA1 and S255IR-SMA3, respectively, also have extremely red spectra, only
several H2 lines can be seen (Figure 23). Therefor we cannot get the spectral
type of these sources.

Figure 23: SINFONI $H$-band (left panel) and $K$-band (right panel) spectra of
the 4 YSOs. The stars are numbered according to Table 11.

####  3.6.4 Cluster membership and the HR diagram

For the late type stars, as discussed in Sec. 3.6.2, most of them have a
surface gravity typical for PMS stars, so they are considered to be the
cluster members. Source #15 and #16 show a luminosity type Giant and Super
Giant, they are unlikely to be associated with this young cluster and since
they show quite high extinction, they are considered to be background stars.

To constrain the mass of the cluster members and the age of the whole cluster,
we construct a HRD to compare the observed parameters with PMS evolutionary
tracks. The derived extinctions allow us to de-redden the $K$-band magnitude,
convert to absolute magnitude with the derived temperature, and plot the
points in a $K$ vs. log(Teff) diagram (Fig. 25, top). We exclude the
background stars #15 and #16 to enable a comparison with the isochrones. The
over-plotted dashed line is the 2 Myr main sequence isochrone taken from
Lejeune & Schaerer (2001). The solid lines in the left panel represent the
evolutionary tracks taken from Da Rio et al. (2009) using the models of Siess
et al. (2000) for 8 different masses: 5, 4, 3, 2, 1.5, 1, 0.5, 0.4 M⊙.
Comparison of the location of the PMS stars with these evolutionary tracks
yields an approximate mass varying from $\sim$ 5 M⊙ for the brightest stars to
0.4 M⊙.

In the bottom panel of Figure 25, the over-plotted solid lines are the PMS
isochrones between 0.1 and 10 Myr. The location of the stars is consistent
with a range of age. Besides source #1, #14 and #4 which lie between the 0.5
and 1 Myr isochrones, most of the objects span the $1-3$ Myr isochrones, with
the more massive objects closer to the 1 Myr isochrone. For the less massive
objects, the spread in age is larger, most likely due to higher uncertainties
in the spectral type determination. The location of the PMS stars suggests an
age of $2\pm 1$ Myr.

Comparison of the location of the PMS stars in the HRD with those of the
Herbig AeBe stars (van den Ancker et al., 1998) shows that the late spectral
type PMS stars are younger than the Herbig AeBe stars and will evolve from
their present G- and K spectral type to late B, A or early F spectral type
when they become main sequence stars. The three relative early spectral type
sources (#1, #2 and #3) already evolved to late B early A spectral type. They
are in the transit phase between our late spectral type sources and the main
sequence.

Figure 24: The stars listed in Table 11 marked on Br$\gamma$ line$+$continuum
map overlaid with H2 line emission contours.

Figure 25: $Top:$ The extinction corrected $K$ vs. log(Teff) HRD. The $K-$band
magnitude has been corrected for the distance modulus. Over-plotted with a
dash line is the 2 Myr main sequence isochrone from Lejeune & Schaerer (2001),
and with the solid lines the pre-main-sequence evolutionary tracks from Da Rio
et al. (2009). $Bottom:$ The same data but over-plotted with solid lines
showing the isochrones from Da Rio et al. (2009).

##  4 Discussion

###  4.1 Different evolutionary stages and triggered star formation?

The SMA and IRAM 30m data together reveal three massive star formation regions
with different evolutionary stages. The SCUBA 850 $\mu$m image (Figure 1)
presents three continuum peaks in the whole S255 complex region, one in each
sub-region, which is S255N, S255IR, S255S, from north to south respectively.
With our high resolution SMA observations, $\sim$2500 AU at the given distance
of 1.59 kpc, we found that all SCUBA 850 $\mu$m sources fragment into several
cores.

Minier et al. (2007) suggests S255S to be at a very young stage without active
star formation, however, our observations show outflows associated with the mm
sources, Furthermore, the virial mass is smaller than the mass obtained from
the SCUBA 850 $\mu$m measurement, which implies that the S255S region will
likely collapse. Since the peak column density is also on the order of the
proposed threshold for high-mass star formation of 1 g cm-2 (Krumholz & McKee,
2008), S255S is at a very early stage of ongoing star formation and may form
massive stars. The single dish continuum properties of S255IR and S255N do not
show much difference, however, our interferometry and NIR observations show us
the different properties of S255IR and S255N. While the large scale NIR
emission shows the existence of the NIR cluster in S255IR and not so many NIR
point sources in S255N, which indicate S255IR is most likely in a more evolved
stage compared to S255N. For the individual sub-cores, more lines are detected
at S255IR-SMA1 than S255N-SMA1, which indicates a higher temperature and more
evolved nature of S255IR-SMA1 compared to S255N-SMA1. Regarding the kinematic
properties, S255N-SMA1 has similar outflow velocities but a much smaller size
of the outflow than S255IR-SMA1, which may also suggest a younger evolutionary
stage of S255N-SMA1. Among other SMA mm continuum cores, S255IR-SMA2,
S255N-SMA2 and S255N-SMA3, which do not have many line emissions are
considered to be at earlier evolutionary stages.

Figure 1 presents the whole star formation region. The young star formation
region S255 complex lies between the evolved H II regions, S255 and S257. From
the morphology of the dust structure and the H II regions, it appears that the
two H II regions pushed gas between them and formed the S255IR dust structure
and triggered the star formation, which has also been suggested by Bieging et
al. (2009); Minier et al. (2007). To inquire that, we studied the velocity map
of our 30m data, however, due to the high noise level at the edge of the map,
we could not find a significant velocity difference between the west and east
edge of the dust structure to prove the interaction between the H II regions
and the dust structure. Chavarría et al. (2008) estimated the dynamical age of
the two H II regions S255 and S257, which is $\sim 1.5\times 10^{6}$ yr,
similar to our NIR cluster in S255IR. So we can not prove the triggering star
formation. However, we witness outflow interaction between S255IR and S255N
(Sec. 3.3). Since S255N is younger compared to S255IR, S255N may be affected
by S255IR.

Further more, the age of the cluster around S255IR we obtained from the
SINFONI data is 2$\pm$1 Myr (Sec. 3.6.4). Chavarría et al. (2008) obtained an
age of the larger-scale cluster of 1 Myr which is consistent with our result.
However, the dynamical age of the high-mass protostar in this region, S255IR-
SMA1, is $\sim 10^{4}$ yr (Table 8). Without the knowledge of the outflow
inclination angle, we could systematically underestimate the age value by a
factor of 2 to 5 (Cabrit & Bertout, 1992). Therefore, S255IR-SMA1 should not
be dynamically older than 105 yr and S255IR-SMA2 is even younger than S255IR-
SMA1. The age difference between the NIR cluster and the massive protostellar
objects indicates that the most massive sources in the cluster form last.

Minier et al. (2007) suggests a collect-and-collapse and triggered star
formation scenario for S255IR, which is that the B-stars (from our
observations they are late B stars, i.e. star #1, #2, #3 in Figure 24) in
S255IR formed through a collect-and-collpase process, and triggered the
formation of NIRS 1 and NIRS 3. The NIR and mm data presented in this paper
show that the cluster detected in the infrared is likely older than the mm
sources detected with our SMA data. While the SMA mm sources lie in the center
of the NIR cluster, the NIR cluster members are more distributed around the mm
sources. Based on this information, we propose an outside-in star formation
scenario, which is that the central gas filament collapses under the pressure
of the two H II regions. The collapse of the clump may start outside-in under
the pressure of the two H II regions, the low mass cores need lower density to
form, and they formed first at the outside region of the clump and collapse to
stars. And then either the low- to intermediate- mass stars may enhance the
instability of the central high-mass cores and potentially trigger the high-
mass star formation in S255IR, or the massive cores could build up slightly
slower and start to collapse afterwards forming the most massive stars in this
region. Our outflow observations show a hint that the energetic outflow from
the YSOs in S255IR may again have affected the star formation in S255N, but
this needs further observations to confirm.

However, the different evolutionary stages of the various regions and even
within each region are quite clear. This suggests a sequential star formation.

###  4.2 Multiple outflows

We detected outflows in all three regions (Sec. 3.3). In the top panels of
Figure 8, the outflow emission which is nearby the SMA continuum sources does
not really follow the NE-SW direction. We suggest this complicated outflow
environment in the nearby region of the continuum sources is due to the
interaction between outflows from NIRS 1 and NIRS 3. The red-shifted lobe to
the south of the continuum source (bottom panels in Figure 8) is most likely
driven by NIRS 1, because the north-south bipolar reflection nebular which is
associated with NIRS 1 follows this direction to the south (Jiang et al.,
2008; Simpson et al., 2009), and the outflow driven by NIRS 3 should be blue-
shifted at this direction.

In Figure 10, the red-shifted gas at the bottom part of line (b) seems to be
associated with S255N-SMA3, however, another blue-shifted feature at the
southern part of the map shows up as the resolution changes. Figure 7 shows
that these components are mixed together with the outflows in S255IR, which
may indicate the interaction between these two regions.

In the youngest region S255S, the velocity of the outflows is much smaller
compared to the other two regions. Table 2 shows that the column density of
this region is still relatively low, S255S is in very young evolutionary stage
and likely just start collapsing. This is the likely reason we did not detect
very energetic outflows like for the other two regions in this region.

###  4.3 Disk candidates in S255IR

NIRS 1, which coincides with S255IR-SMA3, is reported to have a polarization
disk (Jiang et al., 2008), Simpson et al. (2009) also reported a disk,
however, with a slightly different interpretation. Because NIRS 1 is
relatively more evolved, we detected only one unresolved mm continuum source
associated with only the CO isotopologue lines (Sec. 3.2). NIRS 3, which
coincides with S255IR-SMA1, is considered to be a high-mass protostar based on
several signatures (e.g., UCHII region, maser emissions, hot core emissions
and energetic outflows, see Sec. 1 and 3). We detected a rotational structure
coinciding with this source perpendicular to the outflow (Figure 13). The
HCOOCH3 position velocity diagram shows that this rotational structure is not
Keplerian, so it could be a rotating toroid around NIRS 3. However, the C18O
pv-diagram shows a much larger Keplerian-like velocity structure perpendicular
to the outflow (Figure 14). If this source were at a much further distance or
observed with worse resolution, this structure could easily be identified as a
disk. However, with our resolution we resolved two mm source, and we know it
is not a disk. The C18O gas size is $\sim 2\times 10^{4}$ AU and the continuum
source has a size of 1.6$\times 10^{4}$ AU at the given distance of
1.59$\times 10^{3}$ pc, which is similar to the rotational structure described
in Fallscheer et al. (2009). This structure is also much larger than the jeans
length which is $\sim$6 000 AU for a temperature $\sim$20 K and a density
$\sim 1\times 10^{6}$. Our system also exceeds the criteria to maintain a
stable disk described by Kratter et al. (2010). Therefore, we suggest this
structure to be a rotating toroid which fragmented into several sources, and
form a multiple system. Similar structures of fragmenting pseudo-disks have
also been modeled by Krumholz et al. (2009) and Peters et al. (2010).

Following these, we propose a star formation scenario for this large
rotational structure, which is that the rotational elongated dust structure
formed first, and at the central region the massive source NIRS 1 started
forming. However, this large structure is much more massive compared to NIRS 1
and is unstable, then it fragmented into two sources, S255IR-SMA1 and S255IR-
SMA2.

The SINFONI source #17 shows strong Br$\gamma$ and CO emission in $K$-band
(Figure 23), which indicates the existence of a circumstellar disk. Source #17
might also drive the jet like emissions at the northern edge of Figure 19. It
is interesting to find disk signatures toward sources with that different
evolutionary stages within the same forming cluster.

###  4.4 Cores and clumps

Since the single dish has a much larger beam, which can cover all the
continuum sources, and 13CO may be optically thick, the observed velocity
traces the outer layer of the whole clump and shows the mean line-of-sight
velocity. The molecular lines we used to get the velocities for the SMA
observations are all dense gas and disk tracers (Cesaroni et al., 2007), so
what we obtained are the velocities of the high mass cores (Table 6). The
difference between the core velocities and the clump velocities of several km
s-1 (Table 6) are very different from the low-mass core cases (e.g. 0.46 km
s-1 in NGC 1336 (Walsh et al., 2007), 0.17 km s-1 in Oph A (Di Francesco et
al., 2004)). This difference confirms that massive star formation regions are
more turbulent than low-mass ones. It further implies stronger peculiar
motions of the protostars and cores within the clump/cluster gravitational
potential.

##  5 Summary

Combining multi-wavelength data from mm to NIR wavelength, we characterize the
different (proto) stellar populations within the S255 star formation complex.
S255S, S255N and S255IR show different dynamical and chemical properties, not
only at mm wavelength but also at infrared wavelengths, which indicates they
are in different evolutionary stages. With the SMA, IRAM 30m and VLT-SINFONI
observations, we found outflows in all three regions, high velocity collimated
ones in S255IR, high velocity more confined ones in S255N and lower velocity
confined ones in S255S. The multiple outflows we found in S255IR and S255N
suggests a potential interaction between these two regions. From the outflow
maps, we estimated the dynamical age of the outflow driving sources. Although
without the information of the inclination angle this dynamical age could be
underestimated by a factor of 2 to 5 (Cabrit & Bertout, 1992), our mm sources
should not be older than 105 yr.

We detected 25 molecular lines in S255IR, including 7 lines of the
CH3CN(12${}_{k}-11_{k}$) k-ladder with k = 0…6, confirming that the hot core
nature of S255IR-SMA1. Only 10 molecular lines are detected in S255N,
including 2 CH3CN(12${}_{k}-11_{k}$) k-ladder with k = 0, 1, which is
indicative of a younger age and colder temperature. Besides the 3 CO
isotopologue line emissions, only diffuse SO emission is detected in S255S.
This is consistent with different chemical ages.

High-density tracers like CH3CN and HCOOCH3 show rotational structures around
the most prominent high-mass protostar candidates in S255IR and S255N.
Furthermore in S255IR, the C18O presents an elongated rotational structure
with a Keplerian-like velocity gradient perpendicular to the outflow. With a
size of $\sim 2\times 10^{4}$ AU, this structure can not be a disk but may be
a rotational toroid which fragments into several sources.

Near infrared $H$\- and $K$-band integral field spectroscopy observations were
done for S255IR. We identified the excitation mechanism of the H2 emission. We
derived the spectral type of 16 stars, and 14 of them are considered to be the
cluster members. With the knowledge of the spectral type and the SOFI $J$\-
and $K$\- band photometry results (Bik, 2004) of the cluster members, we
constructed an HR diagram to estimate the age of the cluster, which is 2$\pm$1
Myr . This age is consistent with the result Chavarría et al. (2008) obtained.
The age difference between the low-mass cluster and the massive mm cores
indicates different stellar populations in the cluster. This also leads to a
question, do the massive stars in this cluster form last? Our data support the
idea that massive stars form last. We propose the triggered outside-in
collapse star formation scenario for the bigger picture and the fragmentation
scenario for S255IR

The whole picture of the S255 complex suggests triggered star formation,
however, but we can not give hard proof at present. However, the different
evolutionary stages between each region and different stellar populations in
S255IR are consistent with sequential star formation.

###### Acknowledgements.

The authors thank C. J. Cyganowski for providing the VLA 3.6 cm data of S255N,
F. Eisenhauer for providing the data reduction software SPRED, A. Modigliani
for help in the data reduction, R. Davies for providing his routines to remove
the OH line residuals and N. Da Rio for providing the isochrones and
evolutionary tracks, and M. Fang for the discussion. Y.W. acknowledges support
by Purple Mountain Observatory, CAS and Max-Plank-Institute for Astronomy. E.
P. is funded by the Spanish MICINN under the Consolider-Ingenio 2010 Program
grant CSD2006-00070: First Science with the GTC.

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