First Abundance Measurement of Organic Molecules in the Atmosphere of HH 212 Protostellar Disk
aa r X i v : . [ a s t r o - ph . GA ] A p r First Abundance Measurement of Organic Molecules in theAtmosphere of HH 212 Protostellar Disk
Chin-Fei Lee , , Claudio Codella , , Zhi-Yun Li , and Sheng-Yuan Liu ABSTRACT
HH 212 is one of the well-studied protostellar systems, showing the first ver-tically resolved disk with a warm atmosphere around the central protostar. Herewe report a detection of 9 organic molecules (including newly detected ketene,formic acid, deuterated acetonitrile, methyl formate, and ethanol) in the diskatmosphere, confirming that the disk atmosphere is, for HH 212, the chemicallyrich component, identified before at a lower resolution as a “hot-corino”. Moreimportantly, we report the first systematic survey and abundance measurementof organic molecules in the disk atmosphere within ∼
40 au of the central proto-star. The relative abundances of these molecules are similar to those in the hotcorinos around other protostars and in Comet Lovejoy. These molecules can beeither (i) originally formed on icy grains and then desorbed into gas phase or (ii)quickly formed in the gas phase using simpler species ejected from the dust man-tles. The abundances and spatial distributions of the molecules provide strongconstraints on models of their formation and transport in star formation. Thesemolecules are expected to form even more complex organic molecules needed forlife and deeper observations are needed to find them.
Subject headings: stars: formation — ISM: individual: HH 212 — ISM: molecules— ISM: accretion and accretion disk – ISM: jets and outflows. Academia Sinica Institute of Astronomy and Astrophysics, P.O. Box 23-141, Taipei 106, Taiwan;cfl[email protected] Graduate Institute of Astronomy and Astrophysics, National Taiwan University, No. 1, Sec. 4, RooseveltRoad, Taipei 10617, Taiwan INAF, Osservatorio Astrofisico di Arcetri, Largo E. Fermi 5, 50125 Firenze, Italy Univ. Grenoble Alpes, CNRS, Institut de Plan´etologie et d’Astrophysique de Grenoble (IPAG), 38000Grenoble, France Astronomy Department, University of Virginia, Charlottesville, VA 22904, USA
1. Introduction
Accretion disks have been detected in very young protostellar systems, feeding the cen-tral protostars. With the advent of the Atacama Large Millimeter/submillimeter Array(ALMA), we have started to resolve the disks and study their physical processes in greatdetail. HH 212 is one of the well-studied protostellar systems, showing the first vertically re-solved disk with a warm atmosphere around the central protostar (Lee et al. 2017b,c). Thiswarm atmosphere seems to be the hot corino reported recently at a lower angular resolution(Codella et al. 2016, 2018). Hot corinos are the hot ( &
100 K) and compact regions im-mediately around low-mass (sun-like) protostars (Ceccarelli et al. 2007), and rich in organicmolecules including complex organic molecules (COMs, refering to C-bearing species withsix atoms or more, Herbst & van Dishoeck 2009). By determining the connection of the diskatmosphere with the hot corino in HH 212, we aim to determine the origin of the hot corinoand the related physical processes in the innermost region. In particular, the nearly edge-onorientation of the disk in this system provides the best view of the atmosphere, allowing usto study the physical properties of the disk atmosphere, and the formation of the organicmolecules and their role in producing the rich organic chemistry needed for life.HH 212 is a young Class 0 protostellar system deeply embedded in a compact molecularcloud core in the L1630 cloud of Orion at a distance of ∼
400 pc (Kounkel et al. 2017).The central source is IRAS 05413-0104, with a bolometric luminosity of ∼ L ⊙ (updatedfor the new distance) (Zinnecker et al. 1992). The central protostar has a mass of ∼ M ⊙ (Lee et al. 2017c). It drives a powerful bipolar jet (Zinnecker, McCaughrean, & Rayner1998; Lee et al. 2015), which is recently found to be spinning (Lee et al. 2017a). A rotatingdisk must have formed around the protostar in order to launch the jet. Our previous ALMAobservations towards the center indeed showed a spatially resolved nearly edge-on dusty diskwith a radius of ∼
60 au (Lee et al. 2017b). In addition, we also detected a warm atmosphereof the disk with a few organic molecules (Lee et al. 2017c), suggesting that the warm diskatmosphere can be the hot corino reported before at a lower resolution (Codella et al. 2016).Recent observations at a resolution of ∼ ′′ . 15 (60 au) suggested that deuterated water andCH CHO can also reside in the disk atmosphere (Codella et al. 2018). In this paper, wezoom in to the disk region at a higher resolution of ∼ ′′ . 03 (12 au) and higher sensitivity anddetect additional and more complex organic molecules characteristic of a hot corino, withmost of them detected for the first time in the disk atmosphere. Our observations confirmthat the hot corino in HH 212 is indeed the warm disk atmosphere. We will discuss theformation of the organic molecules by comparing their abundances to those in hot corinosaround other low-mass protostars. 3 –
2. Observations
The HH 212 protostellar system was observed with ALMA in Band 7 at ∼ ∼ ∼
98 minutes. Theprojected baselines were ∼ ∼ ∼ − per channelat 346.5 GHz. The primary beam was 17 ′′ . 27. A single pointing was used to map the systemwithin ∼ ′′ of the central source. The maximum recoverable size scale was ∼ ′′ . 4, enoughto map the disk atmosphere without any significant missing flux.The data were calibrated with the CASA package, with quasar J0510+1800 as a pass-band and flux calibrator, and quasar J0541-0211 (a flux of 0.137 Jy) as a gain calibrator.We used a robust factor of 0.5 for the visibility weighting to generate the channel maps witha synthesized beam of 0 ′′ . 036 × ′′ . 03 at a position angle of ∼ − ◦ . We used the line-freechannels to generate a continuum map centered at 356.5 GHz. The channel maps of themolecular lines were generated after continuum subtraction. The noise levels are ∼ − (or ∼ − (or ∼
3. Results
In HH 212, the jet has an axis with a position angle of ∼ ◦ and an inclination angle of ∼ ◦ to the plane of the sky, with its northern component tilted toward us (Claussen et al.1998). The disk is nearly edge-on and exactly perpendicular to the jet axis (Lee et al. 2017b).The systemic velocity is V sys = 1 . ± . − LSR (Lee et al. 2014). In order to facilitateour presentations, we define an offset velocity V off = V LSR − V sys and rotate our maps by 23 ◦ clockwise to align the jet axis in the vertical direction.Figure 1 shows the emission line intensity maps (contours) of nine organic molecules, in-cluding CH OH (methanol, 4 lines), CH DOH (deuterated methanol, 7 lines), CH OH ( Cisotopologue of methanol, 1 line, contaminated by a much weaker CH DCN line), H CCO(ketene, 1 line), CH CHO (acetaldehyde, 22 lines), HCOOCH (methyl formate, 8 lines),t-HCOOH (formic acid in trans state, 1 line), CH CH OH (ethanol, 10 lines), and CH DCN 4 –(deuterated acetonitrile, 5 lines), on the continuum map (color image) of the disk at λ ∼ µ m (or equivalently ∼
347 GHz), obtained from our observations. As seen before at a similarwavelength ( λ ∼ µ m) in Lee et al. (2017b), the continuum map shows a “hamburger-shaped” dusty disk with an equatorial dark lane sandwiched between two brighter featureson the top and bottom. As discussed in that paper, the presence of the equatorial darklane is due to relatively low temperature and high optical depth near the disk midplane. Asdiscussed below, the maps of the molecular emissions are obtained by stacking a number oflines (as indicated above in the parenthesis) in different transitions with a range of upperenergy levels for better detections. Molecular line emissions are only detected in the upper(above the midplane) and lower disk atmosphere, with the emission brighter in the lowerdisk atmosphere below the midplane. No molecular emission is detected toward the diskmidplane, either because the emission is lost in the optically thick dust continuum emissionor because of a lack of emission of these molecules there. All the molecular emissions aredetected within the centrifugal barrier (which has a radius of ∼ ′′ . 11 or 44 au) of the centralprotostar. For CH OH, CH DOH, CH CHO, and CH CH OH, their emissions are clearlydetected in the upper and lower atmosphere with a roughly similar distribution, suggestingthat they are chemically related. In addition, the emission moves closer to the disk mid-plane from CH OH, CH DOH, CH CHO, to CH CH OH, suggesting that the emission ofless abundant molecule (see next section for their abundances) traces deeper into the diskatmosphere, probably due to an optical depth effect. However, it could also be due to achemical stratification in the vertical direction. CH DOH and CH CH OH emission in thelower atmosphere in the outer edge show a structure curving back towards the midplane,following the lower boundary of the dusty disk emission, likely outlining a physical boundaryfor the dusty disk.Figure 2 shows the position-frequency (PF) diagrams obtained with a cut along thelower atmosphere, where the emission is brighter, in order to identify the detections ofvarious molecular lines. The PF diagrams for the upper atmosphere show similar structuresbut fainter and are thus not shown here. As can be seen, many lines are detected, with theiremission detected within ∼ ∼ − within the systemicvelocity) marked by the vertical lines, with one color for each organic molecular species atdifferent transitions. For each line, the diagrams show a roughly linear PF structure with theredshifted emission in the southeast (positive y ) and blueshifted emission in the northwest(negative y ), as seen before in the corresponding position-velocity diagrams for the CH OHand CH DOH lines (Lee et al. 2017c). As discussed in that paper, this linear PF structurelikely arises from a warm rotating ring of the disk atmosphere near the centrifugal barrier.With this feature, we identify lines from the nine organic molecules mentioned above, anda few other simple molecules (e.g., SO, SO and its isotopologue SO O). Other lines from 5 –CO and SiO trace mainly the outflow and jet, and thus do not show such a clear linear PFstructure.Table 1 lists the properties of the organic molecular lines. For each molecular species, westacked the line intensity maps in different transitions, excluding those tentatively detected(marked with T) and blended (marked with B), producing the mean line intensity map shownin Figure 1.
Here we derive the mean excitation temperature and column density in the disk atmo-sphere by fitting the rotation diagram of the molecular lines. This diagram plots the columndensity per statistical weight in the upper energy state in the optically thin limit, N thin u /g u ,versus the upper energy level E u of the lines. Here N thin u = (8 πkν /hc A ul ) W , where theintegrated line intensity W = R T B dv with T B being the brightness temperature.The emission in the lower disk atmosphere is brighter and is thus used to better derivethe mean excitation temperature and column density in the disk atmosphere. Table 1 liststhe integrated line intensities in the lower disk atmosphere measured (with a cutoff of 2 σ )for the reasonably isolated lines detected with more than 3 σ . They are the mean valuesaveraged over a rectangular region that covers most of the emission in the lower atmosphere.Figure 3 shows the resulting diagrams for the six molecules detected with multiple lines. Theblended lines are excluded. With the CH OH lines, we derived an excitation temperatureof ∼ ±
45 K. With the lines from its deuterated species CH DOH, we derived a similartemperature of ∼ ±
19 K. For a check, we also obtained similar temperatures of ∼ ±
32 K and ∼ ±
24 K for the upper atmosphere from the CH OH lines and CH DOHlines, respectively. The resulting column densities of CH OH and CH DOH are listed inTable 2. As discussed later, since the CH OH lines are likely optically thick, the CH OHcolumn derived here is only a lower limit and a more accurate value will be derived usingthe CH OH column density below.From CH OH and CH DOH, we obtain a mean excitation temperature of ∼ ±
50 Kin the lower disk atmosphere, similar to that found before in Lee et al. (2017c). This meantemperature is also consistent with that derived towards the disk at a lower angular resolutionin Bianchi et al. (2017). Assuming this mean temperature for the disk atmosphere, we canestimate the column densities of other molecules with weaker lines. For those moleculesdetected with multiple lines, such as CH CHO, HCOOCH , CH CH OH, and CH DCN, wecan obtain their column densities by fitting their rotation diagrams, as shown in Figures 6 –3(c)-(f). Notice that the column density of CH CHO is estimated here to be (1.5 ± × cm − , with the lower limit consistent with that estimated previously at a lower resolutionwith an excitation temperature of ∼ ±
14 K in Codella et al. (2018). For those detectedwith a single line, such as CH OH, H CCO, and t-HCOOH, we derived their columndensities from their measured integrated line intensity. For CH OH, the only detectedline is contaminated by a weaker line of CH DCN and thus its column density is estimatedafter removing the expected intensity of the CH DCN line. The expected intensity of theCH DCN line is assumed to be given by the mean intensity of other CH DCN lines withsimilar E u and log A ul , and it is estimated to be ∼
21 K km s − , or ∼
18% of the totalintensity there. Also, since the CH OH line is optically thinner than the CH OH lines, wecan improve the CH OH column density by multiplying the CH OH column density bya C/ C ratio of 50 as obtained in the Orion Complex (Kahane et al. 2018). As can beseen from Table 2, the CH OH column density derived this way is about twice that derivedfrom the rotation diagram, suggesting that the CH OH lines are indeed optically thick. Alsoshown in the table are the column densities of NH CHO, D CO, and CH SH measured inLee et al. (2017c), adjusted with the mean excitation temperature here.Based on our disk model that reproduced the dust continuum emission at a similar wave-length of ∼ µ m (Lee et al. 2017b), the dust continuum emission in the disk atmospherehas an optical depth τ .
1. Thus, the derived column densities of the organic moleculescould be underestimated by a factor of e . In that disk model, the disk atmosphere has amean H column density of ∼ × cm − . Hence, the abundances of the molecules, aslisted in Table 2, can be obtained by dividing the column density of the molecules by thismean H column density. The abundances here can be uncertain by a factor a few becausethe mean H column density depends on dust opacity, which can be uncertain by a factor afew.
4. Discussion4.1. Corino-like Disk Atmosphere
Hot corinos with a temperature &
100 K have been detected around low-mass proto-stars and they are rich in organic molecules (including complex organic molecules). A hotcorino has also been recently reported in HH 212 in the inner 100 au with a detection ofCH CHO and deuterated water HDO (Codella et al. 2016, 2018). Now at higher resolution,we find CH CHO to reside in the disk atmosphere. Moreover, we also detect 8 other or-ganic molecules (H CCO, t-HCOOH, CH OH, CH DOH, CH OH, CH DCN, HCOOCH ,and CH CH OH) in the disk atmosphere, with an excitation temperature of ∼ ±
50 K. 7 –Adding 3 other organic molecules (D CO, CH SH and NH CHO) from Lee et al. (2017c), wehave 12 organic molecules, with 9 of them being complex organic molecules, detected in thedisk atmosphere. These results indicate that the hot corino in HH 212 is actually a warmatmosphere of the disk.
We can study the formation of the organic molecules in the disk atmosphere by com-paring their abundances to those estimated in the hot corinos around other Class 0 low-massprotostars. The abundance of CH CHO, which is commonly observed in hot corinos, is es-timated here to be 3.9 × − in HH 212, similar to those found in other hot corinos, e.g.,5.8 × − in IRAS 16293-2422B (Jørgensen et al. 2016), 2.4 × − in B335 (Imai et al. 2016),and 4.2 × − in NGC 1333 IRAS4A2 (L´opez-Sepulcre et al. 2017). To facilitate compari-son, we calculate and compare the abundances of the organic molecules relative to CH CHO.Figure 4 shows the comparison of the relative abundances of 7 organic molecules (excludingdeuterated species and C isotopologue) to those in IRAS 16293-2422B (Jørgensen et al.2016), B335 (Imai et al. 2016), and NGC 1333 IRAS4A2 (L´opez-Sepulcre et al. 2017). Inthis comparison, CH SH is excluded because of no reliable measurement of this molecule inother hot corinos to compare with. Interestingly, the relative abundances of most moleculeshere in HH 212 are similar to those in other hot corinos to within a factor of few, sug-gesting that the formation of these molecules in the disk atmosphere could be similar tothat in other hot corinos. Notice that the abundance of CH SH here is also consistent withthat predicted in the hot corino, which can be as high as 5 × − (Majumdar et al. 2016).Moreover, the abundances here are also similar to those seen in the Class I corino SVS13-A(Bianchi et al. 2019) and to those seen in Comet Lovejoy, which shows similar abundancesto IRAS 16293-2422B (Biver et al. 2015).As discussed in Lee et al. (2017c), the high degree of deuteration in methanol (with[CH DOH]/[CH OH] ∼ .
12) and the detection of doubly deuterated formaldehyde suggestthat the methanol and formaldehyde in the disk atmosphere are originally formed on icygrains and later desorbed (evaporated) into gas phase due to the heating possibly by low-velocity accretion shock near the centrifugal barrier or the radiation of the protostar. Thisheating can also be produced by an interaction with a wind/outflow from an inner region(Lee et al. 2018). The temperature of the disk atmosphere where the organic molecules aredetected is similar to the freeze-out temperature of water, which is ∼
150 K ( ¨Oberg et al.2011). Thus, the detection of deuterated water in the disk atmosphere (Codella et al. 2018)supports that the two organic molecules are evaporated in the disk atmosphere where it is 8 –warm enough to release the water. However, how the other organic molecules are formed isstill an open question and not necessarily on the grain surfaces.In case the organic molecules are formed on the icy grains, they can be formed on theicy grains in the disk, as suggested in the TW Hya protoplanetary disk for the methanol(Walsh et al. 2016), either in situ on the surface or first in the midplane and then brought tothe surface by turbulence (Furuya & Aikawa 2014). Like the hot corinos, it is also possiblethat the molecules already formed on the icy grains in the prestellar core, and then broughtin to the disk surface. CO is frozen onto to the grains at the temperature of ∼
20 K( ¨Oberg et al. 2011). One possible scenario (although to be proven) is that, in the regionssuch as prestellar cores and probably disk midplane where the temperature is below 20 K,CO-rich ices on dust grains can undergo addition reactions with H (and D), O, C, and Natoms accreted from the gas, producing a rich organic chemistry on the grains, as proposedin Charnley & Rodgers (2008).Having said that, some of the organic molecules could also be quickly formed in thegas phase using simpler species released from grains. This could be the case for formamide,which was found to be mainly formed in the gas phase in a young shocked region such asL1157-B1 (Codella et al. 2017). In addition, acetaldehyde can also be formed in gas phase,as suggested in Charnley (2004).These organic molecules are of great importance for forming even more complex organicmolecules such as amino acids and amino sugars, which are the building blocks of life. Ourobservations clearly show that they have formed in a disk or been brought in to a disk inthe earliest phase of star formation and may play a crucial role in producing the rich organicchemistry needed for life. It is also tempting to estimate the alcohol degree in the diskatmosphere. According to Codella et al. (2018), the column density of deuterated water is . × cm − . Assuming [H/D] ∼ [CH OH/CH DOH] ∼ . . × cm − . With the derived column density of ethanol in Table 2, andthe molecular mass of ethanol of 46 and water of 18, the alcohol degree by mass is estimatedto be & Methanol and acetaldehyde have been argued to trace a disk wind in HH 212 (Leurini et al.2016; Codella et al. 2018). Now in observations at higher resolution, they seem to tracemainly a ring of disk atmosphere near the centrifugal barrier, which could in principle beheated by a weak (accretion) shock produced by the rapid decrease of the infall velocity near 9 –the centrifugal barrier. However, since their maps also show small extensions extending outfrom the disk (see Figure 1), they may also trace a wind coming out from the disk surface.Nonetheless, since these extensions appear to be surrounding the SO outflow shell detectedfurther in Lee et al. (2018), they could also be the disk atmosphere pushed away by the SOoutflow shell. Detailed kinematic study of these extensions with the SO outflow shell areneeded to check these scenarios.
5. Conclusions
The nearly edge-on orientation of the disk in HH 212 provides the best view of thedisk atmosphere. Here we have detected 9 organic molecules in the disk atmosphere. Thesemolecules are characteristic of a hot corino and found here to be in the disk atmosphere,confirming that the corino here is a warm disk atmosphere. Adding 3 other organic moleculesfrom our previous study, we have detected 12 organic molecules, with 9 of them being complexorganic molecules, in the disk atmosphere within ∼
40 au of the central protostar. Thesemolecules seem to arise mainly from a ring of disk atmosphere near the centrifugal barrier.Some of them may also trace a wind coming out from the disk surface.The relative abundances of the organic molecules in the HH 212 disk atmosphere aresimilar to those in hot corinos around other low-mass protostars and even to those in CometLovely. It would be interesting to determine whether the hot corinos around other low-massprotostars are also located in their disk atmospheres or not, perhaps through higher reso-lution ALMA observations. In addition, the formation mechanism of the organic moleculescan also be similar to that in those corinos. The organic molecules can originally formed onicy grains, either in the disk or in the prestellar core and then brought in to the disk, andthen desorbed (evaporated) into the gas phase. They can also be quickly formed in the gasphase using simpler species ejected from the dust mantles.We thank the anonymous referee for constructive comments. This paper makes use ofthe following ALMA data: ADS/JAO.ALMA
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This preprint was prepared with the AAS L A TEX macros v5.2.
12 –Fig. 1.— Total emission line intensity maps of nine organic molecules integrated over velocitywithin ∼ − of the systemic velocity, plotted on top of the continuum map (color image)of the dusty disk. As discussed in the text, the maps of the molecular emissions are obtainedby stacking a number of lines (as indicated in the parenthesis) in different transitions with arange of upper energy levels for better detections. The asterisk marks the possible positionof the central protostar. The blue and red arrows indicate the axes of the blueshifted andredshifted jet components, respectively. The contours start at 3 σ with a step of 1 σ . The1- σ noise levels in panels (a) to (i) are 1.96, 1.64, 4.06, 3.91, 0.89, 1.10, 4.23, 1.19, and 1.72mJy beam − km s − . 13 –Fig. 2.— Position-frequency diagrams cut along the lower atmosphere of the disk, in orderto show the detections of various molecular lines. The contours start from 2 σ with a stepof 1 σ of 4 K. The frequency has been corrected for the systemic velocity of 1.7 km s − . Thevertical lines mark the rest frequencies of the lines in different transitions, with one color forone organic molecular species. The lines from simple molecules are marked with the blackvertical lines. The physical properties of the lines are listed in Table 1. 14 –Fig. 3.— Rotation diagrams for molecular lines of the six organic molecules includingCH OH, CH DOH, CH CHO, HCOOCH , CH CH OH, and CH DCN. The diagrams arederived from the line intensities in the lower disk atmosphere listed in Table 1. The errorbars show the uncertainty in our measurements, which are assumed to be 40% of the datavalues. The solid line is a linear fit to the data. In panels (c)-(f), the temperature is fixedat 150 ±
50 K for the fitting in order to derive the column densities from the weak lines. 15 –Fig. 4.— A comparison for the abundances of organic molecules in the HH 212 disk atmo-sphere with those in the hot corinos around other low-mass protostars. As shown in Table2, our measurements have an uncertainty of about 50%. 16 –Table 1. Line Properties from SplatalogueTransition Frequency log( A ul ) E u W a LineQNs (MHz) (s − ) (K) (K km s − ) Listt-HCOOH 15( 2,13)-14( 2,12) 346718.85 -3.331 144.457 147 CDMSH CCO 17( 1,16)-16( 1,15) 346600.45 -3.325 162.789 79 JPLCH OH 16(1)- - 15(2)- 345903.91 -4.044 332.653 234 JPLCH OH 18(-3) - 17(-4), E2 345919.26 -4.136 459.435 200 JPLCH OH 5(4)- - 6(3)- 346202.71 -4.662 115.162 129 m JPLCH OH 5(4)+ - 6(3)+ 346204.27 -4.662 115.162 129 m JPL CH OH 14( 1,13)- 14( 0,14) - + 347188.28 -3.360 254.251 96 † CDMSCH DOH 3(2,1) - 2(1,2), e1 345718.71 -4.373 39.434 51 JPLCH DOH 19(1,19) - 18(2,17), e1 345820.79 -4.531 418.032 B JPLCH DOH 22(4,19) - 22(3,19), e1 345850.48 -3.888 613.625 16 m JPLCH DOH 3(1,2) - 3(0,3), e1 346256.50 -3.699 29.488 152 JPLCH DOH 22(4,18) - 22(3,20), e1 346281.30 -3.874 613.621 16 m JPLCH DOH 21(4,18) - 21(3,18), e1 346419.06 -3.886 566.562 13 JPLCH DOH 21(4,17) - 21(3,19), e1 346783.70 -3.876 566.559 B JPLCH DOH 20(4,17) - 20(3,17), e1 346923.75 -3.886 521.634 85 JPLCH DOH 20(4,16) - 20(3,18), e1 347222.99 -3.878 521.632 B JPLCH DOH 19(4,16) - 19(3,16), e1 347371.16 -3.887 478.841 18 JPLCH DCN 20( 1,20)-19( 1,19) 345685.36 -2.444 179.628 12 JPLCH DCN 20( 0,20)-19( 0,19) 347043.43 -2.438 174.955 27 JPLCH DCN 20( 4,17)-19( 4,16) 347166.47 -2.455 261.246 13 m JPLCH DCN 20( 4,16)-19( 4,15) 347166.48 -2.455 261.246 13 m JPLCH DCN 20( 2,19)-19( 2,18) 347188.29 -2.441 196.562 B JPLCH DCN 20( 3,18)-19( 3,17) 347216.91 -2.447 223.527 B JPLCH DCN 20( 3,17)-19( 3,16) 347219.39 -2.447 223.527 B JPLCH DCN 20( 2,18)-19( 2,17) 347388.21 -2.441 196.615 24 JPL 17 –Table 1—ContinuedTransition Frequency log( A ul ) E u W a LineQNs (MHz) (s − ) (K) (K km s − ) ListCH CHO vt=1, 18(2,17) - 17(2,16), E 346065.34 -2.838 371.350 90 JPLCH CHO v=0, 18(11, 7) - 17(11, 6), E 346697.59 -3.028 430.543 T JPLCH CHO v=0, 18(10, 8) - 17(10, 7), E 346742.00 -2.986 383.352 16 JPLCH CHO v=0, 18(11, 8) - 17(11, 7), E 346754.52 -3.028 430.451 12 m JPLCH CHO v=0, 18(11, 8) - 17(11, 7), E 346755.92 -3.028 430.474 12 m JPLCH CHO v=0, 18(11, 7) - 17(11, 6), E 346755.92 -3.028 430.474 12 m JPLCH CHO v=0, 18(10, 9) - 17(10, 8), E 346763.91 -2.985 383.297 30 JPLCH CHO v=0, 18(9, 9) - 17(9, 8), E 346787.03 -2.950 340.618 B JPLCH CHO v=0, 18(10, 8) - 17(10, 7), E 346805.46 -2.985 383.258 S JPLCH CHO v=0, 18(10, 9) - 17(10, 8), E 346805.46 -2.985 383.258 S JPLCH CHO v=0, 18(9, 10) - 17(9, 9), E 346807.99 -2.950 340.612 S JPLCH CHO v=0, 18(8,10) - 17(8, 9), E 346839.03 -2.920 302.406 23 JPLCH CHO v=0, 18(9,10) - 17(9, 9), E 346849.06 -2.950 340.535 30 m JPLCH CHO v=0, 18(9, 9) - 17(9, 8), E 346849.06 -2.950 340.535 30 m JPLCH CHO v=0, 18(8,11) - 17(8,10), E 346892.18 -2.920 302.343 S JPLCH CHO v=0, 18(8,11) - 17(8,10), E 346893.81 -2.920 302.316 S JPLCH CHO v=0, 18(8,10) - 17(8, 9), E 346893.81 -2.920 302.316 S JPLCH CHO v=0, 18(7,11) - 17(7,10), E 346934.22 -2.896 268.661 B JPLCH CHO v=0, 18(7,12) - 17(7,11), E 346957.55 -2.896 268.606 B JPLCH CHO v=0, 18(7,11) - 17(7,10), E 346957.55 -2.896 268.606 B JPLCH CHO v=0, 18(7,12) - 17(7,11), E 346995.53 -2.896 268.572 24 JPLCH CHO v=0, 18(6,13) - 17(6,12), E 347071.54 -2.875 239.399 33 m JPLCH CHO v=0, 18(6,12) - 17(6,11), E 347071.68 -2.875 239.399 33 m JPLCH CHO v=0, 18(6,12) - 17(6,11), E 347090.40 -2.875 239.397 33 m JPLCH CHO v=0, 18(6,13) - 17(6,12), E 347132.68 -2.875 239.321 33 JPLCH CHO vt=1, 18(4,14) - 17(4,13), E 347182.41 -2.845 400.378 B JPLCH CHO vt=1, 18(5,13) - 17(5,12), E 347216.79 -2.859 420.440 B JPLCH CHO vt=1, 18(5,14) - 17(5,13), E 347251.82 -2.858 419.672 T JPLCH CHO v=0, 18(5,14) - 17(5,13), E 347288.26 -2.858 214.697 53 JPLCH CHO v=0, 18(5,13) - 17(5,12), E 347294.87 -2.858 214.698 115 JPL 18 –Table 1—ContinuedTransition Frequency log( A ul ) E u W a LineQNs (MHz) (s − ) (K) (K km s − ) ListCH CHO v=0, 18(5,13) - 17(5,12), E 347345.71 -2.858 214.640 B JPLCH CHO v=0, 18(5,14) - 17(5,13), E 347349.27 -2.858 214.611 B JPLHCOOCH v=0 28(12,16)-27(12,15) E 345974.66 -3.287 335.433 11 m JPLHCOOCH v=0 28(12,17)-27(12,16) A 345985.38 -3.287 335.434 11 m JPLHCOOCH v=0 28(12,16)-27(12,15) A 345985.38 -3.287 335.434 11 m JPLHCOOCH v=0 28(12,17)-27(12,16) E 346001.61 -3.287 335.430 11 m JPLHCOOCH v=0 28(11,17)-27(11,16) E 346659.86 -3.269 320.394 B JPLHCOOCH v=0 28(11,18)-27(11,17) A 346674.98 -3.269 320.395 22 m JPLHCOOCH v=0 28(11,17)-27(11,16) A 346675.64 -3.269 320.396 22 m JPLHCOOCH v=0 28(11,18)-27(11,17) E 346687.46 -3.269 320.391 22 m JPLHCOOCH v=0 27( 5,22)-26( 5,21) E 347478.25 -3.211 247.252 B JPLHCOOCH v=0 27( 5,22)-26( 5,21) A 347493.96 -3.211 247.256 30 JPLg-CH CH OH 20( 3,18)-19( 3,17) vt=0-0 345648.57 -3.436 242.486 27 JPLg-CH CH OH 20(10,10)-19(10, 9) vt=1-1 346085.56 -3.549 358.843 T JPLg-CH CH OH 20(10,11)-19(10,10) vt=1-1 346085.56 -3.549 358.843 T JPLg-CH CH OH 20( 9,12)-19( 9,11) vt=1-1 346183.19 -3.522 335.566 T JPLg-CH CH OH 20( 9,11)-19( 9,10) vt=1-1 346183.19 -3.522 335.566 T JPLg-CH CH OH 20(11,10)-19(11, 9) vt=0-0 346383.64 -3.580 379.110 5 m JPLg-CH CH OH 20(11, 9)-19(11, 8) vt=0-0 346383.64 -3.580 379.110 5 m JPLg-CH CH OH 20(10,11)-19(10,10) vt=0-0 346424.58 -3.548 353.452 T JPLg-CH CH OH 20(10,10)-19(10, 9) vt=0-0 346424.58 -3.548 353.452 T JPLg-CH CH OH 20( 9,12)-19( 9,11) vt=0-0 346505.34 -3.521 330.252 T JPLg-CH CH OH 20( 9,11)-19( 9,10) vt=0-0 346505.34 -3.521 330.252 T JPLg-CH CH OH 20( 7,14)-19( 7,13) vt=1-1 346565.08 -3.479 296.366 22 m JPLg-CH CH OH 20( 7,13)-19( 7,12) vt=1-1 346565.39 -3.479 296.366 22 m JPLg-CH CH OH 20( 8,13)-19( 8,12) vt=0-0 346620.32 -3.498 309.511 T JPLg-CH CH OH 20( 8,12)-19( 8,11) vt=0-0 346620.33 -3.498 309.511 T JPLg-CH CH OH 20( 7,14)-19( 7,13) vt=0-0 346816.58 -3.478 291.241 12 m JPL 19 –Table 1—ContinuedTransition Frequency log( A ul ) E u W a LineQNs (MHz) (s − ) (K) (K km s − ) Listg-CH CH OH 20( 7,13)-19( 7,12) vt=0-0 346816.92 -3.478 291.241 12 m JPLt-CH CH OH 21( 0,21)-20( 1,20) 346962.59 -3.616 185.843 B JPLg-CH CH OH 20( 6,15)-19( 6,14) vt=0-0 347147.20 -3.461 275.456 13 JPLg-CH CH OH 20( 6,14)-19( 6,13) vt=0-0 347157.99 -3.461 275.457 9 JPLt-CH CH OH 14( 3,12)-13( 2,11) 347350.92 -3.757 99.660 B JPLt-CH CH OH 21( 1,21)-20( 0,20) 347445.52 -3.573 185.853 41 JPLg-CH CH OH 20( 5,16)-19( 5,15) vt=1-1 347473.56 -3.447 267.087 B JPL a : Integrated line intensities (see text for the definition) measured from the lower disk atmo-sphere for the reasonably isolated lines detected with more than 3 σ . They are the mean valuesaveraging over a rectangular region (with a size of 0 ′′ . 2 × ′′ . 05 covering most of the emission) cen-tered at the lower atmosphere. In this column, the line intensities commented with “m” are themean values obtained by averaging over 2 or more lines with similar E u and log A ul for bettermeasurements. T: the lines are tentatively detected with about 2 σ detection and a part of thelinear PF structure. S: the lines are blended with the line(s) of the same molecule at differenttransitions. B: the lines are blended with the line(s) of other molecules. The line intensities hereare assumed to have an uncertainty of 40%. † : CH OH intensity after removing the small contribution from CH DCN (see text). 20 –Table 2. Column Densities and Abundances in the Lower Disk AtmosphereSpecies Excitation Temperature Column Density Abundance † (K) (cm − )CH OH... a ±
45 (6 . ± . × (1 . ± . × − CH OH... b ± c (1 . ± . × (3 . ± . × − CH DOH 148 ±
19 (1 . ± . × (4 . ± . × − CH OH d ± c (2 . ± . × (6 . ± . × − H CCO 150 ± c (3 . ± . × (1 . ± . × − t-HCOOH 150 ± c (2 . ± . × (5 . ± . × − CH DCN 150 ± c (3 . ± . × (9 . ± . × − CH CHO 150 ± c (1 . ± . × (3 . ± . × − HCOOCH ± c (3 . ± . × (8 . ± . × − CH CH OH 150 ± c (2 . ± . × (7 . ± . × − CH SH e ± c (9 . ± . × (2 . ± . × − NH CHO e ± c (1 . ± . × (4 . ± . × − D CO e ± c (3 . ± . × (7 . ± . × − † : Derived by dividing the column densities of the molecules by the H columndensity in the disk atmosphere, which is ∼ × cm − (see text). a : Derived from the rotation diagram in Figure 3a. The value derived thisway is considered as the lower limit of the CH OH column density (see text). b : Obtained by multiplying the CH OH column density by 50, which is the C/ C ratio obtained in the Orion complex (Kahane et al. 2018). The valuederived this way is adopted for the CH OH column density (see text). c : Assuming an excitation temperature (rotational temperature) of 150 ± OH and CH DOH lines. d : Estimated after subtracting the small contribution from CH DCN (seetext). e : Adopted from Lee et al. (2017c), updated for an excitation temperature of 21 –150 ±±