Organic complexity in protostellar disk candidates
Jennifer B. Bergner, Rafael Martin-Domenech, Karin I. Oberg, Jes K. Jorgensen, Elizabeth Artur de la Villarmois, Christian Brinch
DDraft version July 19, 2019
Typeset using L A TEX twocolumn style in AASTeX62
Organic complexity in protostellar disk candidates
Jennifer B. Bergner, Rafael Mart´ın-Dom´enech, Karin I. ¨Oberg, Jes K. Jørgensen, Elizabeth Artur de la Villarmois, and Christian Brinch Harvard University Department of Chemistry and Chemical Biology, Cambridge, MA 02138, USA Harvard-Smithsonian Center for Astrophysics, Cambridge, MA 02138, USA Niels Bohr Institute & Centre for Star and Planet Formation, University of Copenhagen, Øster Voldgade 5-7, 1350 København, Denmark Research Group for Genomic Epidemiology, National Food Institute, Technical University of Denmark, 2800 Kgs. Lyngby, Denmark
ABSTRACTWe present ALMA observations of organic molecules towards five low-mass Class 0/I protostellar diskcandidates in the Serpens cluster. Three sources (Ser-emb 1, Ser-emb 8, and Ser-emb 17) presentemission of CH OH as well as CH OCH , CH OCHO, and CH CO, while NH CHO is detected in justSer-emb 8 and Ser-emb 17. Detecting hot corino-type chemistry in three of five sources represents ahigh occurrence rate given the relative sparsity of these sources in the literature, and this suggests apossible link between protostellar disk formation and hot corino formation. For sources with CH OHdetections, we derive column densities of 10 –10 cm − and rotational temperatures of ∼ OH-normalized column density ratios of large, oxygen-bearing COMs in the Serpens sourcesand other hot corinos span two orders of magnitude, demonstrating a high degree of chemical diversityat the hot corino stage. Resolved observations of a larger sample of objects are needed to understandthe origins of chemical diversity in hot corinos, and the relationship between different protostellarstructural elements on disk-forming scales.
Keywords: astrochemistry – complex organic molecules – hot corinos – low-mass protostars – inter-stellar medium INTRODUCTIONPlanet formation takes place in the gas- and dust-richdisks orbiting young stars. The chemical inventories inthese protoplanetary disks therefore influence the com-positions of nascent planets. It is of particular interestto origins of life studies to understand the chemistryof complex (6+ atom, hydrogen-rich) organic molecules(COMs) in planet-forming regions, since these speciesare considered to be precursors for prebiotic chemistry(Herbst & van Dishoeck 2009; Jørgensen et al. 2012).The process of star (and eventually planet) forma-tion begins in the dense cores of molecular clouds. Theyoungest protostars, termed Class 0, are still deeply em-bedded in their natal envelope. Class I protostars areundergoing envelope infall and accretion onto a circum-stellar disk, and Class II protostars have cleared their en-velope and host Keplerian disks that feed accretion ontothe star (Lada 1987; Andre et al. 1993; Dunham et al.2014). Planet formation was historically thought to take [email protected] place during the Class II stage. In recent years, thehigh sensitivity and spatial resolution of ALMA has en-abled the characterization of complex organic moleculeemission in Class II disks, enhancing our understandingof organic chemistry at this stage ( ¨Oberg et al. 2015;Walsh et al. 2016; Bergner et al. 2018; Loomis et al.2018; Favre et al. 2018). However, there is growing ev-idence suggesting that planet formation begins at ear-lier evolutionary stages (Harsono et al. 2018). Notably,high-resolution (sub-)mm continuum observations arerevealing that dust sub-structure appears ubiquitous inClass II disks (Andrews et al. 2018). One compelling ex-planation for this sub-structure is interactions of large(Neptune- to Jupiter-mass) planets with the disk (Zhanget al. 2018), which would require that the planet forma-tion process begins in younger (Class 0/I) disks. Thisscenario is supported by observations of sub-structure inthe embedded disk HL Tau (ALMA Partnership et al.2015), implying grain growth or even planet formationat an early evolutionary stage.Some deeply embedded protostars are seen to hosta rich and warm gas-phase organic chemistry on small a r X i v : . [ a s t r o - ph . S R ] J u l ( ∼
100 AU) scales. These sources, termed hot corinos,form when temperatures around the protostar exceedthe water ice sublimation point, and ice mantles areliberated into the gas phase (Herbst & van Dishoeck2009). At present, the physical nature of hot corinosis unknown because they are generally unresolved ormarginally resolved in observations. Several structuralelements, including protostellar disks, centrifugal barri-ers, and outflows, occur on similar spatial scales withinthe protostellar inner envelope (Sakai et al. 2014; Leeet al. 2014; Harsono et al. 2014; Yen et al. 2015); with-out high-sensitivity and high-resolution observations itis difficult to disentangle if these elements are relatedor distinct in their physics and chemistry, and which, ifany, are sources of hot corino chemistry.Since protostellar disks are likely the sites whereplanet formation begins, it is of great interest to under-stand whether the organic complexity detected in hotcorinos is related to the material incorporated into thedisk. To date there is no conclusive tie between hot cori-nos and protostellar disks. Some sources with hot cori-nos appear to show velocity structure in their molecularline emission consistent with the presence of a rotatingdisk (Choi et al. 2010; Lee et al. 2014; Codella et al. 2014;Oya et al. 2016, 2017). Still, other hot corino sourcesshow no clear signatures of a rotating disk (Maury et al.2014; Imai et al. 2016; Jacobsen et al. 2018). More ob-servations are needed to understand the chemical andphysical relationships between hot corinos and proto-stellar disks.In this work, we present ALMA observations of fiveprotostellar disk candidates in the Serpens cluster. Forone source, Ser-emb 1, the chemical and physical struc-ture was studied in detail in Martin-Domenech et al.(2019). Here, we aim to characterize the complex or-ganic chemistry in all five sources in order to furtherour understanding of the chemical evolution of proto-stars on small scales. Section 2 describes our sourceselection, line targets, and ALMA observations. Section3 presents the observed morphologies for each source aswell as organic molecule detections. Additionally, we de-rive column densities for CH OH using the populationdiagram method, and for other species by assuming arotational temperature. In Section 4 we discuss the fre-quency of hot corino detection in our sample of diskcandidates. We also compare the CH OH-normalizedcolumn density ratios of organics in the Serpens sourceswith measurements in other hot corinos, low-mass pro-tostellar envelopes, and Solar System comets. OBSERVATIONS2.1.
Observational details
The source sample consists of five low-mass protostarsin the Serpens cluster. Each target source shows non-zero flux at > λ uv distances measured for the 230GHz continuum, which may be due to the presence of aprotostellar disk (Enoch et al. 2011). Source propertiesand candidate disk mass estimates are listed in Table1. For all sources we assume a distance of 436 ± ∼
18 total minutes in two execu-tion blocks, with 42 antennae and baselines spanning 15– 560 m. J1751+0939 was used for bandpass and fluxcalibration, Titan for flux calibration, and J1830-0619for phase calibration. For the higher-frequency setup,each source was observed for ∼
21 total minutes in threeexecution blocks, with 41 antennae and baselines span-ning 15 – 640 m. J1751+0939 was used for bandpasscalibration, Titan for flux calibration, and J1830+0619for phase calibration. This project makes use of datafrom three spectral windows centered on 219.57 GHz,231.49 GHz, and 244.88 GHz. The 219 GHz windowhas a bandwidth of 117 MHz and a channel width of122 kHz ( ∼ ∼ Analysis
Initial pipeline calibration of the ALMA data wasperformed by ALMA/NAASC staff with CASA ver-sions 4.5.3 (lower-frequency setup) and 4.7.0 (higher-frequency setup). An additional 1–2 rounds of phaseself-calibration were performed for each spectral win-dow using the line-free continuum, followed by contin-uum subtraction.In addition to the C O 2–1 line in the 219 GHz spec-tral window, we searched for spectral lines of organicmolecules commonly detected in protostars within the231 GHz and 244 GHz spectral windows. Spectral lineparameters are taken from the JPL (PICKETT et al.1998) and CDMS (M¨uller et al. 2001, 2005) catalogs.For CH OH we searched for lines satisfying A ul > − s − and upper energies <
700 K, and for other COMs weinitially searched for lines with A ul > − s − and up-per energies <
300 K. When generating synthetic spectra(Section 3) we include all lines with A ul > − s − andupper energies <
700 K for all COMs.Image cubes were generated using the tclean taskin CASA version 5.4.1, using Briggs weighting with a
Source R.A. Decl. Class T bol L bol M env Est. M disk a (J2000) (J2000) (K) ( L (cid:12) ) ( M (cid:12) ) ( M (cid:12) )Ser-emb 1 18:29:09.1 0:31:30.9 0 39 [2] 4.1 [0.3] 3.1 [0.05] 0.28Ser-emb 7 18:28:54.1 0:29:30.0 0 58 [13] 7.9 [0.3] 4.3 [0.4] 0.15Ser-emb 8 18:29:48.1 1:16:43.7 0 58 [16] 5.4 b [6.2] 9.4 [0.3] 0.25Ser-emb 15 18:29:54.3 0:36:00.8 I 101 [43] 0.4 [0.6] 1.3 [0.1] 0.15Ser-emb 17 18:29:06.2 0:30:43.1 I 117 [21] 3.8 [3.3] 3.6 [0.4] 0.15 Table 1.
Source properties, taken from Enoch et al. (2009) and Enoch et al. (2011). a Uncertainties are at least ± b Likelyan underestimate.
Spectral window Beam dim. Beam Chan. width Chan. rmscenter frequency (”) PA ( o ) (km s − ) (mJy beam − )Ser-emb 1219 GHz 0.63 × × × × × × × × × × × × × × × Table 2.
Line observation details robust parameter of 0.5. C O was imaged with a ve-locity resolution of 0.25 km s − , and all organic lineswere imaged at the native spectral resolution. Cleanmasks were drawn by hand for detected lines, andfor non-detected lines the 10 σ continuum contour wasused. In addition to hand-masking individual lines, wealso cleaned the full data cubes using the automask-ing task auto-multithresh in tclean . We usedthe standard automasking parameters for short baseline12m line data (sidelobethreshold = 2.0, noisethreshold= 4.5, minbeamfrac = 0.3, negativethreshold = 15.0,lownoisethreshold = 1.5). We have verified the au-tomasking results by comparing the CH OH spectra ex-tracted from the auto-masked and hand-masked images,and find they are consistent. Table 2 shows representa- tive line observation details for each spectral windowused in this work. RESULTS3.1.
Source overview
Figure 1 shows the 1.3 mm dust continuum emission,as well as C O 2–1 and CH OH 5 , – 4 , line emission(see Table 3 for spectral line parameters). All sourcesshow a compact central dust component, while Ser-emb1, Ser-emb 7, and Ser-emb 8 additionally have a moreextended component. The C O morphologies show asimilar pattern, with more extended gas emission in Ser-emb 1, Ser-emb 7, and Ser-emb 8. The sources with ex-tended continuum and C O emission are also the leastevolved, and it is not surprising that we detect a greaterenvelope contribution in these sources.
Molecule Transition Frequency log(A ul ) g u E u Q ( T ) Refs.(GHz) (s − ) (K) 150 K 250 KC O a OH b , – 4 , A 243.916 -4.22 44 49.7 9750 26335 210 , – 9 , A 232.419 -4.73 84 165.410 , – 11 , E 232.946 -4.67 84 190.49 − , – 8 − , E, v t = 1 244.338 -4.39 76 395.618 , – 17 , A 232.783 -4.66 148 446.518 , – 19 , & 243.397 -4.70 296 590.318 , – 19 , A c , – 22 , A 244.330 -4.09 180 636.823 , – 23 , A 243.413 -4.09 188 690.1CH OCH b , – 12 , AA & EE c , – 23 , AA & EE c OCHO a , – 18 , A 233.227 -3.74 78 123.2 59073 [1.20] d d , – 19 , E 244.580 -3.68 82 135.0NH CHO b , – 10 , , – 11 , CO b , – 11 , Table 3.
Spectral line data. a From the JPL catalogue. b From the CDMS catalogue. c Blended transitions. d Numbers in bracketsrepresent the vibrational correction factor to the v T = 0,1 partition function at each temperature. References: [1] Winnewisseret al. (1985) [2] Xu et al. (2008) [3] Endres et al. (2009) [4] Groner et al. (1998) [5] Ilyushin et al. (2009) [6] Plummer et al.(1984) [7] Oesterling et al. (1999) [8] Maeda et al. (2008) [9] Favre et al. (2014) [10] Kryvda et al. (2009) [11] Brown et al. (1990)[12] Guarnieri & Huckaufa (2003) The CH OH 5 , – 4 , line ( E u = 50 K) is detectedabove a 5 σ level only towards Ser-emb 1, Ser-emb 8, andSer-emb 17. Interestingly, the morphology of this tran-sition is different in each source: while all three havea compact central component, Ser-emb 8 also shows ajet-like feature extending out from the central protostar,and Ser-emb 17 shows more amorphous extended emis-sion. This extended emission is only seen for the 50 KCH OH line, and all higher- J lines show emission onlyfrom the compact central component (Figure 2).Our source sample spans a relatively wide range ofbolometric temperatures, bolometric luminosities, enve-lope masses, and estimated disk masses, however thereis no clear relationship between these physical proper-ties and the detection of CH OH in a source (Table 1).For each property, the value of one or both sources withCH OH non-detections falls within the range of valuesfor sources with detections. We note that there arehigh uncertainties on the source luminosities for Ser-emb 8 and 17 and on the candidate disk masses for allsources, so better constraints on these properties mayreveal some trend with CH OH emission. There is alsono clear relationship between the inferred size of the diskand the presence of CH OH emission: in the observa- tions of Enoch et al. (2011), the mm dust emission iscompact and unresolved at spatial scales of ∼
170 AU inSer-emb 7, Ser-emb 15, and Ser-emb 17, but partiallyresolved in Ser-emb 1 and Ser-emb 8. The protostel-lar neighborhood also does not appear to account forCH OH emission: none of the sources show evidence forbinarity in the Enoch et al. (2011) survey, and Ser-emb1 is the only source that is isolated more than ∼
30” fromone or two neighboring protostars. Moreover, Ser-emb 1,7, and 17 are located near to one another in the Serpenscluster B, while Ser-emb 15 is farther away in cluster Band Ser-emb 8 is in the Main cluster (Enoch et al. 2011).The presence of CH OH emission might be influencedby a combination of these and other factors, however,and better characterization of the physical properties ofeach source as well as their local physical environmentis needed. 3.2.
Organic molecule detections
The sources with CH OH detections (Ser-emb 1, Ser-emb 8, and Ser-emb 17) also showed emission from otherorganic molecules. We consider a molecule to be de-tected based on the following criteria. At least one linemust be observed above a 5 σ level in the moment zeromap, and the peak of the extracted spectrum must sim- Figure 1.
Source overview showing the 1.3 mm dust continuum emission (top), C O 2–1 line emission (middle), and CH OH5 , – 4 , line emission (bottom). Continuum contours are drawn at 5, 30, 100, 400 × rms, and line contours are drawn at 5,10, 30 × rms. Color scales are normalized to each individual image, and emission below a 2 × rms threshold is not shown. Thesynthesized beam is shown in the bottom left of each panel. Velocity ranges and rms values for each panel can be found inAppendix A. ilarly be > σ . For each line, we also check for pos-sible neighbors and do not further consider any linesthat are potentially blended with other strong emitters.Lastly, we search the spectra for other lines of each de-tected molecule to ensure that any strong transitionsare not missing. Based on these criteria, CH OCH ,CH OCHO, and CH CO were detected towards Ser-emb1, Ser-emb 8, and Ser-emb 17, while NH CHO was de-tected just towards Ser-emb 8 and Ser-emb 17. We notethat the CH CO detections are based on a single line,however because it is unblended and there are no com-peting line assignments we are confident in the detec-tion. Figure 2 shows moment zero maps for the bright-est transition of each of these organic molecules, alongwith a high-energy ( E u = 165 K) CH OH line. In allcases, the emission is unresolved or marginally resolved.In several sources we also see evidence for emis-sion from HCOOH, CH CHO, C H OH, and possiblyCH OHCHO. However, there are too few lines that areboth unblended and have sufficient SNR to claim detec-tions. Higher resolution data and an increased numberof line targets per molecule are needed to characterizethe emission from these additional species. Table 3 lists the spectral line information for the tran-sitions that we use for measuring column densities. Ad-ditional CH OCHO lines are covered in the frequencyrange of our spectra, however due to line overlap we con-sider only these non-blended lines for analysis. We notethat the CH OCHO partition function from JPL in-cludes the v T = 0,1 states, however at hot corino temper-atures (200–300 K) other vibrational states will becomepopulated. We therefore correct the partition functionwith the vibrational correction factors provided by Favreet al. (2014) via the CMDS catalogue; Table 3 liststhe non-corrected partition function values and vibra-tional correction factors at each temperature. For othermolecules, the catalogs already include sufficient vibra-tionally excited states in the partition function or thevibrational contributions are small at hot corino tem-peratures; the exception is NH CHO for which the vi-brational contribution is not yet available.3.3.
Column densities CH OH For each source, we extract spectra from a single pixelcorresponding to the location of the continuum peak.
Figure 2.
Moment zero maps of organic molecule lines in Ser-emb 1, Ser-emb 8, and Ser-emb 17. Contours are drawn at 5, 10,20, 30 × rms. Color scales are normalized to each individual image, and emission below a 2 × rms threshold is not shown. Thesynthesized beam is shown in the bottom left of each panel. Velocity ranges and rms values for each panel can be found inAppendix A. Velocity-integrated intensities are measured by fitting aGaussian to each spectral feature. For isolated lines wealso include an offset term in the fit to allow slight varia-tions in the baseline. For each source, the line width de-rived for the CH OH 5 , – 4 , transition is adopted as afixed parameter for all other transitions to ensure goodfits for weaker or slightly blended lines. All CH OHspectral lines and Gaussian fits are shown in Figure 3.Integrated intensities can be found in Table 5. Subse-quent uncertainties are propagated based on Gaussianfit uncertainties added in quadrature with a 10% cali-bration uncertainty.Our observations cover 8 CH OH transitions spanninga range of upper state energies from 50 – 690 K. We usethe population diagram method to derive column den-sities and rotational temperatures, taking into accountoptical depths for each line. This treatment is adaptedfrom Goldsmith & Langer (1999) (see also e.g. Taquetet al. (2015); Loomis et al. (2018)) and assumes localthermodynamic equilibrium (LTE) conditions. The total column density N T and rotational temper-ature T R are related to the upper level population N u by: N u g u = N T Q ( T R ) e − E u /T R (1)Here, g u is the upper level degeneracy, Q is the molecularpartition function, and E u is the upper state energy inK. The observed upper level population N u,obs is foundfrom the velocity-integrated surface brightness (cid:82) I ν dv by: N u,obs = 4 π (cid:82) I ν dvA ul hc , (2)where A ul is the transition Einstein coefficient. If thesource does not fill the beam and if the observed lines areoptically thick, we find the true upper level populationfrom N u,obs by: N u = N u,obs C τ Ω a Ω s . (3) Figure 3. CH OH spectral lines in sources where CH OHis detected. Blue lines show the spectra extracted from thecontinuum peak pixel, and shaded regions represent the rms.Red lines show Gaussian fits to the data; a dotted line in-dicates that the feature is not significant above a 3 σ level. C τ is the optical depth correction factor and Ω a Ω s is thebeam dilution factor, where Ω s and Ω a are the sourceand beam solid angles, respectively. C τ is found from: C τ = τ − e − τ , (4)and the optical depth τ is determined by: τ = c A ul N u πν ∆ v ( e hν/kT R − . (5)Here, c is the speed of light, A ul is the Einstein coeffi-cient (lines with a higher A ul are more likely to be op-tically thick), ν is the line frequency, ∆ v is the line fullwidth half-maximum, and k is the Boltzmann constant.For ∆ v we use the FWHM of the CH OH 5 , –4 , linein each source (3.3, 4.0, and 4.9 km/s in Ser-emb 1, 8,and 17, respectively).The degree of beam dilution in our observations is un-certain given that the sources are unresolved or barelyresolved. We therefore solve for column densities usingtwo bounding cases to represent the range of likely val-ues. The maximum source size, corresponding to theminimum beam dilution, is found using the deconvolvedsource sizes (or upper limits) derived in CASA for theCH OH 165 K or 190 K lines (since the 50 K line tracescooler material in addition to the hot corino). This givesbeam dilution factors of 8.7, 15.1, and 9.3 for Ser-emb 1,8, and 17, respectively. The minimum source size, cor-responding to the maximum beam dilution, is found byusing the power-law temperature profile for protostellarenvelopes in Chandler & Richer (2000) to estimate theradius beyond which the temperature falls below 100 K: T ( r ) = 60 (cid:32) r × m (cid:33) − q (cid:32) L bol L (cid:12) (cid:33) q/ K , (6)where q = 2/(4 + β ). Assuming the source luminositiesin Table 1 and β = 1.5, we find maximum beam dilutionfactors of 27.7, 21.0, and 29.7 for Ser-emb 1, 8, and 17,respectively.We fit the observed upper level populations (Equa-tion 2) by generating synthetic upper level populationsfor all detected CH OH transitions, with N T and T R as free parameters. Combining Equations 1 and 3, thesynthetic N u,obs are found from: N u,obs g u = N T Q ( T R ) e − E u /T R C τ Ω s Ω a . (7)We use the affine-invariant MCMC package emcee (Foreman-Mackey et al. 2013) to sample the poste-rior distributions. Additional details on the MCMCfitting can be found in Appendix B. All spectral lineparameters used for this analysis are listed in Table 3.The resulting population diagrams are shown in Fig-ure 4. Table 4 lists the values derived from this fit-ting, along with 1 σ uncertainties. Across the sample,we derive column densities on the order of 10 cm − and rotational temperatures of 200–250 K. Even in themaximum beam dilution cases, in all sources the opticaldepth of the 50 K line is < < OH lines are opti-cally thick, CH OH or CH
OH lines are often usedto derive CH OH column densities (Taquet et al. 2015;Jørgensen et al. 2016). Our observations cover the CH OH line at 231.818 GHz, however we do not detectthis line in any source. Assuming the maximum beamdilution factor, and a C/ C ratio of 70, we obtain3 σ upper limits on the CH OH column density of a few × cm − for our sources. Thus, these non-detectionsare consistent with our derived column densities but donot provide any further constraints.3.3.2. Other organics
For all additional organics, detected lines are fittedwith Gaussians as described previously for CH OH.Again, we adopt the line width of the CH OH 5 , – 4 , transition for all other molecular lines observed towardsa given source. Spectral line fits for all detected linescan be found in Appendix C, and integrated intensitiesare listed in Table 5.We estimate column densities by solving equations 2and 7 assuming optically thin emission and an adoptedrotational temperature. We expect most COMs to sharea roughly similar rotational temperature as CH OH,though some species tend to emit at cooler or warmertemperatures (Jørgensen et al. 2016). We therefore cal-culate column densities assuming the CH OH rotationaltemperatures derived in the previous section ( T M ), aswell as T M ±
75 K. As for CH OH, we calculate therange of column density values assuming minimum andmaximum beam dilution factors. The results are listedin Table 6.Figure 5 shows the complete spectrum extracted fromSer-emb 17 along with synthetic spectra for each or-ganic molecule, calculated with the derived column den-sities and assuming the CH OH rotational temperature.Spectra for all additional sources can be found in Ap-pendix D. Synthetic spectra are calculated for all lineswith A ul > − s − and upper energies <
700 K. The op-tical depths calculated for all COM lines are low ( < DISCUSSION 4.1.
Frequency of hot corino chemistry aroundprotostellar disk candidates
In our survey of five protostellar disk candidates,we see evidence for warm, organic-rich, hot-corino-likeemission from three of five targeted sources. This is ahigh occurrence rate given that only 9 hot corinos havebeen previously identified in the literature: IRAS 16293-2422 (Cazaux et al. 2003), NGC1333 IRAS 4A (Bot-tinelli et al. 2004), NGC1333 IRAS 2 (Jørgensen et al.2005), NGC1333 IRAS 4B (Bottinelli et al. 2007), HH212 (Codella et al. 2016), B335 (Imai et al. 2016), L483(Oya et al. 2017), B1b (Lefloch et al. 2018), and SVS13-A (Lefloch et al. 2018). Our sources were selecteddue to their classification as protostellar disk candidates,suggesting that hot corinos and protostellar disks maybe evolutionarily or structurally linked. We emphasize,however, that higher-resolution kinematic studies areneeded to verify if these sources indeed host disks. In-terestingly, Ser-emb 17 is just the second Class I sourcewith a hot corino detection (Codella et al. 2016; Leflochet al. 2018); while Class 0 sources represent the majorityof hot corino detections, they are clearly not limited tothis early evolutionary stage.Given the quality of these observations, we cannot atpresent put strong constraints on whether the organicmolecule emission we observe originates from a disk.However, we do see interesting hints of structure in theline shapes. Examples of these features are shown in Fig-ure 6. Many organic lines in Ser-emb 17 appear to have adouble-peaked line profile, which can be a signature of arotating disk (Beckwith & Sargent 1993) (though couldalso be an opacity effect). Some lines in Ser-emb 8 andSer-emb 17 also appear to show an inverse P-cygni pro-file, with an absorption feature occurring at red-shiftedwavelengths from the emission peak, indicative of in-fall (Di Francesco et al. 2001; Kristensen et al. 2012).NH CHO line profiles in Ser-emb 8 and 17 also showa slight broadening at red-shifted velocities, suggestingthe presence of a jet/outflow. Strong NH CHO emis-sion in an outflow has been previously reported towardsL1157 (Codella et al. 2017). Clearly, there is evidence foradditional structure in these molecular lines that is notresolved in our observations. Disentangling whether thehot corino, infall, disk, and jet/outflow components aredistinct or overlapping in their chemistry and physics re-quires higher-resolution follow-up observations, and willbe key to understanding the connection between proto-stellar chemistry and disk chemistry.4.2. CH OH column densities
In other hot corinos that have been studied at highangular resolution (NGC1333 IRAS 2A and 4A, IRAS
Min. Ω a Ω s Max. Ω a Ω s N T (10 cm − ) T R (K) N T (10 cm − ) T R (K)Ser-emb 1 1.4 +0 . − . +32 − +0 . − . +32 − Ser-emb 8 5.7 +0 . − . +8 − +0 . − . +8 − Ser-emb 17 2.8 +0 . − . +10 − +0 . − . +9 − Table 4. CH OH column densities and rotational temperatures and 1 σ uncertainties derived from population diagram fitting,assuming minimum and maximum beam dilution factors as described in the text. Figure 4. CH OH population diagrams for Ser-emb 1, Ser-emb 8, and Ser-emb 17. Data and uncertainties are shown in black,and draws from the fit posteriors are shown in blue.
Molecule Line Ser-emb 1 Ser-emb 8 Ser-emb 17Int. intensity Beam dim. Int. intensity Beam dim. Int. intensity Beam dim.(mJy beam − (”) (mJy beam − (”) (mJy beam − (”)km s − ) km s − ) km s − )CH OH 5 , ±
15 0.54 × ±
34 0.55 × ±
29 0.54 × , ± × ±
21 0.57 × ±
16 0.57 × , ± × ±
19 0.57 × ±
15 0.57 × , ± × ±
14 0.57 × ±
12 0.57 × , ±
12 0.54 × ±
16 0.54 × ±
14 0.54 × − , - - 106 ±
14 0.53 × ±
18 0.54 × , - - 120 ±
18 0.53 × ±
16 0.54 × , - - 92 ±
15 0.54 × ±
13 0.54 × OCH , ±
12 0.57 × ±
25 0.57 × ±
16 0.57 × , ±
10 0.54 × ±
11 0.55 × ±
13 0.54 × OCHO 19 , ±
10 0.57 × ±
12 0.57 × ±
12 0.57 × , ± × ±
14 0.55 × ±
14 0.54 × CHO 11 , - - 51 ± × ±
11 0.57 × , - - 64 ±
10 0.54 × ±
11 0.54 × CO 12 , ± × ±
13 0.53 × ±
10 0.54 × Table 5.
Integrated intensities and uncertainties for spectra extracted from the continuum peak pixel. For clarity, only upperstate quantum numbers are used to identify each line; refer to Table 3 for full identifiers. Figure 5.
Full spectrum extracted from the continuum peak pixel in Ser-emb 17 (grey line), along with synthetic spectra of theCOMs studied in this work (colored lines) assuming the CH OH rotational temperature. CH OCH CH OCHO NH CHO CH CON T (10 cm − ) N T (10 cm − ) N T (10 cm − ) N T (10 cm − )Min. Ω a Ω s Max. Ω a Ω s Min. Ω a Ω s Max. Ω a Ω s Min. Ω a Ω s Max. Ω a Ω s Min. Ω a Ω s Max. Ω a Ω s Ser-emb 1 1.1 . . . . . . . . - - 2.5 . . . . Ser-emb 8 2.2 . . . . . . . . . . . . . . . . Ser-emb 17 1.2 . . . . . . . . . . . . . . . . Table 6.
Organic molecule column density estimates. Column densities are calculated assuming two bounding beam dilutionfactors, as described in the text. For each beam dilution factor, column densities are shown for three rotational temperatureassumptions: T M +75 K − K , where T M is the CH OH rotational temperature derived in each source. When multiple lines aredetected, the listed column density is the uncertainty-weighted average of the values calculated for each line.
Figure 6.
Examples of spectra with line profiles suggestive of rotation (a, b), infall (b, c, d), and outflows (e). For clarity, onlyupper state quantum numbers are used to identify each line; refer to Table 3 for full identifiers. OH columndensities are typically on the order 10 cm − , thougha lower value (a few × cm − ) was found in the in-cipient hot corino B1b-S (Marcelino et al. 2018). TheCH OH column densities derived in this work for theSerpens sources are ∼ –10 cm − , consistent withB1b-S but low compared to the other hot corinos. Wehave accounted for beam dilution and optical deptheffects as best as possible given the current obser-vations, but additional observations are needed con-firm the low CH OH column densities in the Serpenssources. Higher-resolution observations that cover ad-ditional CH OH isotopologue lines would enable morerobust constraints on the degree to which beam dilutionand optical depth effects impact our results.If confirmed, the low column densities in the Serpenssources and B1b-S suggest that they are intrinsicallyweaker/smaller hot corinos, which could be due to theirlow source luminosities. This highlights the power ofALMA to detect and characterize new hot corinos thatspan a broad range of physical properties.4.3.
Organic abundances across evolutionary stages
It is of considerable interest to understand how theorganic chemistry changes during the evolutionary pro- gression from protostars to protoplanetary disks to plan-etary systems. At (sub-)millimeter wavelengths, wecan directly probe the gas-phase reservoir of organics,which is also indirectly linked to the ice-phase reser-voir through different desorption processes. In the cold,outer envelopes of low-mass protostars, gas-phase or-ganic molecules likely represent the non-thermal des-orption products of grain-surface chemistry. At the lowdensities and temperatures of these environments, theformation of saturated organics is not efficient in thegas phase and is thought to proceed exclusively in icemantles (Horn et al. 2004; Geppert et al. 2006; Garrod& Herbst 2006). Later, in the hot corino stage, gas-phase molecules emit from much warmer and denser re-gions, where most ice will have sublimated into the gasphase. Gas-phase compositions therefore reflect thermalice desorption and possibly additional gas-phase chem-istry (Herbst & van Dishoeck 2009; Balucani et al. 2015;Skouteris et al. 2017, 2019). Finally, we can best ex-plore COM chemistry at the Class II disk stage usingcometary measurements, which are made in the gas-phase but are believed to probe pristine ices. It is notcurrently clear the extent to which cometary ice com-positions are inherited from the interstellar medium, orreprocessed during the assembly of a planetary system.Protostellar observations suggest that material is altered2by heating/shocking during infall and accretion onto thedisk (Oya et al. 2016), while various lines of evidence in-dicate that solid material (both refractory and volatile)in outer Solar system regions is at least partially inher-ited (Mumma & Charnley 2011; Altwegg et al. 2017).To explore the organic chemistry in different types ofsources, we compare COM/CH OH column density ra-tios. Since CH OH is predicted to serve as a feedstockmolecule for the formation of larger COMs in ice man-tles ( ¨Oberg et al. 2009; Garrod & Herbst 2006; Garrodet al. 2008), the COM/CH OH ratios can be thoughtof as a conversion efficiency. We also note that even inthe maximum beam dilution case the optical depths ofall lines are <
1. This means that uncertainties in thebeam dilution factor will cancel out when comparingCOM/CH OH column density ratios.Figure 7 shows the column density ratios of each or-ganic with respect to CH OH in the Serpens protostellardisk candidates, along with measurements in a sampleof low-mass protostellar envelopes (Bergner et al. 2017);the hot corinos IRAS 16293 B (Jørgensen et al. 2018),L483 (Jacobsen et al. 2018), NGC 1333-IRAS 2A and4A (Taquet et al. 2015), and Barnard 1b-S (Marcelinoet al. 2018) ; and the solar system comets Hale-Bopp,Lemmon, Lovejoy, and 67P summer and winter hemi-spheres (Bockel´ee-Morvan et al. 2000; Crovisier et al.2004; Biver et al. 2014, 2015; Le Roy et al. 2015). For theSerpens sources, organic molecule column densities arecalculated assuming the CH OH rotational temperature T M , and error bars show the values assuming rotationaltemperatures of T M ±
75 K.The NH CHO and CH CO column density ratios withrespect to CH OH are consistent between the Serpenssources and other hot corinos, while the CH OCH andCH OCHO ratios in Serpens are enhanced by at least afactor of ∼
10 compared to the other hot corinos. Thereis also at least an order of magnitude spread in theCH OCH , CH OCHO, and NH CHO ratios across theother hot corinos. We note that some of this variationmay be an artifact of the different angular resolutionsand analysis techniques used to derive column densities.Still, at present it appears that there is a wide rangein the abundances of large, oxygen-bearing molecules indifferent hot corino sources.Several possible mechanisms could contribute to theobserved chemical diversity in hot corinos. Given thathot corino emission is typically not well resolved, it ispossible that different hot corinos host different physical For the B1b-S hot corino, we have multiplied the CH OHcolumn densities listed in Marcelino et al. (2018) by 70 to derivea CH OH column density. components (e.g. infall, jets, accretion shocks, rotatingdisks) that alter what molecules are observed in the gasphase. Alternatively, chemical variation among hot cori-nos could reflect differences in the amounts of time spentin different physical states. For instance, the infall ratedetermines how long ice mantles are heated prior to sub-limation, with slower infall allowing more time for theformation of larger organics in lukewarm ices (Garrodet al. 2008). Time-variable infall chemistry could alsocause large differences in the observed organic inven-tories of different aged sources. Similarly, if gas-phasechemistry is active following ice sublimation in hot cori-nos, then the size and age of the hot corino should influ-ence the observed gas-phase abundances. Interestingly,the Serpens sources are more chemically similar to oneanother than the other hot corinos, suggesting that theconditions in the local birth cloud may play an impor-tant role in setting the protostellar chemistry.While some or all of these explanations could beat play, we emphasize the need for additional high-resolution and high-sensitivity observations. Expandingthe number of hot corinos for which we have resolvedobservations of organic molecule emission will be keyto understanding the variation in hot corino chemistry.Additionally, while we have modeled the CH OH opti-cal depths in the Serpens sources as well as the datawill allow, observations of minor CH OH isotopologuesand/or a wider set of lines is needed to confirm the highCOM to CH OH ratios.As shown in Figure 7, previously studied hot cori-nos are chemically similar both to cold protostellar en-velopes and to Solar system comets, suggestive of chem-ical inheritance throughout these evolutionary stages.However, the Serpens observations show that the hotcorino stage is more chemically variable than previouslyassumed. This means that care must be taken when us-ing protostellar observations to predict the compositionof pre-cometary material in the Solar system. We needa deeper understanding of what drives chemical diver-sity at the hot corino stage, and in turn what sourcesare the best analogs to the young Sun, in order to makeinferences about our chemical history. CONCLUSIONSWe have surveyed the organic molecules CH OH,CH OCH , CH OCHO, NH CHO, and CH CO towardsfive Class 0/I protostellar disk candidates. Based on ouranalysis, we conclude the following:1. Warm organic molecule emission consistent withhot corino chemistry is detected in three of fivetargeted sources. These observations suggest a3
Figure 7.
Comparison of organic molecule column density ratios with respect to methanol at different evolutionary stages. TheSer-emb sources are shown as blue squares. Column densities are derived for the continuum peak position, assuming the CH OHrotational temperature T M ; error bars show the column densities derived for T M ±
75 K. Measured abundances in the cold outerenvelopes of Class 0/I low-mass protostars are shown as pink dots, hot corinos as yellow dots (+ markers indicate IRAS 16293B), and Solar System comets as orange dots. possible association between hot corinos and pro-tostellar disks.2. For the three sources with CH OH detections, wederive column densities of 10 –10 cm − and ro-tational temperatures of ∼ OH-normalized column density ratios ofCH OCH and CH OCHO in the Serpens sourcesand other hot corinos span two orders of magni-tude. Local environmental effects and/or time-dependent warm-up chemistry could contribute tothis chemical diversity.4. High-resolution observations of these and otherhot corino sources are required to disentanglewhether hot corino emission is associated with aprotostellar disk, and to better understand the ori-gins of chemical diversity seen at the hot corinostage. This will be important for connecting Solarsystem comet compositions with earlier evolution-ary stages.ACKNOWLEDGEMENTSThis paper makes use of ALMA data, project code2015.1.00964.S. ALMA is a partnership of ESO (rep-resenting its member states), NSF (USA), and NINS (Japan), together with NRC (Canada) and NSC andASIAA (Taiwan), in cooperation with the Republic ofChile. The Joint ALMA Observatory is operated byESO, AUI/NRAO, and NAOJ. The National Radio As-tronomy Observatory is a facility of the National Sci-ence Foundation operated under cooperative agreementby Associated Universities, Inc.J.B.B acknowledges funding from the National Sci-ence Foundation Graduate Research Fellowship un-der Grant DGE1144152. This work was supportedby an award from the Simons Foundation (SCOL emcee (Foreman-Mackeyet al. 2013),
NumPy (van der Walt et al. 2011),
Mat-plotlib (Hunter 2007),
Astropy (Astropy Collabora-tion et al. 2013), and
SciPy (Jones et al. 2001).4 APPENDIX A. MOMENT ZERO MAP PARAMETERSTable 7 lists the rms values and velocity ranges used to make the moment zero maps shown in Figures 1 and 2.
Ser-emb 1 Ser-emb 7 Ser-emb 8 Ser-emb 15 Ser-emb 17Vel. rms Vel. rms Vel. rms Vel. rms Vel. rms(km/s) (mJy) (km/s) (mJy) (km/s) (mJy) (km/s) (mJy) (km/s) (mJy)C O 2 6.7–10.5 4.1 6.0–12.0 5.4 6.0–10.7 5.6 8.2–12.5 4.4 4.7–10.5 5.7CH OH 5 , OH 10 , OCH , OCHO 19 , CHO 12 , CO 12 , Table 7.
Velocity ranges and rms values used for the integrated intensity maps (Figures 1 and 2) in the main text. For clarity,only upper state quantum numbers are used to identify each line; refer to Table 3 for full identifiers.5
Velocity ranges and rms values used for the integrated intensity maps (Figures 1 and 2) in the main text. For clarity,only upper state quantum numbers are used to identify each line; refer to Table 3 for full identifiers.5 B. POPULATION DIAGRAM DIAGRAM FITTINGFor the MCMC population diagram fits to CH OH, we use a flat prior 10 < N T < cm − and 100 < T r < Figure 8.
Ser-emb 1 rotational diagram MCMC fit results. The corner plot is shown on the left and the walker chain on theright.
Figure 9.
Ser-emb 8 rotational diagram MCMC fit results. The corner plot is shown on the left and the walker chain on theright. Figure 10.
Ser-emb 17 rotational diagram MCMC fit results. The corner plot is shown on the left and the walker chain on theright. C.
SPECTRAL LINE FITSFigures 11–14 show Gaussian fits to the observed lines of each COM, analogous to Figure 3.
Figure 11. CH OCH spectral lines. Blue lines show the spectra extracted from the continuum peak pixel, and shaded regionsrepresent the rms. Red lines show Gaussian fits to the data. Figure 12. CH OCHO spectral lines. Blue lines show the spectra extracted from the continuum peak pixel, and shaded regionsrepresent the rms. Red lines show Gaussian fits to the data.
Figure 13. NH CHO spectral lines. Blue lines show the spectra extracted from the continuum peak pixel, and shaded regionsrepresent the rms. Red lines show Gaussian fits to the data.
Figure 14. CH CO spectra lines. Blue lines show the spectra extracted from the continuum peak pixel, and shaded regionsrepresent the rms. Red lines show Gaussian fits to the data. D. FULL SPECTRAFigures 15–18 show the full spectra extracted from the continuum peak pixel in Ser-emb 1, 7, 8, and 15, analogous toFigure 5. For Ser-emb 1 and 8 (Figures 15 and 17) colored lines show the synthetic spectra of COMs detected in eachsource.9
Figure 15.
Full spectrum extracted from the continuum peak pixel in Ser-emb 1 (grey line), along with synthetic spectra of thedetected COMs (colored lines). Spectra are calculated assuming the CH OH rotational temperature. Figure 16.
Full spectrum extracted from the continuum peak pixel in Ser-emb 7 (grey line). Figure 17.
Full spectrum extracted from the continuum peak pixel in Ser-emb 8 (grey line), along with synthetic spectra of thedetected COMs (colored lines). Spectra are calculated assuming the CH OH rotational temperature. Figure 18.
Full spectrum extracted from the continuum peak pixel in Ser-emb 15 (grey line).
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