Nobeyama 45 m mapping observations toward the nearby molecular clouds Orion A, Aquila Rift, and M17: Project overview

Abstract

We carried out mapping observations toward three nearby molecular clouds, Orion A, Aquila Rift, and M 17, using a new 100 GHz receiver, FOREST, on the Nobeyama 45 m telescope. We describe the details of the data obtained such as intensity calibration, data sensitivity, angular resolution, and velocity resolution. Each target contains at least one high-mass star-forming region. The target molecular lines were ¹²CO (J = 1–0), ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), and CCS (J_N = 8₇–7₆), with which we covered the density range of 10² cm⁻³ to 10⁶ cm⁻³ with an angular resolution of ∼20″ and a velocity resolution of ∼0.1 km s⁻¹. Assuming the representative distances of 414 pc, 436 pc, and 2.1 kpc, the maps of Orion A, Aquila Rift, and M17 cover most of the densest parts with areas of about 7 pc × 15 pc, 7 pc × 7 pc, and 36 pc × 18 pc, respectively. On the basis of the ¹³CO column density distribution, the total molecular masses are derived to be |$3.86 \times 10^{4}\, M_\odot$|⁠, |$2.67 \times 10^{4}\, M_{\odot }$|⁠, and |$8.1\times 10^{5}\, M_{\odot }$| for Orion A, Aquila Rift, and M17, respectively. For all the clouds, the H₂ column density exceeds the theoretical threshold for high-mass star formation of ≳ 1 g cm⁻² only toward the regions which contain current high-mass star-forming sites. For other areas, further mass accretion or dynamical compression would be necessary for future high-mass star formation. This is consistent with the current star formation activity. Using the ¹²CO data, we demonstrate that our data have enough capability to identify molecular outflows, and for the Aquila Rift we identify four new outflow candidates. The scientific results will be discussed in detail in separate papers.

1 Introduction

Star formation not only determines the observed properties of galaxies, but also significantly influences galaxy evolution (Mac Low & Klessen 2004; McKee & Ostriker 2007). What drives and regulates star formation in galaxies? There is little consensus to this apparently simple and fundamental question. Theoretical studies have demonstrated that once self-gravitating objects (cloud cores) form, their gravitational collapse leads to star formation at high rates (Klessen et al. 1998). However, star formation is known to occur at a very low rate in galaxies (Zuckerman & Evans 1974). For example, the total molecular gas mass in our Galaxy is estimated to be 10|$^{9}\, M_\odot$| from CO observations. If all the molecular clouds are converted to stars within a cloud free-fall time which is a few Myr at the typical cloud density of a few thousand cm⁻³, the free fall rate of star formation is calculated to be about 10|$^{3}\, M_\odot$| yr⁻¹, which is about 10³ times larger than the observed star formation rate (McKee & Williams 1997; Murray & Rahman 2010; Robitaille & Witney 2010). More accurate estimations of star formation rates toward individual molecular clouds are basically consistent with the above rough calculation (Krumholz & Tan 2007). The small Galactic star formation rate thus implies that some physical processes make star formation slow and inefficient.

However, it remains uncertain exactly what processes make star formation slow and inefficient. Several have been discussed as slowing down and regulating star formation, e.g., stellar feedback, magnetic field, and cloud turbulence (Shu et al. 1987; McKee & Ostriker 2007; Krumholz et al. 2014). It is thus crucial to characterize the internal cloud structure and physical properties of nearby molecular clouds to understand the roles of these processes in star formation. In our project, we carried out mapping observations toward three nearby molecular clouds using the Nobeyama 45 m telescope, and attempt to address the issues of inefficient star formation.

Star formation processes are often influenced by large-scale events. Protostellar jets and outflows are often extended to 0.1–10 pc-scale (Bally 2016). Expanding bubbles of 1–10 pc generated by stellar winds from intermediate-mass and high-mass stars have been discovered in nearby clouds (Arce et al. 2011; Feddersen et al. 2018; Pabst et al. 2019). Far-ultraviolet (FUV) radiation from high-mass stars sometimes affects the parent clouds at 10 pc scale (Shimajiri et al. 2011; Ishii et al. 2019). Much larger bubbles created by supernovae interact with entire molecular clouds (Frisch et al. 2015). Galactic spiral density waves influence molecular cloud formation and subsequent star formation (Elmegreen 1979). One of the immediate objectives of the present project is to reveal cloud structures and the dynamics of target clouds to attempt to elucidate how such events influence them. In summary, wide-field mapping observations are important to understand the effects of stellar feedback and external events like large-scale shocks because they potentially affect cloud properties and structures at a cloud scale, 1–10 pc. For comparison, we summarize several recent wide-field surveys toward our target clouds in table 1. For Orion A, many molecular-line mapping surveys have been done to date, but many of these surveys were observed only in the northern parts including OMC-1. For the Aquila Rift, many molecular-line mapping surveys have just focused on the two prominent star-forming regions, W|$\, 40$| and Serpens South. As for M17, only a few wide-field surveys that cover at least about a 1 deg² area have been done so far.

Table 1.

Selected recent wide-field surveys toward Orion A, Aquila Rift, and M|$\, 17.^*$|

Telescope/survey	Line/continuum	Resolution	Cloud/key reference
Osaka Pref. 1.85 m	12CO/13CO/C18O J = 2–1	|${2{^{\prime}_{.}}7}$|	Orion A/Nishimura et al. (2015)
Tsukuba 30 cm	12CO J = 4–3	|${9{^{\prime}_{.}}4}$|	Orion A/Ishii et al. (2016)
Harvard-CfA 1.2 m	12CO J = 1–0	|${8{^{\prime}_{.}}4}$|	Orion A/Wilson et al. (2005)
ASTE, NRO 45 m	1.1 m/12CO J = 1–0	36″/21″	Orion A/Shimajiri et al. (2011)
JCMT/GBS	13CO/C18O J = 3–2	17″	Orion A/Buckle et al. (2012)
FCRAO 14 m	12CO/13CO J = 1–0	46″	Orion A/Ripple et al. (2013)
IRAM 30 m	12CO/13CO J = 2–1	11″	Orion A/Berné, Marcelino, and Cernicharo (2014)
ASTE 10 m	12CO J = 3–2	30″	Orion A/Takahashi et al. (2008)
Herschel/HIFI IRAM 30 m	CH+/CO (J = 10–9)/HCN/HCO+ (J = 6–5), ...	12″–27″	Orion A (OMC-1)/Goicoechea et al. (2019)
Herschel-Planck	dust continuum	36″	Orion A/Lombardi et al. (2014)
Spitzer	MIR 3–24 μ m	2″–5″	Orion A/Megeath et al. (2012)
VISTA/VISION	NIR 0.85–2.4μ m	|${0{^{\prime \prime}_{.}}85}$|	Orion A/Meingast et al. (2016)
IN-SYNC	NIR 1.5–1.6 μ m	|${1{^{\prime \prime}_{.}}6}$|	Orion A/Da Rio et al. (2016)
NRO 45 m	N2H+J = 1–0	21″	Orion A/Tatematsu et al. (2008)
NRO 45 m	H13CO+J = 1–0	21″	Orion A/Ikeda, Sunada, and Kitamura (2007)
CARMA + NRO 45 m/CARMA-NRO Orion	12CO/13CO/C18O J = 1–0	8″	Orion A/Kong et al. (2018)
NRO 45 m	12CO/13CO/C18O/N2H+J = 1–0/CCS JN = 76–65	21″–24″	Orion A/this paper
Harvard-CfA 1.2 m	12CO J = 1–0	|${8{^{\prime}_{.}}4}$|	Aquila/Dame Hartmann, and Thaddeus (2001)
Osaka pref. 1.85 m	12CO/13CO/C18O J = 2–1	|${2{^{\prime}_{.}}7}$|	Aquila/Nakamura et al. (2017)
Herschel	dust	36″	Aquila/André et al. (2010)
IRAM 30 m/MAMBO	1.2 mm	11″	Aquila/Maury et al. (2011)
IRAM 30 m	HCN/H13CN/HCO+/H13CO+J = 1–0	40″	Aquila/Shimajiri et al. (2017)
JCMT/GBS	850μ m 12CO J = 3–2	15″/22″	Aquila (W|$\, 40$|⁠)/Rumble et al. (2016)
ASTE 10 m	12CO J = 3–2/HCO+J = 4–3	31″	Aquila (W|$\, 40$|⁠)/Shimoikura et al. (2015)
ASTE 10 m	12CO J = 3–2/HCO+J = 4–3	24″	Aquila (Serpens South)/Nakamura et al. (2011b)
Spitzer	IRAC	2″	Aquila (Serpens South)/Gutermuth et al. (2008)
NRO 45 m	N2H+J = 1–0	24″	Aquila (Serpens South)/Tanaka et al. (2013)
MOPRA	N2H+/H13CN/HCN/HNC/HCO+/H13CO+J = 1–0	40″	Aquila (Serpens South)/Kirk et al. (2013)
NRO 45 m	CCS JN = 43–32/HC3N J = 5–4	37″	Aquila (Serpens South)/Nakamura et al. (2014)
NRO 45 m	12CO/13CO/C18O/N2H+J = 1–0/CCS JN = 76–65	21″–24″	Aquila/this paper
Spitzer	MIR	2″	M17/Povich & Whitney (2010); Povich et al. (2016)
HHT 10 m	12CO/13CO J = 2–1	32″	M17/Povich et al. (2009)
NRO 45 m	12CO/13CO/C18O J = 1–0	20″	M17/Nishimura et al. (2018)
NRO 45 m	12CO/13CO/C18O/N2H+J = 1–0/CCS JN = 76–65	21″–24″	M17/this paper

^∗This is not a complete list of the recent wide-field survey. See also table 1 of Kong et al. (2018) for Orion A.

The main objectives of the present paper are to give a project overview and to describe the details of the observations and calibration of the obtained data. In section 2 we describe how we select our target clouds and lines. The details of the observations are described in section 3. In section 4 we describe how to produce the final data cubes. Then we summarize the quality of our maps in terms of sensitivity, angular resolution, and velocity resolution. In section 5 we present the spatial distributions of the molecular line emissions obtained, and in section 6 we derive the spatial distribution of the ¹³CO and C¹⁸O abundances toward Orion A. We briefly discuss the global molecular gas distributions of the target clouds mainly using ¹²CO and its isotopologues. In section 7 we show the preliminary results of a CO outflow survey toward L 1641 N in Orion A and Serpens South in the Aquila Rift, demonstrating that our data can be used to detect outflows. Finally, we summarize the main results in section 8.

The detailed analysis will be presented in separate papers. For example, Tanabe et al. (2019), Ishii et al. (2019), and Nakamura et al. (2019) describe the results of the outflow survey, cloud structure analysis, and multi-line observations of the OMC-2 FIR 4 region for Orion A, respectively. Shimoikura et al. (2019b) and Kusune et al. (2019) present the detailed cloud kinematic structure and the relationship between filamentary structure and magnetic field, respectively, toward Aquila Rift. For M17, Nguyen Luong et al. (2019), Shimoikura et al. (2019a), and Sugitani et al. (2019) will report the global cloud kinematics, dense core survey, and relationship between cloud structure and magnetic field, respectively. In addition to the three main regions, a few other star-forming regions such as the Northern Coal Sack (Dobashi et al. 2019a) and DR 21 (Dobashi et al. 2019b) were studied during our project. These data were obtained for intensity calibration of the CO lines taken with a new 100 GHz receiver, FOREST.

2 Project overview

2.1 Molecular line data with a 10⁻² pc resolution within distances of a few kpc

In 2011, the state-of-art facility ALMA started science operations. Since then, we can easily conduct observations with much higher sensitivity and higher angular resolution than ever done before. In the ALMA era, one of the important objectives of our project is to make useful datasets to compare with ALMA maps of distant molecular clouds. Even using one of the largest millimeter telescopes, the Nobeyama 45 m telescope, the achieved beam size at 115 GHz is at most ∼15″. On the other hand, we can easily achieve ∼1″ resolution using ALMA. To overcome the disadvantage of the coarse angular resolution, we choose nearby molecular clouds as our targets for the Nobeyama 45 m telescope. Large-scale ALMA mapping observations toward nearby molecular clouds are still limited and are impossible to do because of its small field of view and extremely long observation time. On the other hand, the coarse beams of single-dish telescopes allow us to conduct wide-field mapping observations. In this sense, our Nobeyama mapping project is complementary to the ALMA observations toward distant molecular clouds.

Previous observational studies have often been influenced by the effects of different spatial resolutions when cloud structure and physical properties of molecular clouds are compared. We would like to minimize the effects of the different spatial resolution to better understand the cloud structure and the environmental effects. In our project we therefore aim to obtain maps that can be compared directly with those obtained with ALMA in almost the same spatial resolution.

The number of pointings of the ALMA mosaic observations is limited to below 150 for the 12 m array, which can cover an area of ≲5′ × 5′ with an angular resolution of 1″ in the ALMA band-3 (∼100 GHz). Assuming that we map about 1 deg² area with an angular resolution of 20″ toward regions whose distances are 20 times closer than the areas observed by ALMA with 1″ resolution, both the maps obtained have comparable spatial coverage and spatial resolution. As described below, we thus chose Orion A, Aquila Rift, and M17 as our targets for the Nobeyama 45 m mapping observations. In table 2 we compare the achieved spatial resolutions for several molecular clouds, mainly including our target regions. In our previous Nobeyama 45 m observations we have mapped the nearest molecular clouds (∼140 pc) such as L|$\, 1551$| (Yoshida et al. 2010; Lin et al. 2017) and the ρ Ophiuchus cloud (Maruta et al. 2010; Nakamura et al. 2011a) using the Nobeyama 45 m telescope. These maps have a spatial resolution of about 2000–3000 au, which is comparable to those of the maps of Orion A (414 pc) combined with the CARMA data (Kong et al. 2018). For more distant clouds with distances of a few kpc, ALMA observations can achieve a spatial resolution of a few thousand au at an angular resolution of ∼1″. Thus, our Nobeyama maps of nearby clouds can be directly compared with the maps of the molecular clouds within a few kpc in a comparable spatial resolution of a few thousand au.

Table 2.

Spatial resolutions achieved for the Nobeyama 45 m, CARMA, and ALMA observations.^∗

Telescopes	Angular	Nearest regions	Intermediate regions	Distant regions
	resolution	∼140 pc	∼400 pc	∼3 kpc
		(Taurus, ρ Oph)	(Orion A, Aquila Rift, California)	(M 17, several IRDCs†)
NRO 45 m only	∼20″	2800 au (0.014 pc)	8400 au (0.04 pc)	40000 au (0.2 pc)
CARMA + NRO 45 m	∼8″	—	3300 au	—
ALMA	∼1″	140 au	420 au	3000 au

*The distances are adopted from the following references: nearest regions—Taurus (137 pc; Torres et al. 2007; Loinard et al. 2007), ρ Oph (137 pc; Ortiz-León et al. 2017a); intermediate regions—Orion A (414 pc; Menten et al. 2007; Kim et al. 2008), Aquila Rift (436 pc, Ortiz-León et al. 2017b); distant regions—M17 (2.0 kpc; Xu et al. 2011). The star-forming regions indicated by boldface are the target clouds of the present project. We note that for the M17 region, several clouds with different distances seem to be overlapped along the line of sight (Povich et al. 2016). Here we refer to the ALMA observations just for comparison of the spatial resolutions achieved. In the present paper, we present the data obtained with the Nobeyama 45 m telescope. The detail of the CARMA |$+$| Nobeyama combined data is presented in Kong et al. (2018).

^†Infrared dark clouds.

We have opened all the mapping data we took in this project to the public. The data are available through the Japanese Virtual Observatory (JVO) archival system.¹

2.2 Target lines

We chose the following five molecular lines: ¹²CO (J = 1–0), ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), and CCS (J_N = 8₇–7₆). The ¹²CO molecule is the second most abundant after the hydrogen molecule in molecular clouds, and therefore it can trace the basic cloud structure reasonably well. Rarer isotopologues of ¹²CO such as ¹³CO and C¹⁸O can trace relatively denser gas with densities of 10³–10⁴ cm⁻³. N₂H⁺ has a critical density of ∼10⁵ cm⁻³ and traces cold dense gas, particularly in the prestellar phase, very well (Caselli et al. 2002; Punanova et al. 2016). Thus, we can cover the densities from 10² cm⁻³ to 10⁶ cm⁻³ using these four lines. The main reason for choosing CCS is that it can be obtained simultaneously with ¹³CO (J = 1–0), C¹⁸O (J = 1–0), and N₂H⁺ (J = 1–0), as we mention below. CCS is abundant in the early prestellar phase and traces a similar density range to N₂H⁺. Thus, CCS can provide additional information on the chemical evolution of prestellar dense gas (Suzuki et al. 1992; Marka et al. 2012; Loison et al. 2014; Shimoikura et al. 2018).

Another reason why we mainly chose the CO lines is that the main beam efficiencies of the 45 m telescope in the 100 GHz band (≲ 0.5) are not as superb as those of other telescopes such as the IRAM 30 m and LMT telescopes (see table 3) and thus observations of strong emission lines such as ¹²CO and its isotopologues can be done more efficiently within a limited observation time. Although N₂H⁺ is not as strong as ¹²CO, ¹³CO, or C¹⁸O, a new multi-beam superconductor–insulator–superconductor (SIS) receiver, FOREST, allows us to observe N₂H⁺ line simultaneously with ¹³CO and C¹⁸O. In addition, CCS (J_N = 8₇–7₆) can also be obtained simultaneously with N₂H^{+, 13}CO, and C¹⁸O using a spectral window mode that has become available since the 2016–2017 season.

We also adopted a velocity resolution of ∼0.1 km s⁻¹, so that we can reasonably identify dense cores with a typical internal velocity dispersion of a few tenths km s⁻¹. Our observed emission lines are summarized in table 3.

Table 3.

Observed lines.^∗

Molecule	Transition	Rest frequency	Beam size	ΔV	Main beam efficiency (η)
		(GHz)	(⁠|$^{\prime\prime}$|⁠)	(km s−1)
12CO	J = 1–0	115.271204	14.3 ± 0.4	0.1	41.6 ± 3%
13CO	J = 1–0	110.201354	14.9 ± 0.4	0.1	43.5 ± 3%
C18O	J = 1–0	109.782176	14.9 ± 0.4	0.1	43.7 ± 3%
N2H+	J = 1–0	93.1737637	17 ± 0.5	0.1	50.0 ± 5%
CCS	JN = 87–76	93.870098	17 ± 0.5	0.1	49.7 ± 5%

*The rest frequency of the main hyperfine component of N₂H⁺ is adopted from Pagani, Daniel, and Dubernet (2009). The beam size and main beam efficiencies listed are those measured with FOREST. The errors in the efficiencies are mainly due to the uncertainty of the brightness temperature of the planets used for the measurements. See the NRO web page for more details.

2.3 Target clouds

We chose the target clouds taking into account the following conditions:

1. The main part of the target cloud can be covered by 1 deg² with a few hundred hours using the new 100 GHz four-beam SIS receiver FOREST, including overhead such as pointing observations and flux calibration.

2. The target clouds are relatively well studied, and additional datasets are available.

3. Each target cloud contains regions of ongoing high-mass star formation.

4. The spatial resolution of the Nobeyama 45 m data can be comparable to or smaller than the typical size of dense cores (∼0.1 pc).

Taking the above conditions into account, we chose the following three regions as our targets: (1) Orion A (d ∼ 414 pc; Menten et al. 2007; Kim et al. 2008), (2) Aquila Rift (d ∼ 436 pc; Ortiz-León et al. 2017b), and (3) M17 (d ∼ 2000 pc; Xu et al. 2011).

In this survey we chose nearby clouds that contain the formation sites of O-type and early B-type stars. There are not so many such regions within ∼1 kpc (e.g., Orion A, Aquila Rift, California, Mon R2, and Orion B). Among them, Orion A and the Aquila Rift may be the nearest. The evolutionary stages of the star-forming regions appear to be different in the Aquila Rift and Orion A. For example, star formation lasts at least a few Myr in OMC-1 (Da Rio et al. 2016). On the other hand, in the Serpens South cluster in the Aquila Rift, the fraction of Class O/I protostars is extremely high, suggesting an age within 0.5 Myr. In other words, Orion A contains more evolved star-forming regions than the Aquila Rift. The distances of Orion A and the Aquila Rift are similar (∼400 pc), and thus it is easy to directly compare the cloud structures with the same spatial resolutions for molecular clouds in different evolutionary stages. These are the main reason why we chose these two clouds. The star formation efficiencies of Orion A and the Aquila Rift appear to be similar within a few percent of each other (Evans 2009; Maury et al. 2011), which is typical of nearby star-forming regions (Krumholz & Tan 2007). Thus, we believe that the two regions are representative of nearby star-forming molecular clouds.

The beam size of the Nobeyama 45 m telescope at 100 GHz is ∼15″, corresponding to ∼6200 au and 30000 au at the distances of 414 pc and 2000 pc, respectively. The spatial resolution of M17 is not satisfied with the last condition, but we chose it for the following two reasons.

• The densest part of an infrared dark cloud in M17, M17 SWex, was observed by ALMA, and the N₂H⁺ data are available (see, e.g., Ohashi et al. 2016; Chen et al. 2019), and we can combine the 45 m data with the ALMA data to fill the zero spacing so that we can make maps whose spatial resolution is comparable to the other two targets obtained with the Nobeyama 45 m telescope (∼8000 au).

• Star formation activity in M17 appears to be triggered by a Galactic spiral wave passage (Elmegreen 1979) and thus the M17 data are expected to provide us with a clue to better understand the Galactic star formation process. In addition, M17 is the nearest high-mass star-forming region to the Sagittarius arm. The stellar density at NGC 6811 is the highest at >10³ star pc⁻² after the Carina cluster among the regions in the MYStIX survey of massive star-forming regions in X-ray (Kuhn et al. 2015). M17 is closer to us than the Carina region. Povich and Whitney (2010) proposed that high mass star formation may be delayed in an IRDC region and the mass function of young stars in IRDCs appears to be different from the Salpeter initial mass function. M17 is expected to provide us with key information to understand how high-mass stars form. Therefore, we chose the distant high-mass star-forming region M17.

For Orion A, we combine our Nobeyama 45 m data with the CARMA interferometric data to obtain maps with ∼8″ resolution. The combined maps have spatial resolution comparable to those of the Taurus and ρ Ophiuchus molecular clouds previously obtained with the Nobeyama 45 m telescope, with a spatial resolution of about 3000 au (∼20″) at a distance of 140 pc. Thus, using the Nobeyama data and ALMA data, we can directly compare the internal structure and physical properties of several clouds located at different distances in the range from 140 pc to a few kpc with the same spatial resolutions (see table 2). We expect that the maps obtained will be useful as templates for nearby molecular clouds that can be compared with maps toward more distant molecular clouds, which can be obtained with ALMA, e.g., IRDCs at a few kpc such as the Nessie nebula and G11.11−0.12 (Snake).

We determine the mapping areas by reference to the 2MASS A_V maps (Dobashi et al. 2005; Dobashi 2011), Herschel column density maps, and Spitzer images. Taking the observation time into account, we planned to map the target areas with A_V ≳ a few mag corresponding to a few tenths g cm⁻², so that we can detect almost all the self-gravitating cores in our target clouds—see the A_V threshold derived by André et al. (2010).

Figures 1, 2, and 3 present the mapping areas of Orion A, Aquila Rift, and M17, respectively, overlaid on the Herschel or Spitzer images.

Fig. 1.

Observation areas overlaid on the H₂ column density maps of Orion A. The solid and dashed lines indicate the observation boxes for the ¹²CO+¹³CO set and the C¹⁸O+N₂H⁺+CCS set, respectively. The H₂ column density map has an effective angular resolution of about ∼36″ (see Lombardi et al. 2014). (Color online)

Fig. 2.

Same as figure 1 but for the Aquila Rift. The color image shows the H₂ column density map whose fitting data were downloaded via the Herschel Gould Belt Survey Archive system. The solid and dashed lines are the same as those in figure 1. (Color online)

Fig. 3.

Observation areas overlaid on the Spitzer 8|$\, \mu$|m image of M17. The color image shows the Spitzer 8|$\, \mu$|m image of M17 downloaded from the Glimpse Archival system. The solid and dashed lines are the same as those in figure 1. (Color online)

3 Observations

As described in the next section, we used the molecular line data taken with BEARS for intensity calibration. In this section we first describe the details of the FOREST observations, and then describe the details of the BEARS observations. BEARS is an SIS 25-element focal plane receiver with a frequency coverage of 82–116 GHz, and has been used for many mapping observations since 2000. Therefore, the intensity calibration scheme is well established compared to the new receiver, FOREST, which was also demounted for repair after the first season and reinstalled in the second season. In addition, the surface accuracy of the telescope dish was significantly improved for the first two seasons by applying a holograph to the dish surface. Careful data reduction of the data obtained is crucial to verify the absolute intensity scale of the FOREST data. Thus, we used the data taken with BEARS for intensity calibration. For both observations we adopted the on-the-fly (OTF) mapping mode with a position-switching method using the emission-free positions areas summarized in table 6.

For FOREST, we adopted two frequency sets to observe the target lines. Set 1 is for ¹²CO and ¹³CO, while set 2 is for C¹⁸O, N₂H^{+, 13}CO, and CCS with a spectral window mode. For BEARS, we observed only a single line for the individual observations.

3.1 FOREST

FOREST is a four-beam dual-polarization sideband-separating SIS receiver (Minamidani et al. 2016) and has 16 intermediate frequency (IF) outputs in total; 8 IFs in the upper-sideband and the other 8 IFs in the lower-sideband were used for the molecular line observations. As backends, we used a digital spectrometer based on an FX-type correlator, SAM45, that has 16 sets of 4096 channel arrays. We divided the mapping area into smaller sub-areas whose sizes are summarized in table 4. Then, we carried out OTF observations toward each sub-area. The parameters of the OTF observations are summarized in table 5. Scans of the OTF observations were separated in intervals of |${5{^{\prime \prime }_{.}}17}$|⁠. Thus, individual scans of the four beams of FOREST are overlapped. The positions of the emission-free areas used for the observations are summarized in table 6. We note that the coordinates of the emission-free areas are those of the first FOREST beam.

Table 4.

Sizes of observation boxes for the FOREST observations.^∗

Lines	Scan	12CO/13CO	C18O/	13CO/C18O/
			N2H+	N2H+/CCS
Season		2014–2015/2015–2016/2016–2017	2015–2016	2016–2017
Orion A	x	10′ × 5′	20′ × 5′	20′ × 5′
Orion A	y	5′ × 10′	5′ × 20′	5′ × 20′
Aquila Rift	x	20′ × 10′	—	20′ × 10′
Aquila Rift	y	10′ × 20′	—	—
M17	x	20′ × 10′	20′ × 10′	20′ × 10′
M17	y	10′ × 20′	10′ × 20′	10′ × 20′

*Taking into account the observation schedule of each season, we changed the observation box sizes. For Aquila Rift, we could not obtain y-scan data of C¹⁸O, N₂H⁺, or CCS. In addition, for M17, we significantly reduced the observation area for the C¹⁸O, N₂H⁺, and CCS observations. These incomplete observations are mainly due to a malfunction of the master collimeter driving system that happened in the 2016–2017 season.

Table 5.

Parameters of observations with FOREST.^∗

Box size	10′ × 5′	20′ × 5′	20′ × 10′
Time for scan (s)	15	30	26
Number of ONs per OFF	4	2	3
Separation between scans	|${5{^{\prime \prime }_{.}}17}$|	|${5{^{\prime \prime }_{.}}17}$|	|${5{^{\prime \prime }_{.}}17}$|
Frequency resolution (kHz)	15.26	15.26	15.26

*For all observations, scans of the OTF observations are separated in intervals of |${5{^{\prime \prime }_{.}}17}$|⁠. Thus, individual scans by the four beams of FOREST are completely overlapped. This greatly minimizes the effort needed for flux calibration.

Table 6.

Map coverage and positions of emission-free areas used for the observations and SiO maser objects used for the pointing observations.^∗

Regions	Map coverage	Emission-free areas		Pointing objects (SiO maser line)
		RA (J2000.0)	Dec (J2000.0)	RA (J2000.0)	Dec (J2000.0)
Orion A	|${0{^{\circ}_{.}}7}$| × 2|$^{\circ}$|	|${05}^{\rm h}29^{\rm m}00{^{\rm s}_{.}0}$|	|${-05}^{\circ}25^{\prime}30{^{\prime \prime}_{.}}0$|	|${05^{\rm h}35^{\rm m}14{^{\rm s}_{.}}16}$|	|${-05}^{\circ}22^{\prime}21{^{\prime \prime}_{.}}5$|	Orion KL
Aquila Rift	1|$^{\circ}$| × 1|$^{\circ}$|	|${18^{\rm h}41^{\rm m}19{^{\rm s}_{.}}09}$|	|${+04}{^\circ}12^{\prime}00{^{\prime\prime}_{.}}0$|	|${18}^{\rm {h}}37^{\rm {m}}19{^{\rm {s}}_{.}}26$|	|${+10}{^\circ}25^{\prime}42{^{\prime\prime}_{.}}2$|	V1111-Oph
M17	2|$^{\circ}$| × |${0{^{\circ}_{.}}5}$|	|${18^{\rm h}37^{\rm m}19{^{\rm s}_{.}}26}$|	|${+10}{^\circ }25^{\prime}42{^{\prime\prime}_{.}}2$|	|${18}^{\rm {h}}37^{\rm {m}}19{^{\rm {s}}_{.}}26$|	|${+10}{^\circ }25^{\prime}42{^{\prime\prime}_{.}}2$|	V1111-Oph

^∗These areas and objects were used for both the BEARS and FOREST observations.

Calibration of the observations was done by the chopper wheel technique to convert the output signal into the antenna temperature |$T_{\rm A}^*$|⁠, corrected for atmospheric attenuation. Some details of the observations are also summarized in table 7. The telescope pointing was checked every 1 hr by observing the SiO maser lines from the objects presented in table 6. The pointing accuracy was better than ∼3″ throughout the entire observation.

Table 7.

Summary of the FOREST observations.

Lines	Period	Observation time	T  sys  ∗	Velocity resolution	Noise level
Orion A
12CO	2014 Dec–2016 Dec	150 hr	350–400 K	0.1 km s−1	0.50–1.5 K
13CO	2014 Dec–2017 Mar	150 hr	150–200 K	0.1 km s−1	0.20–0.30 K
C18O	2017 Jan–2017 Mar	150 hr	150–200 K	0.1 km s−1	0.26–0.30 K
CCS	2017 Jan–2017 Mar	150 hr	150–200 K	0.1 km s−1	0.30–0.49 K
N2H+	2017 Jan–2017 Mar	150 hr	150–200 K	0.1 km s−1	0.26–0.30 K
Aquila Rift
12CO	2015 Mar–2017 Mar	50 hr	300–500 K	0.1 km s−1	0.38–0.50 K
13CO	2015 Mar–2017 Mar	150 hr	150–200 K	0.1 km s−1	0.38–0.50 K
C18O	2016 Mar–2017 Mar	120 hr	150–200 K	0.1 km s−1	0.20–0.30 K
CCS	2016 Mar–2017 Mar	120 hr	150–200 K	0.1 km s−1	0.17–0.20 K
N2H+	2016 Mar–2017 Mar	120 hr	150–200 K	0.1 km s−1	0.18–0.22 K
M17
12CO	2015 Mar–2017 Mar	25 hr	300–500 K	0.1 km s−1	0.48–1.7 K
13CO	2015 Mar–2017 Mar	150 hr	150–200 K	0.1 km s−1	0.16–0.80 K
C18O	2016 Mar–2017 Mar	65 hr	150–200 K	0.1 km s−1	0.17–0.34 K
CCS	2016 Mar–2017 Mar	65 hr	150–200 K	0.1 km s−1	0.20–0.27 K
N2H+	2016 Mar–2017 Mar	65 hr	150–200 K	0.1 km s−1	0.15–0.27 K

^∗The values of T_sys are given in the single sideband.

In order to minimize the scanning effects, data with orthogonal scanning directions along the RA and Dec. axes were combined into a single map. We adopted the same gridding convolution function (spheroidal function) as the BEARS data to calculate the intensity at each grid point of the final cube data with a spatial grid size of |${7{^{\prime \prime}_{.}}5}$|⁠.

The FOREST receiver and telescope conditions were not stable for the first two seasons. For example, the receiver was demounted in 2016 June (right after the first season) to improve several internal components. Second, the surface accuracy of the 45 m dish was significantly improved for the first two seasons by applying a holograph to the dish surface. Therefore, the observation conditions were different from season to season. To minimize the effects of these factors in determining the intensity scale of the observed lines, we scaled the data obtained with FOREST by those of BEARS whose intensity calibration method is well established. The detail of the flux calibration of the FOREST data is described in the next section.

In the last season (2016–2017), a new observational capability called a spectral window mode was available, which allows us to obtain more lines simultaneously. We therefore observed ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), and CCS (J_N = 7₆–6₅) simultaneously. In the spectral window mode, we equally divided the whole bandwidth into two, so that the bandwidth and frequency resolution of the spectrometer arrays were set to 31.25 MHz and 15.26 kHz, respectively.

3.2 BEARS

The procedure for the BEARS observations was basically similar to that of FOREST. The details of the BEARS observations are summarized in table 8. Some of the results of the BEARS observations are described in the references listed in the last column of table 8. In brief, we divided the mapping area into many 15′ × 15′ or 20′ × 20′ rectangle sub-areas. The sizes of these sub-areas were determined so as to complete an OTF scan within 1 or 1.5 hr. We carried out mapping observations toward each sub-region in OTF mode (Sawada et al. 2008) using BEARS and 25 sets of 1024 channel auto correlators (ACs) which have a bandwidth of 32 MHz and frequency resolution of 37.8 kHz. The velocity resolutions of the observations for individual lines are listed in table 8.

Table 8.

Summary of the BEARS observations.

Lines	Period	Observation	T  sys  ∗	Velocity resolution	Noise	Observation	Ref.†
		time			level	mode
Orion A
12CO (north)	2007 Dec–2008 May	∼40 hr	250–500 K	0.2 km s−1	0.4 K	OTF	1
12CO (south)	2009 Dec–2010 Jan	∼20 hr	300–600 K	0.5 km s−1	0.52 K	OTF	2
13CO (north)	2013 May	∼50 hr	270–470 K	0.3 km s−1	0.7 K	OTF	3
13CO (south)	2012 Apr–2013 Mar	∼60 hr	210–400 K	0.1 km s−1	1.96 K	OTF	4
C18O (north)	2010 Mar–2013 May	∼100 hr	270–470 K	0.1 km s−1	0.3 K	OTF	3
Aquila Rift (Serpens South)
12CO	2011 Apr–2011 May	∼15 hr	250–500 K	0.5 km s−1	1.3 K	OTF	4
13CO	2011 Apr–2011 May	∼30 hr	210–400 K	0.5 km s−1	0.88 K	OTF	4
C18O	2011 Apr–2014 Apr	∼30 hr	200–400 K	0.1 km s−1	0.9 K	OTF	4

^∗The values of T_sys are given in the double sideband.

^†References. 1: Shimajiri et al. (2011); 2: Nakamura et al. (2012); 3: Shimajiri et al. (2014); 4: This paper.

Calibration of the observations was done by the chopper wheel technique to convert the output signal into the antenna temperature |$T_{\rm A}^*$|⁠, corrected for atmospheric attenuation. At 110 GHz, the half-power beam width was about 15″, which corresponds to about 0.03 pc at a distance of 414 pc. The main beam efficiency (⁠|$\eta _{\rm 45\, m}$|⁠) was about 0.5 at 110 GHz for the corresponding observation season. The telescope pointing was checked every 1 or 1.5 hr by observing the SiO maser sources Orion KL and IRC|$\, +$|00363 for the Orion A and Serpens South observations, respectively, and was better than 3″ during the whole observing period.

We obtained a map by combining scans along two axes that run at right angles to each other. We adopted a convolutional scheme with a spheroidal function to calculate the intensity at each grid point of the final cube data with a grid size of |${7{^{\prime \prime}_{.}}5}$|⁠. This convolutional scheme is the same as that of the FOREST data. The spheroidal function is given by Sawada et al. (2008). The resulting effective resolution was about 21″ at 110 GHz. Finally, we converted the intensities in the antenna temperature scale (⁠|$T_{\rm A}^*$|⁠) into those in the brightness temperature scale (T_mb) by dividing by the main beam efficiencies (η), |$T_{\rm mb} =T_{\rm A}^* /\eta$|⁠. The typical rms noise levels of the final maps are listed in table 8. We note that for Orion A, the coverages of the ¹²CO and ¹³CO maps taken with BEARS are from Dec ∼−5°20′ to ∼−6°30′, which is slightly smaller than the actual mapping area of the FOREST observations. The coverage of the C¹⁸O map was much smaller than those of ¹²CO and ¹³CO. For the other lines (N₂H⁺ and CCS), no OTF data were available. The typical noise levels achieved for the BEARS observations were ∼0.4 K and 0.7 K at 0.1 km s⁻¹ for ¹²CO and ¹³CO, respectively.

4 Flux calibration

The procedure of flux calibration is complicated, as we describe in this section. First, we briefly summarize our flux calibration procedure, and then we describe the details of the procedure in the subsequent subsections.

• 0. All the data were converted to intensity data in |$T_{\rm A}^*$| using the standard NRO data reduction tool, NOSTAR.

• 1. For the ¹²CO and ¹³CO observations of Orion A, the intensities of each emission line and each sub-box were multiplied by the scaling factor |$SF_i^{\rm BEARS, \it l}$|⁠, where the scaling factor |$SF_i^{\rm BEARS, \it l}$| is the ratio of the integrated intensity of the FOREST observations to that obtained with BEARS for the corresponding sub-box i; the superscript l indicates the corresponding line (¹²CO or ¹³CO). For the sub-boxes where BEARS data were not available, the average scaling factor of the corresponding season is adopted. Note that the BEARS data are in the T_mb scale.

• 2. For the other data (in |$T_{\rm A}^*$|⁠), the intensities corrected for the daily variation were obtained by multiplying the intensity data of the corresponding sub-box by the scaling factor of the daily variation SF_ref, where the scaling factor SF_ref is the ratio of the sum of the integrated intensity of the reference area obtained at the the corresponding observation session (⁠|$S_{\rm ref}=\Sigma _i I_{i, \rm ref}$|⁠) to that obtained at the reference date (⁠|$S_{\rm ref}^*= \Sigma _i I_{i, \rm ref}^*$|⁠, |$SF_{\rm ref}=S_{\rm ref}^*/S_{\rm ref}$|⁠). Note that we observed a small reference area at least once during each observation session.

• 3. Then, the intensities corrected for daily variation were divided by the beam efficiency at the corresponding frequency and converted to intensities in the T_mb scale.

We describe the details of the procedure below.

4.1 Orion A

4.1.1 ¹²CO and ¹³CO

The ¹²CO and ¹³CO maps obtained with BEARS cover the northern part of the FOREST observation areas shown in figure 1. The intensity scales of the ¹²CO and ¹³CO data were thus calibrated by comparing with the previous survey data taken by the BEARS receiver (Shimajiri et al. 2011, 2014; Nakamura et al. 2012), which were already corrected into the main beam temperature (T_mb) scale. We obtained the intensities by dividing the FOREST intensities obtained in |$T_{\rm A}^*$| by the scaling factor SF_i determined in each sub-box and each line using the equation

$$\begin{equation} {\rm SF}_i^{\rm BEARS,\it l} = { \displaystyle \sum _{k=1}^{N_i} {W_{k, \rm BEARS,\it l} / W_{k, \rm FOREST,\it l}} \over N_i} , \end{equation}$$

(1)

where |$W_{k, \rm FOREST, \it l}$| is the integrated intensity of a given box in the antenna temperature scale at the kth grid, and |$W_{k, \rm BEARS, \it l}$| is the integrated intensity of the same box in the brightness temperature scale at the kth grid; l is the corresponding line, and N_i is the number of grids in the ith box. We measured the intensity ratio at every grid point and took an average in each sub-box. For almost all the boxes, the dispersions of the scaling factors were within |$3 \, \sigma$| of the mean value of each observation season. For the boxes where the SF values were larger than |$3 \, \sigma$| of the mean value or the boxes where the emission was too weak to measure the SF value, we adopted the value of SF measured immediately before or after the corresponding observation. However, we note that such sub-boxes with large dispersions were rare. The scaling factors so derived are consistent with the results of the measurements of the telescope main beam efficiencies. For example, in the 2015–2016 season, the mean value of the scaling factor was measured to about 2.1, which corresponds to the main beam efficiency of 0.48 at 115 GHz. The actual main beam efficiency of the telescope was measured to be about 0.45.

Finally, we combined the FOREST ¹²CO and ¹³CO data with the BEARS data to reduce the noise levels. The data combination enabled us to typically reduce the noise levels by a factor of 1.5–2, which was crucial for making the CARMA and NRO combined Orion images, so that the noise levels in the uv space were well matched.

4.1.2 C¹⁸O, CCS, and N₂H⁺

Since the map coverage of the C¹⁸O data obtained with BEARS was not too large and we did not have the corresponding N₂H⁺ and CCS data, we applied a different intensity calibration procedure. For the C¹⁸O and N₂H⁺ data we derived the daily scaling factors or daily intensity variations by observing a small area containing FIR 4, which was observed at every observation session once or twice. To determine the daily scaling factors we observed a small area containing FIR 3/4/5 at the date when the wind speed at the telescope site was almost zero and T_sys was close to the lowest value, and we adopted the integrated intensity of the FIR 3/4/5 area at that date as a reference. Immediately after this measurement we also observed a small area of B213 in Taurus to compare the data with those obtained with the IRAM 30 m telescope (Tafalla & Hacar 2015) to check the intensity accuracy. To compare the intensities taken with different telescopes, we smoothed the Nobeyama data to match the map effective angular resolution. Details of the FIR 3/4/5 area are given in Nakamura et al. (2019). For C¹⁸O, we also checked the absolute intensity by using the C¹⁸O map obtained with BEARS toward the northern part of Orion A (Shimajiri et al. 2014, 2015) and we confirmed that our FOREST C¹⁸O intensities agree with the BEARS data within an error of less than 10%. For N₂H⁺, we checked the absolute intensity by using the N₂H⁺ (J = 1–0) fit data obtained with the IRAM 30 m telescope toward B213 in Taurus (Tafalla & Hacar 2015). We confirmed that the intensity of N₂H⁺ (J = 1–0) in B213 obtained with FOREST was only about 5% larger than that of the IRAM value, where we divided the intensity in the antenna temperature scale by the main beam efficiency at 94 GHz. Thus, we consider that the intensity scale of the N₂H⁺ data is reasonably accurate.

For CCS, we simply divided the intensities in |$T_{\rm A}^*$| by the main beam efficiency. The CCS emission is extremely weak for all the targets.

4.2 Aquila Rift and M17

The BEARS maps of the Aquila Rift were only toward the Serpens South region and the map coverages are very small compared to the FOREST mapping areas. Therefore, we did not combine the BEARS and FOREST data. Thus, we applied the same procedure as the Orion A C¹⁸O data calibration to make the ¹²CO, ¹³CO, C¹⁸O, and N₂H⁺ data of the Aquila Rift. After correcting the daily intensity variations of the FOREST data, we compared the FOREST intensities to the BEARS intensities toward Serpens South to determine the scaling factors. The scaling factors computed agreed with those determined with the main beam efficiencies within an error of ∼5%–10%. For M17, we mapped small areas to measure the daily intensity variations and followed the same procedure as for the Aquila Rift.

4.3 Main beam efficiencies of the Nobeyama 45 m telescope

The Nobeyama Radio Observatory measures the main beam efficiencies of the telescope at several frequencies every season, except for the 2016–2017 season which was our last (third) season. Planets such as Mars or Jupiter are often used for the measurements—see the web page of the observatory for details of the measurements. Here, we obtain the main beam efficiencies of the observed lines by fitting the values measured with several receivers on the 45 m telescope.

In figure 4 we show the telescope main beam efficiency measured in the 2015–2016 season. We plotted all the values of the main beam efficiencies measured with the available receivers installed on the 45 m telescope. By fitting the data points with the function

$$\begin{equation} \eta _{\rm A} = \eta _{A0} \times \exp \left[-\left({4 \pi \varepsilon \over \lambda } \right)^2 \right] \end{equation}$$

(2)

we derived the main beam efficiencies at the frequencies of our target molecular lines using a least-squares fit; we obtained η_A0 = 0.539, and the accuracy of the parabola surface is given as ε = 0.151 mm. The derived main beam efficiencies are summarized in the last column of table 3. This procedure is basically the same as the one recommended by the observatory.

We used the obtained main beam efficiencies to convert the intensities measured in the antenna temperature scale into those in the brightness temperature scale, except for ¹²CO and ¹³CO of Orion A.

Fig. 4.

Main beam efficiency of the Nobeyama 45 m telescope as a function of frequency. The crosses indicate the efficiencies measured with the receivers installed on the 45 m telescope. The red line shows the line fitted with equation (2).

5 Global molecular gas distribution

In this section we present the global molecular gas distribution toward the three target clouds. The detailed characteristics, particularly the velocity structures, of the individual clouds will be described in separate papers.

5.1 Average spectra of the three clouds

In figure 5 we present the average spectra of the CO lines for the three clouds. For Orion A, the average spectrum of the ¹²CO emission line has a single peak. There are three peaks or shoulders in ¹³CO and C¹⁸O at around 7 km s⁻¹, 9 km s⁻¹, and 11 km s⁻¹. The Orion A filament has a large velocity gradient toward the northern part, and these components are affected by the global velocity gradient along the filament. However, roughly speaking, the 11 km s⁻¹ component is dominant toward the northern part, while the 7 km s⁻¹ component appears in the southern part. The 9 km s⁻¹ component is strong in the main filamentary structure.

Fig. 5.

Average spectra of ¹²CO, ¹³CO, and C¹⁸O toward (a) Orion A, (b) Aquila Rift, and (c) M17. The blue, red, and black lines indicate the ¹²CO, ¹³CO, and C¹⁸O spectra, respectively. (Color online)

For the other two regions, the spectra have multiple peaks. For Aquila Rift and M17, ¹²CO has at least three major peaks. For Aquila Rift, the C¹⁸O has a single peak at 7.3 km s⁻¹ and thus the dips seen in ¹²CO and ¹³CO are expected to be due to self-absorption (Shimoikura et al. 2018). The weak but distinct component at around 40 km s⁻¹ may be the molecular gas influenced by a supperbubble created by star formation in the Scorpius–Centaurus Association (Breitschwerdt et al. 1996; Frisch et al. 2015). This component is seen in ¹³CO but it is difficult to recognize in C¹⁸O. The ¹²CO profile has a tail between 10 km s⁻¹ and 35 km s⁻¹. This comes from several different components which reside in this region. These components are more prominent in the average spectra of smaller areas and we will discuss them later in a separate paper. These complicated molecular gas distributions may be due to the interaction of molecular clouds with superbubbles (Frisch 1998; Frisch et al. 2015; Nakamura et al. 2017).

For M17, the three components of ∼20 km s⁻¹, ∼40 km s⁻¹, and ∼55 km s⁻¹ are the molecular gas components which belong to the Sagittarius, Scutum, and Norma arms, respectively (Zucker et al. 2015), as discussed by Nguyen Luong et al. (2019), and the main component is the one at 20 km s⁻¹ where NGC 6811 and M17 SWex are located.

5.2 Orion A

Figures 6, 7, 8, 9, and 10 show the integrated intensity maps of Orion A for ¹²CO (J = 1–0), ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), and CCS (J = 8₇–7₆), respectively. For comparison, we have overlaid the contours of the Herschel H₂ column density map on the integrated intensity image in the right panel of each figure. Each map except N₂H⁺ is integrated from 2 km s⁻¹ to 20 km s⁻¹. For N₂H⁺, we integrated the emission from 0 km s⁻¹ to 22 km s⁻¹ so that all seven hyperfine components are summed. Our maps cover a region from OMC-3 to NGC 1999, spanning about 2° in declination. The results of our protostellar outflow survey and cloud structure analysis are given in separate papers (Tanabe et al. 2019; Ishii et al. 2019; H. Takemura et al. in preparation).

Fig. 6.

(a) ¹²CO (J = 1–0) moment-0 map of Orion A, velocity-integrated from 2 km s⁻¹ to 20 km s⁻¹. (b) Same as panel (a), but the contours of the H₂ column density are overlaid on the image. The contour levels are 2.5 × 10²¹ cm⁻², 5.0 × 10²¹ cm⁻², 7.5 × 10²¹ cm⁻², 2.5 × 10²² cm⁻², 5.0 × 10²² cm⁻², 7.5 × 10²² cm⁻², … The effective angular resolution of the ¹²CO map is |${21{^{\prime \prime }_{.}}7}$|⁠. (Color online)

Fig. 7.

(a) ¹³CO (J = 1–0) moment-0 map of Orion A, velocity-integrated from 2 km s⁻¹ to 20 km s⁻¹. (b) Same as panel (a) but the contours of the H₂ column density are overlaid on the image. The contour levels are the same as those of figure 6. The effective angular resolution of the ¹³CO map is |${22{^{\prime \prime }_{.}}1}$|⁠. (Color online)

Fig. 8.

(a) C¹⁸O (J = 1–0) moment-0 map of Orion A, velocity-integrated from 2 km s⁻¹ to 20 km s⁻¹. (b) Same as panel (a), but the contours of the H₂ column density are overlaid on the image. The contour levels are the same as those of figure 6. The effective angular resolution of the C¹⁸O map is |${22{^{\prime \prime }_{.}}1}$|⁠. (Color online)

Fig. 9.

(a) N₂H⁺ (J = 1–0) velocity integrated map of Orion A. The integration range is from 0 km s⁻¹ to 20 km s⁻¹. (b) Same as panel (a), but the contours of the H₂ column density are overlaid on the image. The contour levels are the same as those of figure 6. The effective angular resolution of the N₂H⁺ map is |${24{^{\prime \prime }_{.}}1}$|⁠. (Color online)

Fig. 10.

CCS integrated intensity map of Orion A. We smoothed the CCS image with an effective angular resolution of 32″. The contour levels are the same as those of figure 6. The effective angular resolution of the CCS map is |${24{^{\prime \prime }_{.}}0}$|⁠. (Color online)

Figures 6, 7, and 8 show that the ¹²CO, ¹³CO, and C¹⁸O emissions trace the areas with column density higher than ∼0.5 × 10²² cm⁻², ∼0.75 × 10²² cm⁻², and ∼2.5 × 10²² cm⁻², respectively. N₂H⁺ emission comes from the area with column density larger than ∼5 × 10²² cm⁻² (see figure 9).

In figure 10 we show the integrated intensity map of CCS, where the image was smoothed with an effective angular resolution of 32″ to improve the signal-to-noise ratio. The CCS emission is detected significantly only in the OMC-1 region where the strong C¹⁸O and N₂H⁺ emission is detected. Our CCS map is the first unbiased CCS map of Orion A with the J_N = 8₇–7₆ line. Previous wide-field maps were taken only toward the main filament with other transition lines (J_N = 4₃–3₂ at 45 GHz and J_N = 7₆–6₅ at 81.5 GHz) by Tatematsu et al. (2008, 2014). CCS is known to trace dense gas with densities of 10⁴ cm⁻³, but the abundance of CCS decreases very rapidly due to the destruction. In Orion A, the CCS emission is very weak. Only toward the OMC-1 region we detect the emission with a signal-to-noise ratio of 5 σ. Weaker emission is sometimes seen along the ridge. This weak CCS emission implies that Orion A is a relatively evolved molecular cloud. We note that CCS is detected in the OMC-2 FIR 4 region for much higher-sensitivity observations (Nakamura et al. 2019). The OMC-1 region may be relatively chemically young compared to other parts in Orion A. Recently, Hacar et al. (2017) proposed that the OMC-1 region is gravitationally contracting along the main filament. Such a global infall motion can be recognized in our ¹³CO data (Ishii et al. 2019). If OMC-1 is indeed infalling toward the center, the gas is continuously fed along the main filament in OMC-1. Thus, the significant CCS emission in OMC-1 may come from the material newly fed from outside by the gravititional contraction.

Figure 9 shows the N₂H⁺ map of the Orion A molecular cloud. The map is consistent with the image taken by Tatematsu et al. (2008), who mapped the Orion A filament in a position-switch mode with full-beam sampling, using the BEARS receiver. Our mapping area is much wider than theirs. Also, since we mapped Orion A in the OTF mode, the angular resolution is somewhat better than that of Tatematsu et al. (2008). The N₂H⁺ emission is stronger in the northern part (OMC-1, OMC-2, and OMC-3) and traces the main filamentary structure running in the north–south direction. Our N₂H⁺ map shows that many faint N₂H⁺ cores are distributed outside the main filament as well.

5.3 Aquila Rift

Figures 11, 12, 13, 14, and 15 show the integrated intensity maps of the Aquila Rift for ¹²CO (J = 1–0), ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), and CCS (J = 7₆–6₅), respectively. Each map except N₂H⁺ is integrated from −10 km s⁻¹ to 45 km s⁻¹. For N₂H⁺, we integrated all seven hyperfine components to produce the intensity map.

Fig. 11.

¹²CO (J = 1–0) integrated intensity map of the Aquila Rift. In the panel (b), the contour levels are 1.0 × 10²² cm⁻², 2.5 × 10²² cm⁻², 5.0 × 10²² cm⁻², 7.5 × 10²² cm⁻², 1.0 × 10²³ cm⁻², 1.5 × 10²³ cm⁻². The effective angular resolution of the ¹²CO map is |${21{^{\prime \prime }_{.}}7}$|⁠. (Color online)

Fig. 12.

¹³CO (J = 1–0) integrated intensity map of the Aquila Rift. The contour levels are the same as those of figure 11. The effective angular resolution of the ¹³CO map is |${22{^{\prime \prime }_{.}}1}$|⁠. (Color online)

Fig. 13.

C¹⁸O (J = 1–0) integrated intensity map of the Aquila Rift. The contour levels are the same as those of figure 11. The effective angular resolution of the C¹⁸O map is |${22{^{\prime \prime }_{.}}1}$|⁠. (Color online)

Fig. 14.

N₂H⁺ (J = 1–0) integrated intensity map of the Aquila Rift. The contour levels are the same as those of figure 11. The effective angular resolution of the N₂H⁺ map is |${24{^{\prime \prime }_{.}}1}$|⁠. (Color online)

Fig. 15.

CCS (J = 8₇–7₆) integrated intensity map of the Aquila Rift. The contour levels are the same as those of figure 11. The effective angular resolution of the CCS map is |${24{^{\prime \prime }_{.}}0}$|⁠. (Color online)

The ¹²CO emission is strongest toward the W|$\, 40$| region and the Serpens South cluster. The ¹³CO emission traces an arc-like structure in the W|$\, 40$| region. However, the filamentary structures, particularly toward the Serpens South region, are difficult to recognize in the ¹²CO and ¹³CO images. From the C¹⁸O image, we can vaguely find the filamentary structures in Serpens South. The filamentary structures detected by the Herschel map are prominent toward Serpens South in N₂H⁺

A prominent linear structure is seen in the northeast part of the observed area from W|$\, 40$| in the ¹³CO map. This ¹³CO structure is less prominent in the Herschel column density map. This structure may be created by the interaction of the molecular cloud and the W|$\, 40$| H ii region. Similar linear structures are seen in C¹⁸O at different parts of the mapped area. In figure 16, we compare the 250|$\, \mu$|m Herschel image (gray scale) with the ¹³CO image velocity-integrated from 5.0 km s⁻¹ to 6.6 km s⁻¹ (contours). Weak 250|$\, \mu$|m emission appears to trace the ¹³CO linear structure. See Shimoikura et al. (2019b) for details of the cloud structure and properties of the Aquila Rift. The ¹²CO integrated intensity map indicates the presence of protostellar outflows, particularly in the Serpens South protocluster region (see also Nakamura et al. 2011a; Shimoikura et al. 2015). In section 7 we present the results of the outflow survey toward the Serpens South region. The CCS emission is weak, but significant emission comes along the Serpens South filament. This indicates that the Serpens South filaments are chemically relatively young.

Fig. 16.

The ¹³CO intensity map overlaid on the Herschel 250|$\, \mu$|m image. The integration range of the ¹³CO is from 5.0 km s⁻¹ to 6.6 km s⁻¹. The coutours start at 4.0 K km s⁻¹ with an interval of 4.0 K km s⁻¹. (Color online)

5.3.1 M17

Figures 17, 18, 19, and 20 show the integrated intensity maps of M17 for ¹²CO (J = 1–0), ¹³CO (J = 1–0), C¹⁸O (J = 1–0), and N₂H⁺ (J = 1–0), respectively. Each map except N₂H⁺ is integrated from −10 to 60 km s⁻¹. We could not detect significant CCS emission, and thus we do not show the CCS integrated intensity map. For N₂H⁺ we integrated all seven hyperfine components to make an integrated intensity map.

Fig. 17.

¹²CO (J = 1–0) integrated intensity map of M17. The effective angular resolution of the ¹²CO map is |${21{^{\prime \prime }_{.}}7}$|⁠. (Color online)

Fig. 18.

¹³CO (J = 1–0) integrated intensity map of M17. The effective angular resolution of the ¹³CO map is |${22{^{\prime \prime }_{.}}1}$|⁠. (Color online)

Fig. 19.

C¹⁸O (J = 1–0) integrated intensity map of M17. The effective angular resolution of the C¹⁸O map is |${22{^{\prime \prime }_{.}}1}$|⁠. (Color online)

Fig. 20.

N₂H⁺ (J = 1–0) integrated intensity map of M17. The effective angular resolution of the N₂H⁺ map is |${24{^{\prime \prime }_{.}}1}$|⁠. (Color online)

The ¹²CO and ¹³CO emission is strongest toward the M17 H ii region. The emission lines are spatially extended toward the M17 SWex region, whose dense parts are detected in C¹⁸O. The N₂H⁺ emission is spatially localized and seen as blobs. The detailed cloud structure and kinematics will be discussed in separate papers (e.g., Shimoikura et al. 2019a).

6 Spatial variation of the fractional abundances of ¹³CO and C¹⁸O

Here, using ¹²CO, ¹³CO, and C¹⁸O, we derive physical quantities such as the excitation temperature, fractional abundances, and optical depth toward Orion A. The detailed analysis is presented in Ishii et al. (2019). See separate papers for the results of the other regions.

6.1 Derivation of excitation temperature, optical depth, and column density of CO

Here, we derive excitation temperature, optical depth, and ¹³CO and C¹⁸O column densities toward Orion A. First, we describe how we derive the physical quantities from our data, and then briefly present the spatial distributions of the physical quantities for Orion A.

Assuming that the ¹²CO (J = 1–0) line is optically thick, we derive the excitation temperature T_ex using the following equation (see, e.g., Pineda et al. 2008; Shimajiri et al. 2011; Ishii et al. 2019):

$$\begin{equation} T_{\rm ex} = { 5.53 \over \ln [ 1+ 5.53 / (T_{\rm peak} + 0.819)] } \, {\rm K} , \end{equation}$$

(3)

where T_peak is the maximum intensity along the line of sight. We also assumed that the excitation temperatures of ¹³CO and C¹⁸O are the same as T_ex calculated with the above equation. We note that the excitation temperature of ¹²CO, which is likely to be close to LTE, can be used as a good measure of the gas temperature.

The optical depths and column densities for each channel are derived with the following formulas:

$$\begin{equation} \tau _{\rm ^{13}CO} = - \ln \left\lbrace 1 - {\ T_{\rm ^{13}CO} \over 5.29 [J (T_{\rm ex}) - 0.164]}\right\rbrace , \end{equation}$$

(4)

$$\begin{equation} N_{\rm ^{13}CO} = 2.42 \times 10^{14} \left[ \tau _{\rm ^{13}CO} T_{\rm ex} \Delta V \over 1 - \exp (-5.29 / T_{\rm ex} ) \right] \, {\rm cm^{-2}} , \end{equation}$$

(5)

$$\begin{equation} \tau _{\rm C^{18}O} = - \ln \left\lbrace 1 - {\ T_{\rm C^{18}O} \over 5.27 [J (T_{\rm ex}) - 0.164]}\right\rbrace , \end{equation}$$

(6)

and

$$\begin{equation} N_{\rm C^{18}O} = 2.42 \times 10^{14} \left[ \tau _{\rm C^{18}O} T_{\rm ex} \Delta V \over 1 - \exp (-5.27 / T_{\rm ex} ) \right] \, {\rm cm^{-2}} , \end{equation}$$

(7)

where we assumed that the rotational transition is in the LTE. Here, ΔV is the velocity resolution. The beam filling factors of ¹³CO and C¹⁸O are set to unity. The function J_X(T) is given by [exp (5.29/T) − 1]⁻¹ and [exp (5.27/T) − 1]⁻¹ for ¹³CO and C¹⁸O, respectively. To derive the opacity-corrected total column densities (Pineda et al. 2008), we summed the column density at each channel along the line of sight and multiplied the column densities by the correction factor of τ_i/[1 − exp (−τ_i)], where i is either ¹³CO or C¹⁸O. The opacity-corrected total H₂ column density is then derived from the ¹³CO integrated intensity assuming a constant abundance of 2 × 10⁻⁶ (Dickman 1978). We also derived the ¹³CO and C¹⁸O fractional abundances relative to H₂ by dividing the column densities by the Herschel H₂ column density. See also Ishii et al. (2019) for these derivations for Orion A.

We summarize the molecular gas mass derived from ¹³CO in table 9. The total masses of the observed areas are estimated to be |$2.67 \times 10^{4}\, M_{\odot }$|⁠, |$3.86 \times 10^{4}\, M_{\odot }$|⁠, and |$8.1 \times 10^{5}\, M_{\odot }$| for Orion A, Aquila Rift, and M17, respectively.

Table 9.

Molecular gas mass estimated from ¹³CO toward Orion A, Aquila Rift, and M17.^∗

Name	Region	Mass (⁠|$M_{\odot }$|⁠)
Orion A	total	3.86 × 104
Orion A North	Dec ≥ −05:30:26.0	1.56 × 104
Orion A South	Dec ≤ −05:16:03.5	2.30 × 104
Orion A OMC-1	−05:30:26.0 ≤ Dec ≤ −0.5:16:03.5	7.4 × 103
Aquila Rift	total	2.67 × 104
Aquila East (W|$\, 40$|⁠)	RA ≤ 18:30:39.6	1.30 × 104
Aquila West (Serpens South)	RA ≥ 18:30:39.6	1.37 × 104
M17	total	8.1 × 105
M17	l ≥ 14:26:22.5	3.6 × 105
M17 SWex	l ≤ 14:26:22.5	4.5 × 105

*OMC-1 is part of the north area of Orion A. The distances of 414 pc, 436 pc, and 2.0 kpc are adopted for Orion A, Aquila Rift, and M17, respectively. The excitation temperature of ¹³CO is estimated from the peak intensity of ¹²CO at each pixel.

6.1.1 Orion A

In figure 21 we present the excitation temperature map of Orion A determined by the ¹²CO peak intensity. The excitation temperature ranges from ∼8 K to ∼100 K, and takes its maximum near Orion KL. The main filament has somewhat higher temperatures at ∼50 K in the northern region. The southern part including L 1641 N and NGC 1999 has somewhat lower temperature at 10–20 K. This temperature distribution is consistent with that of Kong et al. (2018) with a finer angular resolution.

Fig. 21.

Excitation temperature map of Orion A. (Color online)

The excitation temperature derived from ¹²CO tends to be somewhat higher than the dust temperature derived from the spectral energy distribution (SED) of the Herschel data. In figure 22 we show the ratio of the excitation temperature derived from ¹²CO to the dust temperature. Here, we smoothed the excitation temperature map with an effective angular resolution of 36″ to match the effective resolution of the dust temperature map. The ratio stays at around two along the main filament, except in OMC-1 where the ratio is about five.

Fig. 22.

Ratio of the excitation temperature derived from ¹²CO to the dust temperature derived from the Herschel data. The ratio is around 1–2 along the main ridge of Orion A, except in OMC-1. (Color online)

Figures 23, 24, and 25 show the spatial distributions of the opacity-corrected ¹³CO column density, the fractional abundances relative to H₂, and the optical depth for ¹³CO (J = 1–0) and C¹⁸O (J = 1–0), respectively. The optical depth of ¹³CO ranges from a few to unity. It is about unity along the filament. Several compact regions such as L 1641 N have larger optical depth at ∼3. On the other hand, the optical depth of C¹⁸O is less than unity for almost all areas. Thus, the C¹⁸O emission is reasonably optically thin for the entire cloud.

Fig. 23.

Opacity-corrected ¹³CO column density map of Orion A. (Color online)

Fig. 24.

(a) ¹³CO fractional abundance and (b) its optical depth maps of Orion A. (Color online)

Fig. 25.

(a) C¹⁸O fractional abundance and (b) its optical depth maps of Orion A. (Color online)

The fractional abundance of ¹³CO ranges from a few × 10⁻⁷ to a few × 10⁻⁶. It tends to be small, being 5 × 10⁻⁷ toward some dense areas such as the main filament of OMC-1/2/3 and L 1641 N. Outside the dense areas, the abundance goes up to a few × 10⁻⁶. For C¹⁸O, the fractional abundance ranges from 1 × 10⁻⁸ to 1.5 × 10⁻⁷. Similar to the ¹³CO abundance, it is small at some dense areas such as OMC-1/2/3 and L 1641 N.

Our map indicates that the fractional abundance varies from region to region within a factor of ∼10 (see figure 26). Toward the dense regions, the ratio seems to approach to ∼5, close to the standard interstellar value. This variation may be partly related to the selective dissociation due to the FUV radiation discussed by Shimajiri et al. (2011) and Ishii et al. (2019). See Ishii et al. (2018) for more details of the analysis.

Fig. 26.

¹³CO-to-C¹⁸O fractional abundance ratio of Orion A. (Color online)

In figure 27a we present the cumulative ¹³CO mass distributions of Orion A. We divide the cloud into three areas: North (Dec ≥ −05:32:56.0), OMC-1 (–05:30:26.0 ≤ Dec ≤ −05:16:03.5), and South (Dec ≤ −05:32:56.0), where the North area includes OMC-1. This indicates that the OMC-1 area contains a fraction of molecular gas with column density higher than ∼1 g cm⁻², the threshold for high-mass star formation proposed by Krumholz and McKee (2008), beyond which further fragmentation of clouds can be avoided to form lower-mass cores, where ∼1 g cm⁻² corresponds to a ¹³CO column density of ∼5 × 10¹⁷ cm⁻², assuming a ¹³CO fractional abundance of 2 × 10⁻⁶ (Dickman 1978). In contrast, in the South, molecular gas with column densities higher than the threshold is deficient. The total ¹³CO mass is estimated to be |$3.9 \times 10^{4}\, M_{\odot }$| with 7.4 × 10³|$M_{\odot }$|⁠, |$1.6\times 10^{4}\, M_{\odot }$|⁠, and |$2.3\times 10^{4}\, M_{\odot }$| for OMC-1, North, and South, respectively.

Fig. 27.

Cumulative column density distributions in (a) Orion A, (b) Aquila Rift, and (c) M17. In Orion A, the area is divided into two (North and South) at Dec = −05:32:56.0. OMC-1 is a part of North and within −05:30:26.0 ≤ Dec ≤ −05:16:03.5. In the Aquila Rift, the area is divided into two (West and East) at RA = 18:30:39.6. In M17, the area is divided into two (H ii and SWex) at l = 14:26:22.5. (Color online)

Below, we repeat the same analysis for the other two regions.

6.1.2 Aquila Rift

Figure 28 shows the spatial distribution of the CO excitation temperature. The excitation temperature is high toward the W|$\, 40$| H ii region. Serpens South has a relatively high excitation temperature of 25 K, but in other parts around Serpens South it is very low at ∼10 K, indicating that the star formation may not be active. In figure 29 we show the ratio of the excitation temperature derived from ¹²CO to the dust temperature. Here, we smoothed the excitation temperature map with an effective angular resolution of 36″ to match the effective resolution of the dust temperature map. The excitation temperature is nearly the same as the dust temperature in the Aquila Rift except in the W|$\, 40$| H ii region where the ratio goes up to about two.

Fig. 28.

Excitation temperature map of Aquila Rift. (Color online)

Fig. 29.

The ratio of the excitation temperature derived from ¹²CO to the dust temperature derived from the Herschel data. The ratio stays at around unity in the molecular clouds. In W|$\, 40$|⁠, it goes up to ∼2. (Color online)

Figures 30, 31, and 32 show the spatial distribution of the ¹³CO column density, the fractional abundances relative to H₂, and the optical depth for ¹³CO (J = 1–0) and C¹⁸O (J = 1–0), respectively. The optical depth of ¹³CO tends to be higher than Orion A, and is sometimes larger than 2–3 toward the Serpens South region. The fractional abundances of ¹³CO and C¹⁸O seem to be low in the filamentary structures seen in the Herschel map. This may indicate that CO molecules are depleted in the cold dense gas in Serpens South. The ratio of ¹³CO and C¹⁸O abundances stays at around 0.1 (see figure 33).

Fig. 30.

The opacity-corrected ¹³CO column density map of Aquila Rift. (Color online)

Fig. 31.

(a) ¹³CO fractional abundance and (b) its optical depth maps of Aquila Rift. (Color online)

Fig. 32.

(a) C¹⁸O fractional abundance and (b) its optical depth maps of Aquila Rift. (Color online)

Fig. 33.

C¹⁸O-to-¹³CO fractional abundance ratio of Aquila Rift. (Color online)

In figure 27b we present the cumulative ¹³CO mass distributions of Aquila Rift. We divide the cloud into two areas: East (Dec ≥ 18:30:39.6) and West (Dec ≤ 18:30:39.6). This indicates that the Eastern area contains a fraction of molecular gas with column density higher than ∼1 g cm⁻², the threshold for high-mass star formation proposed by Krumholz and McKee (2008). In contrast, in the Western area which contains Serpens South, the molecular gas with column densities higher than the threshold is deficient. The total ¹³CO mass is estimated to be |$3.9 \times 10^{4}\, M_{\odot }$| with 7.4 × 10³|$M_{\odot }$|⁠, |$1.6\times 10^{4}\, M_{\odot }$|⁠, and |$2.3\times 10^{4}\, M_{\odot }$| for OMC-1, North, and South, respectively.

6.2 M17

Figure 34 shows the spatial distribution of the CO excitation temperature. The excitation temperature is high toward the M17 H ii region. In contrast, in the infrared dark cloud the excitation temperature stays at 30–40 K. The excitation temperatures are not so low as those of the Serpens South region. Figures 35 and 36 show the spatial distribution of the column density and the optical depth for ¹³CO (J = 1–0) and C¹⁸O (J = 1–0), respectively. The optical depths of ¹³CO and C¹⁸O tend to be higher toward M17 SWex. See Nuygen Luong et al. (2019) for details of the global molecular gas distributions.

Fig. 34.

Excitation temperature map of M17. (Color online)

Fig. 35.

(a) ¹³CO fractional abundance and (b) its optical depth maps of M17. (Color online)

Fig. 36.

(a) C¹⁸O fractional abundance and (b) its optical depth maps of M17. (Color online)

Fig. 37.

C¹⁸O-to-¹³CO fractional abundance ratio of M17. (Color online)

In figure 27c we present the cumulative ¹³CO mass distributions of M17. We divide the region into two areas: the M17 H ii area (l ≥ 14.9) and M17 SWex (l ≤ 14.9). This indicates that the M17 H ii region contains a fraction of molecular gas with column density higher than ∼1 g cm⁻², the threshold for high-mass star formation proposed by Krumholz and McKee (2008). In contrast, for M17 SWex molecular gas with column densities higher than the threshold is deficient. This may be a reason why high-mass star formation is not active in M17 SWex (Povich et al. 2016). The total ¹³CO mass is estimated to be |$3.6 \times 10^{5}\, M_{\odot }$| and |$4.5 \times 10^{5}\, M_{\odot }$| for M17 H ii and M17 SWex, respectively. The M17 SWex area contains about twice as much massive molecular gas than M17 H ii.

7 Molecular outflows

Molecular outflow feedback is an important stellar feedback mechanism (Nakamura & Li 2007, 2014). Molecular outflows can inject significant energy and momentum into molecular clouds. Our ¹²CO data are useful to identify the outflows and gauge how much energy and momentum are injected into the clouds. A molecular outflow survey is one of the main studies we will conduct. Here we briefly present the result of the molecular outflow survey toward small areas in Orion A (L 1641 N) and the Aquila Rift (Serpens South). See Tanabe et al. (2019) for the results of comprehensive outflow surveys toward Orion A.

7.1 Orion A (L 1641 N)

In figure 38 we present a three-color image of the L 1641 N region. The red, blue, and green images show the ¹²CO intensity image integrated from 11 km s⁻¹ to 15 km s⁻¹ (redshifted component), ¹²CO intensity image integrated from −20 km s⁻¹ to 0 km s⁻¹ (blueshifted components), and the Herschel column density image, respectively. In this region, Stanke and Williams (2007) and Nakamura et al. (2012) identified molecular outflows in ¹²CO (J = 2–1) and ¹²CO(J = 1–0), respectively. The high-velocity components they previously identified are recognized in our image. For example, there are several dust cores in this region, shown in green. The redshifted collimated flow running from north to south blows out from the brightest dust core located at the position (RA, Dec) = (5:36:19, −6:22:29).

Fig. 38.

Three-color image of the L 1641 N cluster-forming region with ¹²CO (J = 1–0) intensity integrated from 10 km s⁻¹ to 16 km s⁻¹ (red), ¹²CO (J = 1–0) intensity integrated from 0 km s⁻¹ to 5 km s⁻¹ (blue), and the Herschel column density (green). (Color online)

The result of accurate outflow identification is presented in Tanabe et al. (2019), who found 44 CO outflows in Orion A, 17 of which are new detections. Based on the identified outflow physical parameters, they estimated the momentum injection rates due to the molecular outflows and found that the total momentum injection rate due to the outflows and the expanding shells identified in the ¹³CO data by Feddersen et al. (2018) is larger than the turbulence dissipation rate in Orion A. Thus, the stellar feedback such as molecular outflows and expanding shells driven by stellar winds is an important mechanism to replenish the internal cloud turbulence.

7.2 Aquila Rift (Serpens South)

In figure 39 we present the three-color image of the Serpens South region. Nakamura et al. (2011a) conducted a molecular outflow survey toward the Serpens South cluster in ¹²CO (J = 3–2). The present paper is the first outflow survey using ¹²CO (J = 1–0). The coverage of the image shown in figure 39 is wider than that of Nakamura et al. (2011a). The red, blue, and green images indicate the ¹²CO intensity image integrated from 11 km s⁻¹ to 15 km s⁻¹ (redshifted component), the ¹²CO intensity image integrated from −20 km s⁻¹ to 0 km s⁻¹ (blueshifted components), and the Herschel column density image, respectively. The distribution of the high-velocity components is basically similar to that of ¹²CO (J = 3–2).

Fig. 39.

Three-color image of the Serpens South region with ¹²CO (J = 1–0) intensity integrated from 11 km s⁻¹ to 15 km s⁻¹ (red), ¹²CO (J = 1–0) intensity integrated from −20 km s⁻¹ to 0 km s⁻¹ (blue), and the Herschel column density (green). The squares and circles indicate protostellar core candidates with and without molecular outflows, respectively. The size of the circle is the same as the FWHM beam size of |${21{^{\prime \prime}_{.}}7}$|⁠. (Color online)

By visual inspection, we attempted to identify the high-velocity components that are likely to originate from the molecular outflows toward Herschel protostellar cores in this region. The result of the identification is summarized in table 10. We detected in the ¹²CO (J = 1–0) emission almost all the outflow lobes identified by Nakamura et al. (2011a). In total, we identified 13 outflow driving sources including the 3 tentative detections. From this survey, we identified 4 new outflow sources, all of which are located outside the map of Nakamura et al. (2011a).

Table 10.

Molecular outflow candidates in Serpens South.^∗

No.	RA (J2000.0)	Dec (J2000.0)	Outflows	Classification	Nakamura et al. (2011a)
135	18:28:48.09	−01:38:11.6	N
147	18:29:00.01	−01:42:42.8	N
151	18:29:03.62	−01:39:03.0	BR	C	B12, R7
155	18:29:05.52	−01:41:53.6	R	C	B11, R6
163	18:29:08.34	−01:30:46.8	N
171	18:29:12.68	−01:46:18.4	R	M	new
174	18:29:13.10	−02:03:51.2	N
196	18:29:21.13	−01:37:12.8	N
202	18:29:23.69	−01:38:54.0	R	C	B10, R5
209	18:29:25.54	−01:47:31.5	N
250	18:29:43.44	−01:56:49.9	N
251	18:29:43.84	−02:12:56.0	N
271	18:29:53.06	−01:58:04.5	N
285	18:29:59.67	−02:00:58.7	R	C	R2
289	18:30:00.85	−02:06:57.3	N
290	18:30:01.20	−02:06:09.8	N
292	18:30:01.50	−02:10:25.5	BR	C	B15, R8
297	18:30:03.68	−01:36:29.4	R	C	new
299	18:30:04.19	−02:03:05.5	BR	C	B1, B2, B3, B6, R1, R3, R4?
315	18:30:12.44	−02:06:53.6	B	C	B7
321	18:30:14.93	−01:33:34.9	BR	M	new
323	18:30:16.22	−02:07:16.3	N
326	18:30:17.64	−02:09:59.3	B	C	B14
342	18:30:26.03	−02:10:41.2
347	18:30:27.97	−02:10:59.0
349	18:30:28.98	−01:56:03.2	R	Y	R9
351	18:30:29.28	−01:56:50.6	R	M	new
362	18:30:37.53	−02:08:56.3	N	N

*First column: the number in the Herschel protostellar core catalog of Könyves et al. (2015). Fourthcolumn: B = blueshifted component, R = redshifted component. Fifth column: C = clear, M = marginal. Sixth column: comparison with the identification by Nakamura et al. (2011a); new = new detection (this paper). In the densest part of the Serpens South cluster, outflow components from several different sources are observed. We assigned all such lobes to No. 299.

As discussed by Shimoikura et al. (2015), the Aquila Rift region contains several cloud components with different line-of-sight velocities. The existence of such multiple components sometimes precludes clear identification of high-velocity components, since high-velocity components tend to overlap with different cloud components along the line of sight. More careful inspection of the data cube is needed to fully identify the molecular outflows. We expect that more outflows exist even in the area presented here. We will present a complete outflow survey toward the Aquila Rift in a separate paper.

The scientific results will be reported in more detail in separate papers.

8 Summary

In the present paper we have provided a project overview of the Nobeyama mapping project toward the three nearby molecular clouds Orion A, Aquila Rift, and M17. The main purpose of the present paper is to summarize the complicated observational procedures and flux calibration methods. We summarize the main results of the paper as follows.

1. We conducted wide-field mapping observations toward three nearby molecular clouds, Orion A, Aquila Rift, and M17, in ¹²CO (J = 1–0), ¹³CO (J = 1–0), C¹⁸O (J = 1–0), N₂H⁺ (J = 1–0), and CCS (J_N = 8₇–7₆) using the Nobeyama 45 m telescope.

2. The map coverage is over 1° × 1°. We cover most of the molecular clouds seen in dust emission.

3. We checked the absolute intensities obtained with the new four-beam receiver, FOREST, by comparing the intensities obtained with the previous receiver, BEARS, toward the same areas for ¹²CO, ¹³CO, and C¹⁸O.

4. For N₂H⁺, we compared our results with the intensities of the Taurus molecular cloud obtained with the IRAM 30 m telescope; the fluxes taken with FOREST coincide with those obtained with the IRAM 30 m telescope within an error of 5%.

5. We obtained the column densities of ¹³CO (J = 1–0) and C¹⁸O (J = 1–0) and derived their fractional abundances toward Orion A. Our maps indicate that the fractional abundances depend on the cloud environments, and vary from region to region by a factor of ∼10.

6. The cumulative column density distributions clearly show that only a fraction of the molecular gas has column densities high enough to create high-mass stars for individual clouds.

7. Our maps have sufficient sensitivities to identify the molecular outflows. In particular, in our ¹²CO (J = 1–0) data, we confirmed all the outflows previously detected in ¹²CO (J = 3–2) toward Serpens South, and identified four new outflows in the adjacent region. Using the catalog of the protostars, we identified the driving sources of these CO outflows.

8. Finally, we briefly described results from our project published in separate papers. We revealed the hierarchical structure of Orion A, applying SCIMES and Dendrogram to the ¹³CO cube data. In total, we identified about 80 clouds in Orion A. The abundance ratio of ¹³CO to C¹⁸O varies from region to region, affected by the far-UV radiation (Ishii et al. 2019). We identified 44 ouflows in Orion A, 15 of which are new detections (Tanabe et al. 2019). We estimated the momentum injection rate of the identified outflows and found that they have significant injection momentum rates in the surroundings. Using the data of the OMC-2 FIR 4 region, we characterized the spatial variation of the abundance ratios of several molecules and discussed possible outflow-triggered star formation (Nakamura et al. 2019). Data was taken for the flux calibration of Orion A data. For Aquila, Shimoikura et al. (2019b) reveal evidence of the interaction of the molecular cloud with the expanding H ii region. Kusune et al. (2019) and Sugitani et al. (2019) carried out near-infrared polarization observations toward Aquila (see also Sugitani et al. 2011) and M17, respectively, and reveal that the global magnetic field tends to be perpendicular to the elongation of the molecular clouds. Nguyen Luong et al. (2019) and and Shimoikura et al. (2019b) reveal the global molecular gas distribution in M17, and found that in the IRDC region, the column densities are not dense enough to create high-mass stars, but ongoing cloud–cloud collisions are likely to be forming higher-density regions. Thus, future high-mass star formation is expected.

Acknowledgments

This work was financially supported by JSPS KAKENHI Grant Numbers JP16H05730, JP17H02863, JP17H01118, JP26287030, JP17K00963, JP17H01103, JP18H05441, and JP18H01259. This work was supported by NAOJ ALMA Scientific Research Grant Numbers 2017-04A. 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. We thank the NRO staff for both operating the 45 m and helping us with the data reduction.

Footnotes

1

〈http://jvo.nao.ac.jp/index-e.html〉.