# The Spatial—Kinematic Structure of the Region of Massive Star Formation
S255N on Various Scales

P. M. Zemlyanukha  Institute of Applied Physics, Russian Academy of Sciences,
Nizhnyi Novgorod, Russia  I. I. Zinchenko  Institute of Applied Physics,
Russian Academy of Sciences, Nizhnyi Novgorod, Russia  Lobachevsky State
University, Nizhnyi Novgorod, Russia  S. V. Salii  B.N. Yeltsin Ural Federal
University, Ekaterinburg, Russia  O. L. Ryabukhina  Institute of Applied
Physics, Russian Academy of Sciences, Nizhnyi Novgorod, Russia  Lobachevsky
State University, Nizhnyi Novgorod, Russia  S.-Y. Liu  Institute of Astronomy
and Astrophysics, Academia Sinica, P.O. Box 23-141, Taipei 106, Taiwan

ASTRONOMY REPORTS Vol. 62 No. 5 2018

## Abstract

The results of a detailed analysis of SMA, VLA, and IRAM observations of the
region of massive star formation S255N in CO(2—1), N2H+ (3—2), NH3 (1,1), C18O
(2—1)and some other lines is presented. Combining interferometer and single-
dish data has enabled a more detailed investigation of the gas kinematics in
the moleclar core on various spatial scales. There are no signs of rotation or
isotropic compression on the scale of the region as whole. The largest
fragments of gas ($\approx$0.3 pc) are located near the boundary of the
regions of ionized hydrogen S255 and S257. Some smaller-scale fragments are
associated with protostellar clumps. The kinetic temperatures of these
fragments lie in the range 10— 80 K. A circumstellar torus with inner radius
Rin $\approx$ 8000 AU and outer radius Rout 12 000 AU has been detected around
the clump SMA1. The rotation profile indicates the existence of a central
object with mass $\approx$ 8.5/ sin 2 (i) M⊙ . SMA1 is resolved into two
clumps, SMA1—NE and SMA1—SE, whose temperatures are $\approx$150 K and
$\approx$25 K, respectively. To all appearances, the torus is involved in the
accretion of surrounding gas onto the two protostellar clumps.

##  1 Introduction

Studies of molecular clouds have attracted considerable attention. Much effort
has gone into both simulations and observations of these objects (see, e.g.,
(Mapelli, 2017; Sokolov et al., 2017)). However, there remain relatively few
well studied regions of high-mass star formation. They are encountered more
rarely than the regions of low-mass star formation, and are located at large
distances from the Sun, complicating observations of them. In relation to
detailed studies of processes occurring in their cores, the resolution of
single dishes is insufficient, while interferometers have limited sensitivity
to extended structures.

The object considered in the current study is part of a molecular cloud
(Richardson et al., 1985; Pirogov et al., 2003) located between the two zones
of ionized hydrogen S255 and S257 (Fig. 1). It is supposed that the entire
S254—S258 complex was formed by successive star formation(Bieging et al.,
2007): expansion ofthe HII zones led to the compression of gas in the
molecular cloud. This cloud contains three cores: S255IR, S255S, and S255N.
The first is located in a later stage of evolution (Wang et al., 2011; Ojha et
al., 2011): its radiation is appreciably brighter in the IR, and water-vapor
and Class II methanol maser emission is detected in this region (Zinchenko et
al., 2012; Kurtz et al., 2004). The core S255S is the youngest of the three,
as is testified to by its lower brightness in the IR and in the 1.3 mm lines,
as well as the compactness of the outflow in the core (Wang et al., 2011). A
smaller number of 1.3 mm molecular lines is detected [6], and there are no
signs of massive-star formation; we accordingly did not consider it in the
current study.

The molecular cloud with which S255N (Sh2-255 FIR1, G192.60-MM1) is associated
has been observed on several single-dish radio telescopes (OSO 20 m, IRAM 30
m, NRAO 12 m) (Zinchenko et al., 2009); the core of S255N has been observed
with the SMA and VLA interferometers (Zinchenko et al., 2012). We present here
combined data for the first time. The mass of the S255N core indicated by
single-dish observations is $\sim$300 M⊙, with n${}_{H_{2}}$
$\sim$2$\times$105 cm-3, Tk $\sim$40K, $\Delta$V $\sim$2km/s. The luminosity
of the core is of order 105 L⊙ (Minier et al., 2007). An ultracompact HII
region is observed in this region(Kurtz et al., 1994), as well as H2O and
Class I CH3OH masers. These characteristics of the core provide evidence for
the formation of massive stars in this region. A lack of coincidence in the
positions of the emission peaks in the CO(2—1), HCN(1—0), HNC(1—0), HCO+(1—0),
C18O (2—1), and C34S(5—4) molecular lines has also been demonstrated
(Zinchenko et al., 2009). The presence of hot, massive protostellar clumps in
the core was shown in (Zinchenko et al., 2012; Cyganowski et al., 2007), with
the velocity of half of the cores are different from the velocity of the
ambient gas. The mass of the brightest of these clumps is $\approx$16 M⊙
(Zinchenko et al., 2012) for a distance of 1.78 kpc(Burns et al., 2016). The
presence of two spectral components in the (1, 1) ammonia lines in directions
not coincident with a hot clump was detected in(Zinchenko et al., 2012).
Several high-velocity outflows are present in the core, which also indicates
rich gas kinematics. We attempted in the current study to estimate the motion
of gas on scales fom the size of the core to the sizes of clumps, based on the
collected data. We also used non-standard data-analysis methods that are
robust to a noise to analyze emission that is weak against the background of
the hot-clump emission. We used original methods to detect and identify
kinematic fragments of the core, which were traced in position–velocity
diagrams in a number of molecular lines.

##  2 Observational data and analysis

We considered data obtained on the SMA(Wang et al., 2011; Zinchenko et al.,
2012) and VLA(Zinchenko et al., 2012) interferometric arrays and the 30-m IRAM
telescope (Wang et al., 2011; Zinchenko et al., 2012). Table 1 presents a list
of the observed molecules, the frequencies of their transitions, and the
instruments used. All the data were reduced anew, and self-calibration was
appied to the data of (Zinchenko et al., 2012). The CO(2—1), C18O (2—1),
CH3CN(12—11), and SO(65—54) lines were observed twice in the compact and once
in the extended configuration of the SMA. These data were inverted into a
single image at each line. The width of the synthesized antenna beam was
$\approx$ 1.14′′ at 217 GHz and $\approx$ 3.4′′ for the remaining observations
(the compact configuration of the SMA). The width of the VLA beam was
$\approx$ 2.55′′ (23.7 GHz), and the width of the IRAM beam was 12′′ (217
GHz). The coordinates of the phase center for the SMA observations were
$\alpha$(2000)=06h12m53.800s, $\delta$(2000)=17∘59′22.097′′.

###  2.1 Data reduction

Since one of the main aims of our study was to investigate the kinematics of
the protostellar core on various spatial scales, we considered objects
detected earlier in the 255N core with various extents from 2.5$\times$10-3 pc
to 0.48 pc [11]. The maximum angular sizes to which the instruments used at
the various frequencies are sensitive are 66′′ for the VLA and 25′′ for the
SMA, which corresponds to linear sizes of 0.52 and 0.2 pc at a distance of
1.78 kpc (Minier et al., 2007). These limits hinder the interpretation of
observations over a wide range of spatial scales. An important role is also
played by the loss of flux by the interferometer, which we compensated for by
combining the interferometric and single-dish data. As an example, the profile
of the CO(2—1) line indicated by the SMA observations differs strongly from
the profile observed on the IRAM 30-m telescope in the same direction.
Therefore, we consolidated the SMA data with the observations on the 30-m IRAM
antenna. This was carried out for the N2H+ (3—2), CO(2—1), and SiO(5—4) lines,
as well as the 1.3-mm continuum observations. The spectral-line data were
combined in the visibility domain based on the relative integrated-flux
calibration in the region of overlapping spatial frequencies; the
reconstructed continuum images were consolidated via a Fourier transform using
standard procedures of the MIRIAD package (Sault et al., 1995). The “wings” of
the CO(2—1) line at 3–7 km/s and 11–15 km/s from the single-dish data were
used for the relative flux calibration of the synthesized visibilities at the
frequency of this line.

CO(2—1), C18O(2—1), SO(65—54), DCO+(4—3), DCN(3—2) data were self-calibrated
using the observations in the CH3OH ($4_{2}$-$3_{1}$) line (Zinchenko et al.,
2012). The self-calibration based on this line was fairly effective, since a
single maser source was observed in the telescope field (Schwab, 1984).

###  2.2 Identification of kinematic fragments of the core

The presence of at last two spectral components with different Doppler
velocities in the 255N region was noted earlier (Zinchenko et al., 2012),
however their spatial distribution was not studied. In our current study, we
carried out a careful analysis of the data considered, with the aim of
conducting a detailed study of the spatial—kinematic structure of the core.

Application of the moment method for the identification of kinematic fragments
does not always provide the best solution, especially for lines of molecules
with hyperfine structure, such as NH3 and N2H+, since such estimates can be
shifted due to asymmetry of the line caused by its hyperfine structure. In
addition, in the case of observations of two or more spectral components, the
first moment yields an estimate of their mean velocity. Determining the
spatial boundaries of components precisely is problematic, since they can
interact, overlap, and be represented differently in different molecular
lines, depending on whether or not the conditions in the components are the
same. Therefore, we estimated the minimum number of boundaries from the number
of peaks in a histogram of the line intensities as a function of their
velocities (Fig. 2). Thus, instead of a spatial distribution, we obtained the
distribution of the line intensities in various velocity ranges. This
characterizes the distribution of the density, temperature, or some other
physical parameter, depending on what quantity is traced by the studied line.
We chose three molecular lines for this analysis: CO(2—1), N2H+ (3—2), and NH3
(1,1). We expected to observe signs of fragments in the outer, rarified layers
of the core in the CO line. The emission in NH3 and N2H+ is excited in
quiescent and dense regions of the core. It is possible to obtain more
accurate velocity estimates using these lines, since their transitions display
hyperfine structure. Their spectra are fitted fairly well with the local
thermodynamic equilibrium (LTE) model presented below. The shape of the
histogram (Fig. 2) provides information about the internal structure of the
core – whether it is homogeneous or is appreciably perturbed, and also whether
the gas can be described as forming isolated or overlapping velocity
components (that is, as kinematic fragments of the core), or whether the gas
is involved in some process with a single spectral distribution.

We estimated the positions of the line centers by fitting model line profiles
for the transition to the observed spectra, taking into account hyperfine
structure in an LTE approximation with a Gaussian optical-depth profile:

$T(\nu)=\sum_{n_{g}=1}^{m}(J(T_{ex}^{n_{g}})-J(T_{bg}))(1-e^{-\tau_{n_{g}}(v)}),$
(1)
$\tau_{n_{g}}(v)=\sum_{i=1}^{n_{hfs}}16\pi^{3}\nu_{ref}^{4}\mu^{2}\frac{Sg_{i}N_{u}^{n_{g}}}{3c^{3}J(T_{ex}^{n_{g}})}\phi_{n_{g}i}(v),$
(2)
$\phi_{n_{g}i}(v)=\frac{1}{\sqrt{2\pi}\sigma_{n_{g}}}e^{\frac{-(v-v_{i}-v_{n_{g}})^{2}}{2\sigma_{n_{g}}^{2}}},$
(3)

where $T_{ex}^{n_{g}}$ is the excitation temperature of spectral component
$n_{g}$, $N_{u}^{n_{g}}$ the column density of the molecules in the upper
level for the transition considered, $\sigma_{n_{g}}$ the line width,
$v_{n_{g}}$ the position of the line-component center (the subscript $n_{g}$
denotes the desired parameters of the model used in the fitting),
$J(T)=\frac{h\nu/k}{exp(h\nu/kT)-1}$, $T_{bg}=2.73K$ is the cosmic microwave
background temperature, $\nu_{ref}$ the transition frequency, $\mu$ the dipole
moment, $S$ the line strength, $g_{i}$ the statistical weight of transition i
of the hyperfine structure, $n_{hfs}$ the number of components of the
hyperfine structure (18 for NH3(1,1), 24 for NH3(2,2), 38 for N2H+(3—2))
$v_{i}$ the Doppler velocity of the transition $i$ in the line, and $m$ the
number of spectral components in the line.

The values of $S$,$g_{i}$ and$v_{i}$ were taken from the CDMS database (Müller
et al., 2005). The approximation was applicable for the NH3 (1,1) and (2,2)
and N2H+(3—2) lines. At a number of positions on the map, two peaks were
observed in the spectrum; accordingly, we allowed for the possibility of
fitting lines containing $m$ independent spectral components. In our data
sets, $m$ did not exceed two. The choice between one or two components was
based on a comparison of the rms deviations in the spectrum after subtracting
the model. If the ratio of the deviations for $m$=2 and $m$=1 was less than
1.4, we used the two-component model.

Thus, after fitting a line at each point in the map, we had estimates of the
Doppler shifts of one or two spectral components. However, the number of
kinematic components over the entire map can be appreciably greater than two.
The number of such components can be estimated from a histogram of the
velocity distribution of the line intensities. In the simplest case, the
integrated intensities of the spectral components whose Doppler velocities lay
within a given bin of the histogram were added. There also exist a large
number of methods enabling choice of the bin width based on the data set
considered, two of which we used in our analysis: the rule of Knuth (Knuth,
2006) and the Bayesian block method (Scargle et al., 2013), as realized in the
astropy library (Astropy Collaboration et al., 2013). Figure 2 presents
examples of applying these various methods.

Since a large fraction of a map contains lines with intensities at the
detection threshold (3—5$\sigma$), and the channel width is of order the line
width (5 in a group, 25 in a line, in the case of ammonia), we estimated the
velocity using the k nearest neighbors (kNN) method (Altman, 1992), since
least-squares fitting for these regions either diverged or converged to a
local minimum, with the suppression of one component by the other. Based on
the fitting at each point in the map, regions were chosen whose $\chi^{2}$
values for the LTE approximation were less than unity. The spectra at these
points were used as a reference set of objects relative to which the remaining
regions were considered. The velocities were estimated as

$v_{k}=\frac{\sum_{i=1}^{n}v_{i},d(s_{k},s_{i})^{2}}{\sum_{i=1}^{n}d(s_{k},s_{i})^{2}}$
(4)

where $v_{k}$ is the velocity in a spectrum of interest, $v_{i}$ the velocity
for the known (reference) spectrum, $s_{k}$, $s_{i}$ a spectrum of interest
and a spectrum whose estimated velocity is known, $n=5$ the number of nearby
spectra used to estimate the velocity,
$d(s_{1},s_{2})=\sqrt{\sum_{j=1}^{m}(s_{1_{j}}-s_{2_{j}})^{2}}$ the distance
in intensity space between the two spectra (analogous to the least-squares
residual), $m$ the number of channels, and $s_{i_{j}}$ the intensity of the
spectrum in channel $j$. In this form, the method resembles least squares
fitting, but the algorithm for the descent of the parameters to those yielding
the minimum residual is replaced by a search of n ”similar” spectra from the
reference set. Examples of fitting a spectrum using the least-squares method
and the first nearest neighbor method are presented in Fig. 3. The $\chi^{2}$
value for the LTE approximation is 266 (Jy/beam)2 , and for the first nearest
neighbor method is 72 (Jy/beam) 2. A more detailed analysis of the k nearest
neighbors method will be described in future publications.

###  2.3 Analysis of methanol observations

Estimates of the physical parameters (kinetic temperature and density of
molecular hydrogen, specific column density (column density of molecules per
unit velocity interval), and relative methanol abundance) were carried out
using the database of methanol energy-level populations (Salii, 2006). This
database contains methanol energy-level populations calculated at nodes of a
parameter grid. The parameter values were varied from those corresponding to a
dark molecular cloud to values that may be realized in regions heated by an
embedded object or excited by the passage of a shock. The kinetic temperature
was varied from 10 to 220 K, the density of molecular hydrogen from $10^{3}$
to $10^{9}$ cm-3, the abundance of methanol relative to hydrogen from
$10^{-9}$ to $10^{-6}$, and the specific methanol column density from $10^{8}$
to $10^{13}$ cm-3s. The computations of the methanol level populations in the
database were carried out in the LVG approximation. A model for the excitation
of the methanol molecule developed for the conditions characteristic for star-
forming regions, including regions of massive-star formation, was used in the
computations (Cragg et al., 2005).

##  3 Results

###  3.1 Spatial–kinematic structure of the core

To investigate the kinematics of the gas inside the core, we considered first
and foremost the emission in the N2H+ (3—2)and NH3 (1,1)lines, since these
trace dense, quiescent gas. In addition, both lines contain hyperfine
structure, making it possible to accurately measure Doppler shifts, and
thereby the velocity of the gas. We compared the N2H+ (3—2)data with the data
for the CO(2—1), (65—54), and C18O (2—1)lines in a position–velocity diagram
(see Fig. 9).

The S255N core is part of a more extended structure (Richardson et al., 1985),
which is also observed in our data. The CO(2—1) line emission has no boundary
at the north and south of the map, as can be seen in Fig. 4 (combined data).
The map of the first moment of this line on scales exceeding 20′′ corresponds
to the map presented in (Pirogov et al., 2003). However, a comparison of the
first and second moments of the CO line (Fig. 5) indicates that the line width
grows toward the center of the velocity range in a close to linear shape. The
direction of the velocity variations is close to the direction toward the
ionized regions S255 and S257 from the center of the map, with the mean gas
velocity observed at the map center. There is no clearly distinguished
symmetry axis, such as would be characteristic for the presence of rotation,
as was shown in (Mapelli, 2017). Such a distribution would not be
characteristic for an isotropic collapse. However, the broadening of lines
toward the center of the velocity range may indicate mixing of two portions of
gas driven by the expansion of envelopes.

As can be seen in Fig. 8, the kinematics of the core are very complex. The
presence of two spectral components of ammonia in gas that is not associated
with protostellar clumps was indicated in (Zinchenko et al., 2012), as we also
found in our data. Shifts in the component velocities from point to point are
observed. The overall direction of the velocity gradient lies from the
Southwest to the Northeast and East, from $\sim$6.7 km/s to $\sim$11 km/s.
However, the velocity does not vary uniformly. There are regions where there
is essentially no shift in the line. Figure 7 presents the velocity
distribution of the pixels for maps of the CO(2—1), N2H+(3—2) and NH3(1,1)
lines. The greatest coincidence of peaks in the distributions for the various
lines occurs at a velocity of $\sim$7.8 km/s. The gas with such velocities is
localized primarily in the southwestern part of the map. There is also a
modest peak at $\sim$7.1 km/s located at the southern edge of the map, which
can also be identified in the velocity slice marked 1 in Fig. 9. According to
the data of (Pirogov et al., 2003; Zinchenko et al., 2012), the velocity of
the S255N core is 8.9 km/s, which corresponds to peak 4 in Fig. 6. However,
its position in CO is shifted by $\sim$0.3 km/s relative to its position in
N2H+ and NH3 . Modest peaks are also present at 8.3 km/s and 9.6 km/s. We
denote each kinematic fragment using numbers from 1 to 5, beginning with the
fragment with the lowest velocity, according to the histogram in Fig. 6.

Thus, the gas in the core is distributed inhomogeneously. There exist some
regions within which there are no significant velocity shifts, which we will
refer to as kinematic fragments of the core. Figure 9 presents a frequency
slice through the data cube in the CO(2—1), N2H+(3—2), SO(65—54) and C18O(2—1)
lines along a path passing along the periphery of the core and crossing
through these fragments. This path is shown by the red line in Fig.7. The CO
spectrum slice shape overall repeats the spatial–frequency structure of the
N2H+ and NH3 lines, however, appreciably more compact structures are traced in
the SO and C18O lines. Fragments 1, 3, and 4 are correlated with both of these
lines, while fragment 2 is predominantly correlated with C18O. CO emission at
$\sim$7.9 km/s is observed along the entire path of the slices. We suggest
that this is associated with the diffuse gas surrounding the core observed in
(Richardson et al., 1985). In contrast to CO, the N2H+ and NH3 emission at
this velocity is localized only in the southeastern part of the core. This
region is depicted in Fig. 9 at the beginning and end of the slice path. Two
spectral components in both the N2H+ and NH3 lines are observed a distance
50′′ from the beginning of the path, as was also noted earlier in (Zinchenko
et al., 2012). The red component is not obviously represented in the CO line.
There is likewise no boundary between the fragments in the CO and N2H+ lines
with angular and spectral resolutions of 1.2′′ and 0.5 km/s, apart from the
fifth fragment. It is also noteworthy that, in spite of the different spatial
positions of the emission maxima in the N2H+ and NH3 lines, their
spatial–frequency distributions are otherwise similar to each other. The large
extent of the N2H+ emission can be explained by the addition of the single-
dish data to the interferometric data.

The velocities of the clumps SMA1, SMA2, and SMA3 presented in (Zinchenko et
al., 2012) are close to the gas velocities observed in the N2H+ (3—2)line,
which is not the case for SMA4, SMA5, and SMA6. In addition, a continuous,
elongated filament of gas in the C18O (2—1)line can be traced between clumps
SMA1–3; the velocity gradient along this filament is reflected in the
position–velocity diagram in Fig. 10. It is striking that the DCO+ (4—3) line
is shifted relative to the of the C18O (2—1)emission peak observed in the
clumps. The peak of the DCO+ emission, which is located at the edge of the
map, has coordinates (7′′, –34′′) relative to the phase center, and a full
width at half maximum of 0.7 km/s with its center at 8.8 km/s and an intensity
of $\approx$3 K. The DCO+, CO, and N2H+ spectra are presented in Fig. 11. The
CO line has red wing, suggesting a high-velocity outflow. An IR source with
similar coordinates is presented in (Miralles et al., 1997; Chavarria et al.,
2008), which in all likelihood is identified with an outflow and the DCO peak.
The spatial–velocity structure of the gas is more complex in the vicinity of
the clump SMA1 than in the rest of the core.

###  3.2 Central region

The region of intersection of kinematic fragments 2, 3, and 5 in the vicinity
of SMA1 is of the most interest. The peak of the 1.3-mm continuum emission
also lies in this direction. According to the data presented in (Zinchenko et
al., 2012), the Doppler velocity of the brightest clump, associated with a
peak in the dust emission of S255N—SMA1, is 8 km/s; however, the presence of
two spatially unresolved clumps inside SMA1 is proposed in (Cyganowski et al.,
2007). The existence of a bipolar outflow in the CO(2—1) and SiO(5—4) lines is
also known. Evidence for rotation of the clump SMA1 with a radius of
$\sim$1.5′′, corresponding to 2500 AU at a distance of 1.78 kpc, is presented
in (Wang et al., 2011), but the hypothess that two clumps are present was not
considered. We observed an appreciably more extended structure toward SMA1 in
the ammonia data (Fig. 12) than the structure described in (Wang et al., 2011;
Zinchenko et al., 2012), oriented perpendicular to the outflow, whose velocity
varies from 9 to 10.4 km/s, primarily along the structure. Figure 12 presents
a velocity map together with a position-velocity diagram for this clump in the
C18O (2—1)and DCO+ (4—3) lines with the velocities estimated from the ammonia
line indicated.

The position–velocity diagram in the (1, 1) ammonia line is presented in Fig.
7. Such a velocity profile is characteristic for a Keplerian torus (see, e.g.,
(Higuchi et al., 2015; Beltrán et al., 2014)) with its plane of rotation
facing us, with inner and outer radii R${}_{in}\approx$ 8000 AU and
R${}_{out}\approx$ 12000 AU, and a central mass of M${}_{c}~{}\approx$ 8.5
M${}_{\odot}/sin^{2}(i)$, where is the angle between the torus axis and the
direction toward the observer. The shape of the ammonia emission contours and
the linear radial dependence of the velocity in the inner part of the torus
testify that the inclination of the torus is close to 90∘. We constructed a
position–velocity diagram similarly to the method proposed in the appendix of
(Ohashi et al., 1997), but without taking into account the radial velocity. We
compared the shape of the observed and model contours in the position–velocity
diagram. The sizes of the outer and inner radii can also be traced in the
observed position–velocity diagram, as a linear and Keplerian part of the
diagram (Fig. 7). The inhomogeneous radial distribution of the ammonia
intensity, and also the presence of emission at distances from the clump
exceeding the outer radius of the torus, are noteworthy. Ammonia emission
shifted by 0.8 km/s is observed to the north (the positive direction in Fig.
13).

Figure 13a shows that the C18O (2—1)emission between the clumps SMA1 and SMA3
is continuous spatially and also in frequency. In addition, we found that SMA1
consists of two individual clumps, which we will refer to as SMA1-NE and
SMA1-SW. SMA1-NE is brighter in C18O (2—1)and has a Doppler velocity of 7
km/s, while SMA1-SW has a Doppler velocity of 9.8 km/s. According to the C18O
(2—1)data, their coordinates are $\alpha$(2000)=06h12m53.715${}^{s}\pm$0.005s
$\delta$(2000)=18∘00′27.73${}^{\prime\prime}\pm$0.07′′ abd
$\alpha$(2000)=06h12m53.588${}^{s}\pm$0.007s
$\delta$(2000)=18∘00′26.2${}^{\prime\prime}\pm$0.1′′, respectively.

###  3.3 Physical properties

####  3.3.1 Analysis of methanol observations

Emission in seven methanol lines was detected in S255N: three Class I maser
lines, one Class II maser line, and three quasi-thermal lines (see Table 2).
However, to all appearances, the emission of the Class II maser line is
thermal.

We detected a bright source with the coordinates 06h12m53.70s +18∘00′24.7′ and
with a velocity of 11.2 km/s in the $8_{-1}$–$7_{0}$ and $9_{-1}$–$8_{0}$
methanol transitions. The spectra of this source are shown in Fig. 14. The
intensity of the lines is $\approx$7 Jy/beam, appreciably higher than the
intensity in the direction of the outflow with coordinates 06h12m54.2s
+18∘00′14.9′. This large difference in flux densities suggests that this is a
maser source. Maser emission at the frequencies for these transitions has also
been reported, for example, in (Val’tts, 1998; Hunter et al., 2014; Yanagida
et al., 2014). The maser source is localized at the southeastern boundary of
the clump SMA1 and a bipolar outflow, in both space and frequency, as can be
seen in Fig. 9. We detected only one bright source, rather than the two
reported in (Kurtz et al., 2004). Class I emission is also observed for the
other maser sources reported in (Kurtz et al., 2004). However, the
insufficient sensitivity hinders drawing unambiguous conclusions about the
nature of the molecular excitation. The line is fairly narrow and shows signs
of maser emission in the directions toward points 5, 6, and 7, but the shape
of the spectrum in the direction of the blue part of the bipolar outflow
resembles the observed profiles for the SO(65—54) and SiO(5—4) lines.

The positions of the points where our estimates were carried out are shown in
Fig. 15. In most cases, the maximum methanol line emission is shifted relative
to the center of the clumps presented in (Zinchenko et al., 2012). Methanol
emission peak arranged along the line of the high-velocity outflow can also be
distinguished (Wang et al., 2011). The coordinates of these maxima are
presented in Table 3.

At all peaks line profiles differ from Gaussian. Two spectral peaks at
$\sim$10 km/s and $\sim$6 km/s are clearly visible at positions (1) and (2),
while there is no obvious division into spectral components at positions (3)
and (4). The spectra at positions (1) and (2) were fitted with two Gaussians,
and estimates were obtained separately for each of the components. Only one
component was fitted at positions (3) and (4), at velocities of 7 and 8 km/s,
respectively. Our estimates of the physical parameters of the gas are
presented in Table 4.

####  3.3.2 Other lines

The angular separation of the two emission peaks in the CH3CN (120—110) and
C18O (2—1)lines within SMA1, identified with NE and SW, is 1.5′′, which
corresponds to a linear size in the plane of the sky of $\sim$2700 AU at a
distance of 1.78 kpc(Burns et al., 2016).

Our re-reduction of the data for the (12—11) transitions (Wang et al., 2011;
Zinchenko et al., 2012) leads to a separation of SMA1–NE and SMA1–SW in space
and frequency, as can be seen in Fig. 16. Rotational diagrams (Martin
Hernandez et al., 2008) were used to estimate the kinetic temperatures of the
gas (Tk) for SMA1–NE and SMA1–SW, which were found to be 150$\pm$5K and
25${}^{+5K}_{-10K}$, respectively. We used the first five transitions of the K
ladder for the estimate for SMA1–NE, and the first three transitions for
SMA1–SW. The clump temperatures indicated by the analysis of the methanol data
range from 20K to 100K and from 45K to 100K, respectively, close to the
estimates based on the CH3CN data. The coordinates of the IR source indicated
in (Chavarria et al., 2008) are appreciably closer to SMA1–NE than to SMA1–SW.

We constructed a map of the distribution of Tk based on the (2, 2) and (1, 1)
ammonia lines (Fig. 17). Typical values lie in the range 10-60 K. Figure 13
shows that the ammonia emission is associated with the SW clump. Temperature
estimates for SMA1–SW based on the NH3 data are close to the estimates from
the CH3CN data, and there is no indication of hot gas associated with SMA1–NE
from the ammonia data. The DCN(3—2) and DCO+ (4—3) emission to the south of
the clumps also correlates with the region of low temperatures around SMA1—SW,
but there is no DCN emission to the north. The DCO+ emission toward the north
(Fig. 13) is associated with gas interacting with SMA1–NE, since it lies in
space and frequency inside the part of the region of C18O (2—1)emission where
there is no ammonia. The temperature of the gas in the direction of SMA2
indicated by the ammonia data grows to 54K. A cool region with a temperature
of about 10K with its peak at the coordinates 6h12m53s, 18∘01′01′′ is also
observed at the northern boundary of the map. It is striking that the
temperature grows to the East and West of the ammonia emission peak in
S255N–NH3, consistent with the ratio of the line intensities for these two
transitions at these coordinates. We used the velocity distribution of the
temperatures indicated by the ammonia data to estimate the temperature range
characteristic for each fragment, presented in Table 5. The estimates of the
kinetic temperatures based on the methanol and ammonia data are similar at
points 5, 8, and NH3 of Table 3.

The estimated mass of the filament joining the clumps SMA2, SMA1, and SMA3 is
$\sim$8M⊙. We used the data for the C18O (2—1)line to calculate this mass. The
column density of C18O was obtained using equation (15.28) of (Rohlfs and
Wilson, 2004), and the mass of gas was calculated from the abundance of this
isotope of CO, 1.7$\times$10-7 (Frerking et al., 1982).

##  4 Discussion

The core S255N has a fairly complex kinematic structure. To all appearances,
it is being compressed by expanding HII regions(Bieging et al., 2007). The
region of N2H+ (3—2)emission, which traces dense gas, is located closely
between the radiation-heated dust observed in the IR around the ionized
regions S255 and S257. The large-scale kinematic fragments are oriented
parallel to the boundary of the gas and dust envelopes, but there is no clear
indication of the nature of these fragments. According to (Chavarria et al.,
2008), 13 Class I and II protostellar IR sources were detected in the core.
Seven large protostellar clumps with masses of 1–16 M⊙ were also detected in
the core, which are partially identified with IR sources. The presence of so
many sources provides evidence of active formation of both low-mass and high-
mass stars in the core. The absence of velocity profiles characteristic of
compression suggests that the core has not yet made the transition to the
periphery-collapse phase.

Note that the different spatial scales traced by different molecules have a
mutal kinematic structure. However, several characteristic features are
observed. The outer, dispersed regions of the core traced by the CO (2—1) line
contain a component at a velocity of $\sim$8 km/s, which appears to be related
to unperturbed gas surrounding the core. In addition, three directions for
variations of the velocity can be distinguished from the center toward the
northeast, the southwest, and the southeast. These last two correspond to the
directions toward the ionized regions S255 and S257, suggesting that some
kinematic fragments may come about as a consequence of an external
interaction. Velocity slices for molecules tracing dense gas (NH3 (1,1), N2H+
(3—2), C18O (2—1)) resemble the section in the CO(2—1) line, apart from the
component at $\sim$8 km/s, and also the region around the clumps. However, not
all of the gas fragments can be formed by external factors. In particular, an
extended structure connecting the clumps SMA1, SMA2 and SMA3 and kinematically
related to these clumps can be seen in the C18O line. In addition, two
fragments are observed along the line of sight in the direction of the blue
wing of the bipolar outflow from SMA1 in the N2H+ and NH3 lines; one of these
is not associated with dispersed gas of the core, since it is not represent in
the CO(2—1) line.

The N2H+ NH3 C18O and SO data suggest that the matter in the core is
distributed fairly inhomogeneously, with kinematically distinct interacting
gas components being present. Clumps are not present in all of these, as can
be seen in the velocity in Fig.6.

The dense, compact gas observed in C18O (2—1)has an elongated structure. This
emission is appreciably correlated with the emission in the N2H+ (3—2)line,
but the portion that is not is located in the region of overlap of the
kinematic fragments of the core. Figure 13 shows that SMA1 is surrounded by
matter mainly to the north and south, as can be seen in the position–velocity
diagram. The line spectrum at the currently available resolution varies
smoothly, separating into two components toward SMA1 (NE and SW) – a sign of
interaction of the elongated structure with the clumps. It is also important
that the emission at 8-9 km/s is appreciably weaker at the center of SMA than
in the filament. This may indicate that a large fraction of the filament gas
in the vicinity of the clumps is interacting with this filament (Fig. 9).

The C18O, CH3CN, CH3OH and DCO+ data unambiguously indicate the presence of
two clumps, which were not resolved earlier and have been interpreted as the
object S255N SMA1. We have also continued to use the notation SMA1 when
referring to the most evolved, kinematically rich central region of the core.
It is noteworthy that the mean velocity of the detected giant torus coincides
with the velocity of SMA1–SW, although its size is much larger than the
distance between NE and SW (in the plane of the sky). Compact regions of DCO+
emission that are kinematically related to both clumps are observed. An
appreciably fraction of the DCO+ emission is located outside the clumps, but
inside the torus, correlating with C18O (Fig. 13). The NE clump is associated
with a bipolar outflow (Fig. 18) observed in the wings of CO(2—1), SO(65—54)
and SiO(5—4) lines, and also with a water maser (Cyganowski et al., 2007) and
methanol masers from (Kurtz et al., 2004) and detected by us. Regions of DCO+
emission are present in the gas interacting with NE. This effect is not
observed for SW. It is noteworthy that SW is located practically on the line
of the bipolar outflow from NE. This feature could quite plausibly result due
to the action of the outflow on the gas in the core. However, the relevance of
the torus in this picture remains unclear. Signs of the torus are also
observed in the C18O (2—1)data. Asymmetry of the CO velocity profiles relative
to NE can be traced in Fig. 18.

Analysis of the position–velocity diagram in the C18O, SO, NH3 and DCO+ lines
suggests that the torus traced in the ammonia lines is related to the
accretion of the more extended, dense material of the elongated, filament-like
structure onto the two young protostellar clumps. The position–velocity
diagram in the C18O line in the vicinity of SMA1 indicates an active transfer
of matter between the clumps and the ambient gas in the rotational plane of
the torus, with quasi-Keplerian radial profiles (Figs. 13 a,b,c).

The detected torus is one of several dozen known disk-like objects around
massive protostars (Beuther et al., 2017; Zinchenko et al., 2015; Beltrán and
de Wit, 2016; Quanz et al., 2010; Chini et al., 2006; Sako et al., 2005).
However, a large fraction of known disks have smaller sizes or are located
around less massive protostars. Tori that are larger than the disks and have
appreciable thicknesses are also known (Beltrán et al., 2014; Boley et al.,
2013; Sollins et al., 2005; Beltrán et al., 2011a, b; Higuchi et al., 2015).
The object G28.20-0.05 is closest in terms of its size and physical
parameters. The object we have discovered is large ($\sim$24 000 AU in
diameter, as opposed to $\leq$12 000 AU for G28.20-0.05). The two regions both
host an ultracompact HII region. The temperatures for the two objects
estimated from ammonia lines are similar(Sollins et al., 2005). However,
warmer, chemically rich torus-like rotating structures are also observed
(Beltrán et al., 2011a, b). The presence of a dust component in a similar
object is also known (Quanz et al., 2010). No signs of such a component have
been detected in the case of S255N SMA1 (Miralles et al., 1997; Wang et al.,
2011; Zinchenko et al., 2012).

Theoretical models predict the appearance of torus-like objects that are
gravitationally unstable (Cesaroni et al., 2007), together with an important
role for the late accretion of matter onto a protostar as a transitional
structure (Kratter et al., 2010). The detection of a torus is consistent with
the hypothesis of the formation of massive stars via disk accretion (Balbus
and Hawley, 1998; Krumholz et al., 2009). Observations of the region with
better angular resolution and simultaneously higher sensitivity are needed for
estimates of the accretion rates and detailed studies of the interaction of
the torus and the ambient gas.

##  5 Conclusion

We have considered the kinematic structure of the S255N core on various
spatial scales, from the outer layers of the core to dense clumps in the core.
Our analysis has led to the following discoveries.

  1. 1.

The core does not show characteristic signs of either isotropic collapse or
rotation. Overall, the large-scale motions are in agreement with those
described in (Pirogov et al., 2003), however, the structure of the core is
very inhomogeneous.

  2. 2.

At least five extended parts of the core are observed, within which the gas
velocities are similar.

  3. 3.

Large fragments of the core ($\sim$1/3 the size of the core itself) are
oriented parallel to the boundaries of the HII regions S255 and S257.

  4. 4.

A large fraction of smaller fragments is associated with the SMA1 clump. One
of these fragments is located to the southeast of SMA1 along the line of a
bipolar outflow, and is elongated along this outflow. Another fragment
connecting the clumps SMA3-2-1 is also elongated, and it is fairly massive
($\sim$ 8M⊙).

  5. 5.

We have detected a circumstellar torus rotating around SMA1, with inner radius
R${}_{in}\approx$8000 AU and outer radius R${}_{out}\approx$ 12000 AU. The
rotation profile is characteristic of Keplerian motion around a central mass
of $\sim$ 8.5/$sin^{2}(i)$ M⊙.

  6. 6.

SMA1 is spatially resolved into two separated clumps, NE and SW. They actively
interact with the torus and a filament. One of the clumps is fairly cool
($\sim$25 K), while the other is much hotter ($\sim$150 K). The bipolar
outflow is associated with the hot source.

  7. 7.

A Class I maser source was detected in the CH3OH $4_{2}$–$3_{1}$,
$8_{-1}$–$7_{0}$ and $9_{-1}$–$8_{0}$ lines, with coordinates
$\alpha$(2000)=06h12m53.69s $\delta$(2000)=+18∘00′25.0′ and radial velocity
$\approx$11.2 km/s. This source is located at the northeast boundary of the
bipolar outflow from SMA1–NE.

  8. 8.

The kinetic temperature and number density of the gas in the clumps and
outflows vary in the ranges 15–220 K and 104 – 105 cm-3 . The filling factor
for the sources in the outflows is low, suggesting strong fragmentation of the
gas.

  9. 9.

The temperature of the gas in the core is distributed inhomogeneously both in
terms of the parameters of different fragments and inside the fragments
themselves. Fragments that do not contain clumps have lower temperatures and
smaller temperature ranges.

  10. 10.

We detected a clump with the coordinates (7′′, -34′′ ) relative to SMA1,
together with a related high-velocity outflow.

##  6 Acknowledgments

We thank L.E. Pirogov and A.V. Lapinov for useful discussions of material and
methods. We also thank the referee for thoughtful comments that have helped
improve this paper.

This work was supported by the Russian Foundation for Basic Research (grants
16-32-00873, 16-02-00761, 15-02-06098), the Russian Science Foundation (grant
17-12-01256), decree No. 211 of the Government of the Russian Federation
(contract 02.A03.21.0006), and the Ministry of Education and Science of the
Russian Federation (grant RK AAAA-A17-117030310283-7).

<EMAIL_ADDRESS>
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Table 1: List of analyzed lines. The symbol ∗ denotes data obtained in the extended and compact configurations. SMA+IRAM and SMA∗+IRAM denote the combined data Molecule, transition | Frequency, GHz | Telescope
---|---|---
NH3 (1,1) | 23.69 | VLA
NH3 (2,2) | 23.72 | VLA
CH3OH-E ($5_{1}$—$4_{2}$) | 216.94 | SMA
SiO (5—4) | 217.10 | SMA+IRAM
CO (2—1) | 217.23 | SMA∗+IRAM
DCN (3—2) | 217.23 | SMA
CH3OH-E ($4_{2}$—$3_{1}$) | 218.44 | SMA
C18O (2—1) | 219.65 | SMA∗
CH3CN (12—11) | 220.59 | SMA*
CH3OH-E ($8_{0}$—$7_{1}$) | 220.07 | SMA
CH3OH-E ($8_{-1}$—$7_{0}$) | 229.75 | SMA
CH3OH-E ($3_{-2}$—$4_{-1}$) | 230.02 | SMA
CH3OH ($9_{-1}$—$8_{0}$) | 278.30 | SMA
N2H+ (3—2) | 279.51 | SMA+IRAM
DCO+ (4—3) | 288.14 | SMA
SO (65—54) | 288.51 | SMA∗
C34S (6—5) | 289.20 | SMA
CH3OH-E ($6_{0}$—$5_{0}$) | 289.93 | SMA
Table 2: Methanol lines detected at 216–290 GHz used to estimate physical conditions Transition | Frequency, GHz |  $E_{low}$,K |  $E_{up}$,K | Comments
---|---|---|---|---
$5_{1}$—$4_{2}$ | 216.945 | 37 | 47 | Class II maser
$4_{2}$—$3_{1}$ | 218.440 | 27 | 37 | Class I maser
$8_{0}$—$7_{1}$ | 220.078 | 78 | 88 |
$8_{-1}$—$7_{0}$ | 229.758 | 70 | 81 | Class I maser
$3_{-2}$—$4_{-1}$ | 230.027 | 20 | 31 |
$9_{-1}$—$8_{0}$ | 278.304 | 88 | 102 | Class I maser
$6_{0}$—$5_{0}$ | 289.939 | 40 | 53 |
Table 3: Coordinates of methanol emission peaks N | RA(2000) | DEC(2000) | Nearest object
---|---|---|---
1 | 6:12:53.42 | 18:00:23.65 | Near SMA1
2 | 6:12:53.70 | 18:00:26.90 | 2 and 3111Maser spot detected at 44 GHz (Kurtz et al., 2004)
3 | 6:12:53.86 | 18:00:30.15 | 5 and 6111Maser spot detected at 44 GHz (Kurtz et al., 2004)
4 | 6:12:54.49 | 18:00:43.65 | SMA4
5 | 6:12:53.98 | 18:00:14.40 | SMA5, 7111Maser spot detected at 44 GHz (Kurtz et al., 2004)
6 | 6:12:53.55 | 18:00:21.15 | SMA3
7 | 6:12:53.41 | 18:00:23.15 | SMA1
8 | 6:12:52.88 | 18:00:33.90 | SMA2
NH3 | 6:12:53.23 | 18:00:37.40 | NH3 (1,1) peak
Table 4: Physical parameters of gas from the fitting. 1$\sigma$ parameter range shown in parantheses. At position (4), the density range of 3–9 corresponds to the entire range in the database (model gives no constraints). $ff$ is the beam filling factor. Ntr is the number of model parameters. The column ”!” gives the number of lines with non-zero intensities at the indicated positions N |  VLSR(km/s) |  Tk,K |  lg(Nm/$\Delta$V) |  lg(n)[H2] |  lg(N${}_{CH_{3}OH}$/N${}_{H_{2}}$) | ff |  $\chi^{2}$/Ntr | !
---|---|---|---|---|---|---|---|---
1 | 9.8 | 45(40-100) | 10.50(10.00-11.00) | 6.0(¡7.5) | -6 | 12.45 | 0.90 | 7
1 | 5.6 | 50(¿25) | 10.00( 9.50-10.50) | 6.5(3.0-9.0) | -8 | 1.15 | 0.16 | 7
2111SMA1–NE | 7.6 | 30(¿20) | 11.25(10.00-12.00) | 3.5(¡7.0) | -6 | 11 | 1.9 | 6
2222SMA1–SE | 11.3 | 60(45-100) | 9.75( 9.50-10.00) | 6.0(5.75-6.75) | -6 | 77 | 4.4 | 5
3 | 7.1 | 40(30-160) | 11.25(10.5-12.5) | 6.0(¡7.5) | -6 | 4 | 0.9 | 7
4 | 7.8 | 120(¿45) | 9.50(9.25-10.0) | 6.5(3.0-9.0) | -6 | 69.7 | 4.4 | 4
6 | 9.2 | 85(50-200) | 9.0(8.75-9.25) | 6.0(¿5.5) | -6 | 67.2 | 1.32 | 4
7 | 9.7 | 60(40-120) | 9.75(9.5-10.0) | 6.0(5.5-7.0) | -6 | 65.02 | 5.09 | 6
5 | 7.5 | 165(¿50) | 8.00(¡8.5) | 6.0(5.5-7.5) | -6 | 99.0 | 2.24 | 2
6 | 4.9 | 50(10-220) | 9.25(¡9.0) | 6.0(3-9) | -6 | 55.8 | 0.50 | 5
7 | 5.4 | 55(40-120) | 9.50(9.25-9.75) | 6.0(5.5-7.0) | -6 | 65.41 | 2.22 | 6
8 | 4.7 | 20(¡30) | 11.25(8-13) | 3.5(3-9) | -7 | 9.35 | 1.05 | 3
NH3 | 9.5 | 15(¡25) | 11.25(11-12) | 5.0(3-9) | -8 | 54.8 | 1.06 | 3
NH3 | 7.6 | 15(¡25) | 11.50(11.25-11.75) | 6.5(¿5.5) | -8 | 77.9 | 2.60 | 4
Table 5: Parameters of kinematic fragment of the core. The range of kinetic temperatures in the fragments was obtained from the (1, 1) and (2, 2) transitions of ammonia Fragment number | Velocity, km/s |  Tk, K | Comments
---|---|---|---
1 |  $\sim$7.2 | 10-25 | SMA5
2 |  $\sim$7.8 | 10-40 | SMA1, SMA4
3 |  $\sim$8.3 | 6-30,80 | SMA3
4 |  $\sim$9.1 | 10-53 | SMA2, SMA6
5 |  $\sim$10.3 | 12-23 | Oriented toward the outflow from SMA1 SMA1
Figure 1: Spitzer IRAC three-color image of the region S255N and its
surroundings, composed of images at 8.0 $\mu$m (red), 4.5 $\mu$m (green), and
3.6 $\mu$m (blue). S255N is located in a complex of regions of current and
past star formation (Cyganowski et al., 2007). The color figures here and
below are accessible in the electronic version of the journal. Figure 2:
Example of the velocity distribution of the NH3 (1,1)intensities derived using
various algorithms for choosing the bin width. The blue columns show the
results for the Bayesian block method, the orange bins the results of applying
the Knuth rule, and the green columns results for a bin width of 0.2 km/s.
Figure 3: Observed N2H+ (3—2)spectrum at the point (-7′′ , 7′′) relative to
S225N–SMA1 (histogram), together with fits obtained using the least squares
method (red solid curve) and the kNN method (blue dashed curve). The vertical
lines with lengths proportional to their statistical weights denote hyperfine
components of the (3—2) transition. Figure 4: Map of the first moment of the
CO(2—1) line (background). Contours of the (1, 1) ammonia intensity at levels
of [4, 5.4, 6.8, 8, 9.5] K are superposed. The stars show the positions of
clumps according to the data of (Zinchenko et al., 2012). Figure 5: Dependence
of the second moment (width) on the first moment (velocity) of the CO (2—1)
and N2H+ (3—2)lines. Figure 6: Histograms of the velocity distributions in
various lines. The numbers and vertical dotted lines denote bands
corresponding to the positions of kinematic fragments of the core. Maps of the
first moment were used to construct the CO velocity histogram, and velocity
estimates obtained with the kNN method were used for the remaining lines. The
stars denote the velocities of clumps according to the data of (Zinchenko et
al., 2012). Figure 7: Position–velocity diagram in the NH3 line. A model of a
Keplerian torus with Rout=0.06 pc, Rin=0.04 pc, Mc=8.5 M⊙ is shown in the
background. The contours show the main component of the hyperfine structure of
the ammonia line at levels of [2.8, 3.2, 3.7, 4.3, 5.0] K, and the x’s the
velocities derived using the kNN method in the same line, with the
uncertainties indicated by the horizontal segments. Figure 8: Maps of the
Doppler velocity of the spectral components of the N2H+ (3—2)line derived from
the combined SMA and IRAM data. The left and right images show the velocities
of the blue and red components in the spectra. The superposed contours show
the NH3 intensity at levels of [4., 5.3, 6.7, 8.1, 9.5] K. The white stars
show the positions of clumps according to the data of (Zinchenko et al.,
2012). The red curve shows the path of the spectrogram section in Fig. 9,
beginning at the circle, marked every 10′′ Figure 9: Spectral slice along the
path shown in Fig. 8. The background shows the combined IRAM + SMA intensity
for the CO(2—1) line, the red contours the N2H+ (3—2)intensity at levels of
[6, 9, 12, 16, 23, 31, 43, 58] Jy/beam, the blue dashed contours the emission
in the SO (65—54)line at levels of [0.18, 0.2, 0.35, 0.5, 0.7, 1] Jy/beam, and
the filled contours in shades of blue–violet the intensity in the C18O
(2—1)line at levels of [0.17, 0.23, 0.3, 0.38, 0.49] Jy/beam. The vertical
dot–dashed lines and numbers indicate the positions of kinematic components
identified in Fig. 6. Figure 10: Spectral slice along the clumps SMA3-1-2. The
background shows the CO(2—1) intensity, the filled contours in shades of
blue–violet the C 18 O(2–1) emission at levels of [0.2, 0.3, 0.4, 0.6, 0.8]
Jy/beam, the dashed blue contours the SO (65—54)emission at levels of [0.1,
0.14, 0.18, 0.2, 0.28, 0.35, 0.44, 0.55] Jy/beam, and the red contours the
DCO+ (4—3) emission at levels of [0.37, 0.41, 0.45, 0.49, 0.53, 0.57, 0.62,
0.68] Jy/beam. Figure 11: N2H+ (3—2)(IRAM + SMA), DCO+ (4—3), and CO(2—1)
spectra in the direction of the peak DCO+ emission, (-7′′, 34′′) relative to
the phase center. Figure 12: Map of the Doppler velocity from the fit to the
NH3 (1,1)line (color scale). The contours show the integrated intensity in the
CO(2—1) wings, associated with the bipolar outflow, at levels of [26.5, 35.5]
(solid contours) and [-18.5, -8] (dashed contours). The sizes of the VLA
(outer ellipse) and SMA (inner ellipse) antenna beams are shown in the lower
left-hand corner. The star shows the position of the clump SMA1, and the
$\times$ and + symbols the positions of maser sources according to (Kurtz et
al., 2004) (Class I methanol) and (Cyganowski et al., 2007) (H2O). Figure 13:
Spatial–kinematic structure of the clump SMA1 and its vicinity. (a) map of the
C18O (2—1)intensity (background), with a map of the ammonia velocity shown in
shades of yellow–violet. The black contours shows the DCO+ (4—3) intensity at
levels of [0.4, 0.55, 0.7, 0.85, 1] Jy/beam, and the green lines show the
direction of the slice of the position–velocity diagram along and
perpendicular to the direction of variations in the torus velocity (panels (b)
and (c), respectively). The stars show the positions of clumps. (b) and (c)
Position–velocity diagram along and perpendicular to the torus. The background
shows the C18O emission, the solid black contours the DCO+ emission, and the
dashed orange contours the NH3 (1,1)intensity. (d) Map of the C18O intensity;
the solid red contours show the CO line in the red wing (pale contours are at
levels of [35, 40], bold contours at levels of [16, 26] km/s), and the dashed
blue contours the CO line in the blue wing (the large pale contours are at
levels of [-32.5, -37.5] and the smaller contours at levels of [-5, 2] km/s
Figure 14: Methanol spectra for the transitions $8_{-1}$–$7_{0}$,
$9_{-1}$–$8_{0}$ (lower, shifted by 8 Jy/beam), in the direction of the point
with coordinates 06h12m53.69s +18∘00′25.0′ (J2000) (the maser spot, thin
curves) and 06h12m54.2s +18∘00′14.9′ (J2000) (the bipolar outflow, bold
curves). Figure 15: Map of the first moment of the CO(2—1) line. The contours
show the C18O (2—1)intensity at levels of [0.4, 0.6, 0.8, 1, 1.2] K. The stars
show the positions of clumps according to (Zinchenko et al., 2012), and the
x’s with numbers the positions of the peak methanol emission presented in
Table 3. Figure 16: CH3CN (120-110) position-velocity diagram along the
direction SMA1–NE $\longrightarrow$ SMA1–SW (contours). The background shows
the diagram for the C18O (2—1)line. The contours are drawn for levels of
[0.07, 0.09, 0.12, 0.15, 0.2] Jy/beam. The synthesized beam of the SMA at
220.74 GHz and the width of the spectral channels are shown in the lower left
corner. Figure 17: Map of the distribution of the kinetic temperature from the
(1, 1) and (2, 2) transitions for ammonia. The shades of orange show the
temperature estimates, the background and gray scale contours the C18O
(2—1)intensity at levels of [0.2, 0.4, 0.6, 0.8, 1] Jy/beam, and the blue
contours the integrated DCN intensity ([-3 : 19] km/s). The contours
correspond to levels of [0.7, 1.5, 2.3, 3] Jy/beam m/s. Figure 18:
Position–velocity diagram along the bipolar outflow, denoted by the blue line
in Fig. 13(a). The background shows the combined data in the CO (2—1) line,
the red contours the interferometric data in this same line, the dashed blue
contours the combined data in the SiO (5—4) line, and the filled contours in
shades of blue–violet the data for the C18O (2—1)line. The filled yellow
circles show the positions of methanol masers according to (Kurtz et al.,
2004) and Section 3.3.1, and the filled yellow square the position of the
water maser presented in (Cyganowski et al., 2007).