Nobeyama 45 m mapping observations toward Orion A. III. Multi-line observations toward an outflow-shocked region, Orion Molecular Cloud 2 FIR 4

Abstract

We present the results of mapping observations toward an outflow-shocked region, OMC-2 FIR 4, using the Nobeyama 45 m telescope. We observed the area in ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), CCS (J_N = 8₇–7₆), HCO⁺ (J = 1–0), H¹³CO⁺ (J = 1–0), HN¹³C (J = 1–0), H¹³CN (J = 1–0), DNC (J = 1–0), N₂D⁺ (J = 1–0), and DC₃N (J = 9–8). We detected a dense molecular clump that contains FIR 4/5. We also detected, in the ¹³CO line, blueshifted and redshifted components driven presumably by protostellar outflows in this region. The axes of the FIR 3 and VLA 13 outflows, projected on the plane of the sky, appear to point to the FIR 4 clump, suggesting that it may be compressed by protostellar outflows from Class I sources, FIR 3 and VLA 13. Applying a hyperfine fit of N₂H⁺ lines, we estimated the excitation temperature to be ∼20 K. The high excitation temperature is consistent with the fact that the clump contains protostars. CCS emission was detected in this region for the first time. Its abundance is estimated to be a few × 10⁻¹², indicating that the region is chemically evolved at ∼10⁵ yr, which is comparable to the typical lifetime of Class I protostars. This timescale is consistent with the scenario that star formation in FIR 4 is triggered by dynamical compression of the protostellar outflows. The [HNC]/[HCN] ratio was evaluated to be ∼0.5 in the dense clump and the outflow lobes, whereas it is somewhat larger in the envelope of the dense clump. The small [HNC]/[HCN] ratio indicates that HNC formation was prevented due to high temperatures. Such high temperatures seem to be consistent with the scenario that either protostellar radiation, or outflow compression, or both affected the thermal properties of this region.

1 Introduction

In cluster-forming regions, stellar feedback from forming stars often influences the surrounding environments (McKee & Ostriker 2007; Wang et al. 2010; Krumholz et al. 2014). For example, protostellar outflows can dynamically compress the surrounding gas and/or adjacent dense cores, and promote fragmentation and gravitational contraction, leading to star formation (Shimajiri et al. 2008). In addition, such stellar feedback injects significant momentum in molecular clouds, controlling the efficiency of star formation (McKee 1989; Nakamura & Li 2007, 2014; Nakamura et al. 2012; Arce et al. 2011). Theoretically, the negative effect of the stellar feedback is more important than the positive effect (Li & Nakamura 2006; Nakamura & Li 2007; Krumholz et al. 2014), but the positive effect is sometimes observed to enhance local star formation (Elmegreen & Lada 1978; Sugitani & Ogura 1994; Chauhan et al. 2009). Among the stellar feedback processes, the protostellar outflow feedback can trigger local star formation through the dynamical compression of dense parts in molecular clouds (Sandell & Knee 2001; Yokogawa et al. 2003; Shimajiri et al. 2008).

One of the good examples of the dynamical interaction between the outflow and dense gas is the FIR 4 region in the Orion Molecular Cloud 2 (OMC-2), in which fragmentation seems to be induced by outflow shock from a nearby protostar FIR 3 (Shimajiri et al. 2008; López-Sepulcre et al. 2013). Multi-epoch, multiwavelength Very Large Array (VLA) observations (Osorio et al. 2017) support that the FIR 3 outflow is actually moving toward the FIR 4 region, suggesting that the FIR 3 outflow is interacting with the FIR 4 region (see also González-García et al. 2016). These pieces of morphological evidence strongly support that the outflow is interacting with dense gas in this region. Shimajiri et al. (2008) detected 11 cores based on interferometric 3 mm dust continuum

observations [see also Hacar et al. (2018) for the Atacama Large Millimeter/submillimeter Array (ALMA) observations]. Some of the cores have masses of greater than the thermal Jeans mass in this region. Near-infrared observations suggest that protostars have already been created in this clump. Although the luminosity of the central protostar is not so prominent, the mass of the protostellar envelope was derived to be about 27 M_⊙ (Furlan et al. 2014). Therefore, this protostar is likely to evolve at least into an intermediate-mass star. Since there are several fragments with masses greater than the local Jeans mass, the FIR 4 region may be forming a small star cluster consisting of low-mass and intermediate-mass stars.

Spectral line surveys toward FIR 4 also suggest that FIR 4 is likely to be a hot core (Ceccarelli et al. 2010; Kama et al. 2013). The hot core might be formed by the outflow compression. Alternatively, it might be created before the outflow compression, which promotes fragmentation in the FIR 4 clump. In contrast, very recently, Favre et al. (2018) found that this region has a relatively uniform excitation temperature of c-C₂H₃, which led them to claim that there is no clear evidence of outflow compression, since compression should create a temperature gradient. They proposed that the FIR 3 outflow does not hit FIR 4, and that the spatial overlap of the outflow lobe and the clump in the two-dimensional map may be the apparent coincidence along the line of sight. However, because of the large uncertainty in the estimated temperatures and poor angular resolution, it would be difficult to discuss the temperature distribution in this region accurately based on the data. Higher sensitivity and higher angular resolution observations are needed to further constrain the outflow compression scenario.

Located in the Orion A molecular cloud, the FIR 4 region is one of the closest cluster-forming regions. The distance to the Orion A molecular cloud was estimated to be 414 ± 7 pc by Menten et al. (2007) and 418 ± 6 pc by Kim et al. (2008). These two measurements reasonably coincide with each other. Hereafter, we adopt 414 pc as a representative distance of the FIR 4 region. Some indirect evidence of dynamical interaction by the FIR 3 outflow also comes from recent spectral line surveys (Kama et al. 2013, 2015; Shimajiri et al. 2015b). In the present paper, we characterize the physical and chemical properties of the possible outflow-shocked region, FIR 4, by several molecular line observations, such as N₂H⁺, CCS, HN¹³C, and H¹³CN. In addition, we identify possible protostellar outflow candidates in this region, using the ¹³CO (J = 1–0) and HCO⁺ (J = 1–0) data.

Observations of the FIR 4 region were made part of the Nobeyama Star Formation Legacy Project at the Nobeyama Radio Observatory (NRO) to observe nearby star-forming regions, such as Orion A, Aquila Rift, and M 17. An overview of the project will be presented in a separate paper (Nakamura et al. 2019), and detailed observational results for the individual regions are given in other articles [Orion A: Tanabe et al. (2019) and Ishii et al. (2019); Aquila Rift: Shimoikura et al. (2019b) and Kusune et al. (2019); M 17: Shimoikura et al. (2019a), Sugitani et al. (2019), and Q. Nguyen Luong et al. (in preparation); other regions: Dobashi et al. (2019a, 2019b)]. The FIR 4 region was observed once or twice every day to check the absolute flux calibration of ¹³CO (J = 1–0), C¹⁸O (J = 1–0), and N₂H⁺ (J = 1–0) in Orion A, which were obtained with a new receiver, FOREST, installed on the Nobeyama 45 m telescope. Thus, these molecular line data have excellent sensitivity.

The paper is organized as follows. Section 2 describes the details of our Nobeyama 45 m observations. In section 3, we present the results of our mapping observations. We report the detection of the CCS emission at 93 GHz, which is the first detection of CCS in the OMC-2 region. In section 4, we discuss some characteristics of the internal structure of this region. Using our molecular-line data, we identify the molecular outflow candidates in this region in section 5, and briefly discuss how the outflows interact with the dense clump. Finally, section 6 summarizes the main results of this study.

2 Observations and data

2.1 ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), and CCS (J_N = 8₇–7₆) observations

Fig. 1.

Mapped areas in Orion A. The color image shows the Planck–Herschel H₂ column density created by Lombardi et al. (2014). The areas mapped with FOREST and T70 are indicated with a white and a green square, respectively. We also indicated the regions of OMC-1, OMC-2, and OMC-3.

Table 1.

Observed lines with FOREST.

Molecule	Transition	Rest frequency (GHz)	Effective resolution*	Noise level† (K)
13CO	J = 1–0	110.201354	21|${^{\prime\prime}_{.}}$|7	0.19
C18O	J = 1–0	109.782176	21|${^{\prime\prime}_{.}}$|7	0.16
CCS	JN = 87–76	93.870098	32|${^{\prime\prime}_{.}}$|5	0.04
N2H+	J = 1–0, |$F_1F \rightarrow F_1^{\prime }F^{\prime } =23\rightarrow 12$|	93.1737637‡	23|${^{\prime\prime}_{.}}$|4	0.06

*The CCS data were spatially smoothed to improve the signal-to-noise ratio.

^†The noise levels are measured with a velocity resolution of 0.05 km s⁻¹.

^‡The rest frequency of the main hyperfine component from Pagani, Daniel, and Dubernet (2009). In the last column, we present the rms noise level in T_mb. We observed all six hyperfine components.

We carried out mapping observations in ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), and CCS (J_N = 8₇–7₆) toward the FIR 3/4/5 region in the H ii region Orion A using a new receiver, FOREST (Four-Beam Receiver System on the 45 m Telescope), which is a four-beam, dual-polarization sideband-separating SIS receiver installed on the 45 m telescope of the Nobeyama Radio Observatory. See Minamidani et al. (2016) for details of FOREST. The telescope has a full-width at half-maximum (FWHM) beam size of 18″ at 93 GHz. The beam separation of FOREST is 51|${^{\prime\prime}_{.}}$|7 on the sky. The mapping area is indicated with a white square in figure 1. The observed lines and their parameters are summarized in table 1. The observations were done in the On-The-Fly (OTF) mode (Sawada et al. 2008) during the period from 2016 December to 2017 March. We adopted a spectral window mode that allows us to obtain four lines simultaneously. The FIR 4 maps were obtained for the intensity calibration of the larger Orion A map (see Nakamura et al. 2019 for details). As backends, we used a digital spectrometer based on an FX-type correlator, Spectral Analysis Machine for the 45 m telescope (SAM45), which is 16 sets of 4096 channel arrays. The frequency resolution of all spectrometer arrays was set at 15.26 kHz, which corresponds to 0.05 km s⁻¹ at 93 GHz. The scan interval of the OTF observation is set at 5|${^{\prime\prime}_{.}}$|17, so that individual scans by the four beams of FOREST are fully overlapped.

The temperature scale was determined by the chopper-wheel method. The telescope pointing was checked every 1 hr by observing the SiO maser line from Orion KL (RA [J2000.0], Dec [J2000.0]) = (⁠|${5^{\rm h}35^{\rm m}14{^{\rm s}_{.}}5}$|⁠, |${-5^{\circ}22^{\prime }30{^{\prime\prime}_{.}}4}$|⁠). The pointing accuracy was better than ∼3″ throughout the entire observation. The typical system noise temperature was in the range of 150 K to 200 K in the single-sideband mode at the observed elevation of El = |${30^{\circ}}$|–|${50^{\circ}}$|⁠.

In order to minimize the scanning effects, data with orthogonal scanning directions along the RA (x-scan) and Dec (y-scan) axes were combined on a single map. In total, we combined six x-scan and seven y-scan data on single maps. We adopted a spheroidal function with a width of 7|${^{\prime\prime}_{.}}$|5 as a gridding convolution function to calculate the intensity at each grid point of the final cube data with a spatial grid size of 7|${^{\prime\prime}_{.}}$|5, about a third of the beam size. The resultant effective angular resolution was about 22″ at 110 GHz, corresponding to ∼0.05 pc at a distance of 414 pc.

The main beam efficiencies were 0.50 at 93 GHz and 0.43 at 110 GHz. We divided the intensities by the main beam efficiencies at the corresponding frequencies and made the maps on the brightness temperature scale. More details concerning the observations and data reduction will be given in a separate paper (Nakamura et al. 2018). The CCS image presented in figure 2d was smoothed so as to improve the signal-to-noise ratios of the map, since the CCS emission in the FIR 4 region is very weak.

2.2 HCO⁺ (J = 1–0), H¹³CO⁺ (J = 1–0), HN¹³C (J = 1–0), H¹³CN (J = 1–0), DNC (J = 1–0), N₂D⁺ (J = 1–0), and DC₃N (J = 9–8) observations

Table 2.

Observed lines with T70.

Molecule	Transition	Rest frequency (GHz)	Effective resolution*	Noise level† (K)
HCO+	J = 1–0	89.188526	24|${^{\prime\prime}_{.}}$|9	0.45
H13CO+	J = 1–0	86.754288	33|${^{\prime\prime}_{.}}$|6	0.20
H13CN	J = 1–0, F = 2–1	86.3401764	33|${^{\prime\prime}_{.}}$|6	0.16
HN13C	J = 1–0, F = 2–1	87.090859	33|${^{\prime\prime}_{.}}$|6	0.18
N2D+	J = 1–0	77.107798	32|${^{\prime\prime}_{.}}$|5	0.29
DNC	J = 1–0	76.305697	32|${^{\prime\prime}_{.}}$|5	0.27
DC3N	J = 9–8	75.987149	32|${^{\prime\prime}_{.}}$|5	0.33

*For every line except HCO⁺, the image was smoothed to improve the signal-to-noise ratio. The last column is the average rms noise level of the whole area at a velocity resolution of 0.05 km s⁻¹.

We carried out mapping observations in HCO⁺ (J = 1–0), H¹³CO⁺, (J = 1–0), H¹³CN (J = 1–0), HN¹³C (J = 1–0), DNC (J = 1–0), N₂D⁺ (J = 1–0), and DC₃N (J = 9–8) toward a slightly larger area than is mapped with FOREST, using the T70 receiver installed on the Nobeyama 45 m telescope. The mapping area is indicated in figure 1 with a green square. The observed lines and their parameters are summarized in table 2. The observation procedure was essentially the same as that of the FOREST observations. The observations were done in the OTF mode during the period from 2016 December to 2017 February. We used SAM45 at a frequency resolution of 3.81 kHz, corresponding to ∼0.0125 km s⁻¹ at ∼86 GHz. The scan interval of the OTF observation was set at 5″. The typical system noise temperature was in the range of 150 K to 250 K in the single sideband mode at the observed elevation.

The temperature scale was determined by the chopper-wheel method. The telescope pointing was checked every 1 hr by observing the SiO maser line from Orion KL. The pointing accuracy was better than ∼3″ throughout the entire observation. Adopting a spheroidal function with a width of 7|${^{\prime\prime}_{.}}$|5 as a gridding convolution function to compute the intensity at each grid point, we made final images with a spatial grid size of 7|${^{\prime\prime}_{.}}$|5 by combining an x-scan and a y-scan image on single maps to obtain a map on the antenna temperature scale. The resultant effective angular resolution was about 25″. Finally, we divided the intensities by the main beam efficiencies at the corresponding frequencies to obtain maps on the brightness temperature scale. The main beam efficiencies at 89 GHz, 87 GHz, and 76 GHz for the T70 receiver were about 0.543, 0.544, and 0.549, respectively. Every line except HCO⁺ is relatively weak. To improve the signal-to-noise ratio of the maps, the image presented below for every observed line except HCO⁺ was smoothed, and thus the resultant effective angular resolution was ∼33″ for those images.

3 Spatial distributions of molecular-line emission

Fig. 2.

Total integrated intensity maps of the FIR 4 region. (a) ¹³CO (J = 1–0), (b) C¹⁸O (J = 1–0), (c) N₂H⁺ (J = 1–0), (d) CCS (J_N = 8₇–7₆), (e) HCO⁺ (J = 1–0), (f) H¹³CO⁺ (J = 1–0), (g) HN¹³C (J = 1–0), (h) H¹³CN (J = 1–0), and (i) DNC (J = 1–0) . The integrated intensities are given in K km s⁻¹ on the brightness temperature scale. The squares from north to south indicate FIR 3, FIR 4, and FIR 5 dust cores identified by Chini et al. (1997), and the VLA13 protostar (Reipurth et al. 1999). The positions of the protostars listed in table 3 are shown by the crosses. The numbers in the panels denote the HOPS ID of the protostars. The effective angular resolution of the maps are shown in the lower left-hand corner of each panel.

Table 3.

Protostars in the FIR 3/4/5 region.

Name	RA [J2000.0]	Dec [J2000.0]	Mass	Class	Candidate*	ID†	Note	References‡
HOPS68	|${05^{\rm h}35^{\rm m}24{^{\rm s}_{.}}287}$|	−05°|${08^{\prime }30{^{\prime\prime}_{.}}65}$|	1.7	I	R	C	FIR 2	Aso et al. (2000); Williams, Plambeck, and Heyer (2003);
								Takahashi et al. (2008)
HOPS66	|${05^{\rm h}35^{\rm m}26{^{\rm s}_{.}}928}$|	−05°|${09^{\prime }24{^{\prime\prime}_{.}}40}$|	1.5	Flat				Takahashi et al. (2008) detected near-infrared knots
HOPS370	|${05^{\rm h}35^{\rm m}27{^{\rm s}_{.}}618}$|	−05°|${09^{\prime }34{^{\prime\prime}_{.}}06}$|	2.5	I	BR	C	FIR 3	Aso et al. (2000); Williams, Plambeck, and Heyer (2003); Wu et al. (2005);
								Takahashi et al. (2008); Shimajiri et al. (2008)
HOPS65	|${05^{\rm h}35^{\rm m}21{^{\rm s}_{.}}566}$|	−05°|${09^{\prime }38{^{\prime\prime}_{.}}62}$|	0.2	I
HOPS64	|${05^{\rm h}35^{\rm m}26{^{\rm s}_{.}}928}$|	−05°|${09^{\prime }54{^{\prime\prime}_{.}}51}$|	0.5	?	R	M		This paper
HOPS108	|${05^{\rm h}35^{\rm m}27{^{\rm s}_{.}}015}$|	−05°|${09^{\prime }59{^{\prime\prime}_{.}}59}$|	0.3	0	BR	M	FIR 4	Takahashi et al. (2008) claimed marginal detection
HOPS368	|${05^{\rm h}35^{\rm m}24{^{\rm s}_{.}}717}$|	−05°|${10^{\prime }29{^{\prime\prime}_{.}}78}$|	0.7	I	BR	C	VLA 13	Takahashi et al. (2008)

*R and B indicate the outflow candidates of redshifted and blueshifted components, respectively. Outflow identification is based on the present paper. See section 5 for details.

^†C and M indicate outflows identified as clear and marginal, respectively.

^‡Previous detection.

In this section, we describe the spatial distributions of the molecular-line emission in the observed area. In figure 2, we present integrated intensity maps of the molecular-line emission detected. Recently, Kainulainen et al. (2017) investigated the OMC-2 region with ALMA. Following their table 2, the protostars previously identified in this region are listed in table 3.

3.1 ¹³CO and C¹⁸O

In figures 2a and 2b, we present velocity-integrated intensity maps of ¹³CO and C¹⁸O, respectively. The velocity range of the integration was set at 5.1–15.1 km s⁻¹ for both ¹³CO and C¹⁸O. The distribution of the ¹³CO emission matches reasonably well with that of C¹⁸O. For both the ¹³CO and C¹⁸O maps, the intensities are the strongest near the upper right-hand corner of the observed area (recognized as red areas in figures 2a and 2b), which is a part of the other core traced in the 1.1 mm continuum. [Core No. 320 of Shimajiri et al. (2015a); see also the white box in figure 1.] The ¹³CO emission takes its local maximum near the position of FIR 4 (HOPS108) at (RA [J2000.0], Dec [J2000.0]) = (⁠|${5^{\rm h}35^{\rm m}27{^{\rm s}}}$|⁠, |${-5^{\circ}10^{\prime }00^{\prime \prime }}$|⁠), which coincides with the position of the C¹⁸O local peak. Both the ¹³CO and C¹⁸O emissions trace a dense clump and its envelope which contains FIR 3/4/5. The dense clump corresponds to Core No. 327 of Shimajiri et al. (2015a), and elongates in the north–south direction, which is roughly similar to the large-scale filament axis of the OMC-2 region. The ¹³CO and C¹⁸O images agree very well with those of Shimajiri et al. (2014). The high-resolution ALMA maps also revealed that the main filamentary structure is in the north–south direction (Kainulainen et al. 2017; Hacar et al. 2018).

3.2 N₂H⁺

Figure 2c presents an intensity map of the N₂H⁺ velocity-integrated of 1.6 km s⁻¹ to 20.1 km s⁻¹. In other words, we integrated the emission of all seven hyperfine components.

The strongest emission of N₂H⁺ (J = 1–0) comes near the position of FIR 4 (HOPS108) at (RA [J2000.0], Dec [J2000.0]) = (⁠|${5^{\rm h}35^{\rm m}26{^{\rm s}_{.}}5}, {-5^{\circ}10^{\prime }00^{\prime \prime }}$|⁠). The position of the N₂H⁺ peak agrees well with those of ¹³CO and C¹⁸O. In the following, we refer to the central compact molecular clump detected by N₂H⁺ as the N₂H⁺ clump. The clump contains FIR 4/5. The FIR 3 protostar is located at the northeastern edge of the N₂H⁺ clump. The VLA 13 protostar is located at the southwestern edge of the clump. This N₂H⁺ clump corresponds to Core No. 10 of Tatematsu et al. (2008). The N₂H⁺ clump has a couple of spines, which are indicated with dashed lines in figure 2c. This region was also observed with ALMA in N₂H⁺ (Hacar et al. 2018). The overall distribution of the N₂H⁺ emission agrees well with the ALMA map. For example, the positions of the spines are in good agreement with the structures seen in the ALMA image. The ALMA map revealed that the FIR 4 clump is connected with a couple of small filaments, which can be seen in figure 2c as spines. Such a structure is morphologically reminiscent of the hub-filament structure discussed by Myers (2009), although the size of the region is smaller than those originally discussed by Myers (2009). It is worth noting that the ¹³CO high-velocity component of the VLA 13 outflow anti-correlates with the N₂H⁺ distribution. Similar distributions are recognized in the maps of other high-density tracers (see the next subsections). The other moderately strong N₂H⁺ emission comes from just north of the N₂H⁺ clump, which deviates from the C¹⁸O peak at the upper right-hand corner of figure 2b.

3.3 CCS

In figure 2d we show an integrated intensity map of CCS (J_N = 8₇–7₆). The overall distribution is roughly similar to that of C¹⁸O. This is the first detection of CCS emission in the OMC-2 region. Previous observations by Tatematsu et al. (2008) reported no detection of CCS (J_N = 4₃–3₂), presumably because of a lower sensitivity. The strongest emission comes from the upper right-hand corner of the observed area (west of HOPS68), where the ¹³CO and C¹⁸O emissions are strong. This component is likely to belong to Core No. 320 of Shimajiri et al. (2015a). We can also see another local peak in the northern part just above the N₂H⁺ clump (east of HOPS68). The other strong peaks of CCS are located in the N₂H⁺ clump toward FIR 4 (HOPS108) at (RA [J2000.0], Dec [J2000.0]) = (⁠|${5^{\rm h}35^{\rm m}27{^{\rm s}}}$|⁠, |${-5^{\circ}10^{\prime }00^{\prime \prime }}$|⁠), and just in the southeastern part of VLA 13 (HOPS368).

3.4 HCO⁺ and H¹³CO⁺

In figures 2e and 2f, we present integrated intensity maps of HCO⁺ and H¹³CO⁺. The HCO⁺ distribution elongates in the north–south direction, and its distribution resembles that of ¹³CO and C¹⁸O. However, the position of the local peak at the northern part significantly deviates from those of the ¹³CO and C¹⁸O (east of HOPS68), and is consistent with the N₂H⁺ peak. H¹³CO⁺ emission basically follows the HCO⁺ distribution and is more extended than the N₂H⁺ distribution.

3.5 HN¹³C and H¹³CN

In figures 2g and 2h we present integrated intensity maps of HN¹³C and H¹³CN. H¹³CN has three hyperfine components. We integrated all three components to make the map. Roughly speaking, the spatial distributions of these emission lines are similar to that of H¹³CO⁺, although they are less extended.

The strongest H¹³CN emission comes from the N₂H⁺ clump (at the position of FIR 4). HN¹³C has a few relatively strong peaks at (RA [J2000.0], Dec [J2000.0]) = (⁠|${5^{\rm h}35^{\rm m}27{^{\rm s}}}$|⁠, |${-5^{\circ}10^{\prime }17^{\prime \prime }})$|⁠, (⁠|${5^{\rm h}35^{\rm m}23{^{\rm s}}}$|⁠, |${-5^{\circ}10^{\prime }9^{\prime \prime }}$|⁠), and (⁠|${5^{\rm h}35^{\rm m}25{^{\rm s}}}$|⁠, |${-5^{\circ}8^{\prime }35^{\prime \prime }}$|⁠). The strongest peak is between FIR 4 and FIR 5. A local peak is located between HOPS65 and HOPS368. Another local peak is located near HOPS68. In the N₂H⁺ map, we can see spines stretching from the central N₂H⁺ clump in the southwestern direction. The HN¹³C emission appears to follow the spines seen in N₂H⁺.

3.6 DNC, N₂D⁺, and DC₃N

We did not detect N₂D⁺ and DC₃N emission at a noise level of ∼0.3 K with a velocity resolution of 0.05 km s⁻¹. Below we describe only the distribution of DNC. In figure 2i, we present a velocity-integrated intensity map of DNC. Relatively strong emission of DNC comes between FIR 5 and VLA 13 and the emission extends westward from VLA 13 (HOPS368).

4 Derivation of physical quantities

4.1 N₂H⁺ and CCS abundances

Fig. 3.

Averaged CCS line profile of the area that contains FIR 3/4/5 and VLA 13. We averaged the CCS integrated intensity in the N₂H⁺ clump with ∼0.17 K km s⁻¹ out. The thick solid line is the result of the Gaussian fitting, from which we obtained the peak intensity of 0.068 K, the central velocity of 11.14 km s⁻¹, and the line width of 1.3 km s⁻¹.

Table 4.

Physical parameters of CCS and N₂H⁺ toward FIR 3, FIR 4, FIR 5, and VLA 13.

Parameter	Area	FIR 3	FIR 4	FIR 5	VLA 13
NH2(× 1023 cm−2)	2.3	1.5	2.8	2.5	0.69
NCCS(× 1011 cm−2)	2.8	2.2	3.4	3.1	2.6
fCCS(× 10−12)	2.0	1.5	5.0	1.2	3.8
ΔV (km s−1)	1.3
τCCS	0.12
|${\rm N_{N_2H^+}(\times 10^{14}\:cm^{-2})}$|		5.8	2.9	2.3	1.2
|${\rm T_{ex, N_2H^+} (K)}$|		16.0 ± 0.9	21.8 ± 0.1	23.1 ± 0.3	26.2 ± 1.5
ΔV (km s−1)		1.4 ± 0.1	1.1 ± 0.1	1.1 ± 0.1	1.4 ± 0.1
|${\rm f_{N_2H^+}(\times 10^{-10})}$|		3.8	10.3	8.9	1.7
|${\rm \tau _{N_2H^+}}$|		0.8 ± 0.1	2.9 ± 0.1	1.9 ± 0.1	0.6 ± 0.1

Here, we estimate the abundances of CCS and N₂H⁺ at several positions. Since the CCS emission is weak, we derive only the average value in the small area that contains all four sources FIR 3/4/5 and VLA 13, indicated by the dashed lines in figure 2d. The averaged profile is shown in figure 3. The derived physical quantities are summarized in table 4.

We computed the CCS column densities, assuming that the emission was optically thin. The excitation temperature is adopted to be 5 K. The CCS fractional abundance does not vary significantly from position to position, and stays at about 10⁻¹². Here, we derived the column density of H₂ from the Herschel H₂ map. The rms noise level of the CCS data was about 0.036 K km s⁻¹ for an angular resolution of 32|${^{\prime\prime}_{.}}$|5, corresponding to ∼0.5 × 10¹¹ cm⁻². Therefore, the CCS integrated intensities at the positions of FIR 3/4/5 and VLA 13 varied from 4σ to 7σ. According to chemical evolution calculations by Suzuki et al. (1992) and Marka et al. (2012), such a small abundance can be achieved at the times of 10² yr and 10⁵ yr. According to Marka et al. (2012), the time evolution of the CCS abundance is not so different in the temperature range 10–25 K, whereas in much higher gas temperatures the chemical time scale may be shorter due to a faster chemical reaction (see, e.g., Shimoikura et al. 2018). In the FIR 4 region, the gas temperature is likely to be ∼20 K. Thus, the chemical timescales should be comparable to those in dark clouds with ∼10 K. Since protostars already formed in this region, the evolution time of this region is likely to be ∼10⁵ yr, which is comparable to the typical lifetime of the Class I protostars.

Fig. 4.

Hyperfine fit of N₂H⁺ line profiles at (a) FIR 3, (b) FIR 4, (c) FIR 5, and (d) VLA 13. The red curves show the results of the fitting. The frequencies of the seven hyperfine components are indicated by lines. The main component is indicated by a solid line in each panel.

For N₂H⁺, we applied the hyperfine fit to derive the physical parameters. The results of the hyperfine fits at FIR 3, FIR 4, FIR 5, and VLA 13 are shown in figure 4. We used the same fitting program as that of Tanaka et al. (2013), in which the excitation temperatures and velocity widths are assumed to be identical for the seven hyperfine components. For simplicity, we assumed that only one component exists along the line of sight, although this assumption does not seem to be appropriate at least for FIR 3 or VLA 13, both of which have at least two components. The fractional abundance of N₂H⁺ is relatively high in this region. It is the highest at FIR 4. The excitation temperature of N₂H⁺ is relatively high at ∼20 K. It is the strongest at the position of VLA 13. In the starless cores in the southern part of Orion A, the N₂H⁺ excitation temperatures are as low as 5–10 K (Tatematsu et al. 2014). In the Oph B2 region, where only four protostars are identified, the N₂H⁺ excitation temperatures stay as low as 5 K (Friesen et al. 2010). Other examples of the low excitation temperatures in starless regions can be found in Di Francesco, André, and Myers (2004) toward Oph A and Tanaka et al. (2013) toward Serpens South. These previous studies indicate that in the region where star formation is not active, the excitation temperatures of N₂H⁺ stay as low as 5–10 K. Thus, the high excitation temperature in the FIR 4 region may indicate the existence of protostars in this region. The line widths are about 1 km s⁻¹ for all of the positions, which is comparable to that of CCS.

4.2 The HN¹³C/H¹³CN ratio

Previous observations revealed that the [HNC]/[HCN] ratio varies from region to region. For example, in warm dense gas in OMC-1, the ratio measured much smaller than unity (Goldsmith et al. 1981, 1986). On the other hand, in cold dense gas, Tennekes et al. (2006) derived a somewhat large ratio of 3–4 toward Chamaeleon MM1. Colzi et al. (2018) showed that for high-mass starless cores, the [HNC]/[HCN] ratio stays at around unity, whereas it decreases for more-evolved objects, such as high-mass protostellar cores (∼0.5). Jin, Lee, and Kim (2015) found that the [HNC]/[NCN] ratio increases from infrared dark clouds to ultra-compact H iis. These studies indicate that the [HNC]/[HCN] ratio depends on the physical conditions. The [HNC]/[HCN] ratio may be a useful indicator to guess the physical states of star-forming regions. However, it still remains uncertain which mechanism is more responsible for determining the [HNC]/[HCN] ratio. Here, we examine the spatial distribution of the [HNC]/[HCN] ratio in the FIR 4 region to gain additional information to understand the HNC/HCN chemistry.

Hirota et al. (1998) proposed that the [HNC]/[HCN] ratio depends on the gas temperature, and rapidly decreases as the temperature exceeds the critical value of 24 K, above which a neutral–neutral reaction of HNC + H → HCN + H is responsible for destroying HNC (see also Schilke et al. 1992; Talbi et al. 1996). Colzi et al. (2018) discussed that the small ratio for high-mass protostellar cores may be due to the effect of a selective destruction of HNC by the reactions of HNC + H → HCN + H and HNC + O → NH + CO, suggesting the existence of a threshold temperature for the selective destruction of HNC, although future laboratory experiments are needed to estimate the reaction rate with O at low temperatures in order to understand the temperature dependence of the [HNC]/[HCN] ratio. Loison et al. (2014) pointed out that in the presence of carbon atoms in the gas phase, the C + HNC reaction also prevents the [HNC]/[HCN] ratio from reaching unity. On the other hand, Aguado et al. (2017) discussed the effect of far-ultraviolet (FUV) radiation, which can selectively photodissociate HNC, because the photodissociation cross-section is larger for HNC, leading to smaller [HNC]/[HCN] ratios for a stronger FUV. However, the parameter space has not been fully explored concerning the effects of the FUV radiation. Further investigations would be needed.

Fig. 5.

[HN¹³C]/[H¹³CN] ratio toward the FIR 4 region. In the present paper, we consider that the [HN¹³C]/[H¹³CN] ratio is identical that of [HNC]/[HCN].

In figure 5 we show the spatial distribution of the [HNC]/[HCN] ratio of the FIR 4 region. Here, we assumed that the [HNC]/[HCN] ratio is identical to the [HN¹³C]/[H¹³CN] ratio, and derived the column densities of HN¹³C and H¹³CN by assuming that both emissions are optically thin. In this case, the excitation temperature was set at 20 K to compute the column densities. We note that the column densities are insensitive to the assumed excitation temperature, and they vary only within errors of 30% in the range of from 10 K to 30 K. In the N₂H⁺ clump, the [HNC]/[HCN] ratio is estimated to be ∼0.5. This is smaller than the value at Cha-MMS 1, which has a ratio of 3–4, but not extremely small, such as in the OMC-1 hot core. It is somewhat higher in the envelope of the N₂H⁺ clump. The average ratio in the clump seems to be consistent with the values (∼0.5) in the high-mass protostellar objects derived by Colzi et al. (2018).

As we mentioned above, there are at least two effects discussed for reducing the [HNC]/[HCN] ratio (Hirota et al. 1998; Colzi et al. 2018; Aguado et al. 2017). One is that the gas temperature exceeds the critical temperature for HNC formation. In this case, the achieved [HNC]/[HCN] ratio tends to be very small (≲ 0.1) for higher temperatures (T ≫ 24 K). The average dust temperature of the N₂H⁺ clump is estimated to be about 26 K (Lombardi et al. 2014), which is somewhat larger than, or at least comparable to, the critical temperature of 24 K. The second is the FUV radiation, which can selectively photodissociate HNC, resulting in the smaller [HNC]/[HCN] ratio. Recently, Ishii et al. (2019) estimated the dimensionless FUV strength, G₀, from the Herschel 70 μm data; they found that the G₀ parameter of the FIR 4 region exceeds 100 and peaks at around FIR 3 (∼7000) and VLA 13 (∼1400). If the UV radiation is more responsible for determining the spatial distribution of the ratio, it should be lower at around VLA 13, similarly in the case at FIR 3. In contrast, our derived ratio is somewhat larger at VLA 13 than in the N₂H⁺ clump and FIR 3, and thus seems to be inconsistent with the effect of the FUV radiation. Although a further understanding of the HNC and HCN chemistry is needed, we tentatively conclude that the effect of temperatures comparable to the critical value (T > 24 K) is more responsible for the observed [HNC]/[HCN] ratio.

If the [HNC]/[HCN] ratio is a good indicator of the temperature in the dense regions, our [HNC]/[HCN] image might provide evidence of the outflow–dense gas interaction in this region. When the FIR 3 outflow interacts with the dense molecular gas in the FIR 4 region, dynamical compression heats up the gas. As a result, the [HNC]/[HCN] ratio in the FIR 4 clump may be somewhat reduced. At the positions of the outflow lobes of the FIR 3 and VLA 13 outflows [see the distribution of the outflow lobes in HCO⁺ (figure 8 in the next section) and ¹²CO (Takahashi et al. 2008)], the [HNC]/[HCN] ratio is smaller than the values in the surrounding parts. These features of the [HNC]/[HCN] ratio may be a sign of the outflow–dense gas interaction. In addition, the obtained [HNC]/[HCN] ratio in the FIR 4 clump might be inconsistent with the scenario that FIR 4 is a hot core (Ceccarelli et al. 2010; Kama et al. 2013), since the [HNC]/[HCN] ratio in the OMC-1 hot core is much smaller than the value in FIR 4 (Goldsmith et al. 1981, 1986).

5 Molecular outflows

Fig. 6.

¹³CO and HCO⁺ profiles toward (a) FIR 3, (b) VLA 13, and (c) FIR 4. The blue and red curves indicate the ¹³CO (J = 1–0) and HCO⁺ (J = 1–0) line profiles, respectively. We note that the effective angular resolution of the HCO⁺ profiles is 24|${^{\prime\prime}_{.}}$|9, the original resolution.

Fig. 7.

Distributions of high-velocity components of ¹³CO (J = 1–0). (a) Blueshifted and (b) redshifted components in the FIR 3/4/5 regions. The blueshifted and redshifted components are integrated over 7.5–9.5 km s⁻¹ and 12.5–14.5 km s⁻¹, respectively. (c) The ¹³CO high-velocity components overlaid on the N₂H⁺ integrated intensity image. The positions of the protostars listed in table 3 are indicated with green crosses. The red and blue contours start at 1.5 K km s⁻¹ and 0.6 K km s⁻¹, continuing in steps of 0.5 K km s⁻¹ and 0.6 K km s⁻¹, respectively.

Fig. 8.

Same as figure 7 but for HCO⁺. The velocity range for the integration is the same as that of ¹³CO. The maps taken with T70 are slightly wider than those of FOREST. Therefore, in panel (c), the blueshifted and redshifted emissions are distributed outside the N₂H⁺ map. We smoothed the HCO⁺ data to improve the signal-to-noise ratios of the blueshifted and redshifted components. The resultant effective angular resolution was about 34″. Both red and blue contours start at 0.4 K km s⁻¹ and progress in steps of 0.3 K km s⁻¹. The images are smoothed to improve the signal-to-noise ratios, and the effective resolutions are ∼30″.

In the observed area, seven protostars were previously identified, which are listed in table 3; five of them are classified as Class 0/I. Such protostars often drive powerful molecular outflows. Takahashi et al. (2008) performed an outflow survey in CO (J = 3–2) toward Orion A, and in this region, they detected three outflows: FIR 2 (HOPS68), FIR 3 (HOPS370), and VLA 13 (HOPS368). Here, we identify the outflows driven on the basis of the ¹³CO and HCO⁺ images.

In figure 6, we show the line profiles of ¹³CO and HCO⁺ toward the three representative positions: FIR 3, VLA 13, and FIR 4. Both the ¹³CO and HCO⁺ line profiles at the three positions have high-velocity wings, which may originate from molecular outflows. Figures 7a and 7b show the blueshifted and redshifted components of the ¹³CO emission, respectively. We also show the blueshifted and redshifted ¹³CO emissions overlaid on the N₂H⁺ integrated intensity image in figure 7c. For comparison, in figures 8a, 8b, and 8c, we present the same figures as figures 7a, 7b, and 7c, respectively, but for HCO⁺. The HCO⁺ emission is another good tracer of outflows in OMC-2/3 (Aso et al. 2000).

We adopt a systemic velocity of 11.3 km s⁻¹, which is measured in H¹³CO⁺ (J = 1–0) (Shimajiri et al. 2008). This velocity is in good agreement with the centroid velocity determined from CCS, 11.14 km s⁻¹. First, previous observations suggest that the FIR 3 (HOPS370) protostar is an outflow source, whose axis (NW–SE direction) is roughly on the plane of the sky (Takahashi et al. 2008; Shimajiri et al. 2008). Therefore, both the redshifted and blueshifted components are seen in ¹²CO (J = 1–0) (see figure 1 of Shimajiri et al. 2008). In our ¹³CO maps, the redshifted components (NW–SE direction) are more prominent and the blueshifted component is very vague. The redshifted components are also clearly seen in HCO⁺. The southern tip of the redshifted component appears to reach FIR 4 for both the ¹³CO and HCO⁺ maps, indicating the outflow–dense gas interaction. We note that the velocity range for the integration in ¹³CO is different from that of Shimajiri et al. (2008).

In the ¹³CO map, the redshifted component extending from FIR 3 in the east–west direction is prominent. This component can vaguely be recognized in the HCO⁺ map. It appears to be driven by HOPS 66. On the other hand, the blueshifted component is not seen. It might be overlapped by the FIR 3 component.

In the upper right-hand corner of the figure, we can see blueshifted and redshifted components from the FIR 2 (HOPS 68) outflow. This outflow had already been reported by Takahashi et al. (2008), who suggested that the outflow axis is along the north–south direction. In the ¹³CO maps, both redshifted and blueshifted components are seen in the southern part, and the redshifted component is somewhat stronger. In the HCO⁺ map, we can see a relatively strong redshifted component in the northern part of FIR 2, which is outside the mapping area of the FOREST observations (see figure 8). On the other hand, we do not find any clear blueshifted lobe in HCO⁺. Since both components in ¹³CO are seen in the southern part, the outflow axis may be nearly on the plane-of-sky.

Near FIR 4 and FIR 5, the strong redshifted component, which extends southward from FIR 4, can be seen in the HCO⁺ image. This component might be outflow from FIR 4 (HOPS 64). In fact, the compact blueshifted counterpart is distributed just north of this protostar, although this blueshifted component might be outflow from FIR 3. Takahashi et al. (2008) also detected broad wings toward FIR 4 in CO (J = 3–2). Our detected components in ¹³CO and HCO⁺ may be related to their detected wings.

Finally, we see high-velocity components near VLA 13 (HOPS368). A faint redshifted component is seen in the HCO⁺ emission in the southwestern direction of VLA 13. On the other hand, the blueshifted component is seen just north of VLA 13. The outflow axis appears to be somewhat tilted in the NE–SW direction in the plane-of-sky. In ¹³CO maps, we see both redshifted and blueshifted components in the western part. We note that the axis of the outflow seen in ¹²CO (J = 1–0) also appears to be slightly tilted in the NE–SW direction (see Tanabe et al. 2019), consistent with our interpretation. Since the dense gas is not distributed in the NE–SW direction, there is a possibility that the VLA 13 outflow might blow away the dense gas in this area, and dynamically compress the dense clump of FIR 4. The outflow identification in Orion A using ¹²CO (J = 1–0) is discussed in more detail by Tanabe et al. (2019). It is worth noting that the dense gas distribution anti-correlates with the distribution of any outflow component. However, this anti-correlation is not strong evidence of any outflow interaction.

In total, we confirmed four known outflow candidates, including marginal detection by Takahashi et al. (2008) and one new possible outflow candidate using the ¹³CO and HCO⁺ maps. The results of outflow identification are summarized in table 3. We note that the two marginally detected candidates might be just contamination of nearby outflow components from FIR 3 and VLA 13. To confirm that HOPS64 and HOPS108 are the driving sources of the molecular outflows, higher-angular resolution observations are needed. The scenario that the dynamical compression due to the outflows has triggered star cluster formation in this region (Shimajiri et al. 2008) is not inconsistent with the chemical timescale of 10⁵ yr expected from the CCS abundance. Star formation in FIR 4 may has been triggered by outflow compression. An alternative scenario is that star formation in FIR 4 occurred before the outflow hit the FIR 4 clump, which evolved into a hot core (Kama et al. 2013), and subsequent star formation in FIR 4, or fragmentation, may be triggered by the outflow compression (Shimajiri et al. 2008). Further chemical modeling may give us a hint as to the constraints of the evolution of FIR 4.

6 Summary

We summarize the main results of the present paper as follows.

1. We carried out mapping observations toward an outflow-shocked region, OMC-2, FIR 3/4/5 in ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), CCS (J_N = 8₇–7₆), HCO⁺ (J = 1–0), H¹³CO⁺ (J = 1–0), HN¹³C (J = 1–0), H¹³CN (J = 1–0), DNC (J = 1–0), N₂D⁺ (J = 1–0), and DC₃N (J = 9–8) using a new four-beam receiver, FOREST, and the T70 receiver. We detected every molecular line except N₂D⁺ and DC₃N.

2. We detected faint CCS emission from a dense clump containing FIR 4. This is the first detection of CCS emission in OMC-2. The typical fractional abundance is estimated to be about a few × 10⁻¹².

3. The [HNC]/[HCN] ratio is smaller than unity in the N₂H⁺ clump and the outflow lobes, whereas the [HNC]/[HCN] ratio is sometimes higher in the clump envelope.

4. From the N₂H⁺ hyperfine fit, we found that the excitation temperature in this region is high. The high excitation temperature might be due to the effect of radiation by protostars in this region.

5. We confirmed four known outflow candidates including marginal detection by Takahashi et al. (2008) and one new possible outflow candidate using ¹³CO and HCO⁺ (J = 1–0) emission. Previous studies suggest that the FIR 3 outflow hits the FIR4 clump (Shimajiri et al. 2008; González-García et al. 2016) and triggered star formation in this region (Shimajiri et al. 2008). In the present paper, we proposed that the protostellar outflow from VLA 13 as well as FIR 3 may have compressed the FIR 4 clump, taking into account the spatial distribution of blueshifted and redshifted ¹³CO and HCO⁺ components projected on the plane-of-sky. In addition, there is a possibility that the moderately small [HNC]/[HCN] ratios in the N₂H⁺ clump and the outflow lobes might be due to the dynamical interaction between the clump and the outflow and to the strong FUV radiation, both of which might release the atomic carbon that prevents the formation of HNC.

Acknowledgments

This work was carried out as one of the large projects of the Nobeyama Radio Observatory (NRO), which is a branch of the National Astronomical Observatory of Japan, National Institute of Natural Sciences. SPL and SJL acknowledge support from the Ministry of Science and Technology of Taiwan with Grant MOST 106-2119-M-007-021-MY3. We thank the NRO staff for both operating the 45 m and helping us with the data reduction. This work was financially supported by Grants-in- Aid for Scientific Research (Nos. 17H02863, 17H01118, 26287030). YS received support from the ANR (project NIKA2SKY, grant agreement ANR-15-CE31-0017).