THE DISK-OUTFLOW SYSTEM IN THE S255IR AREA OF HIGH-MASS STAR FORMATION

The formation of high mass stars (more massive than 8–10 M_⊙) is still poorly understood. Studies of this process are currently a "hot topic" of astrophysical research. A major unsolved problem of high-mass star formation is the characterization of accretion disks around young high-mass protostars. To date, most of evidence for the existence of such disks has been indirect (e.g., Beuther et al. 2009). If present, these disks undoubtedly play a crucial role in the star formation process. Nevertheless, the very existence of these disks, much less their physical properties, is not well-established. The importance of understanding these disks cannot be understated. For example, because of the larger masses involved, it is quite possible that high-mass disks are self-gravitating, unlike their Keplerian low-mass counterparts, implying that there may be significant dynamical differences between the high-mass and low-mass star formation processes.

S255IR is a part of the well known massive star-forming complex located between the evolved Sharpless H ii regions S255 and S257. Several authors have estimated the photometric distance to the complex at about 2.5 kpc (Russeil et al. 2007; Chavarría et al. 2008; Ojha et al. 2011). However, Rygl et al. (2010) report a distance of 1.6 kpc based on trigonometric parallax measurements of methanol masers. This seems to be the most accurate distance estimate and we adopt it here, too. The previous studies have shown that S255IR contains a cluster of early B-type stars (Howard et al. 1997; Itoh et al. 2001), several compact H ii regions (Snell & Bally 1986), and a number of H₂ emission features (Miralles et al. 1997). The total mass of the S255IR core has been estimated from single-dish observations at M ∼ 300–400 M_⊙ (Zinchenko et al. 2009; Wang et al. 2011). The estimate of the luminosity has been L ∼ 2 × 10⁴ L_⊙ (Wang et al. 2011).

The interferometric observations by Wang et al. (2011) revealed three continuum clumps and a high-velocity collimated outflow in the S255IR area. Our Paper I (Zinchenko et al. 2012) was devoted to the general structure and kinematics of the complex based on the Submillimeter Array (SMA) observations at 1.3 and 1.1 mm (the synthesized beam sizes were approximately 38 × 30 at 1.3 mm and 29 × 26 at 1.1 mm) as well as VLA ammonia observations and GMRT low frequency continuum data. Our results as well as data by Wang et al. (2011) indicate in particular a presence of a hot rotating core and a spectacular outflow in the S255IR region. However, the angular resolution was insufficient for a more detailed investigation of this core. In addition, a reliable evaluation of physical and chemical properties of this system required observations of additional molecular transitions. Therefore we performed observations of the S255IR area with the SMA at a much higher angular resolution and at higher frequencies.

Here we present observational data for S255IR obtained at 1.3 mm with the SMA in the very extended configuration and at 0.8 mm in the compact configuration. In addition we observed this area with the IRAM 30 m radio telescope in order to obtain short-spacing data complementing the SMA results. We mainly discuss properties of the dense cores and high velocity outflows observed in this area.

2.1. SMA

2.2. IRAM 30 m Radio Telescope

With the SMA we detected several tens of spectral lines in both 350 and 225 GHz bands. A list of these lines, including their frequencies and energy of lower levels, is given in Tables 1, 2. The spectral line parameters are taken from the JPL (Pickett et al. 1998) and CDMS (Müller et al. 2001, 2005) catalogs.

Table 1.  List of Molecular Transitions Observed at the SMA in the Very Extended Configuration in S255IR

Molecule	Transition	Frequency	El
		(GHz)	(K)
12CO	2–1	230.538000	5.532
13CO	2–1	220.398684	5.289
CH3OH	${6}_{1}-{7}_{2}$ A−	217.299202	363.496
	${15}_{6}-{16}_{5}$ A−	217.642677	735.160
	${15}_{6}-{16}_{5}$ A+	217.642678	735.160
	${20}_{1}-{20}_{0}$ E	217.886504	497.919
	${4}_{2}-{3}_{1}$ E	218.440063	34.976
	${8}_{0}-{7}_{1}$ E	220.078561	86.051
	${15}_{4}-{16}_{3}$ E	229.589056	363.420
	${8}_{-1}-{7}_{0}$ E	229.758756	78.076
	${19}_{5}-{20}_{4}$ A+	229.864121	567.565
	${19}_{5}-{20}_{4}$ A−	229.939095	567.561
	${3}_{-2}-{4}_{-1}$ E	230.027047	28.788
	${22}_{2}-{21}_{-3}$ E	230.292196	598.499
	${10}_{2}-{9}_{3}$ A−	231.281110	154.248
	${10}_{2}-{9}_{3}$ A+	232.418521	154.248
	${18}_{3}-{17}_{4}$ A+	232.783446	435.360
13CS	5–4	231.220686	22.194
DCN	3–2	217.238530	10.425
HNCO	${10}_{\mathrm{0,10}}-{9}_{\mathrm{0,9}}$	219.798282	47.471
	${10}_{\mathrm{1,9}}-{9}_{\mathrm{1,8}}$	220.584762	90.916
	${10}_{\mathrm{1,10}}-{9}_{\mathrm{1,9}}$	218.981170	90.569
	${10}_{\mathrm{2,9}}-{9}_{\mathrm{2,8}}$	219.733850	217.739
	${10}_{\mathrm{2,8}}-{9}_{\mathrm{2,7}}$	219.737193	217.739
	${10}_{\mathrm{3,8}}-{9}_{\mathrm{3,7}}$	219.656770	422.417
	${10}_{\mathrm{3,7}}-{9}_{\mathrm{3,6}}$	219.656771	422.417
HC3N	24–23	218.324711	120.504
SO	${6}_{5}-{5}_{4}$	219.949433	24.429
H2CO	${3}_{\mathrm{0,3}}-{2}_{\mathrm{0,2}}$	218.222195	10.483
	${3}_{\mathrm{2,2}}-{2}_{\mathrm{2,1}}$	218.475642	57.608
	${3}_{\mathrm{2,1}}-{2}_{\mathrm{2,0}}$	218.760071	57.613
OCS	18–17	218.903357	89.304
	19–18	231.060983	99.810
HCOOH	${10}_{\mathrm{0,10}}-{9}_{\mathrm{0,9}}$	220.038072	48.061
	${10}_{\mathrm{1,9}}-{9}_{\mathrm{1,8}}$	231.505705	53.355
CH3CN	${12}_{6}-{11}_{6}$	220.594423	315.313
	${12}_{5}-{11}_{5}$	220.641084	236.810
	${12}_{4}-{11}_{4}$	220.679287	172.555
	${12}_{3}-{11}_{3}$	220.709017	122.565
	${12}_{2}-{11}_{2}$	220.730261	86.849
	${12}_{1}-{11}_{1}$	220.743011	65.416
	${12}_{0}-{11}_{0}$	220.747261	58.272

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Table 2.  List of Molecular Transitions Observed at the SMA in the Compact Configuration in S255IR

Molecule	Transition	Frequency	El
		(GHz)	(K)
12CO	3–2	345.795990	16.596
CS	7–6	342.882850	49.372
HCO+	4–3	356.734242	25.682
CH3OH	${13}_{1}-{13}_{0}$ A ${}^{-+}$	342.729796	211.024
	${13}_{-1}-{14}_{-2}$ E	343.599019	607.550
	${18}_{2}-{17}_{3}$ E	344.109039	402.884
	${19}_{1}-{18}_{2}$ A+	344.443433	434.697
	${16}_{1}-{15}_{2}$ A−	345.903916	316.049
	${18}_{-3}-{17}_{-4}$ E	345.919260	442.829
	${13}_{0}-{12}_{1}$ A+	355.602945	193.959
HCN	4–3	354.505473	25.251
	4–3 (0,11c,0)	354.460435	1049.892
	4–3 (0,11d,0)	356.255568	1050.021
H13CN	4–3	345.339769	24.861
HC15N	4–3	344.200109	24.779
HC3N	38–37	345.609016	306.905
	39–38	354.697456	323.492
SO	${8}_{8}-{7}_{7}$	344.310792	70.957
33SO	${8}_{\mathrm{9,8}}-{7}_{\mathrm{8,7}}$	343.086102	61.564
	${8}_{\mathrm{9,9}}-{7}_{\mathrm{8,8}}$	343.087298	61.566
SO2	${13}_{\mathrm{2,12}}-{12}_{\mathrm{1,11}}$	345.338538	76.410
	${12}_{\mathrm{4,8}}-{12}_{\mathrm{3,9}}$	355.045517	93.960
	${13}_{\mathrm{4,10}}-{13}_{\mathrm{3,11}}$	357.165390	105.823
	${15}_{\mathrm{4,12}}-{15}_{\mathrm{3,13}}$	357.241193	132.537
	${11}_{\mathrm{4,8}}-{11}_{\mathrm{3,9}}$	357.387580	82.800
	${8}_{\mathrm{4,4}}-{8}_{\mathrm{3,5}}$	357.581449	55.202
	${9}_{\mathrm{4,6}}-{9}_{\mathrm{3,7}}$	357.671821	15.992
	${7}_{\mathrm{4,4}}-{7}_{\mathrm{3,5}}$	357.892442	47.835
	${6}_{\mathrm{4,2}}-{6}_{\mathrm{3,3}}$	357.925848	41.402
	${17}_{\mathrm{4,14}}-{17}_{\mathrm{3,15}}$	357.962905	162.932
H2CS	${10}_{\mathrm{0,10}}-{9}_{\mathrm{0,9}}$	342.946424	74.132
	${10}_{\mathrm{2,9}}-{9}_{\mathrm{2,8}}$	343.322082	126.830
	${10}_{\mathrm{3,8}}-{9}_{\mathrm{3,7}}$	343.409963	192.613
	${10}_{\mathrm{3,7}}-{9}_{\mathrm{3,6}}$	343.414146	192.556
HCOOH	${16}_{\mathrm{1,16}}-{15}_{\mathrm{1,15}}$	342.521225	127.153
	${15}_{\mathrm{1,14}}-{14}_{\mathrm{1,13}}$	343.952413	119.774
	${16}_{\mathrm{0,16}}-{15}_{\mathrm{0,15}}$	345.030596	126.494
	${16}_{\mathrm{2,15}}-{15}_{\mathrm{2,14}}$	356.137250	141.578

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We present the results in the form of maps as well as spectra and line parameters at selected positions. For continuum observations we give positions, flux densities, and size estimates of the continuum sources.

3.1. Data Analysis

3.2. Millimeter Wave Continuum

Maps of the continuum emission at 0.8 and 1.3 mm are presented in Figure 1. The 1.3 mm map was obtained by combining our new measurements in the very extended configuration with the previous measurements in the compact configuration (Paper I). It shows a rather extended emission and compact cores in the SMA1 and SMA2 clumps.

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Estimates of the continuum source parameters are given in Tables 3, 4. For the measurements at 1.3 mm we indicate the parameters of the compact cores seen in the very extended configuration. Parameters of the more extended components were presented in Paper I. The measured fluxes are several hundreds mJy at 350 GHz and several tens mJy at 225 GHz in the very extended configuration. The deconvolved sizes are from ∼1 arcsec to a few arcseconds at 345 GHz and about 0.3 arcsec for the compact cores at 225 GHz.

Table 3.  Names, Positions, Flux Densities, Deconvolved Angular Sizes, and Position Angles of the Millimeter Wave Continuum Sources Measured at 350 GHz in the Compact Configuration

Name	α(2000)	δ(2000)	${S}_{350}$	${\theta }_{\mathrm{max}}$	${\theta }_{\mathrm{min}}$	P.A.
	(h m s)	(${}^{\circ }$ ${}^{\prime }$ ${}^{\prime\prime }$)	(Jy)	(${}^{\prime\prime }$)	(${}^{\prime\prime }$)	(${}^{\circ }$)
S255IR-SMA1	6:12:54.00	17:59:23.2	0.50	1.4	0.6	−11
S255IR-SMA2	6:12:53.76	17:59:26.1	0.56	2.4	2.1	4
S255IR-SMA3	6:12:53.86	17:59:23.7	0.13	1.3	0.3	44
S255IR-SMA4	6:12:54.01	17:59:12.0	0.25	5.9	3.8	−82

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Table 4.  Names, Positions, Flux Densities, Deconvolved Angular Sizes, and Position Angles of the Millimeter Wave Continuum Sources Measured at 225 GHz in the Very Extended Configuration

Name	α(2000)	δ(2000)	${S}_{225}$	${\theta }_{\mathrm{max}}$	${\theta }_{\mathrm{min}}$	P.A.
	(h m s)	(${}^{\circ }$ ${}^{\prime }$ ${}^{\prime\prime }$)	(Jy)	(${}^{\prime\prime }$)	(${}^{\prime\prime }$)	(${}^{\circ }$)
S255IR-SMA1	6:12:54.010	17:59:23.06	0.058	0.30	0.27	−19
S255IR-SMA2	6:12:53.779	17:59:26.16	0.016	0.49	0.31	−27

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3.3. Basic Properties of the Molecular Emission

Here we give the general description of the observed molecular emission. More detailed information on the relevant species and transitions is presented in the following sections.

The general morphology and kinematics of the molecular emission were described in Paper I. Our new data set contains several tens of molecular transitions. In Figure 2 we present representative maps obtained with the SMA at 0.8 mm.

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These maps confirm that SMA1 is the brightest source of molecular emission in this area. High-excitation lines are observed exclusively toward this clump. Several lower excitation molecular lines are detected also in SMA2. Emission of HCN, SO, and CH₃OH is also observed in SMA4. HCN and CS emission is probably present in the area to the west of SMA2, designated as S255IR-N₂H⁺(1) in Paper I. Since it is observed not in N₂H⁺ only and for simplicity we shall designate it hereafter as SMA2-W.

A new feature, not noticed in our Paper I and in other previous works is a molecular clump to the east of SMA1 with a rather strong emission in the HCN, HCO⁺, CS and SO lines (Figure 2). We shall designate it as SMA1-E. This clump is located near the head of the jet observed in the NIR emission (see below).

Spectra of several representative transitions toward the SMA1 and SMA2 clumps are presented in Figures 3, 4.

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As mentioned in Paper I the main emission peak toward the SMA1 clump is at ${V}_{\mathrm{LSR}}$ ∼ 4–5 km s⁻¹. The line width is $\gtrsim 5$ km s⁻¹. Several lines including CO, HCN, HCO⁺, CS, and SO show broad wings indicative of high-velocity outflow. We discuss this feature in the following sections. A rather strong emission is detected from vibrationally excited HCN with an excitation energy ∼1000 K above the ground level.

The line emission from the SMA1-E clump peaks at about 8 km s⁻¹. The line width is large, ∼5 km s⁻¹.

The emission from the SMA2 clump is observed at about 10 km s⁻¹. The lines are narrow, ∼2 km s⁻¹.

With the very extended array we detected line emission almost exclusively from the SMA1 compact core. No emission was detected in C¹⁸O and SiO indicating an absence of compact structures in the lines of these molecules. The results of these observations are presented and discussed below.

Our data presented in Paper I and here as well as data by Wang et al. (2011) indicate the presence of four continuum clumps in the S255IR area, designated S255IR-SMA1 to S255IR-SMA4. In addition, in Paper I we detected clumps with a rather strong molecular emission without continuum counterparts in the SMA data. One of them, S255IR-N₂H⁺(1), is located close to SMA2 and was observed in the N₂H⁺, NH₃, and several CH₃OH lines. Estimates of their basic physical properties were presented in Paper I. Here we further investigate these objects using the new data set.

4.1. SMA1

SMA1 is the brightest object in this area. As shown in Section 3 most of the observed molecular lines are detected only here. The deconvolved size of the continuum source at 350 GHz measured with the SMA in the compact configuration (Table 3) is close to that found in Paper I. The flux density measured at 350 GHz is only slightly higher than the flux density obtained at 284 GHz (Paper I). It is worth noting that in Paper I we could not separate the SMA1 and SMA3 clumps. However, even if we take the integrated flux of these clumps at 350 GHz, the spectral index in the range 284–350 GHz will be only 1.6 which is inconsistent with a presumably optically thin dust emission at these frequencies. Most probably this implies a significantly larger flux loss at 350 GHz in comparison with the measurements at 284 GHz due to a smaller beam size.

The deconvolved size of the continuum source detected with the SMA in the very extended configuration is about 03 (Table 4) which corresponds to about 500 AU. The observed aspect ratio for this core is close to unity.

4.1.1. Kinematics

Wang et al. (2011) noticed rotation of the core around the axis roughly parallel to the outflow direction. Now, at sub-arcsecond resolution we can investigate the core kinematics on smaller scales. In Figure 5 we present maps of the first moment of the CH₃OH emission in the ${4}_{2}-{3}_{1}$ E line and CH₃CN emission in the ${12}_{3}-{11}_{3}$ line. In this figure we also indicate the axis of the jet previously identified through IR observations (Howard et al. 1997) and locations of the water masers measured with the VLBA (Goddi et al. 2007).

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This figure clearly shows that the core is really rotating around the axis of the jet. The rotation velocities along the line of sight amount to a few km s⁻¹ (it varies from about 3 km s⁻¹ to about 7 km s⁻¹ for CH₃OH ${4}_{2}-{3}_{1}$ E and to about 8 km s⁻¹ for CH₃CN ${12}_{3}-{11}_{3}$). In Figure 6 we present the position–velocity diagram for the CH₃OH emission in the ${4}_{2}-{3}_{1}$ E line along the cut through the core center perpendicular to the jet axis (P.A. = 157°). It clearly shows the velocity gradient which can be consistent with Keplerian rotation but the angular resolution is still insufficient for firm conclusions about the character of this rotation.

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An implicit indication of a probable further increase of the rotation velocity in the innermost part of the core comes from the line width map (Figure 7). It shows a significant increase of the line width in the center which can be related to rotation (the line width increases from about 2 km s⁻¹ at the periphery to about 6 km s⁻¹ in the center).

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It is also worth mentioning an increased line width along the jet axis (Figure 7). It can be probably explained by an increased turbulence caused by the passage of the jet.

The HCO⁺(4–3) and probably HCN(4–3) spectra toward SMA1 (Figure 3) show the redshifted absorption dip which may be suggestive of infall. The CO(3–2) spectrum measured with the SMA has a similar (although broader) feature (Figure 8). At the same time the CO(3–2) spectrum obtained from the combined SMA and IRAM 30 m data shows a peak at these velocities. Most probably it means that this feature is related to an extended component resolved out by the SMA, although an infall cannot be fully excluded.

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4.1.2. Physical Properties

Let us consider physical parameters of this core. We estimated gas kinetic temperature from the CH₃CN and CH₃OH observations. Modeling of the CH₃CN emission (Figure 9) as described in Section 3.1.2 yields a kinetic temperature in the range (74.9–197.6) K (1σ confidence interval) toward the emission peak (Table 5). We have detected a large number (15) of CH₃OH transitions toward SMA1 (Table 1). From the CH₃OH analysis (Section 3.1.1) we obtain ${T}_{\mathrm{kin}}=178$ K with the 165–195 K 1σ confidence interval at the center, ${T}_{\mathrm{kin}}=170-200$ K at 02 to the north, and ${T}_{\mathrm{kin}}=140-165$ K at 02 to the south (Table 6). Therefore there is probably a temperature gradient in the north–south direction. The temperature seems to be anti-correlated with the CH₃OH column density (Table 6, Figure 10).

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Table 5.  Kinetic Temperature, CH₃CN Column Density, Beam Filling Factor, and Line Optical Depth Derived from the CH₃CN Data Obtained in the Very Extended Configuration (HPBW $\approx 0\buildrel{\prime\prime}\over{.} 4$) at the Central Position in SMA1

Parameter	Fit results
Tkin (K)	128.8 (74.9–197.6)
${N}_{{\rm{L}}}({\mathrm{CH}}_{3}\mathrm{CN})$ (1016 cm−2)	2.12 (0.96–5.80)
Beam filling factor	0.23 (0.18–0.36)
$\tau (K=0)$	2.53 (0.93–8.37)

Note. The 1σ confidence intervals are indicated in parentheses.

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Table 6.  Physical Parameters Derived from the CH₃OH Data Obtained in the Very Extended Configuration (HPBW $\approx 0\buildrel{\prime\prime}\over{.} 4$) at Several Positions in SMA1

Offsets	Tkin	${N}_{{\mathrm{CH}}_{3}\mathrm{OH}}/{\rm{\Delta }}V$	${n}_{{{\rm{H}}}_{2}}$	Beam Filling Factor	${N}_{{\mathrm{CH}}_{3}\mathrm{OH}}/{N}_{{{\rm{H}}}_{2}}$
(${}^{\prime\prime }$, ${}^{\prime\prime }$)	(K)	(1012 cm−3s)	(108 cm−3)	(%)
0, +0.2	183 (170–200)	3.6 (2.8–5.0)	3.2	14.8 (13.9–16.4)	${10}^{-6}(\gt {10}^{-7})$
0, 0	178 (165–195)	5.6 (4.0–9.5)	3.2	16.0 (15.2–17.2)	${10}^{-6}(\gt {10}^{-7})$
0, −0.2	153 (140–165)	8.9 (5.0–12.6)	0.2	15.2 (13.3–16.0)	${10}^{-6}(\gt {10}^{-7})$

Note. The 1σ confidence intervals are indicated in the parentheses (the density ${n}_{{{\rm{H}}}_{2}}$ is not constrained by the model as evident from Figure 10).

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We adopt the gas kinetic temperature of 170 K as the average value (it is close to the weighted average of the CH₃OH and CH₃CN results). Assuming the same temperature (170 K) for the dust we obtain a total mass of this hot component of about 0.3 M_⊙. As in Paper I we assume a gas-to-dust mass ratio of 100 and adopt a dust absorption efficiency following Ossenkopf & Henning (1994). The peak gas column density estimated from the continuum data is $N({{\rm{H}}}_{2})\sim 3\times {10}^{24}$ cm⁻². With a size of 500 AU the mean density of hot gas is about 6 × 10⁸ cm⁻³. This is two times higher than our estimate of the SMA1 mean density in Paper I taking into account the difference in the adopted distances here and in Paper I. This increase can be expected since now we consider much smaller scales near the core center. The CH₃OH modeling puts no significant constraints on density. The relative abundance of CH₃OH is ∼10⁻⁶.

Both CH₃CN and CH₃OH observations indicate a beam filling factor of 0.15–0.2. This means that the source is very inhomogeneous on the 04 scale and probably consists of clumps with a size of ≪500 AU and a density much higher than the mean density found above. This higher density estimate does not contradict the CH₃OH data. A similar picture of clumpy medium was inferred from our observations of high-mass star-forming regions on larger scales (Pirogov & Zinchenko 2008; Pirogov et al. 2012). Most probably it reflects turbulence in the cores.

Our data set includes several other tracers of the hot gas. One of the most important is HNCO. We detected HNCO lines in different ${K}_{-1}$ ladders (${K}_{-1}=0,1,2,3$) with excitation energies up to ∼400 K (Table 1). In Figure 11 we present the rotational diagram for the observed HNCO transitions. This diagram was obtained from peak integrated line intensities found by 2D Gaussian fitting of the integrated intensity maps in different HNCO lines. We used the line strengths and dipole moment from the Cologne Database for Molecular Spectroscopy (Müller et al. 2001, 2005).

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This diagram indicates a rotational temperature of 318 ± 70 K (Table 7). However, the ${K}_{-1}=0,1$ transitions can be saturated. The peak brightness temperature in these transitions is close to the brightness temperature in the apparently optically thick lines (e.g., CO and CH₃CN) and our modeling using RADEX (van der Tak et al. 2007) shows that at the derived physical parameters and column density the optical depth in these lines is about unity. In this case the derived rotational temperature would apparently represent an upper limit to the excitation temperature. The deconvolved source size is ∼03 × 02 for the ${K}_{-1}=0$ transition and decreases for higher ${K}_{-1}$ ladders. The map of the integrated intensity in the HNCO ${K}_{-1}=2$ transition is shown in Figure 12 along with the map of the OCS emission which is another tracer of hot gas. The peaks of the HNCO and OCS emission practically coincide with the continuum peak, although distributions of these molecules seem to be somewhat different.

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Table 7.  Fit Results for the Rotational Diagrams Obtained toward SMA1

Molecule	Tkin	$\mathrm{log}{N}_{{\rm{L}}}$	HPBW
	(K)	(cm−2)	(${}^{\prime\prime }$)
HNCO	318(70)	16.34(0.04)	$\approx 0.4$
SO2	146(16)	15.50(0.05)	$\approx 2$
	65(11)	15.35(0.08)	$\approx 2$

Note. The 1σ uncertainties are indicated in parentheses. In the last column the approximate HPBW for the corresponding data is indicated.

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Table 8.  Physical Parameters Derived from the CH₃OH Data Obtained in the Compact Configuration (HPBW $\approx 3^{\prime\prime}$)

Object	Tkin	${N}_{{\mathrm{CH}}_{3}\mathrm{OH}}/{\rm{\Delta }}V$	${n}_{{{\rm{H}}}_{2}}$	Beam Filling Factor	${N}_{{\mathrm{CH}}_{3}\mathrm{OH}}/{N}_{{{\rm{H}}}_{2}}$
	(K)	(1012 cm−3s)	(cm−3)	(%)
SMA2	50 (40–80)	2.0 (1.8–3.5)	$6.3\times {10}^{6}$ ($\leqslant {10}^{8}$)	99.3 ($\geqslant 84$)	${10}^{-8}(\geqslant {10}^{-8})$
SMA2-W	40 (25–65)	3.2 (2.2–4.5)	$1.0\times {10}^{6}$ ($6.3\times {10}^{4}$—$2.5\times {10}^{6}$)	64.0 (63–70)	${10}^{-7}(\geqslant {10}^{-8})$
SMA4	45 (30–90)	3.6 (0.8–7.9)	$5.6\times {10}^{4}$ (103–$6.3\times {10}^{5}$)	99.9 ($\geqslant 20$)	${10}^{-7}({10}^{-6}$–${10}^{-9})$

Note. The 1σ confidence intervals are indicated in parentheses.

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The total HNCO column density derived from the rotational diagram is ∼2 × 10¹⁶ cm⁻². We used the partition function from the Cologne Database for Molecular Spectroscopy (Müller et al. 2001, 2005), too. A comparison with the total gas column density indicates HNCO abundance of $X(\mathrm{HNCO})\sim {10}^{-8}$. This is a rather high value, close to the highest HNCO abundance derived in the survey of massive cores by Zinchenko et al. (2000). However, taking into account the note above about the HNCO lines saturation, this value may need a correction.

In the 350 GHz band we detected many SO₂ lines (Table 2). The corresponding population diagram is shown in Figure 13. It indicates a range of temperatures. For low-excitation transitions the rotational temperature is 65 ± 11 K. For the high-excitation ones it is 146 ± 16 K. The SO₂ observations were performed at a much lower angular resolution (2 arcsec) than in the case of HNCO and apparently include emission from both the hot core and the surrounding cloud. The derived temperatures are consistent with the other estimates for these components.

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Another indicator of the hot environment is the vibrationally excited HCN. We detected the emission in the ${v}_{2}=1$ state, about 1000 K above the ground state. For the ${v}_{2}=2$ emission the upper limit is about five times lower. Following the analysis presented in, e.g., Veach et al. (2013) we obtain an upper limit for the excitation temperature between these states of about 500 K. This is consistent with the other estimates of the gas temperature given above. The critical density for HCN excitation is high (>10¹⁰ cm⁻³, Veach et al. 2013) which is consistent with our density estimates for the hot gas.

4.1.3. A Cold Clump in the Hot Core?

In Figure 14 we plot maps of the DCN $J=3-2$ and ¹³CS $J=5-4$ integrated line emission in the SMA1 core. The emission regions are very compact and the emission peak is shifted from the continuum peak which apparently coincides with the YSO location. The projected distance from the continuum peak is roughly 300 AU.

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The velocities of the DCN and ¹³CS emission are practically the same as those of high excitation lines of other molecules tracing apparently the hot gas. Therefore, most probably the observed DCN emission arises within the disk and since we probably see the disk nearly face-on as discussed above, the physical distance from the center is not much larger than the projected distance. An estimate of the abundances in the LTE approximation gives $X(\mathrm{DCN})\gt {10}^{-11}$ and $X{(}^{13}\mathrm{CS})\gt 3\times {10}^{-11}$ (assuming the total mass of the clump <10 M_⊙ and temperature of 50–100 K). This is a "normal" value for ¹³CS while the derived DCN abundance implies a significant deuteriation as follows from comparison with typical HCN abundances in massive cores (e.g., Zinchenko et al. 2009). As pointed out in Paper I, recent modeling by Albertsson et al. (2013) shows that the DCN/HCN abundance ratio sharply drops at temperatures ≳80 K. It means that the temperature of the DCN-emitting clump should be rather low, much lower than the temperature of the hot gas in the disk. Another possible explanation, as also discussed in Paper I, could be a very young age of the clump, insufficient to reach the steady-state DCN/HCN abundance ratio, but this looks less probable. In principle, interiors of the accretion disk can be well shielded from an external radiation, probably providing necessary low temperatures. In any case this clump deserves further investigation in order to clarify its properties and nature. It is worth mentioning that the mass of this clump cannot be lower than ∼1 M_⊙, otherwise the inferred ¹³CS abundance would be unrealistically high. The virial mass we estimate about the same, hence the clump can be gravitationally bound.

4.2. SMA2

4.3. SMA3

4.4. SMA4

In the S255IR area high velocity emission is detected in lines of CO, SiO and several high density tracers including HCN, HCO⁺, and CS. In Figures 15, 16 we present channel maps of the CO $J=3-2$ emission in the blue and red line wings, respectively. The maps of the integrated line wing emission (Figure 17) show that the CO high velocity emission observed with the SMA looks like a highly collimated bipolar outflow originating near SMA2. There is another more compact component near SMA1.

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On the other hand the bipolar outflow observed in HCN and HCO⁺ is apparently associated with SMA1 while is parallel to the CO flow (Figure 17). It is apparently associated with the jet observed in particular in the Fe ii emission which is also plotted in this figure. It is worth mentioning that the extent of the jet is much smaller than the extent of the outflow seen in CO, although their orientations coincide.

The SMA data alone hint at two parallel outflows with different origins. However, single-dish CO(3–2) observations with the IRAM 30 m telescope show a different picture. Here we see a wider outflow clearly originating at SMA1 (Figure 18). While the emission peaks coincide in the SMA and 30 m maps, near the driving source the CO emission measured with the SMA traces only the northern edge of the outflow. This shows that the SMA map gives a distorted picture and in reality we have here a wider angle outflow originating at SMA1.

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Apparently this means that this northern edge of the outflow contains a relatively bright component with a characteristic scale comparable to the SMA beam. The southern edge of the outflow near the driving source would be more diffuse and resolved out by the SMA. Most probably this implies that the CO emission is formed in a compressed layer surrounding the outflow cavity. For some (unclear) reason in the northern part this layer is more pronounced. It is worth mentioning that the brightest peaks of the high velocity CO emission practically coincide in the SMA and 30 m maps and also coincide with spots of molecular hydrogen emission (Wang et al. 2011) which probably indicate bow shocks. Basic physical parameters of the outflow were estimated by Wang et al. (2011) from their CO(2–1) data. Our new data are consistent with these estimates.

The position–velocity (P–V) diagram for the outflow along the cut indicated in Figure 18, constructed from the 30 m data, is shown in the left panel of Figure 19. The gap at ${V}_{\mathrm{LSR}}\sim 24$ km s⁻¹ is caused by the emission at the reference offset position mentioned above. In the right panel we present the P–V diagram obtained from the SMA CO(3–2) data. The main features of these diagrams coincide. They show non-monotonic dependence of the velocity on offset, which hints at two outflow components at different distances from the driving source. These P–V diagrams are somewhat different from that presented by Wang et al. (2011). However, Wang et al. (2011) plotted this diagram for a different cut, with the position angle of 75°, which makes the direct comparison irrelevant.

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In Figure 20 we plot the P–V diagram for the HCO⁺(4–3) emission as observed with the SMA, along with the part of the CO(3–2) P–V diagram for the same intervals of the offset and velocity. The diagrams are very similar, which means that both molecules trace apparently the same gas.

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5.1. Dense High Velocity Clump

Our data show a strong, compact blue-shifted CS and HCN emission close to the peak of the CO blue-shifted line wing emission. The peak of this CS and HCN emission is marked by cross in Figure 18. The CO(3–2), CO(2–1), HCN(4–3), CS(7–6), HCO⁺(4–3), and N₂H⁺(3–2) spectra toward this position are presented in Figure 21.

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One can see rather strong and broad HCN and CS lines at a central velocity of about −5 km s⁻¹. There is a hint of wings in these lines extending to about −20 km s⁻¹. The peak of the CO emission, especially in the higher $J=3-2$ transition is observed at more negative velocities. There is no detectable HCO⁺(4–3), N₂H⁺(3–2), CH₃CN, and continuum emission (for N₂H⁺ and CH₃CN we use the data from Paper I).

The spatial distribution of the CS(7–6) and CO(3–2) emission integrated in the velocity ranges −13...+1 km s⁻¹ and −19...+1 km s⁻¹, respectively, is shown in Figure 22. The HCN distribution is very similar to that of CS. The deconvolved size of the CS emitting clump is about 11 × 03 which corresponds to 1800 AU × 500 AU. It is apparently located at the head of the stream observed in CO, almost exactly at the jet axis shown in Figure 18.

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At the velocity of the CS and HCN emission peak, the brightness temperatures in the CO(3–2), CO(2–1), and HCN(4–3) lines are practically the same. The frequencies of the CO(3–2), HCN(4–3), and CS(7–6) lines are close to each other and beam parameters for them should be similar, too. However, the comparison with the CO(2–1) is complicated by the significantly different frequencies.

The CO(3–2) line is most probably saturated. Then the HCN(4–3) line should be saturated, too. The CS line is weaker implying either a relatively lower optical depth (but close to unity anyway) or a smaller size of the emission region. A simple modeling using, e.g., RADEX (van der Tak et al. 2007) shows that the column densities of HCN and CS required to explain such optical depths are ≳10¹⁵ cm⁻². The relative abundances of these molecules are ≤10⁻⁸. Therefore the total gas column density in the clump is ≳10²³ cm⁻². Then, using the size estimated above we obtain the gas density n ≳ 3× 10⁶ cm⁻³. This estimate is consistent with the observations of the HCN and CS lines which require at least such densities for excitation.

The large HCN and CS line widths indicate that the clump is gravitationally unbound. The virial mass of this clump estimated in the usual way (e.g., Zinchenko et al. 1994) is ${M}_{\mathrm{vir}}\sim 30$ M_⊙. At the same time, the non-detection of dust emission implies an upper limit for mass orders of magnitudes lower. Therefore, the clump represents a transient entity.

It is worth mentioning the non-detection of HCO⁺ and N₂H⁺ emission. Both molecules can be destroyed by dissociative recombination (e.g., Zinchenko et al. 2009). Therefore their absence can indicate an enhanced ionization in the dense clump.

5.2. Ionized Gas

The surroundings of the SMA1/SMA2 clumps and high-velocity outflow are traced in several molecular lines. One of the most informative is the N₂H⁺(3–2) transition observed with both the SMA and the IRAM 30 m telescope. The combined map of the N₂H⁺(3–2) emission is presented in Figure 18. It shows an absence of N₂H⁺ in the hot core in accordance with our previous findings (Paper I). The overall morphology of the N₂H⁺ emission suggests that it originates in an envelope around the central cores and the outflow lobes.

The distribution of the SiO(5–4) emission is different (Figure 23). It peaks near the SMA1 core. There is also a feature in the area of the blue outflow lobe that may be associated with the N₂H⁺ emission. The CS(7–6) distribution (right panel in Figure 23) seems to be the most uniform one. It shows a rather smooth, almost spherical halo around the SMA1 and SMA2 cores. There is no sign of the outflow influence on the CS distribution. One may suspect that the optical depth in the CS line is too high and we see only the outer regions of the core. However, the SMA1 and SMA2 clumps are well resolved in this map which makes such explanation less probable. Then, the total gas distribution is apparently not significantly affected, too.

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The main goal of this study is the characterization of the outflow and probable accretion disk associated with the massive YSO in the S255IR clump. The data presented in the previous sections shed new light on this system.

Concerning the outflow, one of the main findings is that the SMA interferometric data alone give a rather distorted picture. They hint at two highly collimated parallel outflows with different centers of origin. However, this impression is apparently caused by a significant flux loss in the interferometric measurements. The single-dish CO observations clearly show a single less collimated outflow originating at the SMA1 core. The CO emission retrieved by the SMA near the driving source originates apparently from the northern wall of the outflow cavity. This means that this part of the wall is rather thin and bright. The absence of a noticeable CO emission at the opposite side of the wall implies a more diffuse distribution of the emission here, probably due to the density structure of the surrounding medium.

The question arises of how common this effect is in interferometric studies of outflows. We can easily imagine a situation when an interferometric image will show multiple outflows from a single driving source, while in fact there is a single wide-angle outflow.

The observed P–V diagrams for the outflow (Figures 19, 20) most probably indicate periodic ejections from the driving source. Two events can be traced in the data. They apparently created jet knots at different distances from the central star. The older one is responsible for the extended CO outflow. The peaks of the CO emission coincide with bright H₂ emission spots (Wang et al. 2011) which probably indicate bow shocks at the heads of the jets. The next ejection event created other bow shocks which are seen in particular in the Fe ii emission (Figure 17). This later ejection entrains dense molecular gas observed in the wings of the HCO⁺, HCN and CS lines. It is also traced in the CO emission as can be seen in the CO position–velocity diagram (Figures 19, 20). Another manifestation of this activity is apparently the dense high velocity clump (Section 5.1), which most probably represents dense gas at the head of a bow shock. All the jet knots lay practically on a straight line. Therefore the orientation of the jet does not change significantly with time. The age of the first, extended outflow was estimated at about 7000 years by Wang et al. (2011). Assuming the same ejection velocities for the two events and comparing distances of the bow shocks from the driving source we can conclude that the second ejection happened about 1000 years ago. Therefore the time interval between the ejection events is about 6000 years (the age estimates were done under the usual assumption of an inclination angle for a jet of 45°).

Several entrainment mechanisms are usually discussed for molecular outflows (e.g., Arce et al. 2007; Frank et al. 2014). In our case the most probable one is the jet bow-shock model (e.g., Raga & Cabrit 1993). The main argument in favor of this model is the obvious presence of several bow shocks (traced in the H₂ and Fe ii emission) clearly associated with the high-velocity molecular gas. This model is also consistent with the observed outflow morphology and kinematics. The shape of the P–V diagram differs from the frequently observed "Hubble law" and shows a range of velocities at the largest distances from the driving source. A similar shape was observed in some other outflows which are apparently driven by bow shocks and this type of shape is expected in the theoretical models (e.g., Lee et al. 2000).

The jet is apparently launched from the accretion disk around the central massive young star. As mentioned above, there is a chain of water masers along the jet. For several water maser spots shown in Figure 5, proper motions have been measured (Goddi et al. 2007). The velocities (from ∼10 to ∼30 km s⁻¹ at a distance of 1.6 kpc) are mostly perpendicular to the jet axis with outward components in some cases. The maser velocity pattern is consistent with clockwise rotation around the core center. Velocities of the masers closest to the center imply a dynamical mass of the order of 20 M_⊙. The mass of the central star is estimated to be 27 M_⊙ (Ojha et al. 2011) assuming a distance of 2.5 kpc. At a distance of 1.6 kpc the mass of the star would be 23.5 ± 3.6 M_⊙. This estimate is based on the unpublished value of the star luminosity of (6.8 ± 2.8) × 10⁴ L_⊙ (for the distance of 1.6 kpc). Now we reconsidered the star SED by including data at millimeter wavelengths from Zinchenko et al. (2009). This results in the luminosity of (3.5 ± 0.3) × 10⁴ L_⊙ and mass of 20 ± 2 M_⊙. It is worth mentioning that Pashchenko et al. (2003) estimated the mass of the central star to be (6–7) M_⊙ from their analysis of water maser variability. However, they employed a systemic velocity of 8.3 km s⁻¹ that contradicts our findings which clearly show the systemic velocity to be about 4.8 km s⁻¹. In the following we will take the mass of 20 M_⊙ as the most probable value. This is sufficient to explain velocities of the masers closest to the star. However, maser spots located at larger distances from the center have similar velocities and imply a much larger central mass of the order of 100 M_⊙. This value is inconsistent with all other estimates. The corresponding mean density would be unrealistically high. This means that we probably need another explanation for the maser velocities. One speculation is that it might be jet rotation. Rapid rotation of protostellar jets is expected in theoretical models (e.g., Pudritz et al. 2007) and was indeed observed in some cases (Bacciotti et al. 2002).

Water masers are thought to be excited by interaction of the jet with the surrounding medium. The fact that they are moving in one direction on each side of the central star implies some asymmetry in the system. There may be some misalignment between the jet and the rotation axis of the material in the outer parts of the clump (apparently, in the axially symmetric case the observed maser movements would be symmetric relative to the axis). It can be that in the inner part the orientation of the disk plane is different. However, the difference cannot be large. We see also other indications of the jet interaction with the surrounding medium, in particular, increased line widths along the jet.

Figure 5 shows that the water maser condensations are shifted from the jet axis in direction of their proper motion. This projected shift is of the order of 200 AU. At the velocities mentioned above, the maser condensations would travel this distance in about 30–100 years. From this consideration it looks probable that the event that created these masers happened quite recently.

The disk is apparently strongly fragmented as follows from the derived beam filling factor (∼0.2) for CH₃CN and CH₃OH. The mean gas density in the disk is about 6 × 10⁸ cm⁻³. The density of the fragments should be much higher. Observations of vibrationally excited HCN indicate densities >10¹⁰ cm⁻³ which is consistent with this picture.

The derived mass of the hot gas (∼0.3 M_⊙) is an order of magnitude insufficient to gravitationally bind the core. The virial mass estimated in the usual way (e.g., Zinchenko et al. 1994) from the line width is about 7 M_⊙. There can be also some amount of cold material in the disk, as follows from the DCN observations. However, the total mass of the disk is apparently much lower than the mass of the central star which is estimated to be about 20 M_⊙ (see above). Since the disk mass is significantly lower than the star mass, the disk is probably not self-gravitating and we may expect Keplerian rotation. Unfortunately the achieved angular resolution is not sufficient to reliably measure the disk rotation curve.

For a central mass of 20 M_⊙ the inclination angle of the disk (derived from a comparison of the observed rotation line-of-sight velocity with the expected velocity for this mass) should be small, about 25° (Figure 6), i.e., the disk is seen almost face-on. This is consistent with our maps which do not show a significant elongation of this object. However, this conclusion looks somewhat suspicious because of the well-collimated appearance of the molecular outflow. Observations at a higher angular resolution are needed to clarify the structure and kinematics of this object.

The temperature of the hot gas derived from the CH₃CN, CH₃OH, and SO₂ observations is about 130–180 K. A similar estimate for the temperature of the hot component was obtained by us many years ago from the analysis of the IRAS data (Zinchenko et al. 1990). The HNCO rotational temperature is somewhat higher, about 320 K. However, as mentioned above, this value probably represents an upper limit to the excitation temperature. In addition, the HNCO excitation can be influenced by the FIR radiation. The HNCO excitation in massive cores was discussed by Zinchenko et al. (2000). That analysis shows that in the present case collisional excitation can be effective. The critical density for excitation of the ${K}_{-1}=3$ ladder is of the order of 10¹⁰ cm⁻³. Such density is quite possible in the SMA1 core as follows from the estimates above. However, radiative excitation by FIR emission via the b-type transitions cannot be excluded. With a column density of $N({{\rm{H}}}_{2})$ ∼ 3× 10²⁴ cm⁻² the dust will be optically thick at these frequencies creating a sufficiently strong radiation field. The HNCO excitation especially in the higher ${K}_{-1}$ ladders may reflect the effective temperature of this field.

The highest temperature of the observed molecular material is about 300 K. Taking the source luminosity to be about 3.5 × 10⁴ L_⊙ (see above) this value corresponds to the dust equilibrium temperature at a distance of 200–300 AU from the star. Since we do not see warmer molecular gas, no dust should be present within this radius in a case of no shielding. However, this radius is comparable to the observed size of the core and we do not see any central hole in the molecular distribution. It is then possible that shielded molecular clumps are distributed at smaller radii.

In this respect an interesting feature of our observations is the apparently relatively cold and rather massive clump with a strong DCN emission. It may be gravitationally bound and the mass of the clump is sufficiently high to consider it as a possible low-mass protostar.

Concerning the surroundings of the SMA1 clump, we see that the outflow strongly affects the chemical composition of the medium. The N₂H⁺ molecules are destroyed along the outflow. The SiO distribution also seems to be significantly affected. At the same time there is no noticeable influence on the CS distribution. The density structure is probably not affected either. As mentioned above (Sections 4.2, 4.4) the nearby clumps are probably influenced by shocks that are apparently associated with this outflow.

We presented the results of our observations of the S255IR area with the SMA at 1.3 mm in the very extended configuration and at 0.8 mm in the compact configuration as well as with the IRAM 30 m at 0.8 mm. The best achieved angular resolution is about 0.4 arcsec. The dust continuum emission and several tens of molecular spectral lines are observed. The majority of the lines is detected only toward the S255IR-SMA1 clump. In summary, our main findings are the following.

• 1.  
  The S255IR-SMA1 clump apparently represents a rotating structure (probably a disk) around the young massive star. The achieved angular resolution is still insufficient to establish the character (Keplerian or non-Keplerian) of the rotation. The temperature of the molecular gas reaches 130–180 K. The size of the clump is about 500 AU. It is apparently strongly fragmented as follows from the derived small (∼0.2) beam filling factors for various molecules. The mean gas density is about 6 × 10⁸ cm⁻³. The density of the fragments should be much higher which is confirmed by observations of HNCO and vibrationally excited HCN. The mass of the hot gas is ∼0.3 M_⊙ and the total mass of the clump is significantly lower than the mass of the central star (about 20 M_⊙). The inclination angle of the disk (derived from comparison of the observed rotation line-of-sight velocity with the expected rotation velocity for this mass) should be small, about 25°.
• 2.  
  We detected a strong DCN(3–2) emission near the center of the SMA1 clump. Most probably it indicates the presence of a rather large amount (≳1 M_⊙) of cold (≤80 K) material. This cold clump can be gravitationally bound.
• 3.  
  High velocity emission is observed in the CO line as well as in lines of high-density tracers HCN, HCO⁺, CS, and other molecules. The CO outflow is much more extended than that observed in the lines of high-density tracers. Its morphology obtained from a combination of the SMA and IRAM 30 m data is significantly different from that derived from the SMA data alone. The CO emission detected with the SMA traces only one boundary of the outflow and leads to a rather distorted picture of the outflow structure. The velocity of the CO outflow reaches ∼60 km s⁻¹.
• 4.  
  The outflow is most probably driven by the jet bow shock mechanism. The available data indicate at least two major ejection events with a time interval of several thousand years between them. The direction of the ejections does not change with time. The high-velocity emission in the lines of high-density tracers is associated with the peaks of the Fe ii emission related to the bow shocks caused by rather recent ejections from the SMA1.We detected a dense high-velocity clump apparently associated with one of the bow shocks. It shows a strong emission in the HCN(4–3) and CS(7–6) lines. At the same time there is no detectable HCO⁺(4–3) or N₂H⁺(3–2) emission which can be probably explained by enhanced ionization.
• 5.  
  The proper motions of the water masers excited along the jet imply some misalignment of the jet with the rotation axis of the material in the outer parts of the clump. It could be that the orientation of the disk in the inner and outer parts is somewhat different. However, the difference cannot be large.
• 6.  
  The outflow strongly affects the chemical composition of the surrounding medium. The N₂H⁺ molecules are destroyed along the outflow. The SiO distribution also seems to be significantly affected. At the same time there is no sign of the outflow's influence on the CS distribution. The total gas distribution is apparently not significantly affected either.

This work was supported by the Russian Academy of Sciences (Research program No. 17 of the Department of Physical Sciences), the Russian Foundation for Basic Research (RFBR), the Ministry of Science and Technology (MoST) in Taiwan, and the Department of Science and Technology (DST) of the Government of India in the frameworks of the research grants RFBR 08-02-92001-NSC, RFBR 13-02-92697-Ind, RFBR 15-02-06098, 15-52-45057, and DST-RFBR INT/RUS/RFBR/P-142. S.-Y.L. and Y.-N.S. acknowledge grant support 102-2119-M-001-010-MY2 from MoST. A.M.S. was supported by the Russian Science Foundation (grant No. 15-12-10017). S.V.S. was supported by the Ministry of Education and Science of the Russian Federation (state task No. 3.1781.2014/K). Y.W. was supported by the Swiss National Science Foundation, NSFC 11303097 and 11203081, China. The research is also partly supported by the grant within agreement No. 02.B.49.21.0003 between The Ministry of Education and Science of the Russian Federation and Lobachevsky State University of Nizhni Novgorod and by the Russian Education and Science Ministry Project 3.1252.2014/k. We are grateful to Elena Trofimova for help with line identification and to the anonymous referee for detailed helpful comments. The research has made use of the SIMBAD database, operated by CDS, Strasbourg, France.