The ice composition in the disk around V883 Ori revealed by its stellar outburst
Jeong-Eun Lee, Seokho Lee, Giseon Baek, Yuri Aikawa, Lucas Cieza, Sung-Yong Yoon, Gregory Herczeg, Doug Johnstone, Simon Casassus
aa r X i v : . [ a s t r o - ph . S R ] F e b The ice composition in the disk around V883 Ori revealedby its stellar outburst
Jeong-Eun Lee , Seokho Lee , Giseon Baek , Yuri Aikawa , Lucas Cieza , Sung-Yong Yoon ,Gregory Herczeg , Doug Johnstone , Simon Casassus SchoolofSpaceResearch,KyungHeeUniversity,1732,Deogyeong-Daero,Giheung-gu,Yongin-shi,Gyunggi-do17104,Korea Department of Astronomy, University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033,Japan Facultad de Ingenier´ıa y Ciencias, N´ucleo de Astronom´ıa, Universidad Diego Portales, Av.Ejercito441. Santiago,Chile Kavli Institute for Astronomy and Astrophysics, Peking University, Yi He Yuan Lu 5, HaidianQu, 100871,Beijing,PR China NRC Herzberg Astronomy and Astrophysics, 5071 West Saanich Road, Victoria, BC, V9E 2E7,Canada Departamento deAstronom´ıa,UniversidaddeChile,Casilla36-D Santiago,Chile
Complex organic molecules (COMs), which are the seeds of prebiotic material and precur-sors of amino acids and sugars, form in the icy mantles of circumstellar dust grains butcannot be detected remotely unless they are heated and released to the gas phase. Aroundsolar-mass stars, water and COMs only sublimate in the inner few au of the disk , makingthem extremely difficult to spatially resolve and study. Sudden increases in the luminosity of he central star will quickly expand the sublimation front (so-called snow line ) to larger radii,as seen previously in the FU Ori outburst of the young star V883 Ori . In this paper, wetake advantage of the rapid increase in disk temperature of V883 Ori to detect and analyzefive different COMs, methanol, acetone, acetonitrile, acetaldehyde, and methyl formate, inspatially-resolved submillimeter observations. The COMs abundances in V883 Ori is in rea-sonable agreement with cometary values . This result suggests that outbursting young starscan provide a special opportunity to study the ice composition of material directly related toplanet formation. Observations of comets and asteroids show that the Solar nebula was rich in water and or-ganic molecules. The inventory from the recent Rosetta mission to comet 67P/ChuryumovGerasimenkoincludes many complex organic molecules (COMs) as well as prebiotic molecules. These organ-ics, together with water, could have been brought to the young Earth’s surface by comets . Sincethe ice composition of comets is similar to ices in molecular clouds, it has long been debated ifcometary ices originate in the interstellar matter (ISM). The evolution of volatiles during planetformation may be traced by comparing the abundances from the youngest phases of star and diskformation, the protostellar core phase (when the star and disk are still growing from a circumstel-lar envelope), to the older phases of disk evolution. The Atacama Large Millimeter/submillimeterArray (ALMA) revealed that in some protostellar cores, called hot corinos, the warm gas containsvarious COMs that are also detected in comets
7, 8 . The variation of chemical composition amongcomets , however, suggests that chemical processes could be active even after material from theISM is incorporated to the disk . The chemical reactions, in turn, depend on various physical2arameters in the disk, e.g. ionization rate and dust-gas decoupling, which are under debate
10, 11 .Despite efforts to observe ices in disks, clear detection of ices in the disk midplane is chal-lenging and requires sophisticated analysis that considers the inclination angle and foregroundcontamination . An alternative approach is to observe COMs in the gas phase inside the snowline, defined as the radius outside of which a molecule is predominantly in ice form. The locationof each snow line depends on the heating of the disk by the central star and by the release of grav-itational energy via mass accretion, as well as on the sublimation temperature of the molecule .However, for a typical disk around a solar-mass star and a mass accretion rate of ∼ − M ⊙ yr − ,the COM and water snow lines are located at a radius of only a few AU , too small to be spatiallyresolved even with ALMA. So far, three gaseous COMs, CH OH, CH CN, and HCOOH, havebeen detected in non-bursting protoplanetary disks , with emission that traces non-thermallydesorbed molecules in the outer ( >
10 AU) disk surface, far beyond the snow line. Derivation ofsolid-phase abundances from those observations is not straightforward; it needs detailed under-standing of the non-thermal desorption mechanisms and of gas-phase reactions that could changethe abundances after desorption .The abundances of COMs at the disk locations where icy planetesimals form may instead bebest measured during a protostellar outburst , which viscously heats the disk and extends snowline to a much larger radius. The gas emission from the freshly sublimated COMs should directlytrace the abundances of COMs that were previously in the ice phase. FUors are young stellarobjects that exhibit large-amplitude outbursts in the optical (change in magnitude ∆ m V > mag)3nd undergo rapid increases of accretion rate, by ∼ − to − M ⊙ yr − ) . The high luminosity due to the enhanced accretion rate shifts the snow line to much largerradii . For the FUor outburst V883 Ori studied here, the radius of the water snow line at the diskmidplane is estimated to be ∼
40 AU from 1.3 mm continuum emission ; the snow line at the disksurface may extend to ∼
160 AU. Theoretical studies show that the sublimates are destroyed by thegas-phase reactions in several 10 yr . Since the duration of a typical outburst is usually ≤ ⊙ (see Methods), with ALMA inBand 7 on 8 September 2017 with a resolution of ∼ . ′′ , designed for continuum images, andon 21 February 2018 with a resolution of ∼ . ′′ , designed to measure COM emission. From thelower-resolution observation, we detected many COMs lines, including CH CHO, CH OCHO,and CH COCH , which are robustly detected for the first time in protoplanetary disks (Figure1). CH CHO and CH OCHO were detected contemporaneously and independently in the sametarget . Figure 2 shows the integrated intensity images (contours) and intensity weighted ve-locity images (colors) of emission coadded in many lines of individual COMs from the lower-resolution observation. All COM emission is confined within ∼ ′′ , but the emission peaks areoffset from the continuum peak. We find excitation temperatures for CH OH of 91.5 ± OCHO of 103.3 ± ≤ , see Methods).While the high resolution data was obtained to image the dust continuum, it also containssome COMs emissions. The high-resolution images (Figure 3 and Supplementary Figures 2 and3) reveal that the COM emission originates in a ring region around the water snow line that waspreviously identified from the dust continuum observations. The CH OH emission arises onlybetween 0.1 ′′ -0.2 ′′ , in contrast to the expectation that COMs emission should be bright inside thesnow line. The position-velocity map presented in Figure 3(c) is reproduced well by ring within aKeplerian disk, consistent with the CO emission .Disks are heated by stellar irradiation and viscous accretion
2, 25 . When stellar irradiationdominates, then the disk surface is warmer than the midplane, while the midplane could be warmerwhen the accretion heating dominates. The location where the water sublimates thus varies bothwith the distance from the star and with the height from the midplane, which together defines thesnow surface (Supplementary Figure 5). Because the molecular line opacity tends to be higherthan that of the dust continuum, the 2-D (radial and vertical) temperature distribution is needed toanalyze molecular line emission.Since heating by irradiation and accretion are both important for FUors , we tested two mod-els for the optically thin CH OH line emission: Model A for the disk heated only by irradiationand Model B for the disk heated by an enhanced accretion in the midplane as well as irradiation(see Methods and Supplementary Figure 4 for the details of models). Figure 4 compares the re-sults of the two models with the observed emission distribution of CH OH − . The5bserved emission decreases inside of its snow line, which is well reproduced by Model A and isa consequence of increasing dust continuum optical depths (see the bottom panels of Figure 4 andSupplementary Figures 4) caused by the evaporation of cm-sized icy pebbles (containing micron-sized dust grains) and the re-coagulation of the bare silicates into submm-sized grains . In ModelB, the CH OH emission declines even more steeply inwards due to the self-absorption causedby the negative temperature gradient in the vertical direction. Observations with higher spatialresolution and better sensitivity are needed to discriminate between these two models. Figure 4also compares the total amount of CH OH emission with that produced in the disk surface. Thefraction of emission from the midplane reaches a maximum near the methanol snow line, confirm-ing that the COM compositions in the gas phase originate in the ice compositions in the midplaneregions that are likely forming planetesimals.The average column densities of COMs over their emitting area are derived by fitting the ob-served spectra (Figure 1) with XCLASS (Supplementary Table 1). The abundances with respectto molecular hydrogen are − times higher than those derived from the spatially-extendedCOMs emission of non-bursting disks ( ∼ − − − ) , confirming that our observationsprobe the thermal sublimation of COMs. The D/H and C/ C ratios of CH OH are 0.16 and0.13, which are significantly higher than the ratios in protostellar cores and elemental abundances.While the high D/H ratio could be due to deuterium fractionation in the dense and cold grainsurfaces , the high C/ C ratio indicates that the column density of CH OH is underestimateddue to the high optical depth of the lines, requiring a more detailed treatment than that taken byXCLASS. We thus assume that the actual column density of CH OH is equal to that of CH OH6ultiplied by the elemental abundance of C/ C = 60 (see Methods).We find most COMs abundances relative to CH OH, i.e. their column density ratios, arehigher than those in the hot corino IRAS 16293 B
8, 29, 30 ; the abundances of CH COCH , CH CHO,and CH OCHO are higher than those in IRAS 16293 B by factors of 4 to 16 (Supplementary Table1). The abundances of CH COCH and CH CHO in both V883 Ori and IRAS 16293 B agree withthose in comet 67P/ChuryumovGerasimenko
4, 9 by a factor of a few. Acetone is not detected byby the Rosetta Orbiter Spectrometer for Ion and Neutral Analysis (ROSINA) on board Rosetta butwas estimated to be as abundant as CH CHO in 67P/ChuryumovGerasimenko by the CometarySampling and Composition Experiment (COSAC) , which could be in reasonable agreement withV883 Ori. CH CN abundance in V883 Ori, on the other hand, is lower than in IRAS 16293 B(comets) by a factor of 3 (an order of magnitude). The CH CN/CH OH abundance ratio is alsomuch lower than that in cold vapor in outer radii of other disks ( ∼
16, 17 . These abundancesindicate that formation and destruction of COMs continue after the volatiles are incorporated tothe disk.Our observations and analysis of V883 Ori demonstrate that observations of FUors areunique in revealing fresh sublimates and thus ice composition in protoplanetary disks. AlthoughFUors are rare, they span a range of evolutionary states: some outbursts occur while the systemis still being fed by their envelope, other outbursts occur in systems that have only disks and noenvelope . This diversity may provide us with the leverage to investigate how the chemical evo-lution of complex molecules in protoplanetary disks leaves its imprints on the products of planet7ormation. 8. Herbst, E. & van Dishoeck, E. F. Complex Organic Interstellar Molecules. Ann. Rev. As-tron. Astrophys. , 427–480 (2009).2. D’Alessio, P., Calvet, N. & Hartmann, L. Accretion Disks around Young Objects. III. GrainGrowth. Astrophys. J. , 321–334 (2001). astro-ph/0101443 .3. Cieza, L. A. et al.
Imaging the water snow-line during a protostellar outburst.
Nature ,258–261 (2016). .4. Le Roy, L. et al.
Inventory of the volatiles on comet 67P/Churyumov-Gerasimenko fromRosetta/ROSINA.
Astron. Astrophys. , A1 (2015).5. Altwegg, K. et al.
Organics in comet 67P - a first comparative analysis of mass spectra fromROSINA-DFMS, COSAC and Ptolemy.
Mon. Not. R. Astron. Soc. , S130–S141 (2017).6. Chyba, C. F., Thomas, P. J., Brookshaw, L. & Sagan, C. Cometary Delivery of OrganicMolecules to the Early Earth.
Science , 366–373 (1990).7. Imai, M. et al.
Discovery of a Hot Corino in the Bok Globule B335.
Astrophys. J. Lett. ,L37 (2016). .8. Jørgensen, J. K. et al.
The ALMA Protostellar Interferometric Line Survey (PILS). First resultsfrom an unbiased submillimeter wavelength line survey of the Class 0 protostellar binary IRAS16293-2422 with ALMA.
Astron. Astrophys. , A117 (2016). .9. Mumma, M. J. & Charnley, S. B. The Chemical Composition of Comets − Emerging Tax-onomies and Natal Heritage.
Ann. Rev. Astron. Astrophys. , 471–524 (2011).90. Furuya, K. & Aikawa, Y. Reprocessing of Ices in Turbulent Protoplanetary Disks: Carbon andNitrogen Chemistry. Astrophys. J. , 97 (2014). .11. Schwarz, K. R. et al.
Unlocking CO Depletion in Protoplanetary Disks. I. The Warm Molec-ular Layer.
Astrophys. J. , 85 (2018). .12. Pontoppidan, K. M. et al.
Ices in the Edge-on Disk CRBR 2422.8-3423: Spitzer Spec-troscopy and Monte Carlo Radiative Transfer Modeling.
Astrophys. J. , 463–481 (2005). astro-ph/0411367 .13. Hama, T. & Watanabe, N. Surface Processes on Interstellar Amorphous Solid Water: Ad-sorption, Diffusion, Tunneling Reactions, and Nuclear-Spin Conversion.
Chem. Rev. ,8783–8839 (2013).14. Walsh, C. et al.
First Detection of Gas-phase Methanol in a Protoplanetary Disk.
Astro-phys. J. Lett. , L10 (2016). .15. i¯binfoauthor ¨Oberg, K. I. et al.
The comet-like composition of a protoplanetary disk as revealedby complex cyanides.
Nature , 198–201 (2015). .16. Loomis, R. A. et al.
Detecting Weak Spectral Lines in Interferometric Data through MatchedFiltering.
Astron. J. , 182 (2018). .17. Bergner, J. B., Guzm´an, V. G., bibinfoauthor ¨Oberg, K. I., Loomis, R. A. & Pegues, J.A Survey of CH CN and HC N in Protoplanetary Disks.
Astrophys. J. , 69 (2018). . 108. Favre, C. et al.
First Detection of the Simplest Organic Acid in a Protoplanetary Disk.
Astro-phys. J. Lett. , L2 (2018). .19. Bertin, M. et al.
UV Photodesorption of Methanol in Pure and CO-rich Ices: DesorptionRates of the Intact Molecule and of the Photofragments.
Astrophys. J. Lett. , L12 (2016). .20. Molyarova, T. et al.
Chemical Signatures of the FU Ori Outbursts.
Astrophys. J. , 46(2018). .21. Audard, M. et al.
Episodic Accretion in Young Stars.
Protostars and Planets VI .22. Harsono, D., Bruderer, S. & van Dishoeck, E. F. Volatile snowlines in embedded disks aroundlow-mass protostars.
Astron. Astrophys. , A41 (2015). .23. Nomura, H., Aikawa, Y., Nakagawa, Y. & Millar, T. J. Effects of accretion flow on the chem-ical structure in the inner regions of protoplanetary disks.
Astron. Astrophys. , 183–188(2009). .24. van ’t Hoff, M. L. R. et al.
Methanol and its Relation to the Water Snowline in the Disk aroundthe Young Outbursting Star V883 Ori.
Astrophys. J. Lett. , L23 (2018). .25. Qi, C. et al.
CO J = 6-5 Observations of TW Hydrae with the Submillimeter Array.
Astro-phys. J. Lett. , L157–L160 (2006). astro-ph/0512122 .116. Dullemond, C. P., Hollenbach, D., Kamp, I. & D’Alessio, P. Models of the Struc-ture and Evolution of Protoplanetary Disks.
Protostars and Planets V astro-ph/0602619 .27. Schoonenberg, D., Okuzumi, S. & Ormel, C. W. What pebbles are made of: Interpretation ofthe V883 Ori disk.
Astron. Astrophys. , L2 (2017). .28. M ¨oller, T., Endres, C. & Schilke, P. eXtended CASA Line Analysis Software Suite(XCLASS).
Astron. Astrophys. , A7 (2017). .29. Lykke, J. M. et al.
The ALMA-PILS survey: First detections of ethylene oxide, acetoneand propanal toward the low-mass protostar IRAS 16293-2422.
Astron. Astrophys. , A53(2017). .30. Calcutt, H. et al.
The ALMA-PILS survey: first detection of methyl isocyanide (CH NC) in asolar-type protostar.
Astron. Astrophys. , A95 (2018). .12 orrespondence
Correspondence and requests for materials should be addressed to Jeong-Eun Lee (email:[email protected]).
Acknowledgements
ALMA is a partnership of ESO (representing its member states), NSF (USA) andNINS (Japan), together with NRC (Canada), NSC and ASIAA (Taiwan), and KASI (Republic of Korea), incooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO andNAOJ. J.-E.L. is supported by the Basic Science Research Program through the National Research Founda-tion of Korea (grant no. NRF- 2018R1A2B6003423) and the Korea Astronomy and Space Science Instituteunder the R&D programme supervised by the Ministry of Science, ICT and Future Planning. G.H. is fundedby general grant 11473005 awarded by the National Science Foundation of China. D.J. is supported by theNational Research Council of Canada and by an NSERC Discovery Grant. Y.A. acknowledges support fromJSPS KAHENHI grant numbers 16K13782 and 18H05222.
Author Contributions
J.-E.L., S.L. and G.B. performed the detailed calculations and line fittings usedin the analysis. J.-E.L. wrote the manuscript. All authors were participants in the discussion of results,determination of the conclusions and revision of the manuscript.
Competing Interests
The authors declare that they have no competing financial interests.
The COMs detected in the Cycle 5 ALMA observation.
The grey color is for theobserved spectra. The RMS noises of observed spectra are 0.015 and 0.02 Jy for the spectral win-dows around 330 GHz and the spectral windows around 340 GHz, respectively, and the dashed lineindicates the 3 σ level. Different colors indicate the lines of individual species fitted by XCLASS (see Methods and Supplementary Tables 1 to 3). The spectral window in the second row coversthe series of CH OH 7 k -6 k transitions. For the line fitting, we adopted the temperature of 120 K,hydrogen column density of . × cm − , and the dust properties from the disk model pre-sented in Methods. The identified species and their derived column densities and abundances bythis line fitting process are summarized in Supplementary Table 1. Ethylene oxide ( c -H COCH ),formic acid ( t -HCOOH), and methyl mercaptan (CH SH) are tentatively identified; we detect onlyone transition of each COM. The HDO line at 335.395 GHz (upper state energy, E u = 335 K) isalso marginally detected as marked in the first spectral window, and its estimated column densityis 3.8 × cm − . 14 V eo l c i t y ( k m s - ) Figure 2:
Intensity weighted velocity images (colors) and Integrated intensity maps (contours)of the molecular lines clearly detected in the Cycle 5 ALMA observations.
Multiple lines aredetected in each species except for C O, and they are averaged for a high signal-to-noise ratio.The contours are from 5 σ in step of 3 σ except for CH OH and CH CHO. Their contour stepsare 5 σ . The RMS noise levels are 9.9, 3.5, 3.8, 4.5, 6.1, and 7.3 mJy beam − km s − for C O3-2, CH OH, CH CHO, CH OCHO, CH OH, and CH CN, respectively. The cross indicates theposition of the continuum peak. The observed beam is presented in the lower left panel.15 -20 15 50 85 120 I I N T ( m Jy bea m - k m s - ) -8.0 4.5 17.0 29.5 42.0 I n t en s i t y ( m Jy / bea m ) Figure 3:
The high-resolution images obtained from the Cycle 4 ALMA observation. a,b,
Integrated intensity (I
INT ) maps of CH OH 7 -6 ( a ) and CH OH 12 -12 ( b ). A uv-taperwas applied to make images with high signal-to-noise ratios and a resolution ( ∼ . ′′ ) enoughto resolve clearly the snow line. The lowest contour and subsequent contour step are 5 σ and 3 σ ( σ = − km s − ). The water snow line identified from the analysis of continuumimages in Band 6 is marked with the black solid ellipses. The synthesized beam is plotted in theleft bottom corner. c, The position-velocity map along the white line in b , which follows the disksemi-major axis. ∆ R represents the offset from the continuum peak along the white line. The thingrey contours are for 5 σ and 8 σ ( σ = − ). The thick black solid curves representthe Keplerian motion around the central mass of 1.2 M ⊙ with the inclination of 38.3 ◦ . The blackcontours (20, 40, 60 and 80 % of the peak intensity) illustrate the synthesized emission from thesame simple optically thin disk model for C O 2-1 line , except for the beam size and the fact thatthe emission arises only from 0.1 to 0.2 ′′ . 16 Figure 4:
Models for the CH OH − line. a,b, Radial distribution of the observed ( a )and synthesized ( b ) integrated intensity of CH OH − . Panel a shows the azimuthallyaveraged integrated intensity profile along the de-projected radius r. The sky blue represents the1 σ RMS error. In b , for each model, the solid lines represent the results from the disk modelheated only by irradiation (Model A) while the dotted lines depict the results from the disk modelheated by accretion as well as irradiation (Model B). The black lines show the synthesized intensityprofile from the entire disk in our models, while the red lines represent the results from the samemodels but with CH OH only in the disk surface. The difference between black and red lines isequivalent with the emission from the midplane. c , The gray lines present the fraction of emissionoriginating from the disk midplane ( I Midplane /I Total ). d,e, The two-dimensional distributions oftemperature, T gas ( d ) and optical depth at the line center, τ line ( e ) for model B. z is the height fromthe midplane with units of scale height, H . The white solid and dashed lines in d depict the waterand methanol snow lines, respectively. The black solid and white dashed lines in e describe theheights with the continuum optical depths of 1 and 0.1, respectively.17 ethodsALMA observations V883 Ori was observed using the Atacama Large Millimeter/submillimeter Array (ALMA)during Cycle 4 (2016.1.00728.S, PI: Lucas Cieza) on 2017 Sep. 8 and during Cycle 5 (2017.1.01066.T,PI: Jeong-Eun Lee) on 2018 Feb. 21. For the Cycle 4 observation, a spectral window was centeredat 334.3955GHz (HDO 3 , -4 , ) with a bandwidth of 93.75 MHz and spectral resolution of 488kHz ( δ v ∼ . km s − ). Three windows used for continuum, centered at 333.41 GHz, 345.41 GHz,and 347.41 GHz with a bandwidth of 2 GHz, are not analyzed in this paper. The phase center wasat ( α, δ ) J = (05 h m . s , − ◦ ′ . ′′ ) , and the total observing time was 15.6 minutes.Forty-seven 12-m antennas were used with baselines in the range from 21 m ( ≃
24 k λ ) to 10.6 km( ≃ λ ) to provide the synthesized beam size of . ′′ × . ′′ (PA= . ◦ ) when Briggsweighting (robust = 0.5) is adopted. For the Cycle 5 observation, ten spectral windows wereset to cover many molecular lines such as CH OH, CH OH, H CO + J = 4 → , and C O J = 3 → . Their bandwidths were 468.75 MHz or 117.19 MHz and their spectral resolution was282 kHz ( ∼ − ). The total observing time was 44.53 minutes. The phase center was at ( α, δ ) J = (05 h m . s , − ◦ ′ . ′′ ) . Forty-five 12-m antennas were used with baselinesin the range from 15 m ( ≃
17 k λ ) to 1.4 km ( ≃ λ ) to provide the synthesized beam size of . ′′ × . ′′ (PA= . ◦ ) when Briggs weighting (robust = 0.0) is adopted.We carried out the standard data reduction using CASA 4.7.2 (for Cycle 4) and 5.1.1 (forCycle 5) . For the Cycle 4 data, the nearby quasar J0541-0541 was used for phase calibration,180503-1800 for bandpass calibration, and J0423-0120 for flux calibration. For the Cycle 5 data,J0607-0834 was used for phase calibration while J0510+1800 was used for bandpass and fluxcalibration. For the Cycle 4 data, Briggs weighting with a uv-taper of 1000 k λ was used for themolecular lines to obtain images with a higher signal-to-noise ratio and a resolution ( . ′′ × . ′′ , PA= . ◦ ), enough to resolve the water snow line. The root mean square (RMS) noiselevel of the molecular lines is 4 mJy beam − . For the Cycle 5 data, the RMS noise level for themolecular lines is 7 ∼ − . We did not use a uv-taper for the Cycle 5 data. Distance to V883 Ori
V883 is a member of the L1641 cluster . A cross-match between the sample of L1641members and the Gaia DR2 objects
35, 36 yields 47 stars with parallaxes accurate to at least 5%,have Gaia G -band photometry brighter than 19 mag, are located within 0.5 deg of V883 Ori, andhave a proper motion within 1 mas/yr of the average proper motion of the cluster. (The spatialand proper motion cuts are important because the Gaia DR2 astrometry reveals that the L1641members include at least two distinct clusters located at different distances.) This set of stars hasa distance of pc and a standard deviation of 16 pc, roughly consistent with previous a GaiaDR2 analysis . This standard deviation could be explained by either a real scatter of 12 pc in thedistances or by underestimated errors in the parallax measurements.The Gaia DR2 parallax of V883 Ori yields a distance of ± pc, and a proper motion thatis highly discrepant with the L1641 cluster. If these values are correct, then V883 Ori could not bea member of L1641 or any other known young cluster. Moreover, the Gaia DR2 catalog flags the19strometric solution of V883 Ori as unreliable, which is not surprising for a variable young star.We therefore adopt the median distance to the nearby stars in L1641 of 388 pc for V883 Ori. Withthis updated distance, the midplane snow line is located at 39 AU rather than 42 AU and the massof the central star is 1.2 M ⊙ . Line Identification and Fitting
In Figure 1, all lines were extracted with the circular aperture of 0.6 ′′ , within which the signal-to-noise ratio of CH OH integrated intensity is above 5. The line central velocity at each pixelwithin the aperture, which is affected by the disk rotation, is shifted to the source velocity to reducethe line blending . The source velocity is 4.3 km s − . We identify spectral line transitions (seeSupplementary Table 3) and fit the spectra using the eXtended CASA Line Analysis Software Suite(XCLASS) , which accesses the Cologne Database for Molecular Spectroscopy (CDMS)
39, 40 andJet Propulsion Laboratory (JPL) molecular databases. We still have several unidentified linetransitions. In addition, only one transition was detected for each of c -H COCH , t -HCOOH, andCH SH.XCLASS takes into account the line opacity and line blending in the assumption of local ther-modynamic equilibrium (LTE). The main parameters used to fit lines in XCLASS are the emissionsize, the excitation temperature, the column density of the species, the line full width at half max-imum (FWHM), and the velocity offset with respect to the systemic velocity of the object. Theaperture size (0.6 ′′ ) was adopted as the emission size. There is no velocity offset because all lineswere shifted to the source velocity before combined. All lines are assumed to be Gaussian profiles20ith a FWHM fixed at 2 km s − , the average value of all lines when allowed to vary. The excitationtemperature of 120 K is adopted from the disk model below; the temperature was averaged withina radius of 0.3 ′′ in the model. With these parameters, we used MAGIX (Modeling and AnalysisGeneric Interface for eXternal numerical codes) to optimize the fit and find the best solution.The observational uncertainty is dominated by the absolute calibration error of 10 %. TheCOMs column densities derived from this line fitting are summarized in Supplementary Table1. The error of the column density is calculated by taking absolute calibration uncertainty andfitting error into consideration. The derived column density of CH OH is larger than that fromthe excitation diagram (Supplementary Figure 1) by a factor of 7.6. This may be caused by thecontinuum opacity rather than the optical depths of the CH OH lines, which are optically thinaccording to the XCLASS fitting. The line intensity is scaled by e − τ c , where τ c is the continuumoptical depth, and τ c ∼ . in band 7 at 0.1 ′′ . This continuum optical depth effect is taken intoaccount in XCLASS but not in the excitation diagram.The derived column density of CH OH suggests a very low C/ C ratio of 7.6, which ismuch lower than the typical ISM value (60) . The best-fit column density of CH OH might beunderestimated as the modeled lines at the frequency range of 338.42 to 338.48 GHz are weakerthan the observed ones. If we increase the column density by a factor of 2, those lines are fittedbetter. However, this column density does not fit the rest of lines, and it cannot increase the isotopicratio to the normal value either. This low isotopic ratio between CH OH and CH OH maybe caused by very high optical depth of CH OH, which prevents the CH OH lines from tracing21he matter close to the midplane. Alternatively, the 2-D structures of temperature, density, andabundance could be the main reason. The XCLASS fitting of the spectra in Figure 1 is not perfect;the model spectra deviate from the observation for some other lines as well as those methanollines. It indicates the need for more sophisticated modeling considering the 2-D distributions oftemperature, density, and molecular abundances. We postpone such modeling to future work, sinceour high spatial resolution data cover only a limited number of lines (see Supplementary Table 2).We also list the COM abundances with respect to CH OH to compare with the values in thehot corino IRAS 16293B and in comets. In the calculation of the COMs abundances in V883 Ori,we assume that the true column density of CH OH is the column density of CH OH multipliedby the elemental abundance of C/ C = 60 . The C/ C ratio in CH OH is expected to besimilar to the ratio in CO because CH OH forms mainly via hydrogenation of CO on grain surfaces . Since CO is the dominant carbon carrier, it is reasonable to assume CO/ CO=60. The CO/ CO ratio of 71 ±
15 measured directly from CO ices in star forming regions is also similarto the elemental ratio. The CH OH abundance derived from the CH OH abundance and the C/ C ratio of 60 is then × − , which is in reasonable agreement with the recent study( × − ) .As listed in Supplementary Table 1, the COMs abundances relative to CH OH are higherthan those in the hot corino IRAS 16293 B and similar to those in comets, except for CH CN.This pattern is indicative of chemical evolution in the disk that continues after the ISM material isincorporated to the disk. During the quiescent phase, for example, methyl formate (CH OCHO)22an form via grain-surface reactions of HCO and CH O in the disk midplane . However, it isdifficult to compare directly our observational results with the chemical models because there arestill many uncertainties in the grain surface reactions, such as the branching ratios of radical-radicalreactions. Models
In order to investigate the 2-D (r and z) distribution of CH OH from our high spatial reso-lution data, we construct a simple irradiated disk model to calculate the radiative transfer of the CH OH 12 -12 line in the assumption of LTE. The gas is assumed to be well coupled withdust grains in the region where the CH OH line forms, and thus, the gas temperature is the sameas the dust temperature. We assume that the gas to dust mass ratio is constant as 100 over the entiredisk. The vertical density profile in the disk is given as n H ( r, z ) = N H ( r ) H √ π exp( − z H ) , (1)with the H column density, N H and the scale height, H . For the vertical temperature distribution,we follow the vertical gradient prescription of an irradiated disk model
47, 48 : T ( r, z ) = T atm ( r ) + ( T mid ( r ) − T atm ( r )) cos (cid:18) πz H (cid:19) if z < H = T atm ( r ) if z ≥ H, (2)with the temperature at the atmosphere, T atm and that at the midplane, T mid . For simplicity, weassume that T atm /T mid is 2 based on the dust continuum radiative transfer model for V883 Ori .To find N H ( r ) and T mid ( r ) , we solve a dust continuum radiative transfer for the high reso-23ution ALMA continuum image at Band 6: I ν ( r ) = Z + ∞−∞ B ν ( T ( r, z )) exp( − τ cν ( r, z )) κ cν ( r, z ) dz, (3)with κ cν ( r, z ) = κ abs m H n H ( r, z ) /g d (4) τ cν ( r, z ) = Z z −∞ κ cν ( r, z ′ ) dz ′ . where κ abs is the dust absorption opacity at 1.3 mm ( . g − ) , m H is the molecular hydrogenmass, and g d is the gas to dust mass ratio.A previous study shows that V883 Ori has a very optically thick inner disk and a opticallythin outer disk. The boundary, that is located at 0.1 ′′ , is interpreted as water snow line. In the innerdisk, the dust temperature is relatively well constrained while the dust temperature and the columndensity are degenerate in the optically thin outer disk. Therefore, we analyze the inner and outerdisks separately. At r > ′′ , we assume that T mid ( r ) follows the profile of T mid (0 . ′′ ) × q . ′′ /r ,as used in the irradiated disk . We adopt T mid (0 . ′′ ) = 100 K (ref. ) based on the temperatureestimated in the optically thick region. Given temperature distribution, N H ( r ) can be derived bysolving above equations. The derived column density at . ′′ is ∼ . × cm − and the densityat the midplane is ∼ cm − when H/r is 0.1. At this density, the freeze-out timescale of COMsis shorter than a year if the temperature is lower than the sublimation temperature . To derive thetemperature profile at the inner disk ( r ≤ . ′′ ), N H ( r ) is set to be N H (0 . ′′ ) × ( r/ . ′′ ) − γ .Here, we adopt γ = 1 . , a typical power index of the surface density profile in the disk . Thetemperature distribution in this disk model (Model A, hereinafter) is shown in Supplementary24igure 4.For the synthetic CH OH line emission, the continuum absorption coefficient is consideredas a function of frequency, κ cν ( r, z ) = κ cν ( r, z ) × ( ν/ν ) β with the spectral index of β = 1 in theentire disk. The total absorption coefficient is given by κ ν ( r, z ) = κ cν ( r, z ) + κ lν ( r, z ) , (5) κ lν ( r, z ) = c πν ul A ul . V n H ( r, z ) X mol Q ( T ) " χ l g u g l − χ u exp " − (cid:18) ∆ vσ (cid:19) , where c is the speed of light, ν ul and A ul are the frequency and the Einstein coefficient of thetransition, respectively, χ u ( χ l ) and g u ( g l ) is the level population and statistical weight in theupper (lower) energy state, respectively. ∆ V (= σ × √ ; in cm s − ) is the full width athalf maximum of the line profile, ∆ v is the velocity shift from the source velocity, Q ( T ) is thepartition function, and X mol is the molecular abundance of CH OH. We adopted ∆ V = 2 kms − , which is the average value over all lines, and X mol = 4 × − , which reproduces the observedpeak intensity. When T gas is lower than the sublimation temperature, the molecule is assumed tobe frozen onto grain surfaces, depleted from the gas. The sublimation temperature of CH OH,which fits the radial intensity profile, was 80 K, as marked with the white dashed line in Figure4(d); above the dashed line, X mol = 4 × − , but below the line, X mol = 0 . This lower sublimationtemperature of CH OH than water sublimation temperature is consistent with the lower bindingenergy (5000 K) of CH OH than that of water (5600 K) . The molecular data are adopted fromCDMS database .Finally, the line emission is synthesized using the equation (3), but with the total absorption25oefficient, κ ν ( r, z ) , instead of κ cν . The calculated intensity profile is smoothed to have the sameresolution (0.11 ′′ ) of the observed image. To check whether the molecular line can trace down tothe midplane, we also calculated a model in which the CH OH abundance is artificially set to bezero in the midplane ( | z | ≤ H ) (red lines in Figure 4); the difference between the black and redlines can be considered the emission from the disk midplane. The fraction of emission from thedisk midplane is presented by the gray lines in Figure 4.The spectral index derived from the high resolution Band 7 and Band 6 continuum imagesof V883 Ori is smaller than 2, reaching down to 1, within ∼ . ′′ (Cassasus et al. in prep.). Inthe submillimeter regime, the free-free emission is negligible, and thus, most flux is likely fromthermal emission, whose spectral index must be greater than 2. Note that we can see a deeperregion in Band 6 than in Band 7. Thus, the flux ratio of Band 7 to Band 6 smaller than expectedfrom the isothermal case suggests that the deeper region is hotter. Therefore, we model the diskheated by accretion (at r < . ′′ ) as well as irradiation (Model B).In Model B, to combine the two heating processes, irradiation and accretion, in a simple way,we assume that N H ( r ) and T mid ( r ) are the same as those of Model A, except for the temperaturedistribution within the snow line. Within the snowline ( r < . ′′ ), we assume that the midplaneis vertically isothermal at z ≤ H (accretion heating), while the temperature gradient is given byEquation (2) at z > H (irradiative heating). In Equation (2), which describes the temperatureprofile resulted by irradiation, we assume T mid ( r ) = T mid (0 . ′′ ) × q . ′′ /r as used for r > . ′′ .In Model A at a given radius, one temperature was found to fit the continuum intensity. Therefore,26f at a given radius, the T mid ( r ) calculated by the above equation is lower than T mid ( r ) in ModelA, then a higher midplane temperature (the blue line in the upper right panel in SupplementaryFigure 4) is found to fit the observed intensity. The derived temperature distribution of Model B ispresented in Figure 4(d). 27 ata availability This paper makes use of the ALMA data, which could be downloaded from the ALMAarchive (https://almascience.nao.ac.jp/aq/) with project codes 2016.1.00728.S and 2017.1.01066.T. The datathat support the plots within this paper and other findings of this study are available from the correspondingauthor upon reasonable request.
31. Goesmann, F. et al.
Organic compounds on comet 67P/Churyumov-Gerasimenko revealed byCOSAC mass spectrometry.
Science (2015).32. McMullin, J. P., Waters, B., Schiebel, D., Young, W. & Golap, K. CASA Architecture andApplications. In Shaw, R. A., Hill, F. & Bell, D. J. (eds.)
Astronomical Data Analysis Soft-ware and Systems XVI , vol. 376 of
Astronomical Society of the Pacific Conference Series , 127(2007).33. Allen, L. E. & Davis, C. J. Low Mass Star Formation in the Lynds 1641 Molecular Cloud. InReipurth, B. (ed.)
Handbook of Star Forming Regions, Volume I , 621 (2008).34. Fang, M. et al.
Young Stellar Objects in Lynds 1641: Disks, Accretion, and Star FormationHistory.
Astrophys. J. Suppl. , 5 (2013). .35. Gaia Collaboration et al.
The Gaia mission.
Astron. Astrophys. , A1 (2016). .36. Gaia Collaboration et al.
Gaia Data Release 2. Summary of the contents and survey properties.
Astron. Astrophys. , A1 (2018). .37. Kounkel, M. et al.
The APOGEE-2 Survey of the Orion Star-forming Complex. II. Six-dimensional Structure.
Astron. J. , 84 (2018). .288. Yen, H.-W. et al.
Stacking Spectra in Protoplanetary Disks: Detecting Intensity Profiles fromHidden Molecular Lines in HD 163296.
Astrophys. J. , 204 (2016). .39. M ¨uller, H. S. P., Thorwirth, S., Roth, D. A. & Winnewisser, G. The Cologne Database forMolecular Spectroscopy, CDMS.
Astron. Astrophys. , L49–L52 (2001).40. M ¨uller, H. S. P., Schl¨oder, F., Stutzki, J. & Winnewisser, G. The Cologne Database for Molec-ular Spectroscopy, CDMS: a useful tool for astronomers and spectroscopists.
J. Mol. Struct. , 215–227 (2005).41. Pickett, H. M. et al.
Submillimeter, millimeter and microwave spectral line catalog.
J. Quant. Spectrosc. Ra. , 883–890 (1998).42. M ¨oller, T. et al. Modeling and Analysis Generic Interface for eXternal numerical codes(MAGIX).
Astron. Astrophys. , A21 (2013). .43. Langer, W. D. & Penzias, A. A. (C-12)/(C-13) isotope ratio in the local interstellar mediumfrom observations of (C-13)(O-18) in molecular clouds.
Astrophys. J. , 539–547 (1993).44. Furuya, K., Aikawa, Y., Sakai, N. & Yamamoto, S. Carbon Isotope and Isotopomer Fraction-ation in Cold Dense Cloud Cores.
Astrophys. J. , 38 (2011). .45. Boogert, A. C. A., Blake, G. A. & Tielens, A. G. G. M. High-Resolution 4.7 MicronKeck/NIRSPEC Spectra of Protostars. II. Detection of the CO Isotope in Icy Grain Man-tles.
Astrophys. J. , 271–280 (2002). astro-ph/0206420 .46. Walsh, C. et al.
Complex organic molecules in protoplanetary disks.
Astron. Astrophys. ,A33 (2014). . 297. Dartois, E., Dutrey, A. & Guilloteau, S. Structure of the DM Tau Outer Disk: Probing thevertical kinetic temperature gradient.
Astron. Astrophys. , 773–787 (2003).48. Andrews, S. M. et al.
Resolved Images of Large Cavities in Protoplanetary Transition Disks.
Astrophys. J. , 42 (2011). .49. Cieza, L. A. et al.
The ALMA early science view of FUor/EXor objects - V. Continuum discmasses and sizes.
Mon. Not. R. Astron. Soc. , 4347–4357 (2018). .50. D’Alessio, P., Cant¨o, J., Calvet, N. & Lizano, S. Accretion Disks around Young Objects. I.The Detailed Vertical Structure.
Astrophys. J. , 411–427 (1998). astro-ph/9806060 .51. Lee, J.-E., Bergin, E. A. & Evans, N. J., II. Evolution of Chemistry and Molecular Line Profilesduring Protostellar Collapse.
Astrophys. J. , 360–383 (2004). astro-ph/0408091 .52. Wakelam, V., Loison, J.-C., Mereau, R. & Ruaud, M. Binding energies: New values andimpact on the efficiency of chemical desorption.
Molecular Astrophysics , 22–35 (2017). . 30 upplementary Table 1: Identified COMs Species Formula Column density a X( w.r.t. H ) X( w.r.t. CH OH) b IRAS16293B c Comets d Methanol CH OH 4.00 +0 . − . × × − – – –CH DOH 6.21 +0 . − . × × − × − – – CH OH 5.26 +0 . − . × × − × − – –Acetaldehyde CH CHO 6.40 +0 . − . × × − × − × − × − Methyl Formate CH OCHO 2.37 +0 . − . × × − × − × − × − Acetone CH COCH +0 . − . × × − × − × − –Acetonitrile CH CN 2.45 +0 . − . × × − × − × − × − Ethylene oxide c -H COCH e +0 . − . × × − × − – –Formic acid t -HCOOH e +0 . − . × × − × − – –Methyl mercaptan CH SH e +1 . − . × × − × − – – a The uncertainty of column density corresponds to the 1 σ error. b X( w.r.t CH OH) was calculated using the derived column density of CH OH and the C/ C ratio of 60. c X( w.r.t CH OH) from the refs (8,29,30) d X( w.r.t CH OH) from the ref (4) e Only one transition was detected for these COMs. Therefore, these three COMs are tentatively identified. upplementary Table 2: Cycle 4 Line Identification Line No. Formula Name Frequency [GHz] Transition Einstein-A [log A] E u [K]1 CH CHO vt = 1 Acetaldehyde 334.98085310 (3.48E-5) 18(1, 18) - 17(1, 17), E -2.87593 359.893772 HDCO Formaldehyde 335.09673940 (8.96E-5) 5( 1, 4)- 4( 1, 3) -2.98087 56.248263 CH OH vt = 0 Methanol 335.13357000 (1.3E-5) 2(2)- - 3(1)- -4.57254 44.672664 CH CHO v = 0 Acetaldehyde 335.31810910 (2.83E-5) 18(0, 18) - 17(0, 17), E -2.87112 154.927235 CH CHO v = 0 Acetaldehyde 335.35872250 (2.83E-5) 18(0, 18) - 17(0, 17), E -2.87128 154.852926 CH CHO vt = 1 Acetaldehyde 335.38246150 (3.91E-5) 18(0, 18) - 17(0, 17), E -2.86462 361.484067 CH OH vt = 0 Methanol 335.56020700 (4.0E-5) 12( 1, 11)- 12( 0, 12) - + -3.39358 192.654058 CH OH vt = 0 Methanol 335.58201700 (5.0E-6) 7(1)+ - 6(1)+ -3.78844 78.971839 H C O Formaldehyde 335.81602540 (2.40E-4) 5( 1, 5)- 4( 1, 4) -2.97975 60.23580 upplementary Table 3: Cycle 5 Line Identification Line No. Formula Name Frequency [GHz] Transition Einstein-A [log A] E u [K]spw01 CH CHO v=0 Acetaldehyde 335.35872250 (2.83E-5) 18(0, 18) - 17(0, 17), E -2.87128 154.852922 CH CHO vt=1 Acetaldehyde 335.38246150 (3.91E-5) 18(0, 18) - 17(0, 17), E -2.86462 361.484063 HDO ? Water 335.39550000 (2.6E-5) 3( 3, 1)- 4( 2, 2) -4.58367 335.26718spw14 c-H COCH Ethylene Oxide 336.56139300 (0.00015) 9( 4, 5)- 8( 5, 4) -3.52971 90.313065 CH OH vt=2 Methanol 336.60588900 (1.3E-5) 7(1)+ - 6(1)+ -3.78629 747.41346spw26 (CH ) CO v=0 Acetone 336.62703300 (2.76E-5) 17(16, 1)-16(15, 1) EE -2.80789 142.455897 SO v = 0 ? Sulfur dioxide 336.66957740 (3.9E-6) 16( 7, 9)-17( 6,12) -4.23374 245.114228 (CH ) CO v=0 Acetone 336.70098190 (3.14E-5) 17(16, 2)-16(15, 2) EE -2.80766 142.34420spw39 H S ? Hydrogen disulfide 337.02909220 (6.0E-6) 25(3,22) - 26(2,24) -4.25089 277.6817010 CH OH vt=2 Methanol 337.02957300 (1.8E-5) 7(2) - 6(2) -3.80855 941.3879411 t -HCOOH Formic Acid 337.05330470 (4.07E-5) 11( 3, 9)-11( 2,10) -4.61210 58.2029012 C O Carbon Monoxide 337.06112980 (1.0E-5) J=3-2 -5.63440 32.3532313 CH CHO vt = 2 Acetaldehyde 337.08157220 (0.0001427) 18(1, 18) - 17(1, 17), E -2.88192 526.13974spw414 CH CHO vt = 2 Acetaldehyde 338.31785810 (0.000161) 18(5, 14) - 17(5, 13), E -2.86939 605.1564915 CH CHO vt = 2 Acetaldehyde 338.31906720 (0.000161) 18(5, 13) - 17(5, 12), E -2.86939 605.1565516 SO v=0 ? Sulfur Dioxide 338.32036020 (4.4E-6) 13(2,12)-12(1,11) -3.64436 92.5831317 CH OCHO v=0 Methyl Formate 338.33818400 (0.0002) 27( 8,19)-26( 8,18) E -3.26914 267.1835818 CH OH vt = 0 Methanol 338.34458800 (5.0E-6) 7(-1,7) - 6(-1,6) -3.77807 70.5508319 CH OCHO v=0 Methyl Formate 338.35579200 (0.0001) 27( 8,19)-26( 8,18) A -3.26882 267.1857220 CH OCHO v=0 Methyl Formate 338.39631800 (0.0001) 27( 7,21)-26( 7,20) E -3.25943 257.7450321 CH OH vt = 0 Methanol 338.40461000 (5.0E-6) 7(6) - 6(6), E1 -4.34500 243.7922122 CH OH vt = 0 Methanol 338.40869800 (5.0E-6) 7(0)+ - 6(0)+ -3.76895 64.9814923 CH OCHO v=0 Methyl Formate 338.41411600 (0.0001) 27( 7,21)-26( 7,20) A -3.25931 257.7470324 CH OH vt = 0 Methanol 338.43097500 (5.0E-6) 7(-6) - 6(-6), E2 -4.34266 253.9484325 CH OH vt = 0 Methanol 338.44236700 (5.0E-6) 7(6)+ - 6(6)+ -4.34360 258.6984126 CH OH vt = 0 Methanol 338.45653600 (5.0E-6) 7(-5) - 6(-5), E2 -4.07836 188.9999727 CH OH vt = 0 Methanol 338.47522600 (5.0E-6) 7(5) - 6(5), E1 -4.07807 201.0607728 CH OH vt = 0 Methanol 338.48632200 (5.0E-6) 7(5)+ - 6(5)+ -4.07635 202.8856929 CH OH vt = 0 Methanol 338.50406500 (5.0E-6) 7(-4) - 6(-4), E2 -3.94036 152.8944730 CH OH vt = 0 Methanol 338.51263200 (5.0E-6) 7(4) - 6(4) -3.93970 145.3340631 CH OH vt = 0 Methanol 338.51264400 (5.0E-6) 7(4)+ - 6(4)+ -3.93970 145.3340632 CH OH vt = 0 Methanol 338.51285300 (5.0E-6) 7(2)- - 6(2)- -3.80281 102.7028333 CH OH vt = 0 Methanol 338.53025700 (5.0E-6) 7(4) - 6(4), E1 -3.93781 160.9917834 CH OH vt = 0 Methanol 338.54082600 (5.0E-6) 7(3)+ - 6(3)+ -3.85727 114.7942935 CH OH vt = 0 Methanol 338.54315200 (5.0E-6) 7(3)- - 6(3)- -3.85727 114.7944136 CH OH vt = 0 Methanol 338.55996300 (5.0E-6) 7(-3) - 6(-3), E2 -3.85416 127.70688 OH vt = 0 Methanol 338.58321600 (5.0E-6) 7(3) - 6(3), E1 -3.85584 112.7100938 SO v = 0 ? Sulfur dioxide 338.61180780 (3.5E-6) 20( 1,19)-19( 2,18) -3.54241 198.8775039 CH OH vt = 0 Methanol 338.61493600 (5.0E-6) 7(1) - 6(1), E1 -3.76659 86.0523440 CH OH vt = 0 Methanol 338.63980200 (5.0E-6) 7(2)+ - 6(2)+ -3.80233 102.7175641 CH OH vt = 0 Methanol 338.72169300 (5.0E-6) 7(2) - 6(2), E1 -3.80941 87.2588542 CH OH vt = 0 Methanol 338.72289800 (5.0E-6) 7(-2) - 6(-2), E2 -3.80422 90.9134243 CH OH vt = 0 Methanol 338.75994800 (5.0E-5) 13( 0, 13)- 12( 1, 12) + + -3.66184 205.94501spw544 CH CHO v = 0 Acetaldehyde 346.95755590 (2.84E-5) 18(7, 12) - 17(7, 11), E -2.89625 268.60623CH CHO v = 0 Acetaldehyde 346.95755770 (2.84E-5) 18(7, 11) - 17(7, 10), E -2.89625 268.6062445 H C O Formaldehyde 346.98406710 (0.0002806) 5( 4, 2)- 4( 4, 1) -3.36307 239.7990646 H C O Formaldehyde 346.98409360 (0.0002806) 5( 4, 1)- 4( 4, 0) -3.36307 239.7990647 CH CHO v = 0 Acetaldehyde 346.99553240 (2.82E-5) 18(7, 12) - 17(7, 11), E -2.89625 268.5720948 CH CHO vt = 2 Acetaldehyde 346.99991340 (0.0001791) 18(7, 11) - 17(7, 10), E -2.89745 646.35557CH CHO vt = 2 Acetaldehyde 346.99994090 (0.0001791) 18(7, 12) - 17(7, 11), E -2.89745 646.3555749 CH CHO v = 0 Acetaldehyde 347.07154710 (2.58E-5) 18(6, 13) - 17(6, 12), E -2.87569 239.39974CH CHO v = 0 Acetaldehyde 347.07168440 (2.58E-5) 18(6, 12) - 17(6, 11), E -2.87569 239.3997550 CH CHO v = 0 Acetaldehyde 347.09040150 (2.57E-5) 18(6, 12) - 17(6, 11), E -2.87577 239.3976251 CH CHO v = 0 Acetaldehyde 347.13268590 (2.58E-5) 18(6, 13) - 17(6, 12), E -2.87563 239.3212452 H C O Formaldehyde 347.13389950 (0.0002025) 5( 3, 3)- 4( 3, 2) -3.11271 156.7694453 H C O Formaldehyde 347.14404620 (0.0002025) 5( 3, 2)- 4( 3, 1) -3.11259 156.7700754 CH CHO vt = 2 Acetaldehyde 347.15512480 (0.0001461) 18(4, 14) - 17(4, 13), E -2.84730 573.9299155 CH CHO v = 0 Acetaldehyde 347.16951960 (3.96E-5) 19(0, 19) - 18(1, 18), E -3.66327 171.8148756 CH CHO vt = 1 Acetaldehyde 347.18241320 (2.99E-5) 18(4, 14) - 17(4, 13), E -2.84531 400.3780557 CH OH vt = 0 Methanol 347.18828300 (6.4E-5) 14( 1, 13)- 14( 0, 14) - + -3.36086 254.2518558 CH CHO vt = 1 Acetaldehyde 347.21679780 (2.88E-5) 18(5, 13) - 17(5, 12), E -2.85923 420.4404159 CH CHO v = 0 Acetaldehyde 347.23739280 (7.6E-6) 8(2, 6) - 7(0, 7), E -5.82082 42.5379960 CH CHO vt = 1 Acetaldehyde 347.25182200 (3.08E-5) 18(5, 14) - 17(5, 13), E -2.85890 419.6726361 CH CHO vt = 1 Acetaldehyde 347.28421220 (7.31E-5) 18(10, 8) - 17(10, 7), E -2.98719 586.8846162 CH CHO v = 0 Acetaldehyde 347.28826400 (2.48E-5) 18(5, 14) - 17(5, 13), E -2.85868 214.6975563 CH CHO v = 0 Acetaldehyde 347.29487350 (2.48E-5) 18(5, 13) - 17(5, 12), E -2.85858 214.6983064 CH CHO vt = 1 Acetaldehyde 347.32243660 (5.17E-5) 18(9, 9) - 17(9, 8), E -2.95129 543.8172465 SiO v = 0 Silicon Monoxide 347.33063100 (0.0003475) 8-7 -2.65578 75.0169766 CH CHO v = 0 Acetaldehyde 347.34571040 (2.48E-5) 18(5, 13) - 17(5, 12), E -2.85860 214.6407467 CH CHO v = 0 Acetaldehyde 347.34927830 (2.48E-5) 18(5, 14) - 17(5, 13), E -2.85854 214.61141spw668 CH CHO vt = 1 Acetaldehyde 349.32035150 (3.33E-5) 18(1, 17) - 17(1, 16), E -2.81548 369.2564169 CH CN v = 0 Methyl Cyanide 349.34634280 (2.0E-7) 19(4) - 18(4) -2.61106 281.9877870 CH CN v = 0 Methyl Cyanide 349.39329710 (2.0E-7) 19(3) - 18(3) -2.60218 232.01121spw771 CH DOH Methanol 349.95168460 (7.1E-6) 7(4,4) - 7(3,4), e1 -3.99981 132.07804CH DOH Methanol 349.95183610 (8.6E-6) 14(6,8) - 15(5,10), o1 -4.60584 381.7761072 CH SH v = 0 Methyl Mercaptan 350.00961500 (4.8E-5) 14( 1) + - 13( 1) +A -3.37086 131.1251573 CH DOH Methanol 350.02734890 (8.0E-6) 6(4,3) - 6(3,3), e1 -4.04307 117.08865 H DOH Methanol 350.02776910 (8.0E-6) 6(4,2) - 6(3,4), e1 -4.04307 117.08867spw874 CH CHO v = 0 Acetaldehyde 350.08080620 (7.9E-6) 6(3, 4) - 5(2, 4), A -3.90056 39.7128675 CH DOH Methanol 350.09024240 (9.4E-6) 5(4,2) - 5(3,2), e1 -4.12044 104.24014CH DOH Methanol 350.09038410 (9.4E-6) 5(4,1) - 5(3,3), e1 -4.12044 104.2401576 CH OH vt = 0 Methanol 350.10311800 (5.0E-5) 1( 1, 1)- 0( 0, 0) + + -3.48249 16.8022077 CH CHO v = 0 Acetaldehyde 350.13342960 (2.75E-5) 18(3, 15) - 17(3, 14), E -2.82552 179.2091778 CH CHO v = 0 Acetaldehyde 350.13438160 (2.75E-5) 18(3, 15) - 17(3, 14), E -2.82596 179.17727spw979 CH OH vt = 0 Methanol 350.42158500 (5.0E-5) 8 ( 1 , 7)- 7 ( 2 , 5) -4.15313 102.6155480 CH OCHO v=0 Methyl Formate 350.44225000 (0.0001) 28( 8,21)-27( 8,20) E -3.21991 283.9067581 CH CHO v = 0 Acetaldehyde 350.44577770 (2.74E-5) 18(1, 17) - 17(1, 16), E -2.81490 163.4195282 CH OCHO v=0 Methyl Formate 350.45758000 (5.0E-5) 28( 8,21)-27( 8,20) A -3.21978 283.90950Note. Species with the same numbering are observed as a blended line. upplementary Figure
5: Excitation Diagrams of CH OH and CH OCHO. We plot the logof total number of molecules per degenerate sublevel ( N u /g u ) versus the energy in the upper state( E u ). In order to estimate the column density, we applied the aperture size (0.6 ′′ ) used for theextraction of lines. The column densities derived by the excitation diagrams are smaller than thevalues derived from the XCLASS fitting probably because the dust continuum emission is opticallythick. The error bars on the data points and the temperature and column density errors in the legendindicate the 1 σ error. 36 -20 10 40 70 100 I I N T ( m Jy bea m - k m s - ) Supplementary Figure
6: The same images as Figure 3 except for the CH CHO 18 -17 (left)and H C O 5 -4 (right) lines. These line emission images are also missing flux within the watersnow line. 37 -20 0 20 40 60 I I N T ( m Jy bea m - k m s - ) Supplementary Figure
7: Integrated intensity maps of CH OH 7 -6 (left) and CH OH 12 -12 (right) from the high resolution observation. A uv-taper of 2000 K λ was applied to makeimages with signal-to-noise ratios better than the original images and a resolution ( ∼ . ′′ ), higherthan those in Figure 2(a). The lowest contour and subsequent contour step are 5 and 3 σ ( σ = 5 . mJy beam − km s − ). The identified water snow line is marked with the black solid ellipses.The dotted ellipse describes the projected . ′′ radius. The synthesized beam is plotted in the leftbottom corner. 38 upplementary Figure
8: Radial distribution of temperature (upper) and optical depths (lower)for models. Upper panels show the temperature profiles used for the irradiated disk (Model A) andthe combination of the irradiated disk and the heated midplane by a burst accretion (Model B).The temperature profiles at the atmosphere, z = H , and the midplane are presented with the red,green, and blue lines, respectively. In the lower panels, the solid blue and green lines show theoptical depths at the line center of CH OH 12 -11 at the midplane and z = H , respectively.The dotted lines represent the Band 7 continuum optical depths at the same heights. At very smallradii, the continuum optical depth dominates the line optical depth.39 upplementary Figureupplementary Figure