Analysis of the Herschel/HEXOS Spectral Survey Towards Orion South: A massive protostellar envelope with strong external irradiation

K Tahani; R Plume; E A Bergin; V Tolls; T G Phillips; E Caux; S Cabrit; J R Goicoechea; P F Goldsmith; D Johnstone; D C Lis; L Pagani; K M Menten; H S P Müller; V Ossenkopf-Okada; J C Pearson; F F S van der Tak

doi:10.3847/0004-637X/832/1/12

. Author manuscript; available in PMC: 2019 Dec 16.

Published in final edited form as: Astrophys J. 2016 Nov 11;832(1):12. doi: 10.3847/0004-637X/832/1/12

Analysis of the Herschel/HEXOS Spectral Survey Towards Orion South: A massive protostellar envelope with strong external irradiation

K Tahani ¹, R Plume ¹, E A Bergin ², V Tolls ¹, T G Phillips ², E Caux ^3,⁴, S Cabrit ⁵, J R Goicoechea ⁶, P F Goldsmith ⁷, D Johnstone ⁸, D C Lis ^5,², L Pagani ⁵, K M Menten ⁹, H S P Müller ¹⁰, V Ossenkopf-Okada ¹⁰, J C Pearson ⁷, F F S van der Tak ^11,¹²

PMCID: PMC6914383 EMSID: EMS85105 PMID: 31844334

Abstract

We present results from a comprehensive submillimeter spectral survey toward the source Orion South, based on data obtained with the HIFI instrument aboard the Herschel Space Observatory, covering the frequency range 480 to 1900 GHz. We detect 685 spectral lines with S/N > 3σ, originating from 52 different molecular and atomic species. We model each of the detected species assuming conditions of Local Thermodynamic Equilibrium. This analysis provides an estimate of the physical conditions of Orion South (column density, temperature, source size, & V_LSR). We find evidence for three different cloud components: a cool (T_ex ~ 20 – 40 K), spatially extended (> 60″), and quiescent (ΔV_FWHM ~ 4 km s ⁻¹) component; a warmer (T_ex ~ 80 – 100 K), less spatially extended (~ 30″), and dynamic (ΔV_FWHM ~ 8 km s ⁻¹) component, which is likely affected by embedded outflows; and a kinematically distinct region (T_ex > 100 K; V_LSR ~ 8 km s ⁻¹), dominated by emission from species which trace ultraviolet irradiation, likely at the surface of the cloud. We find little evidence for the existence of a chemically distinct “hot core” component, likely due to the small filling factor of the hot core or hot cores within the Herschel beam. We find that the chemical composition of the gas in the cooler, quiescent component of Orion South more closely resembles that of the quiescent ridge in Orion-KL. The gas in the warmer, dynamic component, however, more closely resembles that of the Compact Ridge and Plateau regions of Orion-KL, suggesting that higher temperatures and shocks also have an influence on the overall chemistry of Orion South.

Subject headings: ISM: abundances, ISM: molecules, ISM: lines and bands, ISM: kinematics and dynamics, ISM: individual (Orion South)

1. Introduction

To date, about 200 different molecular species have been detected in the interstellar medium (Menten & Wyrowski 2011)¹. However, our understanding of the total molecular inventory of individual sources is poor, since few sources have been systematically surveyed in any frequency band due to the large amount of observing time required to perform unbiased spectral surveys (e.g. Blake et al. 1987; Schilke et al. 1997a; Nummelin et al. 1998; Schilke et al. 2001; Comito et al. 2005; Furlan et al. 2006; Tercero et al. 2010; Neill et al. 2014). Therefore, we do not truly understand the origin of the chemical complexity observed in interstellar space. Understanding this complexity is important to comprehend details of the formation of stars, planets and life.

Regardless of how complex chemistry arises in interstellar space, the chemical composition (and subsequent chemical evolution) can, in turn, affect the physical conditions (and subsequent dynamical evolution) of a star forming region (e.g see Herbst & van Dishoeck 2009; Garrod & Herbst 2006; Garrod et al. 2008). For example, the overall molecular (and to some degree atomic) content can play an important role in regulating the gas pressure by changing the temperature of the gas via the process of heating and cooling through line-absorption and emission; (Ceccarelli et al. 1996; Goldsmith & Langer 1978). In addition, molecular ions can affect the strength of coupling between the gas and the magnetic fields (which is related to magnetic turbulent support, e.g. Williams et al. 1998). Thus, there is a complex feedback between the physical and chemical conditions in an interstellar gas cloud that either helps drive the star formation process, or hinders it, and which may help determine the masses of the newly formed stars.

In order to understand the origin of chemical complexity in interstellar space and how this chemistry evolves and affects the process of star formation in the Universe (as well as the formation of planets and pre-biotic chemical species), we require unbiased and complete surveys of spectral lines that span a broad range of wavelengths. These types of datasets are needed so that we can sample a wide variety of molecular and atomic species, as well as obtain multiple emission lines from each of the species, in order to extract the physical conditions in the gas. Fortunately, with the advent of sensitive, high-resolution spectrometers for millimeter/submillimeter wavelengths, especially those developed for space-based observatories, it is now possible to obtain such surveys and to begin to address these issues (e.g. Crockett et al. 2014; Zernickel et al. 2012; Kama et al. 2013; Kaźmierczak-Barthel et al. 2015).

The key project Herschel observations of EXtraOrdinary Sources (HEXOS) (Bergin et al. 2010) was designed to address issues related to the chemical composition of massive star forming regions. HEXOS has obtained spectral line surveys of the Orion-KL, Orion South (hereafter Orion-S), and Orion Bar (Nagy et al. in prep.) regions within the Orion A Molecular Cloud, at high frequencies that are not easily accessible from ground based observatories (480–1900 GHz). Both Orion-KL and Orion-S are relatively nearby (420 pc; Menten et al. 2007) massive star forming regions close to the Orion Nebula. The nearby Trapezium OB stars are the source of high energy photons, which produce Photon Dominated Regions (PDRs) throughout the region. The UV flux (6<E<13.6 eV) in the vicinity of Orion-S is χ = 1.1 × 10⁵χ₀ (Herrmann et al. 1997; Goicoechea et al. 2015), where χ₀ = 2.7 × 10⁻⁴ ergs s⁻¹ cm⁻² sr⁻¹ (Draine & Bertoldi 1996). Observations of Goicoechea et al. (2015) and O’Dell & Harris (2010) suggest that the HII region lies mostly in front of the molecular material, but may wrap behind, at least part of, the Orion-S molecular cloud. That at least some of the Orion-S molecular gas is located in front of ionized material has been convincingly demonstrated by Very Large Array absorption measurements of the H₂CO 6 cm 1₁₀ – 1₁₁ transition (Mangum et al. 1993).

Despite the fact that the far-infrared luminosity of Orion-S (8.5 × 10³ L_⊙; Mezger et al. 1990) is more than an order of magnitude below that of KL, a number of energetic outflows associated with Orion-S suggest ongoing star formation. For example, CO J=2–1 SMA observations (Zapata et al. 2005) revealed a highly collimated bipolar outflow extending ~ 30″ over the velocity range ~ −80 to ~ −26 km s⁻¹ and ~ 22 to ~ 82 km s⁻¹ oriented NW-SE. The sub-millimeter continuum source with a deconvolved size ≤ 0.6″ and an integrated flux of 116.2 ± 9.0 mJy at 1.3 mm is well centered on the bipolar outflow axis, α₂₀₀₀ = 05^h35^m13.550^s, δ₂₀₀₀ = −05°23′59.14″. In addition, another quite extended, collimated, low-velocity (5 km s ⁻¹) CO outflow has been observed, oriented NE-SW (Schmid-Burgk et al. 1990), and a low-velocity (10 km s ⁻¹) bipolar SiO (5–4) outflow with a length ~ 30″ (oriented NE-SW) has been reported by Ziurys et al. (1990). Four other SiO outflows are also listed by Zapata et al. (2006).

Despite the presence of star formation activity, as indicated by the IR luminosity and molecular outflows, BIMA observations of a few selected species by McMullin et al. (1993) suggest that the chemistry of Orion-S resembles that of the Orion-KL quiescent ridge and has fewer, narrower and weaker lines than KL. These observations may imply that Orion-S is a more quiescent and younger star forming region, in which the star formation activity has not had time to significantly alter the dynamics and chemistry of the region. This idea is also consistent with dynamical ages from outflow observations in Orion-S (i.e. Schmid-Burgk et al. 1990; Bally et al. 2000; Zapata et al. 2005). Assuming no projection effects, the maximum corresponding dynamical age for the largest outflow is found to be less than 45000 years which is still remarkably young (Schmid-Burgk et al. 1990). The dynamical age for all the other outflows can be shown to be less than 5000 years (Bally et al. 2000; Zapata et al. 2005). A more detailed comparison between Orion-S and Orion-KL is, therefore, of great interest, since the two regions presumably formed under similar conditions, but could have very different chemical abundances possibly based on differences in their ages, densities, temperatures, radiation fields, etc.

In this paper we present a comprehensive study of the Herschel/HIFI spectral survey toward Orion-S. The observations presented here were obtained as part of HEXOS and span over 1.2 THz of frequencies, mostly not accessible from the ground. In §2 we present our observations and data reduction methods, including the removal of off-position contamination, line identification, and Gaussian fitting of the spectral features. Our results (including LTE modeling of each individual species) together with a chemical comparison of Orion-S are presented in §3. The conclusions are provided in §4.

2. Observations & Data Reduction

The data presented in this paper were taken with the Heterodyne Instrument for the Far-Infrared (HIFI) (de Graauw et al. 2010), one of three instruments aboard the Herschel Space Observatory (Pilbratt et al. 2010). HIFI operated over the frequency range 480–1900 GHz (with two gaps: one at 1280–1430 GHz, due to the switch between SIS and HEB detectors (Roelfsema et al. 2012), and one at 1540–1570 GHz, which was an observational time saving strategy since this frequency range was expected to have few transitions). HIFI was separated into 14 different bands (1a, 1b, …, 7b). Each receiver band had independent channels for horizontal (H) and vertical (V) polarizations, each with its dedicated Wide Band Spectrometer (WBS) having a native spectral resolution of 1.1 MHz (Roelfsema et al. 2012). Bands 1–5 were observed with a LO redundancy of 6, whereas Bands 6 and 7 used a redundancy of 2. Redundancy refers to the number of observations of each sky frequency with different LO settings. For example, redundancy of 6 means that each frequency in the band was observed with 6 different LO settings. This redundancy was necessary in order to distinguish lines originating from the upper and the lower sidebands (Comito & Schilke 2002). A redundancy of 2 was sufficient for bands 6 and 7 due to the relatively lower density of transitions at these high frequencies. The central position of Orion-S was α₂₀₀₀ = 5^h35^m13.44^s, δ₂₀₀₀ = –5°24′08.1″. All observations were taken in Dual Beam Switch (DBS) mode using the Fast Chop option (>0.5 Hz chop frequency).

We used the hifiPipeline task in HIPE 9.0 for all data reduction. The hifiPipeline task is a pre-compiled script in HIPE used to process level 0 data to any higher level (e.g. 0.5, 1.0, etc.). See Ott (2010) for a description of the various data products. Spurious spectral features were removed and fully calibrated, double side band (DSB) spectra were deconvolved into single sideband (SSB) spectra (e.g. level 2.5). Additional details on data reduction and observational parameters can be found in Bergin et al. (2010) and Crockett et al. (2010). After processing by HIPE 9.0, the H and V polarizations were coadded (except for Band 4a which, due to processing errors specific to this band, had much noisier V polarization data that were excluded) and then the spectra were Hanning smoothed by two to sixteen channels (see Table 1) to improve the signal to noise ratio. Results are provided in Table 1, which shows typical values for the 1σ rms noise, system temperature, velocity and frequency resolutions after smoothing, and the number of channels smoothed for each band. The 1σ rms noise is calculated from line-free regions of the spectrum immediately adjacent to the lines. A noise range is provided since the noise is not uniform across the bands.

Table 1. Data Smoothing and Noise Characteristics.

	Smoothing Channels	Freq. Res. (MHz)	Vel. Res. (km s⁻¹)	T_sys (K)	1σ rms (K)
band 1	2	2.2	1.0−1.4	~ 100	~ 0.02 − 0.05
band 2	4	4.4	1.6−2.1	~ 150	~ 0.03 − 0.07
band 3	4	4.4	1.4−1.7	~ 200	~ 0.04 − 0.09
band 4	8	8.8	2.4−2.8	~ 400	~ 0.1 − 0.2
band 5	8	8.8	2.1−2.4	~ 1000	~ 0.1 − 0.3
band 6	16	17.6	3.1−3.7	~ 1300	~ 0.3 − 2.0
band 7	16	17.6	2.8−3.1	~ 1300	~ 0.4 − 4.0

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All data in this paper are presented on the T_A temperature scale and, for subsequent analysis, were converted to T_mb using the main beam efficiencies in Müller et al. (2014)². The final data after deconvolution and spectral smoothing (bands 1a–7b) are shown in Figures 1-4, in which the T_A range is in Kelvin and frequency in MHz. The strongest lines are labeled in each band, and in order to make the residual noise recognizable and comparable from one band to another the intensity is fixed to 15 K for all bands. Note that certain broad features, like the one near 790 GHz, are most likely due to excess noise, since individual observations show quite a few noise spikes in these spectral ranges.

Fig. 1 — HEXOS/HIFI spectral scans of (from top to bottom) band 1a, band 1b, band 2a and band 2b after Hanning smoothing. Resolutions, noise levels, and smoothing factors for each band are listed in Table 1. Baselines are not subtracted and some of the strongest lines are labelled.

Fig. 4 — Same as Figure 1 except for band 7a, and band 7b. The higher noise level in band 7 is due to the HEB mixers which produce higher noise in comparison with the SIS mixers used in the first five bands.

2.1. Removal of Off-Position Contamination

As described in the previous section, the HEXOS Orion-S spectral survey was observed in DBS mode. Since the observations used the chopper in this mode, the reference positions were fixed to ~3 arc minutes from the target position, with an angle set by observatory constraint. In a crowded field like the Orion Molecular Cloud region, it is very likely that the reference position is not free of emission or absorption for some or all molecular lines detected. Typically, emission (absorption) in the reference position appears as an absorption (emission) like feature in the final spectrum. Since the case of absorption in the reference beam is rare, due to a low continuum flux, we will consider only emission. Figure 5 an image of the dust emission at a wavelength of 250 μm obtained with the Herschel/SPIRE instrument. The positions of the Orion-S observations and the two reference observations for each of the 14 HIFI bands are overlaid. The diameters of the circles shown represent the FWHM of the individual beams for the center frequencies of each HIFI band. It is apparent that a few of the reference observations (on the east side) were located near the Orion Bar region, making reference beam line contamination very likely. In addition, we even captured emission in the lower-J ¹²CO transitions in the opposite chopping direction.

Fig. 5 — HEXOS Orion-S Observations: The circles indicate the beam positions with diameters corresponding to the FWHM at the center frequency of the HIFI bands. Circles near the center of the image (at l ~ 83.81°) indicate the position of the spectral scan observations of Orion-S. Circles to the left and to the right (at l ~ 83.86° and 83.76° respectively) indicate the off position observations. The background shows the *Herschel*/SPIRE 250 μm dust emission in the Orion-KL region (white regions indicate saturated pixels). From the location of the beam circles near the Orion Bar, it is apparent that some of the observations see emission in at least one reference beam.

The identification of potentially contaminated lines was first performed after the deconvolution by checking the line profile of the detected lines. Once lines were identified, we performed additional tests by subtracting the Level 0.5 nod2 reference spectrum from the associated nod1 reference spectrum of the scans that cover the frequency ranges of these lines. If the resulting spectrum showed only noise we assumed no emission in the reference spectra (we never experienced the case that emission in both reference beams cancelled out perfectly, which would hide this problem). If the resulting spectrum showed an emission line, there was emission in the nod1 reference beam; and if the spectrum showed an absorption feature, there was emission in the nod2 reference beam.

In order to remove the emission in the reference beams, we used the Herschel/HIPE hifiPipeline task to create the Level 0.5 product, in which the reference observations are still separated from the target observations. A custom HIPE script then extracted all affected reference scans and the neighboring scans taken before and after the affected scans that have slightly changed LO-settings such that the contaminating emission is at sufficiently different intermediate frequencies (IF). The repair is based on the assumption that the bandpass of all observations changes only very little with small changes in the LO-setting. The primary change is in the amplitude, while the shape of the bandpass changes negligibly. Thus, we could use the neighboring reference scans to repair the contaminated reference flux.

The first step of the repair was to determine the IF-frequency interval [f_l, f_h], covering the reference beam emission and two smaller, abutting intervals [f_l – Δ_l, f_l] and [f_h, f_h + Δ_h], indicated by the green areas in Fig 6, for scaling. The next step included extracting and averaging the flux of the neighboring scans over the frequency range [f_l – Δ_l, f_h + Δ_h] (Ref 1 and Ref 2 in Figure 6). Then, to properly scale the averaged flux to replace the contaminated flux, we determined the ratios of the original reference flux to the averaged flux over the two green intervals, interpolated the corresponding values over the interval containing the contaminated flux (the white area in Figure 6), and calculated the new reference flux by multiplying the averaged flux with the just determined ratio over the entire frequency range [f_l – Δ_l, f_h + Δ_h]. This new reference flux (New Ref) now replaced the original reference flux (Orig Ref). From here on, we continued to use the Herschel/HIPE hifiPipeline task to create the Level 1.0 and higher products.

Fig. 6 — Removing the reference beam emission in the HEXOS Orion-S spectral scan: The blue solid line shows the emission in a reference beam. The red dashed and the green dotted lines show the references for an earlier scan and a later scan with slightly changed LO-settings, but covering the same IF interval. Averaging and scaling these reference scans using just the frequency ranges marked with green background results in the new reference spectrum, which replaces the old reference spectrum only in the frequency range shown (the total spectrum is still 4 GHz wide or ~1 GHz for each of the 4 HIFI WBS sub-bands). For display purposes we subtracted the new reference from all scans causing it to appear as a straight line.

The lines that needed repairs are: B1a ([CI]), B2b (H₂S, C¹⁸O, ¹³CO, C¹⁷O), B3a (CO, [CI], CH⁺), B3b (C¹⁸O, ¹³CO, C¹⁷O, CO), B4b (H₂S, C¹⁸O, ¹³CO), B5a (C¹⁷O, CO, C¹⁸O, ¹³CO), B7a (CO), and B7b (CO, [CII]). Figure 7 shows an example of how the repair recovered the true line profile of the [CII] 158 μm line.

Fig. 7 — Example of the repaired [CII] line at 1900.526 GHz. The red dashed line shows the original result with emission in both reference beams and the blue solid line shows the result with the emission in the reference beams removed.

2.2. Line Identification

We used CASSIS³, a Java based software package designed to analyze astrophysical spectroscopic data to perform the line identification and modeling. Our line identification procedure involved two main steps. First we visually identified the strongest (well above 5σ signal-to-noise) and best known emission lines in the spectrum (e.g. from CO, CS, HCO⁺, HCN, H₂O, etc.) and some of their isotopologues utilizing the JPL⁴ (Pearson et al. 2005) and CDMS⁵ (Müller et al. 2005) spectral line databases. These databases include tabulated values of the central frequency error for each transition. Although in some cases, the difference between the listed centroid frequency of a particular transition from these two catalogues is larger than their given error bars, the observed line width of the transition usually compensates for this ambiguity and makes the identification robust. Many of the strongest identified species are shown in Figures 1-4. In these cases, line blending (e.g. the appearance of more than one transition/species at a single frequency) is not considered a problem since the emission from the well-known species will invariably overwhelm the weak emission from a less well-known and, presumably, lower abundance blended line.

Once the strongest emission lines were accounted for, we examined all other spectral lines in our data that had a signal-to-noise ratio above 3σ (in peak intensity). We first performed the line identification via visual inspection of each spectral feature in each HIFI band and compared the transition’s frequency to those listed in the databases. From the possible database entries we investigated all species with transitions that fell within a Doppler velocity range of 5.5 to 8.5 km s⁻¹ (i.e. within ±1.5 km s⁻¹ of the assumed central velocity of Orion-S). Within this velocity range we examined a smaller sub-sample of possible spectral lines with upper state excitation energies (E_up) less than 1500 K. If a single database entry from this sub-sample matched the observed spectral feature, we considered this to be a tentative identification. In order to confirm or reject this tentative identification, we then searched for other predicted transitions of the selected species in all of the HIFI bands. If we saw other spectral features in the data that matched the predicted frequencies, we accepted the initial line identification as likely correct. If we did not, then the initial line identification was still considered to be only tentative, since we realize that the absence of other predicted transitions may be due to special excitation conditions. Therefore, in both cases, in order to finally confirm or reject our initial line identifications, Local Thermodynamic Equilibrium (LTE) modeling was performed (described in detail in Section 3), which allowed us to determine if all the observed spectral features from the tentatively detected species could be theoretically reproduced under uniform excitation conditions.

Our LTE models explored excitation conditions with T_ex ≤ 1500 K (where T_ex is the excitation temperature - equal to the kinetic temperature in LTE), E_up ≤ 1500 K, and total species column density ≤ 10¹⁷ cm⁻². If the LTE model produced emission at the frequency of the spectral feature then the line identification was considered confirmed. If not, the species was assumed to have been incorrectly identified and a new identification for that spectral feature was sought. Note, at this stage we are simply trying to produce some visible model emission at the frequency of the spectral feature and not trying to fit or replicate the observed spectral line profile This will be performed in a subsequent stage described in Section 3.1. If a spectral feature could not be reproduced by an LTE model of any species, or if there was no database entry at the frequency of the observed spectral feature, that feature was listed as an unidentified line (32 lines in total). Visual inspection of the original, DSB spectra indicates that all of these features are “ghosts” (i.e. artifacts of the deconvolution routine). A list of all identified species is given in Table 2. A frequency ordered list of all spectral features above 3σ in intensity (as well as their peak intensity) is given in Table 3. Ghost lines are identified as “ghost”. In total we identified 52 different species (including isotopologues) which are responsible for 685 transitions (including the blended lines) in the HIFI spectra. It is, of course, possible that additional species and transitions exist in Orion-S, but at intensities too weak to be detected. This will be addressed in Section 3.3.2.

Table 2. Identified species in Orion-S.

atoms	di-atomic molecules	multi-atomic molecules	Ionized species
C-atom	CO	CCH	C⁺
	¹³CO	DCN	CH⁺
	¹³C¹⁸O	HNC	CO⁺
	C¹⁸O	HCN	DCO⁺
	C¹⁷O	H¹³CN	H¹³CO⁺
	CS	HC¹⁵N	HC¹⁸O⁺
	¹³CS	HDO ^a	HCO⁺
	C³⁴S	o−H₂O	HCS⁺
	CH	p−H₂O	N₂H⁺
	CN	$o - H_{2}^{18} O$ ^a	SH⁺
	HCl ^a	$p - H_{2}^{18} O$ ^a
	H³⁷Cl ^a	o−H₂S
	HF ^a	p−H₂S
	NO ^a	$H_{2}^{33} S$
	SiO	$H_{2}^{34} S$
	SO	H₂CS
		o−H₂CO
		p−H₂CO
		o−NH₃
		p−NH₃
		A−CH₃OH
		E−CH₃OH
		CH₃OCH₃
		NH₂
		SO₂

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JPL database used for these species. For all other species CDMS database is used.

Table 3. All observed lines above the 3σ noise level in order of increasing frequency.


Band 1a

Frequency (GHz)	T_A (K)	Species

480.2699	0.7	A−CH₃OH
481.5056	0.7	A−CH₃OH
481.9167	0.3	C³⁴S
482.2179	0.1	A−CH₃OH
482.2833	0.8	E−CH₃OH
482.9598	1.4	E−CH₃OH
483.0808	0.2	A−CH₃OH
483.1418	1.7	E−CH₃OH
483.3898	0.3	E−CH₃OH
483.3983	0.1	A−CH₃OH
483.4623	0.2	A−CH₃OH
483.4728	0.2	E−CH₃OH
483.5393	0.2	A−CH₃OH
483.5533	0.3	A−CH₃OH
483.5668	0.3	E−CH₃OH
483.5818	0.1	A−CH₃OH
483.6868	0.6	E−CH₃OH
483.7633	0.3	A−CH₃OH
484.0058	0.8	A−CH₃OH
484.0238	0.7	E−CH₃OH
484.0718	0.4	E−CH₃OH
484.2703	0.2	SO₂
485.2638	0.9	A−CH₃OH
486.9419	0.9	A−CH₃OH
487.5319	0.6	A−CH₃OH
487.6639	0.2	H₂CS
489.0379	0.9	A−CH₃OH
489.7509	4.5	CS
490.5970	0.2	HDO
491.5520	0.7	A−CH₃OH
491.9335	0.2	SO₂
491.9690	3.1	o−H₂CO
492.1615	4.7	[CI]
492.2795	2.0	A−CH₃OH
492.7841	0.1	CH₃OCH₃
493.7000	1.4	A−CH₃OH
493.7350	1.4	A−CH₃OH
494.4820	0.6	A−CH₃OH
494.7781	0.2	SO₂
495.1741	1.2	E−CH₃OH
496.9226	0.1	E−CH₃OH
497.8296	0.4	A−CH₃OH
501.5897	0.4	A−CH₃OH
503.0142	0.1	$H_{2}^{34} S$
504.2008	0.2	DCO⁺
504.2948	1.3	E−CH₃OH
504.6783	0.1	SO
505.5658	1.5	o−H₂S
505.7633	0.3	A−CH₃OH
505.8343	1.7	p−H₂CO
506.1548	0.2	E−CH₃OH
506.7728	0.2	H₂CS
506.8273	0.2	DCN
508.5369	0.1	¹³CS
508.7069	0.1	SO₂
509.0924	0.1	A−CH₃OH
509.1469	0.8	p−H₂CO
509.2939	0.3	HDO
509.5654^b	0.8	E−CH₃OH
		+ o−H₂CO
509.8314	0.3	p−H₂CO
510.1559	0.9	o−H₂CO
510.2389	0.9	o−H₂CO
510.3459	0.3	A−CH₃OH
510.9114	0.1	HC¹⁸O⁺
511.0904	0.2	SO₂
511.5024	0.1	SO₂
511.7166	0.1	CH₃OCH₃
511.9455	0.2	HCS⁺
513.0775	0.7	p−H₂CO
513.3610	0.1	H₂CS
513.3698	0.1	CH₃OCH₃
514.8535	0.7	SO
515.1700	0.2	A−CH₃OH
515.3335	0.2	A−CH₃OH
515.8230	0.1	Ghost
516.2616	0.2	HC¹⁵N
516.3361^b	0.7	SO
		+ H₂CS
517.3551	0.9	SO
517.9706	0.4	H¹³CN
520.1801	1.2	E−CH₃OH
520.4612	1.1	H¹³CO⁺
520.7297	0.1	A−CH₃OH
520.8817	0.1	SiO
522.4057	0.1	H₂CS
523.1002	0.0	E−CH₃OH
523.2752	0.3	E−CH₃OH
523.4827	0.1	¹³C¹⁸O
523.9727	2.2	CCH
524.0352	1.8	CCH
524.2682	0.1	E−CH₃OH
524.5847	0.1	E−CH₃OH
524.6673	0.1	E−CH₃OH
524.7418	0.2	E−CH₃OH
524.8053	0.2	E−CH₃OH
524.8628	0.2	E−CH₃OH
524.9098	0.2	E−CH₃OH
524.9483	0.1	E−CH₃OH
525.0548	0.1	E−CH₃OH
525.6663	2.5	o−H₂CO
526.0363	0.1	SH⁺
526.0453	0.1	SH⁺
526.5233	0.1	A−CH₃OH
526.5483	0.1	E−CH₃OH
526.7878	0.1	E−CH₃OH
527.0548	0.4	A−CH₃OH
527.1738	0.1	E−CH₃OH
527.6608	0.1	E−CH₃OH
528.1813	0.1	E−CH₃OH
528.6833	0.1	E−CH₃OH
529.1419	0.2	E−CH₃OH
529.2904	0.1	SO₂
529.5409	0.2	E−CH₃OH
529.8679	0.2	E−CH₃OH
529.9739	0.2	SO₂
530.0699	0.2	A−CH₃OH
530.1234	0.3	C³⁴S
530.1849	0.6	E−CH₃OH
530.0244	0.1	E−CH₃OH
530.3174	0.4	E−CH₃OH
530.4559	0.5	E−CH₃OH
530.5499	0.5	E−CH₃OH
530.6124	0.5	E−CH₃OH
530.6484	0.5	E−CH₃OH
530.8244	0.2	E−CH₃OH
531.0804	1.2	E−CH₃OH
531.3199	1.3	A−CH₃OH
531.6384	0.2	A−CH₃OH
531.7159	7.5	HCN
531.8709	0.4	A−CH₃OH
531.8924	0.3	A−CH₃OH
532.0334	0.4	E−CH₃OH
532.0709	0.2	E−CH₃OH
532.1349	0.3	A−CH₃OH
532.3239	0.1	CH₃OCH₃
532.4669	0.6	E−CH₃OH
532.5684	0.4	E−CH₃OH
532.7214	2.4	CH
532.7909	1.0	CH
533.3810	0.1	Ghost
535.0610	14.6	HCO⁺
536.1925	0.6	A−CH₃OH
536.7585	2.1	CH
536.7805	0.5	CH
536.7925	0.9	CH
538.5716	2.2	A−CH₃OH
538.6891	3.4	CS
539.2806	0.1	E−CH₃OH
540.4656	0.2	H₂CS
540.9236	0.1	E−CH₃OH
541.7536	0.2	SO₂
541.8162	0.1	SO₂
542.0022	1.3	A−CH₃OH
542.0832	1.4	A−CH₃OH
543.0777	1.2	E−CH₃OH
543.8987	0.1	HNC
545.0437^b	0.2	E−CH₃OH + A−CH₃OH
545.1032	0.2	Ghost
545.8872	0.2	Ghost
546.2488	0.2	A−CH₃OH
547.3003	0.2	Ghost
547.6768	0.3	$H_{2}^{18} O$
548.8313	8.1	C¹⁸O
549.2998	0.1	SO₂
549.5506	0.1	CH₃OCH₃
549.5533	0.1	NO
550.6564	0.0	A−CH₃OH
550.9254	29.4	¹³CO
551.1874	0.3	NO
551.5159	0.1	Ghost
551.5344	0.3	NO
553.1474	1.0	E−CH₃OH
554.0569	0.2	E−CH₃OH
554.5784	0.1	HCS⁺
555.6670	0.1	SO₂
556.9370	7.2	o−H₂O
557.1260	0.1	H₂CS
558.0870	0.3	SO
558.3465	0.4	E−CH₃OH
558.9681	2.5	N₂H⁺
559.3216	0.4	SO
560.1781	0.5	SO
560.2496	0.2	E−CH₃OH
560.2716	0.2	E-CH₃OH
560.3001	0.2	Ghost


Band 1b

Frequency (GHz)	T_A (K)	Species

556.9370	8.4	o−H₂O
558.0875	0.5	SO
558.3460	0.5	E−CH₃OH
558.9681	3.4	N₂H⁺
559.3206	0.5	SO
560.1781	0.6	SO
561.7136	3.0	C¹⁷O
561.9001	2.7	o−H₂CO
564.2522	0.1	SiO
566.7312	1.3	CN
566.9477	1.6	CN
567.2612	0.2	Ghost
567.5953	0.1	SO₂
568.2358	0.2	E−CH₃OH
568.4353	0.2	CH₃OCH₃
568.5673	1.3	E−CH₃OH
568.7853	0.2	A−CH₃OH
570.2434	0.1	CH₃OCH₃
572.4974	3.3	o−NH₃
572.8984	0.2	E−CH₃OH
574.1399	0.1	H₂CS
574.8084	0.1	SO₂
574.8694	0.3	A−CH₃OH
576.2050^b	0.4	DCO⁺
576.2660	62.0	CO
576.3835	0.2	Ghost
576.7095	1.1	p−H₂CO
578.0080	0.4	E−CH₃OH
578.2165	0.2	C³⁴S
579.0860	1.4	A−CH₃OH
579.1520	0.8	E−CH₃OH
579.1995	0.1	DCN
579.4605	0.9	A−CH₃OH
579.8590	0.2	A−CH₃OH
579.9225	1.4	A−CH₃OH
580.0600	0.1	A−CH₃OH
580.1760	0.2	A−CH₃OH
580.2135	0.2	A−CH₃OH
580.3696	0.3	E−CH₃OH
580.4446	0.2	E−CH₃OH
580.5026	0.2	A−CH₃OH
580.9036	0.4	E−CH₃OH
581.0916	0.3	E−CH₃OH
581.6131	0.5	p−H₂CO
582.3841	0.2	o−H₂CO
582.7246	0.2	p−H₂CO
583.1456	0.6	o−H₂CO
583.3096	0.5	o−H₂CO
584.4511	1.9	A−CH₃OH
584.8227	0.5	A−CH₃OH
587.4537	0.4	p−H₂CO
587.5702	0.2	SO₂
587.6167	2.3	CS
589.1653	0.1	CH₃OCH₃
589.8703	0.1	CO⁺
590.2793	1.2	A−CH₃OH
590.4418	1.1	A−CH₃OH
590.7522	0.1	CH₃OCH₃
590.7923	0.9	E−CH₃OH
591.8218	0.1	H₂CS
593.9424	0.1	SO₂
593.9624	0.1	SO₂
597.2084	0.1	HCS⁺
598.5485	-0.1	Ghost
599.9280	0.3	HDO
600.3320	1.7	o−H₂CO
600.9065	0.1	¹³CS
601.2570	0.4	SO
601.8521	0.1	E−CH₃OH
602.2346	0.8	E−CH₃OH
602.2750		HC¹⁵N
602.2916	0.4	SO
603.0236	0.5	SO
604.2641	0.2	H¹³CN
604.3696	0.2	SO₂
605.8801	0.1	E−CH₃OH
607.1762	0.8	H¹³CO⁺
607.2167	0.3	E−CH₃OH
607.6112	0.1	SiO
607.7982	0.2	H₂CS
608.0984	0.1	CH₃OCH₃
609.7087	0.1	CH₃OCH₃
611.2688	1.2	CCH
611.3308	1.0	CCH
611.5518	0.1	SO
613.0798	0.1	SO₂
616.9814	1.2	E−CH₃OH
620.3035	5.6	HCN
620.7045	0.6	o−H₂O
622.5705	0.1	A−CH₃OH
622.6605	0.3	A−CH₃OH
622.7755	0.2	E−CH₃OH
624.1801	0.2	E−CH₃OH
624.2081	13.3	HCO⁺
624.9616	1.0	H³⁷Cl
624.9751	1.4	H³⁷Cl
624.9856	0.6	H³⁷Cl
625.7521	0.4	E−CH₃OH
625.7601	0.1	A−CH₃OH
625.8996	2.1	HCl
625.9166	2.8	HCl
625.9296	1.5	HCl
626.0896	0.2	SO₂
626.3521	0.1	C³⁴S
626.4996	0.2	H₂CS
626.5141	0.1	A−CH₃OH
626.5566	0.2	E−CH₃OH
626.6271	1.5	A−CH₃OH
627.0203	0.1	CH₃OCH₃
627.1036	0.1	Ghost
627.1721	0.6	E−CH₃OH
627.5602	0.7	A−CH₃OH
628.0522	0.2	A−CH₃OH
628.1422	0.2	¹³C¹⁸O
628.3302	0.2	E−CH₃OH
628.4482	0.2	E−CH₃OH
628.4717	0.3	A−CH₃OH
628.5142	0.2	A−CH₃OH
628.5272	0.2	A−CH₃OH
628.6627	0.1	CH₃OCH₃
628.6982	0.3	E−CH₃OH
628.8182	0.2	E−CH₃OH
628.8682	0.2	A−CH₃OH
629.1417	1.6	A−CH₃OH
629.3232	0.4	E−CH₃OH
629.6517	0.3	E−CH₃OH
629.9222	1.7	A−CH₃OH
631.2847	0.2	Ghost
631.7038	1.7	o−H₂CO
632.1918	0.2	SO₂
633.4248	0.4	A−CH₃OH
634.5118	0.6	HNC
635.8673	0.3	A−CH₃OH
636.3394	0.5	A−CH₃OH
636.3979	0.4	A−CH₃OH
636.4209	0.6	A−CH₃OH


Band 2a

Frequency (GHz)	T_A (K)	Species

626.5191	0.7	A−CH₃OH
626.6271	1.3	A−CH₃OH
627.1711	0.6	E−CH₃OH
627.5607	0.5	A−CH₃OH
629.1417	1.0	A−CH₃OH
629.3242	0.3	E−CH₃OH
629.6482	0.3	E−CH₃OH
629.9232	1.2	A−CH₃OH
631.7038	1.3	o−H₂CO
633.4233	0.3	A−CH₃OH
634.5113	0.5	HNC
636.2514	0.2	A−CH₃OH
636.2749	0.3	A−CH₃OH
636.3059	0.3	A−CH₃OH
636.3349	0.3	A−CH₃OH
636.3669	0.4	A−CH₃OH
636.3949	0.4	A−CH₃OH
636.4209	0.3	A−CH₃OH
636.5199	0.3	A−CH₃OH
636.5339	1.3	CS
638.2804	0.7	E−CH₃OH
638.5259	1.0	A−CH₃OH
638.8194	0.9	A−CH₃OH
644.3252	0.2	NH₂
644.3790	0.4	SO
645.2576	0.4	SO
645.8756	0.4	SO
645.9282	0.1	CH₃OCH₃
647.0831	0.2	p−H₂CO
647.6130	0.2	CH₃OCH₃
648.1956	0.2	DCO⁺
649.5402	0.2	E−CH₃OH
651.2997	0.1	SO₂
651.4352	0.3	NO
651.6177	0.6	E−CH₃OH
651.7742	0.4	NO
652.0972	2.8	N₂H⁺
653.9703	0.4	p−H₂CO
655.2128	0.2	o−H₂CO
655.6353	0.2	p−H₂CO
656.1653	0.6	o−H₂CO
656.1723	0.4	E−CH₃OH
656.4663	0.5	o−H₂CO
658.5544	7.7	C¹⁸O
661.0659	30.3	¹³CO
662.2105	0.4	p−H₂CO
664.8193	0.1	CH₃OCH₃
665.2485	0.2	SO₂
665.4440	1.3	E−CH₃OH
670.3602	0.2	SO₂
670.4242	0.3	A−CH₃OH
672.5657	0.2	SO₂
672.8362	0.1	A−CH₃OH
672.9037	0.3	E−CH₃OH
673.4172	0.3	E−CH₃OH
673.7472	1.4	A−CH₃OH
674.0107	2.4	C¹⁷O
674.8113	1.1	o−H₂CO
674.9913^b	1.4	A−CH₃OH
		+ E−CH₃OH
675.0558	0.2	H₂CS
675.1353	0.5	E−CH₃OH
675.6138	0.5	A−CH₃OH
675.7748	0.8	E−CH₃OH
676.2143	0.3	A−CH₃OH
676.4618	0.1	SO₂
676.7523	0.3	A−CH₃OH
676.8298	0.3	A−CH₃OH
677.0138	0.3	E−CH₃OH
677.7113	0.4	E−CH₃OH
678.2528	0.2	E−CH₃OH
678.7688	0.2	Ghost
678.7864	1.6	A−CH₃OH
680.0429	0.6	CN
680.2599	0.8	CN
681.5310	-0.1	Ghost
681.9919	0.4	A−CH₃OH
685.4375	1.0	CS
685.5050	0.6	E−CH₃OH
686.7325	0.7	A−CH₃OH
687.0220	0.1	$H_{2}^{34} S$
687.1515	0.2	$H_{2}^{33} S$
687.2270	0.8	A−CH₃OH
687.3040	2.3	p−H₂S
687.4571	0.4	SO
688.2021	0.3	SO
688.7286	0.3	SO
690.5551	0.2	H¹³CN
691.4706	64.4	CO
693.8787	0.4	H¹³CO⁺
697.1083	0.1	Ghost
697.1448	0.2	E−CH₃OH
698.5458	0.6	CCH
698.6083	0.6	CCH
701.3709^b	1.3	o−H₂CO
		+ E−CH₃OH
704.2609	0.2	Ghost
704.4134	0.2	CH₃OCH₃
704.9400	0.5	Ghost
705.1840	0.3	E−CH₃OH
705.4000	0.2	Ghost
705.9605	0.4	Ghost
706.6255	0.3	Ghost
707.3180	0.4	Ghost
707.7905	0.4	E−CH₃OH
708.4735^b	0.5	o−H₂S
	0.0	$+ H_{2}^{33} S$
708.7085	0.4	Ghost
708.8741	3.2	HCN
709.2771	0.3	Ghost
713.3422	10.0	HCO⁺
713.9837	0.9	E−CH₃OH
716.9402	0.5	p−H₂CO
718.1618	0.3	A−CH₃OH
718.4378	0.9	E−CH₃OH
719.6653	1.1	A−CH₃OH
719.7293	1.0	Ghost
719.9548	1.5	Ghost
720.0118	0.9	E−CH₃OH
720.4418	1.2	A−CH₃OH
721.0118	0.2	E−CH₃OH
723.0404	0.4	E−CH₃OH
723.6219	0.4	A−CH₃OH
724.1224	0.7	E−CH₃OH
724.3489	0.4	A−CH₃OH
725.0954	0.9	Ghost
716.9412	0.6	p−H₂CO


Band 2b

Frequency (GHz)	T_A (K)	Species

718.1623	0.3	A−CH₃OH
718.4378	1.0	E−CH₃OH
719.6668	1.2	A−CH₃OH
720.4433	1.4	A−CH₃OH
721.0093	0.3	E−CH₃OH
723.0429	0.5	E−CH₃OH
723.2824	0.3	E−CH₃OH
723.6199	0.5	A−CH₃OH
724.1219	0.9	E−CH₃OH
724.3409	0.2	A−CH₃OH
725.1089	0.3	HNC
725.1279	0.3	A−CH₃OH
725.3169	0.2	E−CH₃OH
726.0540	0.2	E−CH₃OH
726.2100	0.4	p−H₂CO
726.9025	0.3	E−CH₃OH
728.0535	0.3	o−H₂CO
728.5860	0.3	p−H₂CO
728.5960	0.2	p−H₂CO
728.8650	1.5	A−CH₃OH
729.2120	0.5	o−H₂CO
729.7260	0.4	o−H₂CO
730.5006	0.4	SO
730.5206	0.4	A−CH₃OH
731.1416	0.3	SO
731.5981	0.3	SO
732.4336	0.6	A−CH₃OH
732.7761	0.2	¹³C¹⁸O
734.2706	0.4	$H_{2}^{34} S$
734.3286	0.8	CS
734.8952	0.6	A−CH₃OH
735.6762	0.6	A−CH₃OH
736.0317	4.3	o−H₂S
737.3407	0.4	p−H₂CO
737.6272	0.2	Ghost
741.2263	0.2	E−CH₃OH
742.2408	0.2	H₂CS
745.2119	1.8	N₂H⁺
747.3035	0.3	p−H₂S
749.0735	0.8	o−H₂CO
751.5560	0.4	E−CH₃OH
751.6766	0.3	NO
752.0321	7.9	p−H₂O
752.1076	0.6	E−CH₃OH
752.1366	0.4	E−CH₃OH
752.1716	0.3	E−CH₃OH
752.3111	0.3	E−CH₃OH
753.4161	0.2	HDO
753.8671	0.1	E−CH₃OH
754.2226	0.2	E−CH₃OH
762.6378	0.9	E−CH₃OH
763.8823	0.3	E−CH₃OH
763.9533	0.9	A−CH₃OH
764.5828	-0.5	Ghost
764.8119	0.2	E−CH₃OH
765.5134	0.2	E−CH₃OH
765.9404	0.2	o−H₂S
766.0309	0.3	E−CH₃OH
766.3974	0.4	E−CH₃OH
766.6489	0.5	E−CH₃OH
766.7124	1.2	A−CH₃OH
766.7624	0.8	E−CH₃OH
766.8114	0.4	E−CH₃OH
766.9094	0.5	E−CH₃OH
766.9614	0.5	E−CH₃OH
766.9844	0.5	E−CH₃OH
768.2529	5.6	C¹⁸O
768.5404	0.2	E−CH₃OH
770.8855	0.4	E−CH₃OH
770.8980	0.9	o−H₂CO
771.1825	27.1	¹³CO
771.5770	0.4	A−CH₃OH
772.4425	0.2	A−CH₃OH
772.4545	0.6	E−CH₃OH
773.2611	0.2	A−CH₃OH
773.4226	0.3	A−CH₃OH
773.5136	0.3	SO
773.8896	0.2	E−CH₃OH
773.9481	0.2	A−CH₃OH
774.0666	0.2	SO
774.3331	0.3	E−CH₃OH
774.4541	0.3	SO
775.5996	0.4	E−CH₃OH
779.0072	0.5	A−CH₃OH
779.0322	0.4	E−CH₃OH
779.3822	1.3	A−CH₃OH
780.5672	0.3	H¹³CO⁺
783.0028	0.5	A−CH₃OH
783.1993	0.6	CS
784.1793	0.3	A−CH₃OH
785.8058	0.3	CCH
785.8679	0.2	CCH
786.2829^b	1.5	C¹⁷O
		+ p−H₂CO
790.9360	0.7	Ghost
793.3410	0.4	CN
793.5480	0.3	CN
794.5211	0.2	Ghost
794.8206	0.2	Ghost
797.4306	1.7	HCN
798.3106	0.2	p−H₂CO


Band 3a

Frequency (GHz)	T_A (K)	Species

802.2430	0.3	E−CH₃OH
802.2810	0.3	o−H₂CO
802.4590	7.1	HCO⁺
803.1130	0.3	o−H₂CO
806.6501	59.4	CO
807.8661	0.6	A−CH₃OH
809.3431	6.8	[CI]
811.4452	0.6	E−CH₃OH
812.5522	0.9	A−CH₃OH
813.5442	0.2	A−CH₃OH
815.0723	0.8	E−CH₃OH
815.6943	0.1	HNC
816.0153	0.2	E−CH₃OH
816.4953	0.2	SO
816.9743	0.2	SO
817.3143	0.2	SO
818.6674	0.2	E−CH₃OH
819.4844	0.3	A−CH₃OH
820.5024	0.2	A−CH₃OH
820.7644	0.4	E−CH₃OH
821.4774	0.2	A−CH₃OH
821.7014	0.2	A−CH₃OH
821.8694	0.2	E−CH₃OH
822.5475	0.3	E−CH₃OH
823.0845	0.6	o−H₂CO
824.3525	0.2	E−CH₃OH
824.7255	0.2	E−CH₃OH
825.2795	0.3	E−CH₃OH
827.4516	0.2	A−CH₃OH
829.8906	0.7	A−CH₃OH
830.3506	1.0	A−CH₃OH
831.0487	0.3	A−CH₃OH
832.0627	0.5	CS
832.7567	0.4	A−CH₃OH
834.6517	0.2	E−CH₃OH
834.7517	0.1	E−CH₃OH
834.8397	0.2	E−CH₃OH
834.9017	0.1	E−CH₃OH
834.9577	0.2	E−CH₃OH
835.0037	0.1	E−CH₃OH
835.1358	3.0	CH⁺
838.3088	0.8	N₂H⁺
840.2779	0.6	o−H₂CO
851.4161	0.7	A−CH₃OH
851.9151	0.3	NO
853.5112	0.3	E−CH₃OH
855.1542	0.3	p−H₂CO
857.9613	0.7	A−CH₃OH


Band 3b

Frequency (GHz)	T_A (K)	Species

860.4618	0.5	E−CH₃OH
863.3639	0.6	E−CH₃OH
863.4269	0.2	E−CH₃OH
867.3250	0.3	A−CH₃OH
869.0370	0.5	E−CH₃OH
869.9821	0.3	A−CH₃OH
870.1091	0.2	E−CH₃OH
870.2811	0.2	p−H₂CO
873.0361	0.2	CCH
873.7811	0.1	o−H₂CO
875.3692	0.2	o−H₂CO
876.6462	0.1	o−H₂CO
877.9232	3.4	C¹⁸O
878.2273	0.7	A−CH₃OH
879.0153	0.2	A−CH₃OH
880.9043	0.3	CS
881.2703	21.6	¹³CO
881.4223	0.3	A−CH₃OH
881.7833	0.8	A−CH₃OH
885.9664	1.2	HCN
890.1265	0.3	E−CH₃OH
891.5586	5.5	HCO⁺
893.6387	-0.7	HDO
894.6146	0.5	A−CH₃OH
896.8077	0.3	o−H₂CO
898.5247	0.9	C¹⁷O
898.9777	0.2	Ghost
900.9658	0.1	E−CH₃OH
902.9388	0.5	A−CH₃OH
902.9838	0.2	SO
905.3959	0.2	E−CH₃OH
905.3959	0.2	E−CH₃OH
906.5939	0.2	CN
906.8029	0.2	CN
907.4303	0.2	NH₂
909.5120	0.3	o−H₂CO
909.7400	0.4	E−CH₃OH
910.8110	0.2	E−CH₃OH
911.6440	0.5	E−CH₃OH
912.1100	0.8	E−CH₃OH
916.1771	0.5	p−H₂O
916.6502	0.3	E−CH₃OH
917.2702	0.4	E−CH₃OH
917.4082	0.2	E−CH₃OH
921.7973	59.5	CO
921.9873	0.5	E−CH₃OH
923.5853	0.3	p−H₂CO
926.5564	0.5	A−CH₃OH
926.8944	0.3	A−CH₃OH
929.7315	0.3	CS
930.2035	0.3	A−CH₃OH
931.3885	0.4	N₂H⁺
933.6945	0.6	A−CH₃OH
937.4826	0.4	A−CH₃OH
947.4759	0.3	A−CH₃OH
952.5422	-0.3	NH₂
952.5740	-0.8	NH₂
952.6268	-0.3	NH₂


Band 4a

Frequency (GHz)	T_A (K)	Species

960.4732	0.6	E−CH₃OH
965.4513	0.4	E−CH₃OH
974.4895	0.7	HCN
974.6895	0.3	A−CH₃OH
974.8795	0.5	A−CH₃OH
980.0316	0.4	A−CH₃OH
980.6376	4.3	HCO⁺
986.1018	0.6	A−CH₃OH
987.5618	2.2	C¹⁸O
987.9298	8.5	p−H₂O
991.3259	17.6	¹³CO
991.5839	0.5	A−CH₃OH
993.1019	1.7	o−H₂S
1002.7782	0.8	p−H₂S
1006.1242	0.4	E−CH₃OH
1008.8183	0.6	E−CH₃OH
1010.7323	0.6	C¹⁷O
1013.5664	0.4	E−CH₃OH
1023.1986	0.6	A−CH₃OH
1036.6990	1.1	SO₂
1036.9070	66.9	CO
1039.0150	0.5	A−CH₃OH
1057.1194	0.5	E−CH₃OH


Band 4b

Frequency (GHz)	T_A (K)	Species

1057.1239	0.5	E−CH₃OH
1062.9821	1.0	HCN
1069.6982	2.8	HCO⁺
1072.8303	1.0	o−H₂S
1092.4668	0.4	A−CH₃OH
1097.1589	1.4	C¹⁸O
1097.3689	3.4	o−H₂O
1101.3450	12.5	¹³CO
1101.6990	-0.7	$p - H_{2}^{18} O$
1105.3691	0.4	E−CH₃OH
1113.3472	4.0	p−H₂O


Band 5a

Frequency (GHz)	T_A (K)	Species

1113.3477	3.5	p−H₂O
1119.8339	0.4	A−CH₃OH
1122.9060	0.3	C¹⁷O
1151.4506	0.6	HCN
1151.7555	0.9	A−CH₃OH
1151.9826	51.0	CO
1152.9207	0.6	A−CH₃OH
1153.1267	4.4	o−H₂O
1153.5527	0.6	E−CH₃OH
1158.7328	1.7	HCO⁺
1162.7149^b	0.6	A−CH₃OH
		+ E−CH₃OH
1162.9129	4.4	o−H₂O
1168.1190	0.5	A−CH₃OH
1168.4524	-1.1	p−NH₃
1196.0197	0.5	o−H₂S
1206.7279	0.7	C¹⁸O
1207.6540	0.5	p−H₂O
1211.3260	6.8	¹³CO
1214.8529	-1.0	o−NH₃
1215.2457	-1.3	p−NH₃
1228.7924	1.3	p−H₂O
1232.4685	-1.2	HF


Band 5b

1228.7922	1.7	p−H₂O
1232.4692	-1.2	HF
1239.9125	0.4	HCN
1247.7406	0.8	HCO⁺
1267.0091	44.6	CO


Band 6a

1496.9192	42.0	CO


Band 6b

1611.7859	37.3	CO
1669.9169	7.7	o−H₂O


Band 7a

1726.5976	26.9	CO


Band 7b

1841.3367	20.1	CO
1900.5261	69.5	[CII]

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In some cases, our LTE modeling resulted in a particular spectral feature being reasonably explained by a superposition of lines from more than one species/transition (i.e. blended lines). In order to account for the possible effects of line blending, we performed modeling of all transitions of the two species in all HIFI bands. After obtaining good fits to the unblended transitions, we were able to determine how much each species contributed to the blended feature. An example of a blended line is shown in Figure 8, in which the data are shown by the black histogram, the solid red line indicates a C¹⁷O transition, the solid blue line indicates an H₂CO transition and the solid green line is the superposition of these two model results. Note that the solid green line in Figure 8 does not represent a multicomponent Gaussian fit but rather the LTE modeling required to reproduce the line profiles (see Section 3.1). Given the relatively few spectral lines in Orion-S, significant blending was only a problem in 6 of 685 lines detected. These blended lines are indicated by a “b” superscript in Tables 3 and 4. In Table 3 both species are listed. Blended lines were excluded from the modeling analysis presented in Section 3.

Fig. 8 — Example of a blended line at 786.3 GHz, produced by C¹⁷O and H₂CO. The red line shows the LTE modeled synthetic spectrum for C¹⁷O and the blue line shows that for H₂CO (Table 5 and 6). The green line is the superposition of these two components. Data are shown by the black histogram.

Table 4. Gaussian fits to lines above 5σ.

[CI]

	Transition ^2s+1L_\|L+S\|	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	³P₁ → ³ P₀	492161.28	5.2	7.6	4.7	26.1	23.62
	³P₂ → ³ P₁	809342.91	6.9	7.2	4.2	30.8	62.46

[CII]

	Transition ^2s+1L_\|L+S\|	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	²P_3/2 → ² P_1/2	1900527.55	69.1	8.6	3.8	279.8	91.21

CCH

	Transition N_J,F₁	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	6_6.5,7 → 5_5.5,6	523972.07	2.2	6.6	4.5	10.7	88.02
	6_5.56 → 5_4.5,5	524034.40	1.9	6.7	4.0	8.1	88.04
	7_7.5,8 → 6_6.5,7	611267.06	1.2	6.8	5.2	6.4	117.36
	7_6.5,6 → 6_5.5,5	611329.64	1.0	7.0	4.9	5.0	117.38
	8_8.5,9 → 7_7.5,8	698544.47	0.6	7.4	5.9	3.7	150.88
	8_7.5,7 → 7_6.5,6	698607.43	0.5	7.0	5.2	2.8	150.91
	9_9.5,10 → 8_8.5,9	785801.64	0.3	7.9	8.2	2.3	188.59
	9_8.5,8 → 8_7.5,7	785864.19	0.2	7.0	5.8	1.3	188.62

CH

	Transition N_K,J,F₁	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	1_1,1.5,1 → 1_−1,0.5,1	532721.73	2.4	7.1	4.2	10.8	25.73
	1_1,1.5,1 → 1_−1,0.5,0	532791.46	1.0	8.4	4.8	4.9	25.73
	1_−1,1.5,2 → 1_1,0.5,1	536759.12	2.2	8.2	4.0	9.2	25.73
	1_−1,1.5,1 → 1_1,0.5,1	536779.69	0.6	8.3	3.8	2.2	25.76
	1_−1,1.5,1 → 1_1,0.5,0	536793.54	1.0	8.3	4.1	4.2	25.76

CH⁺

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	1 → 0	835135.84	4.0	8.3	4.8	20.4	40.08

A−CH₃OH

	Transition $J_{+ K, v_{t}^{π}}$	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	3_2,0⁺ → 3_1,0⁻	480269.83	0.7	6.7	3.7	2.7	51.64
	2_2,0⁺ → 2_1,0⁻	481505.55	0.7	6.1	3.3	2.4	44.67
	10_0,0⁺ → 9_0,0⁺	483141.47	1.5	6.6	4.3	6.9	127.6
	10_2,0⁻ → 9_2,0⁻	483388.95	0.3	6.8	4.9	1.4	165.35
	10_3,0⁺ → 9_3,0⁺	483553.40	0.2	5.5	6.4	1.7	177.46
	10_3,0⁻ → 9_3,0⁻	483566.43	0.3	6.5	3.9	1.2	177.46
	10_2,0⁺ → 9_2,0⁺	483762.05	0.2	6.6	5.7	1.4	165.4
	2_2,0⁻ → 2_1,0⁺	484005.49	0.7	6.5	4.2	3.0	44.67
	3_2,0⁻ → 3_1,0⁺	485263.94	0.8	6.6	4.0	3.5	51.64
	4_2,0⁻ → 4_1,0⁺	486941.55	0.9	6.5	4.1	3.7	60.92
	10_1,0⁻ → 9_1,0⁻	487532.70	0.6	6.5	4.4	2.9	143.28
	5_2,0⁻ → 5_1,0⁺	489037.53	0.8	6.6	4.1	3.5	72.53
	6_2,0⁻ → 6_1,0⁺	491551.87	0.7	6.3	4.1	2.9	86.46
	4_1,0⁺ → 3_0,0⁺	492279.27	2.0	6.7	4.2	8.8	37.55
	5_3,0⁺ → 4_2,0⁺	493699.96	1.3	6.5	3.7	5.1	84.62
	5_3,0⁻ → 4_2,0⁻	493734.65	1.4	6.4	3.7	5.3	84.62
	7_2,0⁻ → 7_1,0⁺	494482.35	0.6	6.5	4.3	2.6	102.7
	8_2,0⁻ → 8_1,0⁺	497829.07	0.4	6.4	4.3	2.0	121.27
	9_2,0⁻ → 9_1,0⁺	501589.39	0.4	6.7	4.0	1.5	142.15
	10_2,0⁻ → 10_1,0⁺	505762.88	0.3	6.4	4.5	1.3	165.35
	11_2,0⁻ → 11_1,0⁺	510345.79	0.2	6.4	4.7	1.0	190.86
	16_0,0⁺ → 15_1,0⁺	515170.91	0.2	6.6	4.9	1.1	315.21
	11_1,0⁺ → 10_1,0⁺	527054.34	0.4	6.5	3.9	1.6	166.37
	11_0,0⁺ → 10_0,0⁺	531320.09	1.2	6.6	4.3	5.4	153.1
	11_2,0⁻ → 10_2,0⁻	531636.16	0.2	7.1	6.9	1.7	190.87
	11_4,0⁻ → 10_4,0⁻	531869.53	0.4	7.5	6.2	2.3	233.52
	11_3,0⁻ → 10_3,0⁻	531893.14	0.3	6.9	4.8	1.4	202.98
	11_2,0⁺ → 10_2,0⁺	532133.99	0.3	6.2	4.1	1.1	190.94
	11_1,0⁻ → 10_1,0⁻	536192.13	0.6	6.4	3.7	2.4	169.01
	5_1,0⁺ → 4_0,0⁺	538571.19	2.1	6.7	4.4	9.7	49.06
	6_3,0⁺ → 5_2,0⁺	542002.05	1.2	6.4	3.7	4.7	98.55
	6_3,0⁻ → 5_2,0⁻	542082.94	1.2	6.5	3.7	4.7	98.55
	12_1,0⁺ → 11_1,0⁺	574869.74	0.3	6.3	4.2	1.4	193.96
	2_2,0⁻ → 1_1,0⁻	579085.75	1.3	6.5	3.7	5.2	44.67
	12_0,0⁺ → 11_0,0⁺	579460.65	0.9	6.5	4.4	4.0	180.91
	2_2,0⁺ → 1_1,0⁺	579922.27	1.3	6.5	4.1	5.7	44.67
	12_3,0⁺ → 11_3,0⁺	580176.81	0.2	6.6	4.2	0.8	230.83
	12_3,0⁻ → 11_3,0⁻	580212.19	0.2	7.3	5.3	1.1	230.83
	12_2,0⁺ → 11_2,0⁺	580503.31	0.2	6.3	4.5	0.9	218.8
	6_1,0⁺ → 5_0,0⁺	584450.72	2.5	6.6	4.3	11.3	62.87
	12_1,0⁻ → 11_1,0⁻	584823.31	0.4	6.6	5.3	2.5	197.08
	7_3,0⁺ → 6_2,0⁺	590278.73	1.0	6.5	3.9	4.3	114.79
	7_3,0⁻ → 6_2,0⁻	590441.53	1.0	6.4	3.8	4.1	114.79
	13_1,0⁺ → 12_1,0⁺	622660.02	0.3	6.4	4.3	1.4	223.85
	3_2,0⁻ → 2_1,0⁻	626627.34	1.4	6.5	3.9	5.5	51.64
	13_0,0⁺ → 12_0,0⁺	627559.43	0.6	6.5	5.4	3.4	211.03
	13_2,0⁻ → 12_2,0⁻	628052.23	0.2	6.8	6.0	1.2	248.84
	13_3,0⁺ → 12_3,0⁺	628471.51	0.2	6.2	6.3	1.3	260.99
	13_4,0⁻ → 12_4,0⁻	628513.72	0.2	6.3	3.5	0.7	291.53
	13_3,0⁻ → 12_3,0⁻	628524.83	0.2	7.1	5.3	1.1	261
	13_2,0⁺ → 12_2,0⁺	628868.08	0.2	7.5	5.1	1.1	248.98
	3_2,0⁺ → 2_1,0⁺	629141.49	1.4	6.5	4.1	6.2	51.64
	7_1,0⁺ → 6_0,0⁺	629922.01	1.6	6.7	4.6	7.9	78.97
	13_1,0⁻ → 12_1,0⁻	633424.48	0.4	6.3	4.9	2.1	227.48
	7_4,0⁻ → 7_3,0⁺	636336.97	0.3	7.2	4.1	1.2	145.33
	4_4,0⁻ → 4_3,0⁺	636422.01	0.5	6.0	2.8	1.5	103.56
	3_2,0⁻ → 2_1,0⁻	626627.67	1.1	6.3	4.0	4.8	51.64
	13_0,0⁺ → 12_0,0⁺	627560.25	0.5	6.1	3.9	2.0	211.03
	3_2,0⁺ → 2_1,0⁺	629141.29	1.0	6.6	4.2	4.2	51.64
	7_1,0⁺ → 6_0,0⁺	629921.68	1.1	6.8	5.0	5.9	78.97
	13_1,0⁻ → 12_1,0⁻	633425.22	0.2	6.0	5.2	1.3	227.48
	7_4,0⁻ → 7_3,0⁺	636335.86	0.3	7.7	4.4	1.4	145.33
	6_4,0⁺ → 6_3,0⁻	636365.90	0.3	6.1	4.4	1.6	129.09
	5_4,0⁻ → 5_3,0⁺	636395.60	0.3	6.1	4.6	1.6	115.16
	4_4,0⁺ → 4_3,0⁻	636420.70	0.3	6.6	3.5	1.0	103.56
	8_3,0⁺ → 7_2,0⁺	638524.78	0.9	6.4	4.4	4.3	133.36
	8_3,0⁻ → 7_2,0⁻	638818.88	0.9	6.5	4.4	4.0	133.36
	14_1,0⁺ → 13_1,0⁺	670424.18	0.2	6.3	6.8	1.6	256.02
	4_2,0⁻ → 3_1,0⁻	673747.18	1.3	6.5	4.2	5.8	60.92
	8_1,0⁺ → 7_0,0⁺	674991.12	1.3	6.8	4.5	6.2	97.38
	14_0,0⁺ → 13_0,0⁺	675613.75	0.5	6.5	5.5	2.8	243.45
	4_2,0⁺ → 3_1,0⁺	678786.36	1.4	6.5	3.9	5.9	60.93
	14_1,0⁻ → 13_1,0⁻	681991.47	0.3	6.3	6.9	2.2	260.21
	9_3,0⁺ → 8_2,0⁺	686733.14	0.7	6.3	4.0	2.8	154.25
	9_3,0⁻ → 8_2,0⁻	687225.61	0.7	6.6	4.7	3.6	154.25
	9_1,0⁺ → 8_0,0⁺	719665.76	1.0	6.6	5.2	5.7	118.08
	5_2,0⁻ → 4_1,0⁻	720442.77	1.1	6.5	3.9	4.6	72.53
	15_0,0⁺ → 14_0,0⁺	723621.75	0.4	6.0	4.8	1.8	278.18
	9_1,0⁺ → 8_0,0⁺	719665.58	1.1	6.7	4.6	5.4	118.08
	5_2,0⁻ → 4_1,0⁻	720442.56	1.3	6.5	4.1	5.7	72.53
	15_0,0⁺ → 14_0,0⁺	723620.68	0.4	6.4	6.1	2.8	278.18
	5_2,0⁺ → 4_1,0⁺	728863.72	1.4	6.5	4.3	6.3	72.53
	15_1,0⁻ → 14_1,0⁻	730520.33	0.3	6.7	5.1	1.5	295.27
	10_3,0⁺ → 9_2,0⁺	734895.06	0.5	6.5	5.1	2.8	177.46
	10_3,0⁻ → 9_2,0⁻	735674.32	0.6	6.5	5.0	3.0	177.46
	10_1,0⁺ → 9_0,0⁺	763953.80	0.8	6.8	5.5	4.5	141.08
	6_2,0⁻ → 5_1,0⁻	766711.66	1.0	6.5	3.8	4.0	86.46
	16_1,0⁻ → 15_1,0⁻	779006.31	0.4	7.9	4.2	1.9	332.65
	6_2,0⁺ → 5_1,0⁺	779381.81	1.2	6.4	4.3	5.3	86.46
	11_3,0⁺ → 10_2,0⁺	783002.74	0.4	6.6	4.1	1.9	202.98
	11_3,0⁻ → 10_2,0⁻	784178.53	0.4	6.6	4.8	1.8	202.99
	11_1,0⁺ → 10_0,0⁺	807866.57	0.6	6.8	5.2	3.1	166.37
	7_2,0⁻ → 6_1,0⁻	812551.63	0.8	6.5	4.4	3.8	102.7
	4_4,0⁻ → 3_3,0⁻	829891.72	0.7	6.9	4.6	3.1	103.56
	7_2,0⁺ → 6_1,0⁺	830351.20	0.9	6.3	4.2	4.0	102.72
	12_3,0⁺ → 11_2,0⁺	831048.06	0.3	5.9	4.4	1.4	230.83
	12_3,0⁻ → 11_2,0⁻	832755.31	0.3	6.4	5.3	1.8	230.83
	12_1,0⁺ → 11_0,0⁺	851415.12	0.6	7.0	4.7	3.0	193.96
	8_2,0⁻ → 7_1,0⁻	857960.16	0.6	6.6	3.8	2.4	121.27
	5_4,0⁺ → 4_3,0⁺	878227.11	0.6	6.8	4.5	2.7	115.16
	13_3,0⁺ → 12_2,0⁺	879015.02	0.3	6.3	5.9	1.6	260.99
	13_3,0⁻ → 12_2,0⁻	881420.97	0.3	6.9	6.6	1.8	261
	8_2,0⁺ → 7_1,0⁺	881783.47	0.7	6.5	4.5	3.6	121.29
	13_1,0⁺ → 12_0,0⁺	894614.92	0.4	6.8	4.7	1.9	223.84
	13_1,0⁺ → 12_0,0⁺	894614.92	0.4	6.8	4.7	1.9	223.84
	9_2,0⁻ → 8_1,0⁻	902936.66	0.4	6.4	5.3	2.5	142.15
	6_4,0⁺ → 5_3,0⁺	926555.20	0.4	7.0	5.3	2.5	129.09
	9_2,0⁺ → 8_1,0⁺	933694.55	0.6	6.5	4.6	2.7	142.19
	14_1,0⁺ → 13_0,0⁺	937479.55	0.3	6.6	5.8	2.0	256.02
	7_4,0⁻ → 6_3,0⁻	974878.74	0.4	6.4	6.4	2.9	145.33
	10_2,0⁺ → 9_1,0⁺	986100.37	0.5	6.2	5.1	3.0	165.4
	8_4,0⁺ → 7_3,0⁺	1023196.73	0.5	7.2	4.4	2.2	163.9

E−CH₃OH

	Transition J_{±K, v_t}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	10_0,0 → 9_0,0	482283.13	0.7	6.4	3.8	2.8	132.71
	10_−1,0 → 9_−1,0	482959.74	1.4	6.4	3.9	5.7	125.25
	10_1,0 → 9_1,0	483687.10	0.5	6.5	3.9	2.2	140.83
	10_2,0 → 9_2,0	484023.82	0.6	6.6	4.5	3.0	142.08
	10_−2,0 → 9_−2,0	484072.13	0.3	6.8	5.0	1.6	145.73
	7_0,0 → 6_−1,0	495173.84	1.2	6.6	3.9	4.8	70.18
	7_1,0 → 6_0,0	504294.53	1.2	6.4	3.6	4.7	248.24
	11_1,0 → 10_2,0	506153.78	0.2	6.6	5.1	1.0	166.37
	10_2,0 → 9_1,0	509564.94	0.7	6.8	5.2	3.9	142.08
	2_−2,0 → 1_−1,0	520179.58	1.1	6.7	4.1	5.0	24.96
	14_−1,0 → 13_0,0	523275.25	0.3	6.5	4.9	1.4	241.04
	8_3,0 → 8_2,0	530123.52	0.3	6.9	5.4	1.7	123.38
	11_0,0 → 10_0,0	530184.87	0.5	6.6	4.6	2.5	158.15
	7_3,0 → 7_2,0	530316.84	0.3	6.6	4.9	1.8	104.81
	6_3,0 → 6_2,0	530455.52	0.4	6.5	4.1	1.9	88.56
	5_3,0 → 5_2,0	530549.77	0.4	6.7	4.6	2.2	74.63
	4_3,0 → 4_2,0	530611.06	0.4	6.6	4.3	2.0	63.03
	3_3,0 → 3_2,0	530647.07	0.4	7.1	5.0	2.1	53.74
	11_−1,0 → 10_−1,0	531080.10	1.1	6.5	4.3	4.9	150.74
	11_1,0 → 10_1,0	532032.12	0.4	6.6	5.2	2.2	166.37
	11_2,0 → 10_2,0	532467.35	0.5	6.4	4.6	2.7	167.63
	11_−2,0 → 10_−2,0	532567.79	0.3	6.4	5.2	1.7	171.29
	8_0,0 → 7_−1,0	543076.99	1.1	6.6	4.5	5.2	88.72
	8_1,0 → 7_0,0	553147.29	0.9	6.4	3.6	3.6	96.73
	11_2,0 → 10_1,0	558345.87	0.4	6.3	4.3	1.9	167.63
	11_2,0 → 10_1,0	558345.39	0.5	6.6	4.5	2.3	167.63
	3_−2,0 → 2_−1,0	568566.84	1.2	6.6	4.3	5.6	31.93
	15_−1,0 → 14_0,0	572900.17	0.2	6.2	5.0	1.2	275.74
	12_0,0 → 11_0,0	578007.58	0.4	6.4	4.6	1.9	185.89
	12_−1,0 → 11_−1,0	579151.93	0.7	6.5	4.4	3.5	178.53
	12_{1,0 →} 11_1,0	580369.39	0.3	6.6	4.5	1.5	194.22
	12_2,0 → 11_2,0	580903.85	0.4	6.4	4.9	2.2	195.51
	12_−2,0 → 11_−2,0	581092.64	0.3	6.5	4.8	1.3	199.18
	9_0,0 → 8_−1,0	590791.90	0.8	6.5	4.1	3.7	109.56
	9_1,0 → 8_0,0	602234.23	0.8	6.4	4.0	3.4	117.62
	12_2,0 → 11_1,0	607217.09	0.3	6.4	5.2	1.7	195.51
	4_−2,0 → 3_−1,0	616980.70	1.2	6.6	4.1	5.1	41.22
	16_−1,0 → 15_0,0	622775.50	0.2	6.1	5.0	1.0	312.73
	13_0,0 → 12_0,0	625750.66	0.3	6.4	5.4	1.7	215.92
	13_−1,0 → 12_−1,0	627171.60	0.6	6.5	4.9	3.0	208.64
	13_{1,0 →} 12_1,0	628698.28	0.2	6.1	5.0	1.1	224.39
	13_{3,0 →} 12_3,0	628817.62	0.2	6.3	4.0	0.8	251.06
	13_2,0 → 12_2,0	629322.36	0.4	6.7	5.0	2.0	225.71
	13_−2,0 → 12_−2,0	629653.22	0.2	6.3	7.8	1.9	229.4
	13_−1,0 → 12_−1,0	627171.35	0.5	6.6	4.4	2.3	208.64
	13_2,0 → 12_2,0	629321.48	0.3	7.1	6.0	1.6	225.71
	13_−2,0 → 12_−2,0	629651.72	0.2	7.0	5.2	1.1	229.4
	10_{0,0 →} 9_−1,0	638280.27	0.7	6.7	4.4	3.1	132.71
	10_{1,0 →} 9_0,0	651618.53	0.6	6.5	4.6	2.9	140.83
	13_2,0 → 12_1,0	656166.61	0.6	8.1	6.4	3.8	225.71
	5_−2,0 → 4_−1,0	665443.13	1.2	6.7	4.3	5.2	52.83
	14_0,0 → 13_0,0	673416.97	0.3	6.6	4.9	1.4	248.24
	14_−1,0 → 13_−1,0	675135.73	0.4	6.5	6.0	2.8	241.04
	3_3,0 → 2_2,0	675773.95	0.7	6.8	4.9	3.6	53.74
	14_1,0 → 13_1,0	677011.82	0.2	7.5	8.3	2.1	256.88
	14_2,0 → 13_2,0	677711.91	0.3	6.1	4.9	1.7	258.24
	11_0,0 → 10_−1,0	685505.14	0.5	7.0	5.9	3.2	158.15
	6_−2,0 → 5_−1,0	713983.43	0.9	6.6	4.1	3.9	66.76
	6_−2,0 → 5_−1,0	713983.43	0.9	6.6	4.1	3.9	66.76
	4_−4,0 → 3_−3,0	718436.93	0.8	6.7	4.0	3.5	103.22
	15_−1,0 → 14_−1,0	723040.58	0.4	6.9	5.4	2.0	275.74
	4_3,0 → 3_2,0	724122.89	0.6	6.5	5.7	3.4	63.03
	4_3,0 → 3_2,0	724122.10	0.8	6.8	5.3	4.2	63.03
	12_0,0 → 11_−1,0	732433.68	0.4	6.4	5.1	2.2	185.89
	12_1,0 → 11_0,0	751552.50	0.4	6.4	5.8	2.3	194.22
	7_−2,0 → 6_−1,0	762636.47	0.7	6.7	4.4	3.4	83.02
	9_−3,0 → 9_−2,0	766028.92	0.3	6.7	5.2	1.4	159.27
	8_−3,0 → 8_−2,0	766396.97	0.4	6.6	5.9	2.3	138.38
	7_−3,0 → 7_−2,0	766648.47	0.4	6.6	4.5	1.9	119.81
	5_−4,0 → 4_−3,0	766761.68	0.8	6.6	3.3	2.9	114.82
	6_−3,0 → 6_−2,0	766811.56	0.4	6.6	2.6	1.1	103.56
	5_−3,0 → 5_−2,0	766908.51	0.4	6.8	4.2	1.9	89.63
	4_−3,0 → 4_−2,0	766961.54	0.5	6.4	3.4	1.6	78.03
	3_−3,0 → 3_−2,0	766983.75	0.4	6.4	3.1	1.2	68.74
	5_3,0 → 4_2,0	772453.48	0.5	7.2	7.1	4.0	74.63
	13_0,0 → 12_−1,0	779030.27	0.4	7.2	6.0	2.2	215.92
	8_−2,0 → 7_−1,0	811444.68	0.5	7.0	6.0	3.4	101.6
	6_−4,0 → 5_−3,0	815071.72	0.7	6.6	4.3	3.1	128.75
	6_3,0 → 5_2,0	820764.66	0.4	6.2	4.4	1.8	88.56
	9_−2,0 → 8_−1,0	860459.44	0.4	7.0	5.5	2.4	122.5
	7_−4,0 → 6_−3,0	863365.78	0.5	6.7	5.2	2.9	145
	7_3,0 → 6_2,0	869039.06	0.5	6.6	5.2	2.6	104.81
	10_−2,0 → 9_−1,0	909738.07	0.3	6.9	4.7	1.6	145.73
	8_−4,0 → 7_−3,0	911643.69	0.4	6.5	4.5	2.1	163.56
	3_−3,0 → 2_−2,0	912109.65	0.7	6.6	5.1	3.5	68.74
	8_3,0 → 7_2,0	917269.03	0.4	7.3	4.9	1.9	123.38
	4_−3,0 → 3_−2,0	960473.04	0.6	6.4	5.3	3.1	78.03

CN, v=0

	Transition N_J,F₁	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	5_4.5,5.5 → 4_3.5,4.5	566729.48	1.3	7.3	5.3	7.3	81.59
	5_5.5,5.5 → 4_4.5,4.5	566946.10	1.6	7.6	5.6	9.8	81.64
	6_5.5,6.5 → 5_4.5,5.5	680046.29	0.5	7.5	5.4	3.0	114.23
	6_6.5,6.5 → 5_5.5,5.5	680263.15	0.7	7.5	5.4	4.1	114.29
	7_6.5,7.5 → 6_5.5,6.5	793337.05	0.3	7.3	6.9	2.2	152.38
	7_7.5,7.5 → 6_6.5,6.5	793552.12	0.3	7.6	6.4	1.8	152.38

CO

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

Main	5 → 4	576267.41	52.5	7.3	6.3	352.5	82.98
Wing	5 → 4	576267.63	7.7	7.2	19.1	156.7	82.98
Main	6 → 5	691472.71	48.2	7.2	5.8	299.0	116.16
Wing	6 → 5	691472.71	11.9	7.2	17.6	222.7	116.16
Main	7 → 6	806653.11	46.6	6.5	5.0	248.5	154.88
Wing	7 → 6	806651.38	7.8	7.2	22.4	185.5	154.88
Main	8 → 7	921801.29	45.2	6.5	5.1	244.5	199.11
Wing	8 → 7	921799.21	7.3	7.2	23.8	184.5	199.11
Main	9 → 8	1036911.45	45.8	7.3	6.0	292.2	248.88
Wing	9 → 8	1036911.85	11.7	7.2	18.0	223.7	248.88
Main	10 → 9	1151987.56	39.3	6.5	4.6	194.1	304.17
Wing	10 → 9	1151987.37	5.5	6.5	23.0	134.4	304.17
Main	11 → 10	1267014.10	31.2	7.1	5.0	166.0	364.97
Wing	11 → 10	1267016.60	10.8	6.5	15.0	173.1	367.97
Main	13 → 12	1496920.46	29.8	7.5	4.8	152.8	503.14
Wing	13 → 12	1496925.41	8.7	6.5	18.2	167.9	503.14
Main	14 → 13	1611790.99	28.5	7.5	4.9	148.6	580.5
Wing	14 → 13	1611792.67	4.9	7.2	21.9	114.6	580.5
Main	15 → 14	1726597.84	20.4	7.8	3.5	75.5	663.36
Wing	15 → 14	1726601.59	6.2	7.2	14.0	92.3	663.36
Main	16 → 15	1841340.61	14.0	7.8	4.0	59.4	751.73
Wing	16 → 15	1841344.53	4.0	7.2	15.0	63.2	751.73

¹³CO

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	5 → 4	550925.82	30.9	7.3	4.8	157.6	79.33
	6 → 5	661066.67	30.8	7.3	4.7	155.2	111.05
	7 → 6	771182.86	28.0	7.5	5.1	152.0	148.06
	8 → 7	881271.54	22.6	7.4	4.7	112.0	190.36
	9 → 8	991327.01	17.5	7.7	4.7	87.8	237.94
	10 → 9	1101346.80	13.1	7.8	4.6	64.4	290.79
	11 → 10	1211327.07	6.7	7.6	4.3	31.0	348.93

C¹⁸O

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	5 → 4	548831.12	8.5	6.8	4.1	36.4	79.02
	6 → 5	658553.65	8.0	6.9	3.7	31.5	110.63
	7 → 6	768252.15	5.8	7.0	4.1	25.4	147.5
	8 → 7	877922.54	3.5	6.9	5.0	18.7	189.64
	9 → 8	987559.09	2.2	7.6	5.9	13.8	237.03
	10 → 9	1097163.80	1.5	6.8	3.2	5.1	289.69

C¹⁷O

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	5 → 4	561713.01	3.0	6.7	3.7	11.9	80.88
	6 → 5	674009.48	2.4	6.6	3.7	9.3	113.23
	7 → 6	786282.9^b	–	6.5	4.5	7.6	150.96
	8 → 7	898523.38	0.9	7.0	4.1	4.0	194.09
	9 → 8	1010731.41	0.4	7.2	3.5	1.5	242.59
	10 → 9	1122902.89	0.2	7.4	4.2	1.1	296.49

¹³C¹⁸O

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	5 → 4	523484.35	0.1	6.9	4.1	0.6	75.37
	6 → 5	628141.57	0.1	6.6	2.6	0.3	105.52

CS,v=0

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

Main	10 → 9	489750.85	4.1	7.3	4.3	18.8	129.29
Wing	10 → 9	489750.92	0.6	6.7	7.8	4.6	129.29
Main	11 → 10	538688.84	3.3	7.2	4.7	16.5	155.15
Wing	11 → 10	538688.92	0.3	7.0	16.9	4.5	155.15
Main	12 → 11	587616.41	2.2	7.1	4.9	11.2	183.35
Wing	12 → 11	587619.04	0.2	6.5	12.9	2.4	183.35
Main	13 → 12	636532.25	1.0	7.2	5.3	5.8	213.9
Wing	13 → 12	636531.41	0.2	8.0	15.0	3.9	213.9
Main	14 → 13	685435.93	1.0	7.0	6.0	6.6	246.79
Wing	14 → 13	–	–	–	–	–	246.79
Main	15 → 14	734325.78	0.6	6.7	5.1	3.1	282.04
Wing	15 → 14	734327.70	0.2	6.5	20.4	4.8	282.04
Main	16 → 15	783201.36	0.3	7.2	5.0	1.4	319.62
Wing	16 → 15	783203.40	0.3	7.0	10.3	2.9	319.62
Main	17 → 16	832061.55	0.4	6.6	6.1	2.9	359.56
Wing	17 → 16	–	–	–	–	–	359.56
Main	18 → 17	880903.85	0.3	7.2	9.0	2.9	401.83
Wing	18 → 17	–	–	–	–	–	401.83

¹³CS

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	11 → 10	508535.83	0.1	7.1	6.1	0.5	146.46

C³⁴S

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	10 → 9	481916.67	0.3	6.8	4.6	1.4	127.23
	11 → 10	530071.15	0.2	7.5	5.3	1.0	152.67
	12 → 11	578215.88	0.2	7.6	7.3	1.3	180.42

DCN,v=0

	Transition N_k	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	7₀ → 6₀	506825.63	0.2	6.9	5.2	1.0	97.3

DCO⁺

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	7 → 6	504200.27	0.2	7.1	4.4	0.8	96.80
	8 → 7	576205:0^b	–	–	–	–	124.45
	9 → 8	648195.15	0.2	6.1	2.1	0.4	155.56

o–H₂S

	Transition J_Ω,Λ	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	2_2,1 → 2_1,2	505565.22	1.5	7.0	5.0	8.2	59.59
	3_1,2 → 3_0,3	708470.60	0.4	6.9	7.5	3.5	116.99
	2_1,2 → 1_0,1	736034.21	4.5	7.4	5.9	27.9	35.32
	3_0,3 → 2_1,2	993107.06	1.7	5.8	6.3	11.4	82.99
	2_2,1 → 1_1,0	1072837.04	0.9	7.7	6.0	5.7	59.59

p–H₂S

	Transition J_Ω,Λ	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	2_0,2 → 1_1,1	687303.34	2.4	7.1	5.1	12.7	54.7
	3_1,3 → 2_0,2	1002777.30	0.6	7.4	5.6	3.8	102.82

$H_{2}^{34} S$

	Transition N_{K_a,K_c}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	2_1,2 → 1_0,1	734269.67	0.3	6.9	4.1	1.3	55.01

H₂CS

	Transition J_{K_a,K_c}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	14_1,13 → 13_1,12	487664.03	0.2	6.8	4.5	0.8	188.8
	15_1,15 → 14_1,14	506771.50	0.2	6.6	5.2	1.1	207.86
	15_0,15 → 14_0,14	513361.94	0.1	6.6	4.0	0.4	197.44
	15_2,13 → 14_2,12	516336.1^b	–	–	–	–	250.71
	15_1,14 → 14_1,13	522403.63	0.1	6.9	5.1	0.9	213.87
	16_1,16 → 15_1,15	540465.31	0.1	6.5	3.5	0.5	233.79
	17_1,17 → 16_1,16	574140.12	0.1	6.8	4.7	0.5	261.35

HF

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	1 → 0	1232469.17	-1.1	8.7	3.2	-7.0	59.15

o–H₂CO

	Transition J_{K_a,K_c}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	7_1,7 → 6_1,6	491968.45	3.1	7.0	4.4	14.7	91.15
	7_5,3 → 6_5,2	509562.12^b	–	–	–	–	376.68
	7_3,5 → 6_3,4	510155.94	0.9	6.9	4.8	4.4	188.73
	7_3,4 → 6_3,3	510238.04	0.9	6.9	4.4	4.0	188.73
	7_1,6 → 6_1,5	525666.11	2.5	6.8	4.6	12.2	97.63
	8_1,8 → 7_1,7	561899.30	2.6	7.0	4.4	12.5	118.12
	8_5,3 → 7_5,2	582382.77	0.2	6.6	6.3	1.3	404.63
	8_3,6 → 7_3,5	583144.85	0.6	6.9	5.2	3.2	216.71
	8_3,5 → 7_3,4	583308.64	0.6	7.0	4.6	2.8	216.73
	8_1,7 → 7_1,6	600331.09	1.7	6.7	4.5	8.1	126.44
	9_1,9 → 8_1,8	631703.48	1.5	6.7	4.6	7.3	148.43
	9_5,5 → 8_5,4	655212.67	0.2	6.7	5.4	1.1	436.08
	9_3,7 → 8_3,6	656166.83	0.6	6.0	6.2	3.9	248.2
	9_3,6 → 8_3,5	656465.35	0.5	6.6	4.6	2.3	248.23
	9_1,8 → 8_1,7	674810.63	1.1	6.6	4.8	5.4	158.83
	10_1,10 → 9_1,9	701370.13	1.3	7.1	5.1	7.1	182.1
	10_5,6 → 9_5,5	728052.29	0.2	7.5	5.1	1.0	471.02
	10_3,8 → 9_3,7	729213.26	0.4	6.7	4.5	2.0	283.2
	10_3,7 → 9_3,6	729725.77	0.3	6.7	5.9	2.1	283.25
	10_1,9 → 9_1,8	749072.76	0.8	6.7	4.5	3.8	194.78
	11_1,11 → 10_1,10	770895.15	0.7	7.4	7.7	5.7	219.09
	11_1,10 → 10_1,9	823084.31	0.5	6.4	5.1	2.8	234.28
	12_1,12 → 11_1,11	840277.03	0.5	6.5	4.8	2.5	259.42
	12_1,11 → 11_1,10	896807.59	0.3	6.2	4.9	1.7	277.32
	13_1,13 → 12_1,12	909511.04	0.3	5.9	6.0	2.0	303.07

p–H₂CO

	Transition J_{K_a,K_c}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	7_0,7 → 6_0,6	505834.17	1.7	6.7	4.0	7.2	97.44
	7_2,6 → 6_2,5	509146.74	0.8	6.7	4.2	3.4	144.93
	7_4,4 → 6_4,3	509830.59	0.2	6.4	4.2	1.0	286.17
	7_2,5 → 6_2,4	513076.79	0.7	6.7	4.5	3.2	145.35
	8_0,8 → 7_0,7	576709.19	1.1	6.6	4.2	4.9	125.12
	8_2,7 → 7_2,6	581612.27	0.5	6.8	4.3	2.3	172.84
	8_4,5 → 7_4,4	582723.54	0.2	6.7	5.7	1.2	314.14
	8_2,6 → 7_2,5	587454.12	0.5	6.8	4.8	2.4	173.55
	9_0,9 → 8_0,8	647082.69	0.9	6.6	4.1	3.7	156.18
	9_2,8 → 8_2,7	653971.43	0.4	6.4	4.6	1.9	204.23
	9_4,6 → 8_4,5	655640.45	0.2	6.8	6.0	1.2	345.6
	9_2,7 → 8_2,6	662209.75	0.4	6.7	4.5	1.9	205.33
	10_0,10 → 9_0,9	716939.12	0.4	6.7	5.6	2.4	190.58
	10_2,9 → 9_2,8	726209.22	0.3	6.6	6.5	2.0	239.08
	10_2,8 → 9_2,7	737342.68	0.3	7.0	5.3	1.7	380.57
	11_0,11 → 10_0,10	786282.9^b	–	6.5	4.5	7.6	228.32

o–H₂O

	Transition J_{J_{K₋₁,K₊₁}}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

Main	1_1,0 → 1_0,1	556935.64	6.2	7.2	7.8	51.9	26.73
Wing	1_1,0 → 1_0,1	556939.86	1.4	4.9	34.3	52.1	26.73
Main	5_3,2 → 4_4,1	620704.48	0.5	5.3	2.0	1.1	697.84
Wing	5_3,2 → 4_4,1	620704.48	–	–	–	–	697.84
Main	3_1,2 → 3_0,3	1097364.47	1.5	7.1	5.1	8.2	215.2
Wing	3_1,2 → 3_0,3	1097371.43	2.0	5.2	14.8	31.2	215.2
Main	3_1,2 → 2_2,1	1153127.12	1.6	6.9	5.3	8.9	215.203
Wing	3_1,2 → 2_2,1	1153133.21	2.6	5.3	19.2	53.7	215.2
Main	3_2,1 → 3_1,2	1162913.04	2.9	6.6	4.3	13.4	271.01
Wing	3_2,1 → 3_1,2	1162914.76	1.4	6.2	18.6	28.4	271.01

p–H₂CO

	Transition J_{J_K₋₁,_K₊₁}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

Main	2_1,1 → 2_0,2	752033.42	5.5	6.9	5.8	33.6	136.94
Wing	2_1,1 → 2_0,2	752034.00	2.5	6.7	21.8	58.8	136.94
Main	4_2,2 → 3_3,1	916174.82	0.3	5.9	6.9	2.0	454.34
Wing	4_2,2 → 3_3,1	916174.82	–	–	–	–	454.34
Main	2_0,2 → 1_1,1	987927.13	6.1	6.9	7.5	48.6	100.85
Wing	2_0,2 → 1_1,1	987938.36	2.4	3.5	26.2	66.2	100.85
Main	1_1,1 → 0_0,0	1113349.92	1.9	5.1	4.5	9.1	53.43
Wing	1_1,1 → 0_0,0	1113352.60	1.8	4.4	22.3	42.8	53.43
Main	4_2,2 → 4_1,3	1207638.86	0.3	7.0	13.6	4.2	454.34
Wing	4_2,2 → 4_1,3	1207638.86	–	–	–	–	454.34
Main	2_2,0 → 2_1,1	1228791.19	1.2	6.4	5.4	7.2	195.91
Wing	2_2,0 → 2_1,1	1228791.19	–	–	–	–	195.91

$o- H_{2}^{18} O$

	Transition J_{K₋₁,K₊₁,v}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	1_1,0,0 → 1_0,1,0	547676.36	0.2	7.0	5.6	1.43

HDO

	Transition J_{K₋₁,K₊₁}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	2_0,2 → 1_1,1	490597.00	0.2	6.8	3.9	0.7	66.43
	1_1,0 → 1_1,1	509293.66	0.3	6.3	3.4	1.0	46.76
	2_1,1 → 2_0,2	599927.69	0.2	6.5	2.9	0.7	95.23

HCl

	Transition N_J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	1_1.5 → 0_1.5	625901.60	2.1	7.8	4.4	9.9	30.04
	1_2.5 → 0_1.5	625918.76	2.8	7.9	4.1	12.0	30.04
	1_0.5 → 0_1.5	625932.01	1.4	7.9	4.7	7.0	30.04

H³⁷CI

	Transition N_J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	1_1.5 → 0_1.5	624962.48	1.0	8.3	3.1	3.1	29.99
	1_2.5 → 0_1.5	624975.93	1.4	8.1	3.8	5.5	29.99
	1_0.5 → 1_1.5	624986.44	0.6	8.8	5.4	3.6	29.99

HCN

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

Main	6 → 5	531716.10	7.4	7.1	5.4	43.0	89.32
Wing	6 → 5	531718.60	0.5	6.0	25.3	14.0	89.32
Main	7 → 6	620303.70	5.1	7.2	5.3	28.4	119.09
Wing	7 → 6	620305.69	0.8	5.9	17.1	14.3	119.09
Main	8 → 7	708876.27	2.5	7.4	5.2	13.8	153.11
Wing	8 → 7	708878.31	0.9	6.2	16.0	15.0	153.11
Main	9 → 8	797432.27	0.9	7.6	4.9	4.9	191.38
Wing	9 → 8	797435.00	0.8	6.3	13.5	11.2	191.38
Main	10 → 9	885969.75	0.6	7.4	5.7	3.9	233.9
Wing	10 → 9	885973.43	0.6	6.3	16.0	10.2	233.9
Main	11 → 10	974486.55	0.6	6.5	5.5	3.5	280.67
Wing	11 → 10	–	–	–	–	–	280.67
Main	13 → 12	1151448.54	0.5	5.9	6.2	3.0	386.95
Wing	13 → 12	–	–	–	–	–	386.95

H¹³CN

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	6 → 5	517969.24	0.4	7.3	5.0	2.1	87.01
	7 → 6	604266.38	0.2	7.8	7.5	1.7	116.01

HNC

	Transition J_{K_a,K_c}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	6_0,0 → 5_0,0	543897.76	1.2	7.0	4.2	5.4	91.37
	7_0,0 → 6_0,0	634510.41	0.5	7.2	4.5	2.2	121.82

HCO⁺

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	6 → 5	535061.20	15.3	7.2	5.5	90.3	89.88
	7 → 6	624208.34	13.8	7.0	5.7	83.7	119.84
	8 → 7	713341.48	10.5	6.9	5.3	59.4	154.08
	9 → 8	802458.39	7.4	6.9	5.2	41.0	192.59
	10 → 9	891558.14	5.6	6.7	5.2	31.2	235.38
	11 → 10	980636.88	4.2	6.9	5.3	23.8	282.44
	12 → 11	1069695.54	2.5	6.5	5.2	14.0	333.78
	13 → 12	1158729.89	1.6	6.3	4.7	7.9	389.39
	14 → 13	1247737.53	0.9	6.4	5.0	5.0	449.27

H¹³CO⁺

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	6 → 5	520460.32	1.1	6.7	3.9	4.6	87.43
	7 → 6	607175.53	0.7	6.7	3.7	2.7	116.57
	8 → 7	693877.57	0.4	6.4	3.4	1.6	149.87
	9 → 8	780563.19	0.2	6.9	5.5	1.3	187.33

N₂H⁺

	Transition J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	6 → 5	558967.76	2.7	6.3	0.1	7.2	93.9
	7 → 6	652097.23	2.7	6.3	0.0	6.7	125.19
	8 → 7	745211.80	1.6	6.3	0.1	3.7	160.96
	9 → 8	838309.70	0.9	6.2	0.2	1.8	201.19
	10 → 9	931388.75	0.5	6.0	0.3	1.3	245.89

NH₃

	Transition N_K,v	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	Δ V_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

Main	1_0,0 → 0_0,1	572498.04	3.3	7.1	5.9	20.8	27.48
Wing	1_0,0 → 0_0,1	572496.26	0.2	7.5	12.0	2.1	27.48
Absorption	2_1,1 → 1_1,0	1215245.71	-0.7	6.5	5.1	-3.9	80.45

NO

	Transition N_ΛF₁F₂	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	6_−1,5.5,5.5 → 5_1,4.5,4.5	551187.85	0.3	6.7	4.6	1.5	84.15
	6_1,5.5,4.5 → 5_−1,4.5,3.5	551533.80	0.3	7.1	4.8	1.4	84.25
	6_1,6.5,6.5 → 6_−1,5.5,5.5	651433.16	0.3	6.8	5.8	1.7	115.42
	6_−1,6.5,5.5 → 6_1,5.5,4.5	651773.27	0.4	6.9	4.3	1.7	115.53

SO

	Transition N_J	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	4₃ → 1₂	504676.68	0.1	6.8	3.8	0.5	28.68
	12₁₁ → 11₁₀	514853.39	0.7	7.2	6.7	4.7	167.59
	12₁₂ → 11₁₁	516335.63	0.6	7.1	7.6	5.1	174.22
	12₁₃ → 11₁₂	517354.09	0.8	7.3	6.4	5.7	165.78
	13₁₂ → 12₁₁	558087.15	0.4	7.3	5.9	2.4	194.37
	13₁₃ → 12₁₂	559319.68	0.4	7.0	6.6	2.7	201.07
	13₁₄ → 12₁₃	560178.46	0.5	7.1	5.7	2.7	192.66
	14₁₃ → 13₁₂	601257.88	0.4	7.3	9.6	3.7	223.23
	14₁₄ → 13₁₃	602291.95	0.3	7.5	9.0	3.2	229.97
	14₁₅ → 13₁₄	603021.65	0.4	7.0	8.2	3.5	221.61
	5₄ → 2₃	611551.86	0.1	7.3	3.7	0.4	38.58
	15₁₄ → 14₁₃	644377.53	0.3	7.7	9.0	3.3	254.15
	15₁₅ → 14₁₄	645253.67	0.3	7.6	9.0	2.8	260.94
	15₁₆ → 14₁₅	645874.95	0.3	7.5	10.5	3.8	252.6
	16₁₅ → 15₁₄	687457.56	0.3	7.1	12.6	3.4	287.15
	16₁₆ → 15₁₅	688204.39	0.2	7.1	10.5	2.7	293.97
	16₁₇\| → 15₁₆	688733.81	0.3	7.8	14.9	4.0	285.66
	17₁₆ → 16₁₅	730500.26	0.3	7.2	7.5	2.2	322.2

SO₂

	Transition N_{K_a,K_c}	Frequency MHz	T_A (K)	V_lsr (km s⁻¹)	ΔV_FWHM (km s⁻¹)	∫T_AdV (K km s⁻¹)	E_u (K)

	13_3,11 → 12_2,10	484270.75	0.2	7.1	4.0	0.6	105.83
	7_4,4 → 6_3,3	491933.52	0.1	7.7	5.4	0.8	65.01
	12_3,9 → 11_2,10	494779.07	0.2	7.4	3.3	0.7	93.96

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Superscipt “b” means it is a blended line and excluded.

2.3. Line Profiles

Although all lines above the 3σ S/N level in intensity were identified, Gaussian fitting and subsequent modeling was only performed on lines that were above the 5σ noise level (where the noise is calculated from line-free regions of the spectrum immediately adjacent to the line). Gaussian fits were obtained using the Levenberg-Marquardt algorithm as implemented in the CASSIS software package and was done independently from LTE modeling, the latter of which will be described in Section 3.1. A linear or a second order baseline was fit to the data prior to Gaussian fitting but was not removed so we could include the continuum in subsequent modeling (important for absorption lines).

In most cases, a single component Gaussian fit to a specific species could reasonably reproduce the observed lines. However, in some cases a two component Gaussian fit was needed, one component being narrow (ΔV_FWHM = 3–5 km s⁻¹) and the other broad (ΔV_FWHM = 7–14 km s⁻¹); the latter could be the effect of a hot core, an outflow, or a shock. Figure 9 provides an example of a species that needed only a single component fit (¹³CO), whereas Figure 10 shows the spectra for HCN, a typical example of a species requiring a two component fit In addition, ¹²C¹⁶O line profiles are the only ones among 52 identified species, which clearly had a non-Gaussian shape, probably due to self absorption. A few other transitions are seen in absorption rather than emission and are listed in Table 3 with negative intensities.

Fig. 9 — One component LTE modeling for ¹³CO. Black histogram shows the data. Resulting model spectra are shown in red. LTE model parameters are provided in Table 5.

Fig. 10 — Two component LTE modeling for HCN. Black histogram shows the data. Models of the narrow and the broad components are shown by the red and blue lines respectively. The superposition of the two components is shown in green. LTE model parameters are provided in Table 6.

Table 4 shows the results of the Gaussian fitting for each species. The reported result for each individual line profile is the best Gaussian fit. We should note that, usually, the signal to noise ratio is lower, as we move to higher frequencies. In case of a two component fit the narrower component is referred as “main” and the broad component is referred to as “wing”. The first column in Table 4 is the transition quantum number (an explanation of the quantum numbers is provided on the CDMS and JPL websites). The centroid frequency of the fitted Gaussian profile is listed in the second column. T_A and V_LSR are respectively the observed antenna temperature and centroid velocity of the corresponding Gaussian fit. ΔV_FWHM is the “Full Width Half Maximum”, and ∫T_AdV (km s⁻¹) is the integrated line intensity.

Note that Gaussian fitting was performed on each line separately (i.e. we did not utilize a single set of Gaussian parameters to fit all transitions simultaneously). This implies that each transition of a given species can have slightly different V_LSR and ΔV_FWHM. This effect is best demonstrated using methanol as an example, since it has the largest number of transitions. The mean values of V_LSR and ΔV_FWHM for methanol A & E combined is 6.6±0.3 km s⁻¹ and 4.7±0.9 km s⁻¹ respectively, where the errors are the 1σ standard deviations about the mean. The calculated scatter about the mean, however, is an intensity dependent parameter. Figures 11 and 12 plot V_LSR and ΔV_FWHM vs T_A (> 5σ) for all the fitted methanol A & E transitions in Table 4, and clearly show that the scatter in both fitted parameters decreases with increasing T_A. For transitions with T_A < 0.5 K, <V_LSR >= 6.59 ± 0.39 km s⁻¹ and < ΔV_FWHM >= 4.99 ± 0.92 km s⁻¹. Whereas for transitions with T_A > 0.5 K, <V_LSR >= 6.57 ± 0.23 km s⁻¹ and < ΔV_FWHM > = 4.39 ± 0.67 km s⁻¹. This suggests that most of the observed scatter in these parameters is not due to the emission itself but it is due to our Gaussian fitting procedure, which clearly is subject to larger errors for weaker lines. This behaviour is predicted by Porter et al. (2004) who show that the error in V_LSR and ΔV_FWHM from Gaussian fitting increases with decreasing signal to noise. The green and the blue lines on Figures 11 and 12 denote, respectively, the 1σ and 3σ theoretical error envelope. These were calculated from equation A.1 in Porter et al. (2004) assuming < ΔV_FWHM >~ 5 km s⁻¹ and T_rms ~ 0.1 K (typical values for methanol) and illustrate this effect quite clearly.

Fig. 11 — Plot of the measured V_LSR vs T_A derived from independent Gaussian fits to each of the A-CH₃OH and E-CH₃OH transitions listed in Table 4. The green and blue lines are the 1σ and 3σ (respectively) theoretical error envelope for the V_LSR determined from Gaussian fitting of noisy lines predicted by Porter et al. (2004). Curves are calculated assuming < Δ_FWHM >~ 5 km s⁻¹ and T_rms ~ 0.1 K (typical values for methanol).

Fig. 12 — Plot of the measured ΔV_FWHM vs T_A derived from independent Gaussian fits to each of the A-CH₃OH and E-CH₃OH transitions listed in Table 4. The green and blue lines are the 1σ and 3σ (respectively) theoretical error envelope for the Δ_FWHM determined from Gaussian fitting of noisy lines predicted by Porter et al. (2004). Curves are calculated assuming < Δ_FWHM >~ 5 km s⁻¹ and T_rms ~ 0.1 K (typical values for methanol)

HIFI data obtained in beam switching mode generally provide quite a good measure of the continuum. Therefore, in each band, we have integrated the emission over the entire frequency range to obtain the line+continuum emission. Summation over the integrated intensities, listed in Table 4, of all the transitions in each band provides the corresponding total line emission. Comparing the two provides the line to continuum ratio, which is interesting for the interpretation of broadband continuum images of star forming regions. The advantage of our data is that the line and continuum emission are measured in the same beam with the same instrument and, therefore, there are no complications that arise from cross-calibration between different instruments or beam sizes. Figure 13 plots the band integrated continuum emission (red triangles), the integrated line emission in each band (blue triangles) and, on the right y-axis, the percentage line to continuum ratio (green circles). The figure shows that the line to continuum ratio is ~ 3 – 1% in bands 1a–2a and drops to less than ~ 0.5% in the higher bands. These are both smaller than the ~ 10% seen in Orion-S in the 300 GHz band (Groesbeck 1995) suggesting that the line to continuum ratio generally decreases with increasing frequency. The dramatic drop in the line to continuum ratio between Band 1a to 2b is due to two factors. Over this frequency range, the continuum emission rises by a factor of a few while, at the same time, the number of spectral lines and their corresponding intensity drops by a factor of a few. The red line is a power law fit to the red triangles using a modified black body in which the Planck function is multiplied by κ_o(ν/ν_o)^β, where κ_o is the dust mass opacity coefficient. The best fitted value for β is 1.0. This value is consistent with the behaviour of dust in other studies of star forming regions (e.g. Shetty et al. 2009).

Fig. 13 — Plot of the continuum emission integrated over each band (red triangles), integrated line emission in each band (blue triangles), and the line to continuum ratio percentage (green circles). The red line is a power law fit to the continuum emission using a modified black body function (see text for details).

3. Results & Discussion

3.1. LTE Modeling

LTE modeling assumes that the gas is in Local Thermodynamic Equilibrium, meaning that the density is sufficiently high that collisions dominate the excitation. The LTE modeling capability implemented in CASSIS has 5 input variables, N_t, T_ex, V_LSR, ΔV_FWHM, Ω, where N_t is the total column density, T_ex is temperature, and Ω is the size of the emitting region (which couples to the variable HIFI beam sizes to take into account beam dilution effects). Note that by definition, under LTE conditions, the excitation temperature that determines the relative populations of the upper and the lower level of a spectral line, the rotation temperature that describes the populations of all the rotational levels of one species and the gas kinetic temperature are all identical. Each combination of these variables produces a Gaussian model spectrum for each transition of the selected species. Note that, unlike the Gaussian fitting procedure which fits the V_LSR and ΔV_FWHM to each line separately, for the LTE modeling we obtain a single average value of V_LSR and ΔV_FWHM for all transitions of a given species.

In order to find the set of parameters which produce synthetic spectra that best fit the observed spectral line profiles, we used a Markov Chain Monte Carlo (MCMC) method implemented in CASSIS (e.g. Guan & Krone 2007). The MCMC method randomly picks a seed in the five dimensional parameter space (that we call the X₀ state). Then it randomly chooses one of the nearest neighbors (called the X₁ state), as specified by a variable step size, which is calculated for each iteration. The χ² of the new state is calculated and if p = χ²(X₀)/χ²(X₁) > 1 then the new state is accepted. If p = χ²(X₀)/χ²(X₁) < 1 this new state might still be accepted with a certain acceptance probability. If the new state is rejected the X₀ state will remain, and another random nearby state will be picked as the X₁ state. Having a finite probability to accept a new position even if the χ² is worse, ensures we do not converge directly to a local minimum, but instead forces better sampling of the full parameter space. The code runs with several initial random states and, usually, when the variance among different clusters of states is smaller than the variance of each cluster, it is assumed to have converged to the correct solution (Hastings 1970; Roberts et al. 1997). When the code approaches convergence, it calculates a number of models and χ² values in a tight cluster surrounding the “best” solution. This allows us to calculate a median value for each fitted parameter and its statistical standard deviation, which are listed in Table 5.

Table 5. LTE modeling results for species requiring one component fits.

	N (cm⁻²)	T_ex (K)	ΔV_FWHM (km s⁻¹)	ΔΩ (″)	V_lsr (km s⁻¹⁾
[CI]	1.0 ± 0.2 ×10¹⁸	48.3 ± 8.8	4.2	64.9 ± 14.7	7.8
[CII]	1.4 to 1.9 ×10¹⁸	200.0 to 500.0	4.3	90.0	8.4
CCH	8.9 ± 2.1 ×10¹⁴	36.5 ± 5.9	4.1 ± 0.1	65.7 ± 12.1	7.2 ± 0.0
CH	1.8 ± 0.4 ×10¹⁴	37.7 ± 15.2	4.2 ± 0.3	69.8 ± 8.1	8.0 ± 0.0
CH⁺	2.9 to 4.2 ×10¹³	30.0 to 300.0	4.6	90.0	8.0
CN	2.6 ± 1.3 ×10¹⁴	29.1 ± 3.9	4.3 ± 0.1	67.6 ± 7.8	7.6
¹³CO	1.4 ± 0.4 ×10¹⁷	88.8 ± 14.4	4.2 ± 0.7	79.6 ± 7.8	7.1
C¹⁸O	3.5 ± 0.5 ×10¹⁶	61.9 ± 1.6	3.9 ± 0.0	69.3 ± 13.3	7.1
C¹⁷O	1.4 ± 0.4 ×10¹⁶	53.3 ± 1.9	3.6	64.7 ± 19.9	7.2 ±0.0
¹³C¹⁸O	8.1 ± 1.6 ×10¹⁴	54.3 ± 1.3	4.5	62.6 ± 12.0	7.1
DCN	4.9 ± 1.9 ×10¹²	38.1 ± 5.4	6.1 ± 0.2	60.4 ± 15.4	7.0
DCO⁺	2.4 ± 1.3 ×10¹²	39.9 ± 2.3	3.0 ± 0.1	54.0 ± 15.3	6.7
H₂CS	6.2 ± 0.7 ×10¹³	85.0 ± 2.8	4.5	68.0 ± 11.0	6.7
HCl	1.6 ± 0.6 ×10¹⁴	32.2 ± 5.5	3.9 ± 0.1	52.5 ± 24.1	7.4
H³⁷Cl	6.7 ± 3.4 ×10¹³	52.5 ± 15.6	3.7	53.3 ± 27.1	7.6
HNC	2.4 ± 0.3 ×10¹³	28.2 ± 1.1	4.3 ± 0.1	67.2 ± 3.0	7.3
H¹³CO⁺	7.1 ± 1.8 ×10¹²	39.7 ± 1.3	3.8 ± 0.0	58.6 ± 14.4	7.0
N₂H⁺	1.6 ± 0.5 ×10¹³	46.6 ± 1.9	2.6 ± 0.0	57.9 ± 17.2	6.3
NO	0.3 to 1.1 ×10¹⁶	20.0 to 200.0	5.8	90.0	7.2
SO₂	1.5 ± 0.1 ×10¹⁴	141.3 ± 8.4	6.7 ± 0.1	64.9 ± 9.6	7.5

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Despite the fact that we identify all transitions above 3σ, for our modeling we only utilize transitions above 5σ, while neglecting the blended lines. However, when exploring the validity of our models we also investigate frequency regions where potential transitions of the selected species exist in the molecular line databases, but were not detected above 3σ. This ensures that our models do not produce synthetic spectra where no transitions are actually observed. At the beginning of the procedure we usually let all five parameters vary. However, frequently we were able to find good solutions for the V_LSR, and FWHM after the first convergence of the code. Therefore, on subsequent runs, we fixed the V_LSR and FWHM and allowed the other three parameters to vary. This significantly speeds up the computational time of subsequent runs. Once we obtain a good fit, we run the code five to ten more times to ensure that different runs converge to the same solution within the error bars. In some cases, after running the code five to ten times, the scatter of the converged solutions is larger than the standard deviation of any of the individual solutions. In this case, in Tables 5 and 6 we report the average of the multiple runs (i.e. we take the median value of each run and compute the average between runs) and the standard deviation of the solutions about this new average. In all cases we let the source size (Ω) vary up to 90″ twice that of the largest HIFI beam (~ 45″). However, if the source size is larger than the largest beam it is essentially unconstrained (although there is some sensitivity to source sizes that are larger than the beam since the beam is Gaussian in shape and not a tophat profile). In such cases, the source is simply considered to be extended in nature. Table 5 provides the results of our MCMC χ² fitting of the spectral lines listed in Table 4 (i.e. those with S/N > 5σ). Column 1 is the species, column 2 lists the median total column density of the species and the standard deviation, column 3 is the excitation temperature, column 4 the FWHM line width, column 5 is the source size (Ω), and column 6 is the median LSR velocity. In cases where the error is not listed, the parameter was fixed in the MCMC fitting routine. Table 5 shows that the column density uncertainties range from 10–50%. To ensure that our assumption of LTE is valid, we ran the same MCMC fitting procedure using the non-LTE (RADEX) models implemented in CASSIS for a handful of species (N₂H⁺, SO₂, ¹³CO, and DCO⁺). In all cases, the non-LTE column densities are consistent and within the reported error bars of the LTE models, as shown in Table 5. Given that the errors in Table 5 are the statistical uncertainties on the LTE solutions, to account for the possibility that LTE is not always a good approximation, we suggest that uncertainties on the high side of this range are probably appropriate.

Table 6. LTE modeling results for species requiring two component fits.

	N₁ (cm⁻²)	T_ex1 (K)	ΔV_FWHM₁ (km s⁻¹)	ΔΩ₁(″)	V_lsr1 (km s⁻¹)	N₂ (cm⁻²)	T_ex2 (K)	ΔV_FWHM₂ (km s⁻¹)	ΔΩ₂(″)	V_lsr2 (km s⁻¹)

	Narrow Component					Broad Component
CS	6.5 ± 1.9 × 10¹⁴	36.5 ± 1.5	4.1	66.8 ± 10.8	7.2	7.3 ± 1.4 ×10¹³	108.3 ± 2.4	9.5 ± 0.5	34.5 ± 8.2	7.0
¹³CS ^a	1.4 ± 0.3 ×10¹³	36.5	4.1	74.0 ± 13.6	7.2	1.6 ± 0.2 ×10¹²	109.6 ± 1.3	9.5	34.5	7.0
C³⁴S ^a	4.5 ± 0.7 ×10¹³	34.5	4.1	60.5 ± 8.9	7.2	3.9 ± 0.2 ×10¹²	111.5 ± 3.5	9.5	34.5	7.0
o–H₂S	1.0 ± 0.2 ×10¹⁵	24.5 ± 0.7	3.6	66.1 ± 14.1	7.0	1.3 ± 0.2 ×10¹⁴	85.8 ± 18.3	8.8 ± 1.2	39.9 ± 8.1	7.2
p–H₂S	9.4 ± 1.5 ×10¹⁴	24.3 ± 1.1	3.6	65.4 ± 13.5	7.0	1.5 ± 0.2 ×10¹⁴	74.7 ± 9.9	7.4 ± 1.1	52.5 ± 10.5	7.1 ± 0.1
$H_{2}^{34} S$ ^a	3.2 ± 0.6 ×10¹³	25.1 ± 0.5	3.6	51.7 ± 7.7	7.0	3.0 ± 0.7 ×10¹²	94.0 ± 15.1	9.0 ± 0.7	48.6 ± 10.0	7.2
o–H₂CO	1.5 ± 0.2 ×10¹⁴	47.7 ± 4.3	4.0	69.3 ± 7.2	6.7	3.8 ± 0.6 ×10¹³	153.6 ± 15.0	9.0 ± 0.3	47.7 ± 10.0	7.2 ± 0.2
p–H₂CO	2.0 ± 0.4 ×10¹⁴	44.8 ± 3.0	4.0 ± 0.1	75.2 ± 10.7	6.7	6.7 ± 1.6 ×10¹³	163.0 ± 27.2	9.0 ± 1.0	46.0 ± 9.3	7.3 ± 0.2
HCN	2.2 ± 0.4 ×10¹⁴	34.3 ± 5.0	4.4	64.4 ± 6.5	7.2	4.8 ± 0.1 ×10¹³	66.8 ± 5.2	13.4 ± 0.6	41.1 ± 6.2	6.6 ± 0.1
H¹³CN	5.1 ± 0.7 ×10¹²	30.5 ± 1.2	4.4	74.8 ± 7.4	7.2	3.6 ± 1.9 ×10¹²	79.1 ± 7.8	11.0 ± 0.4	37.0 ± 10.8	6.6
HCO⁺	7.7 ± 0.7 ×10¹³	68.9 ± 1.8	4.3	74.8 ± 8.2	7.0	1.5 ± 0.4 ×10¹³	69.5 ± 4.9	12.0 ± 0.8	40.1 ± 12.1	7.1 ± 0.1
NH₃ p & o	1.4 ± 0.3 ×10¹⁴	20.0	4.2	80.3 ± 5.0	7.0	8.4 ± 5.4 ×10¹³	35.8 ± 2.4	10.0	34.9 ± 11.2	7.2
SO	1.0 ± 0.4 ×10¹⁵	33.7 ± 5.0	3.8 ± 0.5	60.2 ± 17.8	6.8 ± 0.2	2.1 ± 1.1 ×10¹⁴	121.8 ± 13.6	10.7 ± 0.5	34.8 ± 8.9	7.2 ± 0.3

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Calculation based on the main isotopologue

Figure 9 provides an example of one component modeling for ¹³CO. The apparent shift in centroid velocity between the data (black histogram) and the LTE model (red Gaussian curve) is seen in a number of species and can also be seen by comparing the tabulated V_LSR listed in Tables 4 and 5. These apparent shifts of a few tenths of a km s⁻¹ are caused by the fact that the spectral lines are not perfectly Gaussian in shape; in some cases possibly due to optical depth effects. In addition, the MCMC routine optimizes a number of free parameters to obtain the best overall physical model which fits all spectral lines simultaneously; as opposed to the Gaussian fitting routine, which simply fits a mathematical Gaussian profile to each spectral line separately. Thus, the median V_LSR and ΔV_FWHM determined from the LTE modeling may not perfectly match the actual V_LSR and ΔV_FWHM of any individual transition.

Although one component modeling usually results in a remarkably good fit to the observations (e.g C¹⁸O, CH, CCH, etc), there are some cases in which a second (broad) component is necessary to properly reproduce the observations. This is independent of the broad line wings seen in the transitions of some species that required a two component Gaussian fitting as mentioned in Section 2.3. In some cases, even species that were well fit by a single Gaussian component required two component LTE modeling, one with a narrow line width, and the other with a broad line. This is because, in these cases, there is no single combination of model parameters (notably T_ex and N_tot) that could reproduce the intensities of all transitions simultaneously. The two component LTE modeling implementation in CASSIS uses a two slab model; from the perspective of the observer, component 1 is the front slab and component 2 is the slab located behind it. The code allows component 1 to absorb emission from component 2. The results of the two component MCMC LTE modeling are listed in Table 6. Figure 10 provides an example of a species for which two component modeling was required. To see if this could be the result of non-LTE effects, we also attempted to model these species using the RADEX code (van der Tak et al. 2007) as implemented in CASSIS. For the RADEX modeling we used the identical parameter range as for the LTE modeling. RADEX, however, invokes one additional free parameter, namely the H₂ volume density which we allowed to range from 10² to 10¹⁰ cm⁻³. In all cases, the RADEX modeling was unable to produce a good fit to all transitions unless a second physical component was included.

Note that ¹²CO is not presented in either Table 5 or Table 6 due to the presence of self-absorption from foreground material which complicated the modeling procedure. We do, however, model the broad shock/outflow component separately for the analysis presented in Section 3.3.2.

3.2. Comments on Individual Species

The detected species listed in Table 2 can be related to a variety of physical processes that exist in the ISM such as: shocks, UV irradiation by nearby OB stars, and hot core chemistry. In this section we discuss some specific molecules in the context of these physical processes.

3.2.1. Tracers of UV Irradiation

Given the high UV flux in the Orion-S region (χ = 1.1 × 10⁵χ₀; Herrmann et al. 1997), it is not surprising that we detect a wide variety of UV tracers.

[CI] & [CII]

[CI] and [CII] are the fine structure lines of neutral atomic and singly ionized carbon. Both have been seen over large regions of the ISM. [CI] is known to trace PDRs at the UV illuminated surfaces of GMCs (Papadopoulos et al. 2004; Plume et al. 1999; Tielens & Hollenbach 1985), and [CII] is a tracer of the interface between the diffuse warm ionized medium and the outermost surface of GMCs (Velusamy et al. 2012). [CII] is also thought to be a tracer of CO “dark gas” (Langer et al. 2010). We have detected all [CI] and [CII] transitions accessible to HIFI, i.e. both of the [CI] ground-state fine-structure transitions: ³P₁ → ³P₀ and ³P₂ → ³P₁ (V_LSR ~ 7.5 km s ⁻¹), and the single [CII] transition: ²P_3/2 → ²P_1/2 (V_LSR ~ 8.6 km s ⁻¹), toward Orion-S. With only one transition of [CII], we modeled the column density assuming that the excitation temperature in the PDR was between 200–500 K. The velocity of [CII] is considerably different from the velocity of the dense, quiescent cloud component of Orion-S as traced by C¹⁸O, CS, DCO⁺, HCO⁺, etc. (e.g. ~ 7 km s ⁻¹; see Figure 14). This suggests that [CII] is tracing a kinematically distinct component of Orion-S; most probably photoevaporating material moving away from the molecular clump surfaces (e.g. Goicoechea et al. 2015). [CI], however, does not have a velocity that is dramatically different from the quiescent cloud component and is, in fact, similar to that of C¹⁸O (Figure 14). This is probably due to the fact that neutral atomic carbon exists slightly deeper in the cloud (A_V > 3 – 4) where it is still mixed with molecular material (see e.g. Hollenbach & Tielens 1997; Mookerjea et al. 2012).

Fig. 14 — Mean V_LSR for each species derived from the Gaussian fits (Table 4). Red triangles indicate the mean V_LSR for species fit by a single Gaussian component. In cases requiring two component Gaussian fits, the narrow component is indicated by blue circle and the broad component is indicated by a green square. Error bars reflect the range in fitted V_LSR values provided in Table 4.

CH⁺, CH & CCH

CH⁺, CH & CCH are often associated with PDRs, with the former two also being tracers of “CO-dark molecular gas” (Nagy et al. 2013; Gerin et al. 2010). In addition, CH⁺ and SH⁺ can also form via turbulent chemistry in the diffuse ISM (Godard et al. 2012). Transitions above 5σ detected toward Orion-S for these species are listed in Table 4. Enough transitions of CH and CCH were detected above 5σ that we could model the emission from these species (Table 5), both of which were well fit by 1 component models. For CH⁺, we provide a range of column densities for a range of excitation temperatures between 30–200 K. From the Gaussian fits in Table 4, both CH and CH⁺ have similar kinematics (V_LSR > 8.0 km s ⁻¹), whereas the CCH has V_LSR ~ 7.2 km s⁻¹ (see Figure 14). This suggests that CH and CH⁺ trace the same region as the [CII] emission i.e. the UV illuminated surface of the cloud, although possibly a deeper and denser region of the PDR as suggested by Pan et al. (2001). CCH, which has a velocity closer to that of the quiescent gas, likely arises from deeper layers in the cloud (Nagy et al. 2015). The formation pathways for these species may help clarify these velocity differences. For example, CH⁺ forms by an endothermic reaction: C⁺ + H₂ → CH⁺ + H (Federman et al. 1996). The formation of CH follows after a hydrogen abstraction reaction with CH⁺ to form ${CH}_{2}^{+}$ and a subsequent dissociative recombination. Since these two species are closely linked to the C⁺ abundance, through one or two steps in the reaction network, it makes sense that they would be linked physically and, therefore, kinematically. The formation of CCH, however, involves additional steps in the reaction chain, starting with the formation of $C_{2} H_{2}^{+}$ followed by dissociative recombination to form CCH (e.g. Wootten et al. 1980). Since this requires additional reactions involving molecular material, this species is probably more closely linked to the denser molecular gas.

SH⁺ & CO⁺

SH⁺ & CO⁺ are also species thought to trace regions with enhanced UV fields (Nagy et al. 2013). We detect two weak (< 5σ) hyperfine components of SH⁺ in Orion-S, which are too weak to be fitted or modeled but which, interestingly, are seen in emission rather than the usual absorption line profiles seen in the diffuse ISM (Godard et al. 2012). Although we do not report the Gaussian fit parameters of SH⁺ or CO⁺ due to the weakness of the transitions, inspection of their lines suggests V_LSR of ~ 8.5 km s⁻¹ which is virtually identical to the [CII] velocity. Although SH⁺ can form via turbulent chemistry in the diffuse ISM, given the strength of the UV field in Orion-S, it is likely that the main formation pathway is S⁺ + H₂ → SH⁺ + H. Therefore, like [CII], SH⁺ probably also originates in the PDR at the surface of the cloud. The same is true of CO⁺ which has a similar V_LSR (8.5 km s ⁻¹) as SH⁺ and [CII], and like CH⁺ forms directly from C⁺ via the reaction OH + C⁺ → CO⁺ + H.

CN & HCN

CN & HCN have both been detected in Orion-S. While both molecules are good tracers of warm dense gas, the CN/HCN abundance ratio is suggested to be an indirect measure of the UV field (e.g. Fuente et al. 1993); i.e. if the ratio is significantly larger than 1, then the UV field is thought to be enhanced. We have identified the N = 5–4, N = 6–5 and N = 7–6 transitions of CN above the 5σ noise level toward Orion-S and have modeled these transitions using a narrow component (see Table 5). The transitions of HCN (J = 6–5 to 13–12), however, exhibit the characteristic broad line wings that required two component Gaussian fitting and LTE modeling (see Table 6). Figure 10 shows the LTE model fit to our HCN observations. Since the fits are constrained by the rms noise in each spectrum, the higher frequency transitions (which tend to have much larger noise) appear to be less well fit than the lower frequency/lower noise transitions. They are, however, still acceptable fits to the data within the given noise levels. Comparing the CN column density with the narrow component of HCN, we obtain a CN/HCN abundance ratio of 1.2±0.6 indicating a moderately enhanced UV field. Given the high critical densities of these transitions (≥ 10⁸ cm⁻³) it is unlikely that they originate at the UV illuminated cloud surface. Instead, both their critical densities and the CN/HCN abundance ratio of 1.2 suggest that they arise deeper in the cloud (around A_V > 5; Fuente et al. 1993) and, therefore, are more closely associated with the dense molecular gas. This is also borne out by their velocities, which are similar to the dense, quiescent cloud component (Figure 14).

3.2.2. Complex Organic Molecules and precursors

Complex organic molecules are often associated with hot core chemistry. Unlike Orion-KL in which a plethora of complex organic molecules were detected (Crockett et al. 2014), in Orion-S we only detect a handful of molecules that might be considered complex.

CH₃OH

Methanol is an asymmetric top molecule, whose internal rotation results in two distinct symmetry species A–CH₃OH and E–CH₃OH. In total we observed 359 methanol transitions above 3σ toward Orion-S, 170 A–CH₃OH and 189 E–CH₃OH. 198 of the lines were above the 5σ noise level, 111 A–CH₃OH and 87 E–CH₃OH. While methanol is known to be a good temperature probe (e.g. Beuther et al. 2005; Wang et al. 2011), detailed modeling of methanol is beyond the scope of this paper and will be the subject of future work.

H₂CO

Formaldehyde is another commonly used tracer of gas temperature (e.g. Mangum & Wootten 1993), in which the two hydrogen atom spins separate the molecule into distinct ortho and para species. Transitions of H₂CO above the 5σ noise level are listed in Table 4. We needed two component LTE modeling for both the ortho and para H₂CO molecules since one component models could not simultaneously reproduce all observed transitions. The modeling (Table 6) results in low temperatures for the narrow component (T_ex ~ 45 – 50 K) and higher temperatures (T_ex ~ 150 – 165 K) for the broad component. The large linewidths and the fact that the estimated source sizes are quite large (> 45″) may indicate that the high temperature H₂CO emission arises from shocks in the outflows rather than from a “hot core” region. Ortho to para ratios in the narrow and broad components are 0.8±0.1 and 0.6±0.1 respectively. These low values are also consistent with our results for H₂S (see below). The spin temperatures associated with the ortho and para species are 7±1 K and 6±1 K respectively, suggesting that if formaldehyde formed under LTE conditions that the formation temperature was very low.

CH₃OCH₃

Dimethyl ether is a complex molecule detected toward Orion-S. Since this molecule has no transition above 5σ, we only report it as a detection in Table 2.

The lack of complex organic molecules suggests that if hot cores exist in Orion-S they are still in their infancy and have not had time to either expand dynamically or develop chemically. This is not surprising given the very small size of the embedded submillimeter continuum sources detected in the region (Zapata et al. 2005). In addition, if these submillimeter continuum sources are indeed hot cores they are approximately 10 times smaller than the Orion-KL hot core. Therefore, the beam dilution in Orion-S would be a 100 times worse. Thus, any transitions arising from the Orion-KL hot core that have an intensity less than a few K in the survey of Crockett et al. (2014) would be undetectable in our survey if they originate from the considerably smaller region in Orion-S. Alternatively, it is possible that Orion-S is not a massive star forming region at all and, therefore, there are no hot cores in this region. Observations with higher spatial resolution or at lower frequencies (with associated lower excitation temperatures) would help address this issue by revealing the presence of more complex organics.

3.2.3. Pure Shock Tracers

SiO

SiO abundances can be enhanced by more than two order of magnitudes in hot and shocked regions (e.g. Iglesias & Silk 1978; Martin-Pintado et al. 1992a) and SiO emission is often used as a tracer of molecular outflows since the SiO emission traces the outflow material itself, rather than the dense protostellar core (Martin-Pintado et al. 1992b). This is believed to be due to Si-bearing dust grains being shattered by the outflow, followed by a rapid gas-phase reaction with free oxygen to produce SiO (e.g. Schilke et al. 1997b; Gusdorf et al. 2008a,b). A number of outflows have already been identified in Orion-S by Zapata et al. (2006). In our data, although we could not identify any SiO emission above the 5σ level, we have identified three SiO transitions above the 3σ level (J=12–11, 13–12, 14–13, v=0) at the velocity of 6.2 km s⁻¹ (similar to the quiescent gas). With additional spectral smoothing (to a velocity resolution of ~ 4 km s⁻¹) it is clear that these transitions are real, with S/N > 5σ. These transitions are quite broad (ΔV ~ 20–30 km s⁻¹), which is reasonable considering the observed characteristics of the SiO outflows as seen by Zapata et al. (2006). Given the existence of such high-J transitions with excitation energies above ~ 150 K, this indicates the presence of at least a small amount of hot shocked SiO in Orion-S.

3.2.4. Tracers of Quiescent Gas

CO,¹³CO, C¹⁸O, C¹⁷O, & ¹³C¹⁸O

For all carbon monoxide isotopologues, excluding ¹²C¹⁶O itself, one component Gaussian fitting and LTE modeling match the observations remarkably well. ¹²C¹⁶O itself, however, exhibits a broad line wing (clearly tracing an out-flow/shock) and, due to the presence of self-absorption, was not modeled. The existence of an outflow is not visible in any of the ¹²C¹⁶O isotopologue transitions due to their lower abundances. The LTE modeling of ¹³CO is shown in Figure 9. For all ¹²C¹⁶O isotopologues (except ¹³C¹⁸O) we see transitions from J=5–4 to 11–10 above the 3σ level. For ¹³C¹⁸O the highest transition we detect above 3σ is J=7–6. The higher J transitions are buried in the larger noise of the higher frequency HIFI bands. For ¹²C¹⁶O, however, we detect lines up to J = 16–15. For the isotopologues, the typical V_LSR is approximately 7 km s ⁻¹, indicating that these species trace the quiescent gas in the cloud. The V_LSR of the main isotopologue, however, is often a bit higher than this, probably due to the Gaussian fits being skewed by the presence of self-absorption in the spectra, or due to the fact that with its high opacity ¹²C¹⁶O may be tracing a different physical region of the cloud. Interestingly, in Table 5, a correlation can be seen between the CO isotopologues and the derived excitation temperature; with the more optically thick species, which trace the cloud surface (e.g. ¹³CO) having a higher temperature than the optically thin ones which preferentially trace the interior (e.g. ¹³C¹⁸O). This suggests that the external UV field is responsible for much of the heating in Orion-S (see also Tauber & Goldsmith 1990). This is different than the usual case of isolated star formation, in which the gas is predominantly heated internally by the process of gravitational collapse and the formation of an embedded protostar.

Deuterium-Bearing Molecules

Deuterated species are subject of considerable interest in the ISM, since the D/H ratio in molecular clouds can be considerably enhanced over the cosmic value of ~ 10⁻⁵. In Orion-S we detect only a few deuterated species: DCN, DCO⁺, and HDO, which all have velocities similar to that of the quiescent gas. Enhanced deuteration can occur because fractionation reactions involving deuterium are favoured in low-temperature environments associated with pre-stellar cores, the resultant deuterated molecules can freeze onto grains, and then be released back into the gas phase when star formation activity begins to heat the natal gas (e.g. Ceccarelli et al. 2007). Thus, deuterated species such as DCN and HDO can trace the chemical history of the gas. In Orion-S we found the DCN/HCN column density ratio to be 0.02±0.01, suggesting considerable enhancement in cold gas. The DCO⁺/HCO⁺ column density ratio is 0.03±0.02. While the DCO⁺ abundance can be enhanced in cold gas via H₂D⁺, Parise et al. (2009) have shown that deuteration can also occur in the gas phase of warm regions like the Orion Bar via the CH₂D⁺ ion. Although we detected three HDO transitions, it was the only species for which we were not able to find any models that converged to a good solution. Therefore, there is no way to give even a rough estimate for the D/H ratio in water.

N₂H⁺

While N₂H⁺, J = 1–0, is often associated with cold, dense gas, we detect N₂H⁺ transitions from J = 6–5 to 10–9. LTE modeling indicates excitation temperatures of ~ 47 K, suggesting that even the dense gas in Orion-S is quite warm. Previous observations of CH₃C₂H in Orion-S (Bergin et al. 1994) confirm this idea. The V_LSR of N₂H⁺ (Figure 14) also suggests that it originates from the quiescent gas. The upper limit for the column density of N₂D⁺ with the same excitation condition as found for N₂H⁺ is 5×10¹¹ cm⁻². This provides a rough estimate of the D/H ratio of < 0.03.

3.2.5. Tracers of Both Shocked and Quiescent Gas

As previously mentioned, there are a number of species, for which we had to invoke two component LTE modeling in order to fit the observed transitions (see Table 6). Narrow spectral components are usually associated with quiescent gas, whereas broader spectral components trace more dynamic gas that is often associated with shocks. This suggests that species listed in Table 6 can simultaneously exist in both quiescent and shocked gas components. This is not surprising, since Bachiller & Pérez Gutiérrez (1997) show that, in the bipolar outflow L1157, while some species are clearly quiescent gas tracers, many species exist in both components. Of these latter species, their abundances in the shocked gas are often an order of magnitude or more higher than their abundances in the quiescent gas. We will explore the issue of abundances further in Section 3.3.2. Here, however, we will briefly discuss some of the species listed in Table 6 as possible tracers of both shocked and quiescent gas.

H₂O

While H₂O is not listed in Table 6, it is an important molecule in the ISM and has been the subject of a number of important studies using the Herschel Space Observatory in both shocked and unshocked gas. Both the ortho and para forms of H₂O were detected in Orion-S, as well as one transition of $o- H_{2}^{18} O .$ The H₂O transitions required two component Gaussian fitting due to the presence of a broad line wing in the spectra (Table 4). The modeling of water is a complex affair and is beyond the scope of this paper. However, Choi et al. (2014) modeled the ortho and para $H_{2}^{18} O$ in Orion-S and found LTE column densities of 2×10¹¹ cm⁻² and 2×10¹² cm⁻² respectively, which suggests an ortho to para ratio of 0.1, indicating that it is unlikely that water formed under LTE conditions. Their non-LTE analysis of the data, however, brings the ratio up to a factor of 2. Choi et al. (2014) also show that the ortho to para ratio is ~ 0.3 in the nearby Orion Bar. Both values are well below the usual value of 3, which indicates non-LTE formation mechanism for water in both Orion-S and the Bar, possibly due to photodesorption from dust grains.

H₂S & $H_{2}^{34} S$

H₂S is an asymmetric rotor which has ortho and para spin modifications. It is considered to be a tracer of high temperature grain surface chemistry (e.g. Watson & Walmsley 1982). Similar to SO, despite the fact that we fit the H₂S transitions with a single Gaussian in Table 4, H₂S also required a second physical component in order to obtain a good χ² fit from the LTE modeling process. As Table 6 shows, we modeled the H₂S emission from the ortho and para spin modifications separately. In both cases, the narrow component has a linewidth of ~3.6 km s ⁻¹, a low excitation temperature (24 K), and is fairly extended (emission extending beyond the Herschel beam) whereas the broad component has a larger linewidth (~ 8 km s ⁻¹), is warmer (~ 80 K), and yet is still fairly extended (> 35″). The ortho to para ratio in the narrow component is 1.1±0.3 and in the broad component is 0.9±0.1 indicating a spin temperature of 9±2 K. These values are consistent with those determined for H₂O in Orion-S by Choi et al. (2014) and for Formaldehyde (above). Similarly, this low ortho to para ratio suggests either a very low formation temperature for H₂S or that non-LTE effects had an important role in its formation. To model $H_{2}^{34} S,$ we coupled its single ortho transition with those of the common isotopologue. The isotopic ratio is another possible free parameter in CASSIS, which assumes that the other parameters of both isotopologues are identical. The isotopic ratio converged to 31±9. Note, however, that there is only one weak transition of $H_{2}^{34} S$ that we used to determine this ratio. For comparison, typical values for the ³²S/³⁴S ratio in galactic molecular clouds are ~ 19±8 (Lucas & Liszt 1998) which is consistent with the solar value of 23 (Anders & Grevesse 1989).

CS

CS is a well-known tracer of dense gas due to its high critical density (Plume et al. 1992). We observe many transitions of CS, from J = 10–9 to J = 19–18. Like CO, CS requires two components to be successfully fit by a Gaussian profile and modeled; one narrow (ΔV_FWHM ~ 4.1 km s ⁻¹), extended (Ω = 67″), and cool (T = 37 K), and the other broader (ΔV_FWHM ~ 10 km s ⁻¹), moderately extended (Ω = 35″), and warm (T = 108 K). The broad component is not seen in the spectra of the CS isotopologues but was needed to successfully model the transitions. ¹³CS and C³⁴S were modeled in a fashion similar to that of $H_{2}^{34} S$ (i.e. the transitions of the isotopologues modeled simultaneously with those of the common isotope, and found to have isotopic ratio of 46±17 and 14±5 respectively). The Gaussian fitted (Table 5) and modeled velocity (Table 6) of both components is ~ 7 km s⁻¹, suggesting that both components originate in the same material (Figure 14).

HCN

The transitions of HCN (J = 6–5 to 13–12) exhibit the characteristic broad line wing that required both two component Gaussian fitting and LTE modeling (see Table 6). As is usual for other species, the broad component is hotter (67 K vs 34 K), broader (13.4 km s⁻¹ vs 4.4 km s⁻¹), and less spatially extended (41″ vs 64″) than the narrow component. While the broad component’s V_LSR is lower than that of narrow component (6.6 km s⁻¹ vs 7.2 km s⁻¹) both are consistent with the systemic velocity of Orion-S (Figure 14).

HCO⁺

HCO⁺ is another well-known tracer of both dense molecular gas and outflows. We detected the J = 6–5 to J = 13–12 transitions of HCO⁺ in Orion-S, which span a wide range of physical conditions (E_up= 90 K, n_crit ~ 3.2 × 10⁷ cm⁻³ to E_up=389 K, n_crit ~ 6.0 × 10⁸ cm⁻³). HCO⁺ has some of the strongest lines seen in our survey of Orion-S. Despite the fact that we fit the HCO⁺ transitions with a single Gaussian in Table 4, HCO⁺ required a second physical component in order to obtain a good χ² fit from the LTE modeling process (see Table 6). Both components have a velocity of ~ 7 km s⁻¹ and are fairly warm (~ 69 K), suggesting a common origin.

NH₃

We detect two transitions above the 5σ level in Orion-S: one in emission (1_0,0 – 0_0,1) and one in absorption (2_1,1 – 1_1,0). In fact, we detect 3 additional transitions of NH₃ (all in absorption) but since they were just below the 5σ level, we do not report them in Table 4. Modeling these transitions simultaneously requires two components: a cold (T ~ 20K), quiescent (ΔV ~ 4.2 km s ⁻¹) layer of gas in front of a warmer (T ~ 36K), broader (ΔV = 10 km s ⁻¹) component. The V_LSR values of both NH₃ components are consistent with the systemic velocity of the cloud (Figure 14). The presence of absorption lines provides additional evidence for the existence of two components in Orion-S (one warm and one cool). The fact that the low energy transition is seen in emission, whereas the higher energy transition is seen in absorption may be related to the beam size and the strength of the continuum. At 572 GHz the continuum is weaker than it is at 1215 GHz (see Figure 13) and may be too beam-diluted to see absorption. However, at higher frequencies, the beam couples better to the source and absorption may become more prevalent. This suggests that NH₃ may not come from a high density, hot region which is consistent with our conclusion that there are no hot cores in Orion-S.

SO₂

SO₂ is also often a tracer of shocks, since it can freeze onto grain mantles at early evolutionary times when the gas is cold and dense, and later be returned to the gas phase by shocks (e.g. Millar & Herbst 1990; Esplugues et al. 2013). Toward Orion-S we detected a number of SO₂ transitions above the 5σ level (see Table 4). Given the fact that all the observed SO₂ lines are relatively weak, one component modeling matches the observations remarkably well. LTE modeling for this species (Table 5) shows that the SO₂ is fairly warm (T_ex ~ 150 K), broad (ΔV ~ 6.7 km s⁻¹), and extended (~ 65″). The velocity of SO₂ is also similar to the velocity of the quiescent gas (Figure 14).

SO

In contrast with SO₂, and despite the fact that we fit the SO transitions with a single Gaussian in Table 4, SO required a second physical component in order to obtain a good χ² fit from the LTE modeling process (see Table 6). The narrow component has a line width of 3.8 km s⁻¹, a moderately low temperature (34 K), and is extended beyond the Herschel beam. Despite the fact that the broad component (ΔV_FWHM = 11 km s⁻¹) is much warmer (122 K), it is still quite extended in size (35”).

The large linewidths (7–13 km s⁻¹), high temperatures (70–150 K), and extended size (> 30″) determined for the second physical components of these species suggest that the embedded outflows seen in Orion-S (Zapata et al. 2005; Schmid-Burgk et al. 1990; Ziurys et al. 1990) have affected a large volume of the region both thermally and dynamically. Whether or not these shocks have affected the chemistry of the gas will be examined in Section 3.3.

3.3. Chemical Comparison with Orion-KL

One of the main goals of this project is to explore the chemical differences and similarities between Orion-S and Orion-KL. As mentioned in Section 1, a detailed comparison between the chemical abundances in Orion-S and Orion-KL is useful, since both regions presumably formed under similar conditions, but could have developed very different chemical abundances based on differences in their ages, densities, temperatures, radiation fields etc. As part of the HEXOS survey, and in a direct analogue to our study, Crockett et al. (2014) have observed the same frequency range in Orion-KL using the same instrument. Therefore, we have a perfectly matched database, with which to compare our results. In this section we will compare the chemistries of these two regions. All of the Orion-KL data are taken from the HEXOS survey of this region as listed in Crockett et al. (2014) using the column densities derived from their XCLASS LTE modeling of the data.

3.3.1. Chemical Abundances in Orion-S

Our LTE modeling produces column densities but, to obtain chemical abundances, we must scale each column density by the H₂ column density, which is not known directly. Therefore, we use C¹⁸O as a proxy for the H₂ column density. The C¹⁸O transitions we observe in Orion-S are optically thin and have only a single narrow component (Table 4) that is well fit by a column density of 3.5 × 10¹⁶ cm⁻² (Table 5). To convert this to an H₂ column density we use a C¹⁸O:H₂ conversion factor of 1.7×10⁻⁷ (Goldsmith et al. 1997) to obtain N_H₂ = 2.1×10²³ cm⁻².

Other species have been modeled using both a narrow and broad component (Table 6), the latter possibly indicative of gas affected by shock. To compute the chemical abundances in the broad component of species listed in Table 6, we need an estimate of the C¹⁸O column density specifically present in this broad component. Since the broad component is too optically thin to be detected in C¹⁸O or even ¹³CO, we rely on ¹²CO instead for this purpose. In this case, we attempt to model the broad (i.e. line wing) component of our observed ¹²CO lines by fitting three Gaussian components to each transition. The first two Gaussians fit the “main” component of the asymmetric CO profile (see Table 4) and the third fits the outflow (ΔV ~ 18 km s⁻¹). This third component is then modeled via our LTE procedure to estimate the physical parameters of the ¹²CO outflow (T_ex ~ 200 K; N(CO) ~ 7.1 × 10¹⁶ cm^−2; Ω > 30″). Dividing this column density by the canonical ¹²CO:C¹⁸O abundance ratio of 500:1 provides a C¹⁸O column density of 1.4 × 10¹⁴ cm⁻² for the broad component in Orion-S; a factor of ~ 250 smaller than the C¹⁸O column density in the narrow component as measured by directly modeling our C¹⁸O observations. Using the same C¹⁸O:H₂ scaling relationship as above we obtain N_H₂ = 8.2×10²⁰ cm⁻²

Therefore, dividing the modeled column densities of the broad components of the species in Table 6 by 8.2×10²⁰ cm⁻² gives the abundance (with respect to H₂) of all species in Orion-S that are possibly affected by shocks. Dividing the modeled column densities of the rest of the species in Table 5, as well as the narrow component of the species in Table 6, by 2.1×10²³ cm⁻² gives the abundance (with respect to H₂) of all species in Orion-S that likely originate in quiescent gas. The results are provided in Table 7 and illustrated in Figure 15, which clearly show that in the broad component (green squares) the abundances are enhanced by a factor of 10–100 with respect to the narrow component (red triangles). Abundance enhancements of this magnitude indicate classic shock behavior (Bachiller & Pérez Gutiérrez 1997). This suggests that shock chemistry is playing an important role in Orion-S.

Table 7. The abundance of species with respect to H₂.

Species	Narrow comp wrt H₂	Broad comp wrt H₂	Abundance enhancement factor
CCH	4.32×10⁻⁹
CN	1.26×10⁻⁹
C¹⁷O	6.80×10⁻⁸
¹³C¹⁸O	3.93×10⁻⁹
DCN	2.38×10⁻¹¹
HCl	7.77×10⁻¹⁰
H³⁷Cl	3.25×10⁻¹⁰
HNC	1.17×10⁻¹⁰
NO	3.40×10⁻⁸
H¹³CO⁺	3.45×10⁻¹¹
H₂CS	3.01×10⁻¹⁰
SO₂	7.29×10⁻¹⁰
H₂S	9.42×10⁻⁹	3.29×10⁻⁷	35
$H_{2}^{34} S$	1.46×10⁻¹⁰	3.41×10⁻⁹	23
SO	4.86×10⁻⁹	2.47×10⁻⁷	51
¹³CS	6.80×10⁻¹¹	1.88×10⁻⁹	28
C³⁴S	2.19×10⁻¹⁰	4.59×10⁻⁹	21
NH₃	6.80×10⁻¹⁰	9.88×10⁻⁸	145
H₂CO	1.70×10⁻⁹	1.24×10⁻⁷	73

Open in a new tab

Fig. 15 — The abundance (with respect to H₂) of species listed in Tables 5 & 6. Open red triangles indicate the abundance ratio for species fitted by a single component LTE model in Orion-S. In cases requiring two component LTE models for the Orion-S data, the abundance ratio for the narrow component is indicated by solid red triangles and that for the broad component is indicated by solid green squares. Therefore, the dotted line connects species/components that likely trace quiescent gas, whereas the dashed line connects species/components that may trace shocked gas.

3.3.2. Species Common to Both Orion-KL and Orion-S

Crockett et al. (2014) detected ~ 13,000 lines from 39 different molecules (79 species if one includes all the isotopologues). This is considerably more than the 685 lines from 52 species (including isotopologues) that we have detected in Orion-S. In addition, the lines in Orion-KL are typically an order of magnitude stronger than those seen in Orion-S. A more interesting comparison, however, is to examine the abundances of species common to the two sources.

Crockett et al. (2014) produce column densities for many of the detected species in Orion-KL. To compare with the chemical abundances in Orion-S (Section 3.3.1), we must also scale by the H₂ column density in Orion-KL. For Orion-KL we use the C¹⁸O column densities derived by Plume et al. (2012) as a proxy, which breaks down the results for each of the four known kinematic components: the Hot Core (V_LSR ~ 4 − 6 and ΔV ~ 7 − 12 km s⁻¹), the Plateau (V_LSR ~ 7 − 11 and Δ_VLSR ≥ 20 km s⁻¹), the Compact Ridge (V_LSR ~ 7 − 9 and ΔV ~ 3 − 6 km s⁻¹), and the Extended Ridge (V_LSR ~ 8 − 10 and ΔV ~ 2 − 4 km s⁻¹) (Blake et al. 1987). The H₂ column density can then be produced using the same C¹⁸O:H₂ conversion factor of 1.7×10⁻⁷.

Producing abundances in this way does depend on the assumptions regarding the C¹⁸O:H₂ abundance ratio. However, by dividing the abundance of a given species in Orion-KL by the abundance of the same species in Orion-S, we eliminate the C¹⁸O:H₂ abundance ratio altogether and are essentially normalizing to the C¹⁸O column density in each region. This does, of course, assume that C¹⁸O abundances are the same in both sources, which may be reasonable based upon similarities between the observed C¹⁸O:C¹⁷O ratios (e.g. 2.5 in Orion-S, 3.0 in the Hot Core, 6.5 in the Compact Ridge, 3.5 in the Plateau, and 2.3 in the Extended Ridge). These ratios are also consistent with those found by Ladd (2004).

Therefore, we are essentially creating the following ratio:

\frac{X_{K L}}{X_{S}} = \frac{{(\frac{N_{i}}{N_{C^{18} O}})}_{K L}}{{(\frac{N_{i}}{N_{C^{18} O}})}_{S}}

(1)

where N_i is the column density of species i and N_C¹⁸O is the column density of C¹⁸O. The subscripts KL and S refer to this ratio in Orion-KL and Orion-S respectively. Given the four distinct kinematic components of Orion-KL, we create this ratio for the Hot Core (HC), the Plateau (P), the Compact Ridge (CR), and the Extended Ridge (ER) separately and, again, use different values for the C¹⁸O abundance in Orion-S depending on whether the species in question has a narrow or broad spectral line profile.

Figure 16 shows the comparison between Orion-S and the Orion-KL Hot Core. Note that, by common use, the term “hot core” refers to a dense, warm region surrounding a central high mass protostellar object that dominates its energetics (Kurtz et al. 2000). It has been argued that the eponymous hot core in the Orion-KL region does not fulfil this criterion Zapata et al. (2011). Rather, these authors suggest that this region is rather powered by the aftermath of the explosion caused by a stellar merger event (Bally & Zinnecker 2005). Regardless of this, in the present paper we are comparing the chemical abundances of Orion-S with those in what is traditionally referred to as the “hot core” component of Orion KL. The x-axis indicates the species and the y-axis shows the ratio as calculated from Equation 1. Open red triangles indicate molecules for which one component LTE models in Orion-S were sufficient. In cases where we required two components to model the Orion-S data, the solid red triangles indicate the ratio for the narrow component and the solid green squares indicates the ratio for the broad component. The dotted line connects species/components that likely trace quiescent gas, whereas the dashed line connects species that have a broad component. Not all species detected and modeled in Orion-S are represented in this figure. This is due to the fact that Crockett et al. (2014) did not model all their detected species (notably the atomic species), nor did they provide column densities for species in which the lines were optically thick (e.g. CO, ¹³CO, HCO⁺, CS) in Orion-KL. Note also that not every species listed in Figure 16 has a symbol associated with it (e.g. CN, HCl, SO, etc.). This is because emission from these species were not attributed to the HC, but to one or more of the other kinematic components in Orion-KL. Error bars are calculated from the statistical uncertainties determined from our LTE modeling of Orion-S (~ 10−50%; see Table 5 and 6) and with the assumption of 10% error bars of the reported column densities in Orion-KL Crockett et al. (2014) which includes the effects of calibration errors, pointing errors, etc. However, to account for the possibility that LTE is not a good approximation in either Orion-S or Orion-KL we add an additional 40% error to the column densities. This value is based on a comparison of LTE versus non-LTE column density calculations for the Orion-KL Extended Ridge (Crockett et al. 2014).

Fig. 16 — Comparison of the abundances of species detected in the Orion-KL Hot Core to those in Orion-S as given by equation 1 in Section 3.3.2. Symbols are the same as described in Figure 15.

Inspection of Figure 16 clearly shows that the abundances of species in the Orion-KL HC are significantly higher than those in the narrow component of Orion-S (dotted line in Figure 16). Except for CCH, C¹⁷O, and H¹³CO⁺, the abundances in the HC are ≳ 10 times larger than those in Orion-S. Examining the abundance ratios in the narrow component, we obtain $〈 \frac{X_{K L}}{X_{S}} 〉 = 135 (SD = 260)$ where SD is the standard deviation about the mean. The large standard deviation simply reflects the enormous scatter in the ratios (note that the y-axis in Figure 16 is on the log scale). Although still not a good match, the disagreement is smaller for species that have a broad component, (dashed line in Figure 16). In this case we obtain $〈 \frac{X_{K L}}{X_{S}} 〉 = 6 (SD = 12) .$ Given the lack of complex molecules noted in Section 3.2.2 and the poor match in the abundances between the Orion-KL HC and Orion-S, this suggests that the gas detected in this study of Orion-S does not originate in a hot core.

Figure 17 shows that the abundances of species in the Orion-KL CR are also higher than those in the narrow component of Orion-S $(〈 \frac{X_{K L}}{X_{S}} 〉 = 23; SD = 45)$ but the agreement is better than it is for the HC. Again, the match is better to the broad component $(〈 \frac{X_{K L}}{X_{S}} 〉 = 1; SD = 2)$ of the two component fits in Orion-S (dashed line) than it is to the narrow component.

Fig. 17 — Same as for Figure 16 except for the Orion-KL Compact Ridge

Figure 18 shows the comparison between Orion-S and the Orion-KL Plateau region. The match between abundances here is clearly better than it is for the HC or the CR with $〈 \frac{X_{K L}}{X_{S}} 〉 = 14; (SD = 22)$ for the narrow component (dotted line) and 1 (SD = 2) for the broad component (dashed line). Note that although SO and SO₂ are often associated with shocked gas they do not appear in Figure 18. This is because these molecules were optically thick and Crockett et al. (2014) did not provide column densities.

Fig. 18 — Same as for Figure 16 except for the Orion-KL Plateau.

The best agreement with molecular abundances in the narrow component of Orion-S is with the Extended Ridge of Orion-KL (Figure 19) where we obtain $〈 \frac{X_{K L}}{X_{S}} 〉 = 7 (SD = 14) .$ For the broad component (dashed line) we obtain $〈 \frac{X_{K L}}{X_{S}} 〉 = 0.3 (SD = 0.5) .$

Fig. 19 — Same as for Figure 16 except for the Orion-KL Extended Ridge.

Given that the best match to the abundances in the narrow component of Orion-S is the ER of Orion-KL, it seems as though these species/components do indeed trace quiescent gas. In particular, it probably is the same gas out of which both star forming regions have been formed. The broad component of Orion-S, however, seems better matched to the CR and Plateau of Orion-KL. Figures 16 to 19 only have a few broad component points and, therefore, it is difficult to make any strong statistical arguments based on these data alone. However, this evidence along with the chemical abundance analysis presented in Section 3.3.1 provide fairly strong support for the idea that shocks have also had an influence on the chemistry of Orion-S.

3.3.3. Species Detected in Orion-KL but not in Orion-S

It is well known that Orion-KL has an incredibly rich molecular chemistry (e.g. Schilke et al. 1997a; Comito et al. 2005; Olofsson et al. 2007; Persson et al. 2007; Leurini et al. 2006; Tercero et al. 2005). However, it is possible that the species that were detected in Orion-KL, but not in Orion-S, exist in the latter source, but at levels too weak to be detected. Some of these species might be observable with ALMA at lower frequencies. In this Section, we explore this possibility by providing upper limits for the abundances of all the species detected in Orion-KL by Crockett et al. (2014) but not detected in Orion-S.

Modeling was accomplished by fixing T_ex, ΔV_FWHM, Ω, and V_LSR and finding the column density that produced transitions whose intensities were < 3σ across all the HEXOS bands. Since we were able to model all of our species using, at most, two components, we determine two different column density upper limits: one assuming that the undetected emission arises from the narrow component and assuming that it originates in the broad component. For the narrow component, we used fixed values of T_ex = 40K, ΔV_FWHM = 4 km s ⁻¹, Ω = 60″, and V_LSR = 7.1, which were found to be typical for the narrow component (see Tables 5 and 6). For the broad component, we used fixed values of T_ex = 80K, ΔV_FWHM = 8 km s ⁻¹, Ω = 40″, and V_LSR = 7.1, which were found to be typical for the broad component (see Table 6). Results are listed in Table 8. Column 1 is the species name, column 2 the upper limit total column density assuming the gas arises in the narrow component, column 3 the upper limit total column density assuming the gas originates in the broad component, and column 4 lists the maximum upper state energy for the model (i.e. no transitions with E > E_up were modeled). Different values of E_up were used for different molecules to keep the number of modeled transitions to a reasonable value.

Table 8. Column density upper limit for species in Orion-KL not detected in Orion-S.

Species	N_{t, T_ex=40K}	N_{t, T_ex=80K}	E_u
¹⁵NH₃	7.0 × 10¹¹	3.0 × 10¹²	700
²⁹SiO	2.0 × 10¹²	1.0 × 10¹²	700
³⁰SiO	3.0 × 10¹²	2.0 × 10¹²	300
C₂H₃CN	N/A	N/A	300
C₂H₅CN	N/A	N/A
C₂H₅OH	N/A	N/A
C³³S	6.0 × 10¹²	4.0 × 10¹²	700
CH₂DOH	6.0 × 10¹³	6.0 × 10¹³	300
CH₂NH	7.0 × 10¹²	2.0 × 10¹³	300
${CH}_{3}^{13} CN$	3.0 × 10¹⁴	2.0 × 10¹³	700
CH₃CN	7.0 × 10¹⁴	5.0 × 10¹³	700
CH₃CN, v₈ = 1	N/A		700
¹³CH₃CN	4.0 × 10¹⁴	2.0 × 10¹³	700
CH₃OCHO	2.0 × 10¹⁵	2.0 × 10¹⁵	300
CH₃OD	N/A	N/A	700
¹³CH₃OH	N/A	N/A	300
D₂O	5.0 × 10¹¹	2.0 × 10¹²	300
$H_{2}^{13} CO$	1.5 × 10¹⁵	1.5 × 10¹⁴	700
$H_{2}^{17} O$	5.0 × 10¹¹	1.5 × 10¹²	300
H₂CCO	1.5 × 10¹⁵	2.0 × 10¹⁴	700
H₂O, v₂	N/A	N/A	700
H¹³CN, v₂ = 1	N/A	N/A	700
HC₃N	2.0 × 10¹⁷	5.0 × 10¹⁴	700
HC₃N, v=0	4.0 × 10¹⁷	6.0 × 10¹⁴	700
HCN, v₂ = 1	N/A	N/A	700
HCN, v₂ = 2	N/A	N/A	700
HD¹⁸O	5.0 × 10¹¹	3.0 × 10¹²	300
HDCO	7.0 × 10¹²	9.0 × 10¹²	300
HN¹³C	8.0 × 10¹¹	1.0 × 10¹²	300
HN¹³CO	5.0 × 10¹²	1.0 × 10¹³	300
HNC, v₂ = 1	N/A	N/A	700
HNCO	5.0 × 10¹²	1.0 × 10¹³	300
NH₂CHO	3.0 × 10¹³	3.0 × 10¹³	700
NH₂D	4.0 × 10¹²	8.0 × 10¹²	300
NH₃, v₂	N/A	N/A	700
NS	1.5 × 10¹³	1.0 × 10¹³	300
O₂	2.0 × 10¹⁷	4.0 × 10¹⁷	700
OCS	3.0 × 10¹⁷	3.0 × 10¹⁵	700
OD	1.0 × 10¹⁴	1.5 × 10¹³	700
OH	3.0 × 10¹⁵	1.5 × 10¹⁴	300
SiS	9.0 × 10¹⁴	6.0 × 10¹³	700
SO₂, v₂ = 2	N/A	N/A	700
³³SO	3.0 × 10¹³	1.5 × 10¹³	700
³⁴SO	5.0 × 10¹³	2.0 × 10¹³	300
³³SO₂	2.0 × 10¹³	4.0 × 10¹³	300
³⁴SO₂	1.5 × 10¹³	3.0 × 10¹³	200

Open in a new tab

There are also a number of species that we detected (S/N > 3σ), but did not model, since their S/N was < 5σ. Using the same assumptions as for the undetected species we provide upper limits for the abundances for these species in Table 9. The column density limits listed in Table 9 are approximately an order of magnitude or more smaller than the column densities of the same species detected in Orion-KL (Crockett et al. 2014).

Table 9. Column density upper limit for unmodeled species in Orion-S.

Species	N_{t, T_ex=40K}	N_{t, T_ex=80K}
CH₃OCH₃	1.5 × 10¹⁴	3 × 10¹⁴
CO⁺	2 × 10¹²	1.5 × 10¹²
$H_{2}^{18} O$	2 × 10¹²	3 × 10¹²
HC¹⁵N	1 × 10¹²	2 × 10¹²
HCS⁺	2.5 × 10¹³	1 × 10¹³
NH₂	7 × 10¹²	1 × 10¹³
SH⁺	4 × 10¹²	7 × 10¹²
SiO	1 × 10¹³	4 × 10¹²

Open in a new tab

A possible question is whether the species listed in Table 8 could actually exist in Orion-S even though they are undetected. We inspect the excitation conditions of these species, examining E_up of all possible transitions, to determine whether they are detectable given the reported noise of the HIFI bands. A portion of these species have E_up much greater than 500 K (e.g. H₂O v₂, H¹³CN v₂ = 1, HC₃N, HC₃N v=0, HCN v₂ = 1, HCN v₂ = 2, OCS, etc.). Based on the analysis done in this paper, Orion-S can barely excite species with E_up > 500 K. Therefore, transitions of these species would not be observed above the noise, even if they were present. In addition, there are some complex organic species with transition of E_up < 100 K which we also did not detect (e.g. CH₂NH, CH₂DOH, CH₂NH, CH₃OCHO, etc.). This is not surprising given the absence of hot core chemistry in Orion-S.

4. Summary and Conclusion

We have presented results from a comprehensive spectral survey toward Orion South, taken with the HIFI instrument aboard the Herschel Space Observatory covering the frequency range 480 to 1900 GHz with a resolution of 1.1 MHz. We detected 685 spectral lines with S/N > 3σ originating from 52 different molecular and atomic species. Using the CASSIS spectral line analysis software package, we modeled each of the detected species assuming conditions of Local Thermodynamic Equilibrium. Based on this modeling, we found evidence for three different cloud components: a cool (T_ex ~ 20 − 40 K), spatially extended (> 60″), and quiescent (ΔV_FWHM ~ 4 km s ⁻¹) component; a warmer (T_ex ~ 80 − 100 K), less spatially extended (~ 30″), and more dynamic (ΔV_FWHM ~ 8 km s ⁻¹) component, which is likely affected by embedded outflows; and a kinematically distinct region dominated by emission from species which trace UV irradiation. Indirect evidence to support the existence of the first two components can be inferred from McMullin et al. (1993) who mapped the region in a few spectral lines (SiO, H¹³CO⁺, SO₂, CH₃OH, and HC₃N) with the BIMA array. Their H¹³CO⁺ and HC₃N data confirm the existence of a fairly extended (~1′) quiescent (FWHM ~3 km s⁻¹) component, whereas the SO₂ and CH₃OH data reveal a smaller emitting region (~20″) of warm gas (~75 K). While the spectra for the latter two species are too weak to determine line widths, their SiO data reveal a similarly small region (offset by only a few arc seconds from the SO₂ and CH₃OH emission peaks) with broad line widths (~7 km s⁻¹). In addition, McMullin et al. (1993) reports column densities of SO₂ and H¹³CO⁺ of < 2×10⁻¹⁰ and 4×10⁻¹¹ respectively which compare favourably to the values reported in Table 7. Finally, while there are no higher resolution observations to confirm the existence of the third component (i.e. the UV irradiated region), since CO⁺ is only ever detected in PDRs, its presence in our data strongly suggests that such a component must exist.

We also presented a comprehensive chemical abundance comparison between Orion-KL and Orion-S; two star forming regions that potentially formed from the same natal molecular gas but are at different evolutionary stages. Based on a paucity of complex molecules in Orion-S, we found little chemical evidence for the existence of a significant “hot core” component. This is likely due to the fact that either the hot cores associated with the embedded star formation have either not had sufficient time to develop chemically, or that they are simply too small for their line emission to be detected in the large Herschel beam, or that Orion-S is not a massive star forming region and hot cores massive enough to produce their characteristic rich spectra simply do not exist. The presence of a number of UV tracers such as [CII], [CI], CH, CH⁺, SH⁺, CO⁺, and the fact that transitions of these species have velocities that are 1-1.5 km s⁻¹, higher than those of the quiescent gas, suggest that these species arise from a kinematically distinct PDR; most likely the UV illuminated surface of the cloud. The best match to the chemical abundances in the cooler, quiescent gas in Orion-S is with the quiescent extended ridge of Orion-KL, indicating that most of the gas in Orion-S is still quiescent as well, and relatively unaffected by higher temperature or UV driven chemistry. The best agreement with the warmer, broad component of Orion-S is with the Orion-KL Plateau and Compact Ridge regions, suggesting that shocks have had an influence on the overall chemistry in Orion-S.

Fig. 2 — Same as Figure 1 except for band 3a, band 3b, band 4a and band 4b.

Fig. 3 — Same as Figure 1 except for band 5a, band 5b, band 6a and band 6b. The higher noise level in band 6 is due to the HEB mixers which produce higher noise in comparison with the SIS mixers used in the first five bands.

Acknowledgements

HIFI was designed and built by a consortium of institutes and university departments from across Europe, Canada and the United States under the leadership of SRON Netherlands Institute for Space Research, Groningen, The Netherlands and with major contributions from Germany, France and the US. Consortium members are: Canada: CSA, U.Waterloo; France: CESR, LAB, LERMA, IRAM; Germany: KOSMA, MPIfR, MPS; Ireland, NUI Maynooth; Italy: ASI, IFSI-INAF, Osservatorio Astrofiscio di Arcetri-INAF; Netherlands: SRON, TUD; Poland: CAMK, CBK; Spain: Observatorio Astronmico Nacional (IGN), Centro de Astrobiologa (CSIC-INTA). Sweden: Chalmers University of Technology - MC2, RSS & GARD; Onsala Space Observatory; Swedish National Space Board, Stockholm University - Stockholm Observatory; Switzerland: ETH Zurich, FHNW; USA: Caltech, JPL, NHSC. We also need to acknowledge the support by the Deutsche Forschungs-gemeinschaft (DFG) via the collaborative research grant SFB 956, project C1 & C3, as well as the ERC and the Spanish MINECO for funding support under grants ERC-2013-Syg-610256 and AYA2012-32032. Support for this work was provided, in part, by a National Sciences and Engineering Research Council of Canada (NSERC) grant to R. Plume and K. Tahani and by NASA through an award issued by JPL/Caltech. This work was carried out in part at the Jet Propulsion Laboratory, which is operated for NASA by the California Institute of Technology.

Facilities: Herschel.

Footnotes

Also see: http://www.astro.uni-koeln.de/cdms/molecules and http://www.astrochymist.org/astrochymist_ism.html

http://herschel.esac.esa.int/twiki/bin/view/Public/HifiCalibrationWeb#HIFI_performance_and_calibration.

CASSIS was developed by IRAP-UPS/CNRS. See: http://cassis.irap.omp.eu

⁴

http://spec.jpl.nasa.gov/ftp/pub/catalog/doc/catintro.pdf

⁵

http://www.astro.uni-koeln.de/cdms/catalog

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PERMALINK

Analysis of the Herschel/HEXOS Spectral Survey Towards Orion South: A massive protostellar envelope with strong external irradiation

K Tahani

R Plume

E A Bergin

V Tolls

T G Phillips

E Caux

S Cabrit

J R Goicoechea

P F Goldsmith

D Johnstone

D C Lis

L Pagani

K M Menten

H S P Müller

V Ossenkopf-Okada

J C Pearson

F F S van der Tak

Abstract

1. Introduction

2. Observations & Data Reduction

Table 1. Data Smoothing and Noise Characteristics.

Fig. 1.

Fig. 4.

2.1. Removal of Off-Position Contamination

Fig. 5.

Fig. 6.

Fig. 7.

2.2. Line Identification

Table 2. Identified species in Orion-S.

Table 3. All observed lines above the 3σ noise level in order of increasing frequency.

Fig. 8.

Table 4. Gaussian fits to lines above 5σ.

2.3. Line Profiles

Fig. 9.

Fig. 10.

Fig. 11.

Fig. 12.

Fig. 13.

3. Results & Discussion

3.1. LTE Modeling

Table 5. LTE modeling results for species requiring one component fits.

Table 6. LTE modeling results for species requiring two component fits.

3.2. Comments on Individual Species

3.2.1. Tracers of UV Irradiation

[CI] & [CII]

Fig. 14.

CH+, CH & CCH

SH+ & CO+

CN & HCN

3.2.2. Complex Organic Molecules and precursors

CH3OH

H2CO

CH3OCH3

3.2.3. Pure Shock Tracers

SiO

3.2.4. Tracers of Quiescent Gas

CO,13CO, C18O, C17O, & 13C18O

Deuterium-Bearing Molecules

N2H+

3.2.5. Tracers of Both Shocked and Quiescent Gas

H2O

H2S & H234S

CS

HCN

HCO+

NH3

SO2

SO

3.3. Chemical Comparison with Orion-KL

3.3.1. Chemical Abundances in Orion-S

Table 7. The abundance of species with respect to H2.

Fig. 15.

3.3.2. Species Common to Both Orion-KL and Orion-S

Fig. 16.

Fig. 17.

Fig. 18.

Fig. 19.

3.3.3. Species Detected in Orion-KL but not in Orion-S

CH⁺, CH & CCH

SH⁺ & CO⁺

CH₃OH

H₂CO

CH₃OCH₃

CO,¹³CO, C¹⁸O, C¹⁷O, & ¹³C¹⁸O

N₂H⁺

H₂O

H₂S & $H_{2}^{34} S$

HCO⁺

NH₃

SO₂

Table 7. The abundance of species with respect to H₂.