Evidence for magnetic activity at starbirth: a powerful X-ray flare from the Class 0 protostar HOPS 383
Nicolas Grosso, Kenji Hamaguchi, David Principe, Joel Kastner
AAstronomy & Astrophysics manuscript no. 38185 c (cid:13)
ESO 2020June 5, 2020 L etter to the E ditor Evidence for magnetic activity at starbirth: a powerful X-ray flarefrom the Class 0 protostar HOPS 383
Nicolas Grosso , Kenji Hamaguchi , , David A. Principe , and Joel H. Kastner Aix-Marseille Univ, CNRS, CNES, LAM, Marseille, Francee-mail: [email protected] CRESST II and X-ray Astrophysics Laboratory NASA / GSFC, Greenbelt, MD, USA Department of Physics, University of Maryland, Baltimore County, Baltimore, MD, USA Massachusetts Institute of Technology, Kavli Institute for Astrophysics and Space Research, Cambridge, MA, USA Center for Imaging Science, School of Physics & Astronomy, and Laboratory for Multiwavelength Astrophysics,Rochester Institute of Technology, Rochester, NY, USAReceived 16 April 2020 / Accepted 2 May 2020
ABSTRACT
Context.
Class 0 protostars represent the earliest evolutionary stage of solar-type stars, during which the majority of the system massresides in an infalling envelope of gas and dust and is not yet in the central, nascent star. Although X-rays are a key signature ofmagnetic activity in more evolved protostars and young stars, whether such magnetic activity is present at the Class 0 stage is stilldebated.
Aims.
We aim to detect a bona fide Class 0 protostar in X-rays.
Methods.
We observed HOPS 383 in 2017 December in X-rays with the
Chandra
X-ray Observatory ( ∼
84 ks) and in near-infraredimaging with the Southern Astrophysical Research telescope.
Results.
HOPS 383 was detected in X-rays during a powerful flare. This hard ( E > ∼ ∼ × erg s − in the 2–8 keV energy band, a level at leastan order of magnitude larger than that of the undetected quiescent emission from HOPS 383. The X-ray flare spectrum is highlyabsorbed ( N H ∼ × cm − ), and it displays a 6.4 keV emission line with an equivalent width of ∼ . Conclusions.
The detection of a powerful X-ray flare from HOPS 383 constitutes direct proof that magnetic activity can be presentat the earliest formative stages of solar-type stars.
Key words. stars: flare – stars: individual: HOPS 383 – stars: low-mass – stars: magnetic field – stars: protostars – X-rays: stars
1. Introduction
Low-mass objects that have evolved beyond the Class 0 stageare conspicuous X-ray emitters (Dunham et al. 2014), that is,Class I protostars with remnant envelopes and massive ac-cretion disks, and Class II and III pre-main sequence starswith and without circumstellar disks (T Tauri stars). Theirhigh luminosities ( ∼ − erg s − ) compared to the solar max-imum ( ∼ erg s − ) and their intense flaring activity in X-rays make them appear as extremely magnetically active youngsuns (Feigelson & Montmerle 1999; Güdel & Nazé 2009). InClass 0 protostars (André et al. 1993, 2000), the hydrostaticcore is deeply embedded within its envelope and the molecularcloud, making its detection di ffi cult at most wavelengths (Gia-rdino et al. 2007). The most deeply-embedded X-ray sources re-ported in star-forming regions (Hamaguchi et al. 2005; Getmanet al. 2007 and references therein; Kamezaki et al. 2014) haveevolved beyond the Class 0 stage or their bolometric luminos-ity is not accurate enough for a robust conclusion (Appendix A).Moreover, any non-thermal radio emission from putative mag-netic activity at Class 0 protostars is very likely absorbed by thebases of their ionized outflows (Güdel 2002; Dzib et al. 2013). The object of this study in the Orion Molecular Cloud 3(OMC-3) was identified as a protostar by a Spitzer survey(Megeath et al. 2012) and included in the Herschel Orion pro-tostar survey (HOPS; Furlan et al. 2016) as source 383. The sub-millimeter to bolometric luminosity ratio of HOPS 383, 1.4%,(Safron et al. 2015) well surpasses the (0.5%) threshold for bonafide Class 0 designation (André et al. 1993, 2000). Notably,HOPS 383 is the first Class 0 protostar known to have undergonea mass-accretion-driven eruption (Safron et al. 2015), whichpeaked by 2008 and ended by 2017 September (Appendix B).
2. Observations
We observed HOPS 383 three times with the
Chandra
X-ray Ob-servatory (Weisskopf et al. 2002) from 2017 December 13 to 14with simultaneous near-infrared imaging on 2017 December 14using the 4.1 m Southern Astrophysical Research (SOAR) tele-scope (Krabbendam et al. 2004) in Chile.We used
Chandra ’s Advanced CCD Imaging Spectrometer(ACIS-I) in the very faint mode with a frame time of 3.141 s(Appendix C.1).
Chandra
ACIS-I has on-axis a full-width half-maximum (FWHM) angular resolution of 0 (cid:48)(cid:48) . Article number, page 1 of 12 a r X i v : . [ a s t r o - ph . H E ] J un & A proofs: manuscript no. 38185 a
32s 5h35m30s 28s - m s s Right Ascension (J2000) D e c li na t i on ( J ) b SMZ 1-2A SMZ 1-2B
CXOU J053529 HOPS 383 b - m s - m s Right Ascension (J2000) D e c li na t i on ( J ) c1.8" HOPS 383 SMZ 1-2BCXOU J053529 c . s - m . s Right Ascension (J2000) D e c li na t i on ( J ) JVLA @ 3.0 cm HOPS 383 JVLA-SE @ 4.0 cmJVLA-NW @ 4.0 cmCXOU J053529
Fig. 1.
Chandra detection of HOPS 383.
Panel a : the red and blue 0 (cid:48)(cid:48) . × (cid:48)(cid:48) .
492 pixels display logarithmically the 0.5–2 keV and 2–7 keV X-rays, respectively, which were detected from 2017 December 13 to 14 during three ACIS-I exposures, totaling 83 877 s. The logarithmic grayscale K -band image was obtained on 2017 December 14 with the Spartan infrared camera on the SOAR telescope. North is up and east is left. The pinkpluses and diamond, the green pluses, and cyan crosses are near- and mid-infrared sources (Megeath et al. 2012), H -shocked emission (Stankeet al. 2002), and X-ray sources, respectively. The white dashed box shows the field-of-view of Fig. 1b. Panel b : Spartan H -filter image obtainedon 2017 December 14 in linear grayscale. The white contours show the CO blueshifted outflow of HOPS 383 (Feddersen et al. 2020). The bluepluses are radio sources (Galván-Madrid et al. 2015; Rodríguez et al. 2017). The blue and green arrows show the direction of the thermal-jetcandidate and the proper motion of the H emission knot SMZ 1-2B, respectively. The white dashed box shows the field-of-view of Fig. 1c. Panel c : enlargement around HOPS 383 shows the mid-infrared (pink), radio (blue), and X-ray (cyan) positional error circles. spectral resolution of 261 eV at 6.4 keV. Data reduction is de-scribed in Appendix C.2.The SOAR Spartan infrared camera (Loh et al. 2012) is com-posed of 4 CCDs of 2 , × ,
048 pixels. We selected the wide-field mode with a single-detector field-of-view of 2 (cid:48) . × (cid:48) . (cid:48) . × (cid:48) .
04 edge-to-edge) and a pixel scale of 0 (cid:48)(cid:48) . K -band and H narrow-band filters. We used a five-pointdithering pattern where, in the first exposure, HOPS 383 wasput near the center of the northeastern detector (det3), whichwas then moved by 30 (cid:48)(cid:48) from its initial position sequentiallytoward the south, north, west, and east in the four other expo-sures. This sequence was repeated until the requested exposurewas achieved (Appendix D.1). Data reduction is described in Ap-pendix D.2.
3. Near-infrared results
The near-infrared nebulosity that was prominent during the out-burst (Safron et al. 2015) is not detected in our SOAR K -bandimage (Fig. 1a), indeed, it vanished by 2015 December 30 (Fis-cher & Hillenbrand 2017). The H -narrow filter image shows anemission knot of shocked molecular hydrogen at ∼ (cid:48)(cid:48) southeastfrom the mid-infrared location of HOPS 383. We identify this H source with the H emission knot SMZ 1-2B (Stanke et al. 2002)that was observed on 1997 September 13, yielding a proper mo-tion of ∼ (cid:48)(cid:48) . ± (cid:48)(cid:48) southeast in 20.25 yr (Fig. 1b). The correspond-ing velocity is ∼ (95 ± × (sin i / sin 69 . ◦ ) − × ( d /
420 pc) km s − ,where d is the distance to HOPS 383 and i is the inclination an- gle (90 ◦ corresponds to an edge-on view; Table 1 of Furlan et al.2016). This is typical of the velocity values observed in the out-flows from Class 0 protostars (Bontemps et al. 1996; Reipurth &Bally 2001). Assuming a constant velocity, the angular distanceof ∼ (cid:48)(cid:48) between HOPS 383 and this emission knot correspondsto a kinematical timescale of ∼ ±
100 yr. The proper motiondirection is consistent with the orientation of the outburst neb-ula, the CO blueshifted outflow (Feddersen et al. 2020), andthe position angle of the binary radio counterpart of HOPS 383at 4 cm (Galván-Madrid et al. 2015). We note that the unre-solved 3 cm counterpart (Rodríguez et al. 2017) corresponds tothe faint northwest component at 4 cm, JVLA-NW, which sug-gests that JVLA-NW would be the base of the radio thermal jetlaunched by HOPS 383, whereas the brighter southeast compo-nent at 4 cm, JVLA-SE, would be an emission knot along thisthermal jet (Fig. 1c).
4. X-ray results
We detect a hard ( E > Chandra observation ofthe OMC-2 and 3 region obtained in 2000 January at an o ff -axisangle of 6 (cid:48) . Chan-
Article number, page 2 of 12icolas Grosso et al.: a powerful X-ray flare from the Class 0 protostar HOPS 383 a D e c li n a t i o n ( J ) - d m s s s s b E v e n t e n e r g y ( k e V ) R e s p o n s e ( p r o b a b ili t y ) C u m u l a t i v e p r o b a b ili t y
54% C.I.(±0.74 ) c C u m u l a t i v e c o un t s C o un t r a t e ( c o un t k s ) Silverman'srule-of-thumbh rot /hour=0.241Adaptive kernelfrom h rot with =0.5Cross-validationh cv /hour=0.873Adaptive kernelfrom h cv with =0.5Mirrored rise phase Fig. 2.
X-ray flare of HOPS 383.
Panel a : the sky position of the X-ray events detected during the
Chandra observation on 2017 December 13.The cyan cross and circle are the X-ray position and positional error of HOPS 383, respectively. The dashed and solid contours encompass 90%(position determination) and 96% (spectrum extraction) of 1.49 keV point source emission, respectively. The diamonds mark the events from the6.4 keV emission line.
Panel b : event energy versus time of arrival. The gray stripe is centered at 6.4 keV and corresponds to a confidence intervalof 54% centered at this emission line observed with ACIS-I (inset).
Panel c : cumulative counts and various smoothed light curves. The dashedcurve is the mirrored rise phase of the red curve.
Panels b and c : the vertical dotted line at ∼ dra exposure that we obtained in 2017 (Fig. 2a). The first pho-ton is detected ∼ ∼ ∼ − in ∼ ∼ / or shocks in out-flows (Güdel & Nazé 2009), but the former produces plasmasthat are not bright and hot enough to emit the observed hard X-rays and the latter is constant on week timescales (Kastner et al.2005), which is much longer than the hour variability that weobserved. Therefore, this burst of X-rays indicates magnetic ac-tivity from HOPS 383 that is associated with JVLA-NW. We modeled the X-ray spectrum with isothermal coronal plasmaemission su ff ering from intervening absorption (Appendix C.5).The best-fit model has an excess and a deficit relative to the ob-served spectrum below and above 6.4 keV, respectively, suggest-ing an additional component at 6.4 keV. This deviation is likelyassociated with the emission line arising from neutral or low-ionization iron. Therefore, we included in the model a Gaussianemission line centered at 6.4 keV. We performed a Markov chainMonte Carlo analysis (MCMC) to compute the probability den-sity functions of the physical parameters, from which we de-termined the median parameter values and 90% confidence in-tervals. We find that CXOU J053529 is highly embedded witha hydrogen column density ( N H , X ) of 7 . + . − . × cm − . Thisresult for N H is consistent with the large median energy of the 0.5–8 keV counts (MedE = , given the log N H ver-sus MedE relations found for YSOs in the Orion nebula cluster(ONC; Fig. 8 of Feigelson et al. 2005) and M17 (Fig. 10a ofGetman et al. 2010).We detected the 6.4 keV emission line with a high proba-bility of 98.9%. Our spectral analysis suggests that we have de-tected five photons from the Fe 6.4 keV emission line; these pho-tons are likely the closest in energy to 6.4 keV. Indeed, we ob-served five photons with energies that are densely packed aroundthe predicted emission feature at 6.4 keV (diamonds in Fig. 2binset). Moreover, the spatial distribution of these Fe line photonsis point-like (Fig. 2a), suggesting all of the observed Fe fluo-rescence emission arises from within the Chandra point spreadfunction, that is to say, a radius from HOPS 383 that is lowerthan 105 au at the assumed distance to the protostar. The iron lineequivalent width, EW = . + . − . keV, is large (even at the lowerlimit) relative to what was expected from the possible emissionprocesses. The neutral iron can be ionized by photons or elec-trons that have energies larger than the Fe K-shell ionization po-tential of 7.112 keV. The contributions of these two ionizationchannels to the observed X-ray line in accreting young stars andClass I protostars (Güdel & Nazé 2009; Hamaguchi et al. 2005;Czesla & Schmitt 2010; Hamaguchi et al. 2010; Pillitteri et al.2019) are not firmly settled, but photoionization is less energeti-cally challenging (Czesla & Schmitt 2010). The plasma temper-ature is not well constrained mainly due to the low count rate, buta hot plasma temperature is favored ( kT = . + . − . keV), whichis consistent with a magnetic flare and the photoionization ofiron. For epoch 2017, the degradation of the ACIS e ff ective area at lowenergy due to the contamination of the optical blocking filter likelydoes not significantly modify the median energy of a source as hardas HOPS 383. Article number, page 3 of 12 & A proofs: manuscript no. 38185 a z ( a u ) X i =69.5(66.4 72.5 ) 6040200204060 z ( R )
60 40 20 0 20 40 60r (R )18 16 14 12 10 8log(Density / g cm ) b z ( a u )
18 16 14 12 10 8log(Density / g cm ) c z ( a u ) CXO angular resolution0.5
210 au @ 420 pc19.0 18.5 18.0 17.5 17.0 16.5 16.0log(Density / g cm ) d Distance along the line-of-sight from the X-ray source (au)10 N H ( c m ) accretion disk bipolar cavityinfalling envelope cloud i =72.5 i =69.5 i =67.0 i =66.4 Fig. 3.
Probing the gas and dust toward HOPS 383 with the X-ray spectrum.
Panels a-c : 2D density distribution from small to large spatial scalesand viewing angle from the model of HOPS 383 (Furlan et al. 2016). Dots indicate regions with dust.
Panel d : cumulative hydrogen columndensity along the line-of-sight from the X-ray source. The dashed lines show the contribution of the accretion disk, the bipolar cavity, the infallingenvelope, and the cloud. The data point with an error bar (90% C.I.) is the hydrogen column density that we obtained from the X-ray spectrummodeling.
The mean value of the absorption-corrected luminosity duringthe flare time interval is 1 . + . − . × erg s − in the 2–8 keVenergy band. Multiplying this value by the flare peak to flaremean count rate ratio (2.39) provides the flare peak value of ∼ . + . − . × erg s − . For comparison purposes, the maximumpeak value of the 84 X-ray flares with typical shapes observedfrom YSOs in the ONC (Getman et al. 2008) is 8 . × erg s − in the 0.5–8 keV energy band; only ∼
34% of these flares havepeak values that are higher than the HOPS 383 one. During ourobservations, the non-detection of HOPS 383 outside the flaretime interval implies that its quiescent absorption-corrected lu-minosity is lower than 2 . × erg s − in the 2–8 keV en-ergy band (Appendix C.6). Therefore, the flare that we detectedpeaked at least 21 + − times above the quiescent level. The lowerlimit on the total energy released in X-rays during this flare, as computed from the mean luminosity times the flare duration, was2 . + . − . × ergs.
5. Discussion
Starting from a model of HOPS 383 consisting of an infallingenvelope with bipolar cavities carved by outflows and a small(5 au radius) accretion disk (Fig. 3 a–c), constrained by the post-outburst spectral energy distribution (SED) using a Monte Carloradiative transfer code (Furlan et al. 2016), we computed the pre-dicted N H toward the central protostar (Appendix E). Due to thehigh inclination angle in this model, the line-of-sight interceptsthe upper layers of the accretion disk, which produce the bulkof N H ; moreover, this value is extremely sensitive to the view-ing angle (Fig. 3d). The observed N H , X favors a slightly lowerinclination. However, the available grid of protostellar models(Furlan et al. 2016) does not include a lower-mass accretion diskor a model with no accretion disk at all, which may reproduce Article number, page 4 of 12icolas Grosso et al.: a powerful X-ray flare from the Class 0 protostar HOPS 383 both the SED and N H , X . These model trials illustrate the impor-tance of X-ray constraints on N H to the overall consistency ofprotostar modeling.In the optically thin case at 6.4 keV, that is to say, aslong as the H column density by a cold absorber with asolar elemental abundance ( N (cid:48) H ) is lower than ∼ cm − ,the iron line EW produced by photoionization is EW = . Ω / π )( N (cid:48) H / cm − ) eV, where Ω is the subtended angleof the irradiated material from the X-ray irradiating source (Tsu-jimoto et al. 2005). Since Ω / π <
1, the large observed value of EW in HOPS 383 rules out the optically thin condition. In theoptically thick case of the photoionization of photospheric ironby an X-ray flare, the iron line EW that is computed by Monte-Carlo radiative transfer simulations is limited to ∼
130 eV forsolar iron abundance and can only be increased by factors < EW can be obtained in both optical depth conditions when the pho-toionizing X-ray source is partly eclipsed (Drake et al. 2008). Inthe model, the X-ray source is located on the obscured region ofthe star and near the base of the accretion funnels (Fig. 3a). Thissource irradiates the accretion funnels, the upper layers of the ac-cretion disk (Fig. 3b), and the circumstellar envelope (Fig. 3c).
6. Conclusions
Our detection of an X-ray flare from HOPS 383 provides conclu-sive evidence that strong magnetic activity is present at this bonafide Class 0 protostar. The resulting X-ray irradiation contributesto the ionization of the base of the outflow (Shang et al. 2004),and the magnetic reconnection that triggers powerful flares, sim-ilar to the one we detected, most likely drives energetic parti-cle ejections (Feigelson et al. 2002). Such protostellar cosmicrays have been proposed to collide with refractory dust grainslocated at the inner edge of the accretion disk, inducing spal-lation reactions that could yield short-lived radionuclei, such asare observed in the refractory inclusions of chondritic meteorites(Gounelle et al. 2013; Sossi et al. 2017). Mass ejections havesimilarly been invoked in the case of the bona fide Class 0 pro-tostar OMC-2 FIR 4 to explain the production of free electronsvia collisions in the envelope, thereby enhancing the abundancesof molecular ions (Ceccarelli et al. 2014). The Class 0 stage( ∼
10 000 yr) has an ∼
10 times shorter duration than the Class Istage; however, during the earliest stage, half of the envelopemass is accreted, building the central star and the accretion disk.Therefore, the determination of the internal irradiation level inClass 0 protostars is paramount for the understanding of proto-stellar chemistry.Longer observations are required to determine the flaringactivity level and to collect more photons from such deeply-embedded, nascent stars. Athena X-IFU (Barret et al. 2018),which is scheduled to be launched at the beginning of the 2030s,will improve the X-ray throughput and spectral resolution at6.4 keV, enabling a potential leap in our knowledge of the on-set of protostellar magnetic activity.
Acknowledgements.
We thank the anonymous referee for the careful readingof the manuscript. We are grateful to Joshua Wing (
Chandra
X-ray Center;CXC) and Steve Heathcote (Director of the Cerro Tololo Inter-American Ob-servatory) for the scheduling of the simultaneous
Chandra and SOAR observa-tions. K.H., D.P., and N.G. are thankful to Jonathan (Jay) Elias (SOAR Director)and Patricio Ugarte (Observer Support) for technical support during the Spar-tan observation. N.G. thanks Patrick Broos and Leisa Townsley for Acis Extractsupport. This work was supported by “the Programme National de PhysiqueStellaire” (PNPS) of CNRS / INSU co-funded by CEA and CNES. Support forthis work was provided by the National Aeronautics and Space Administration(NASA) through
Chandra
Award Numbers GO7-18012A and B issued by the CXC, which is operated by the Smithsonian Astrophysical Observatory for andon behalf of the NASA under contract NAS8-03060. The scientific results re-ported in this article are based on observations made by the
Chandra
X-ray Ob-servatory (ObsIDs 18927, 20882 and 20883). This research has made use ofsoftware provided by the CXC in the application packages CIAO. SOAR is ajoint project of the Ministério da Ciência, Tecnologia, e Inovação (MCTI) daRepública Federativa do Brasil, the U.S. National Optical Astronomy Observa-tory (NOAO), the University of North Carolina at Chapel Hill (UNC), and Michi-gan State University (MSU). This publication makes use of data products fromWISE, which is a joint project of the University of California (Los Angeles) andthe Jet Propulsion Laboratory / California Institute of Technology (JPL / Caltech),and NEOWISE which is a project of the JPL / Caltech. WISE and NEOWISEare funded by the NASA. This work has made use of data from the Euro-pean Space Agency (ESA) mission
Gaia (https: // / web / gaia,processed by the Gaia
Data Processing and Analysis Consortium (DPAC,https: // / web / gaia / dpac / consortium). Funding for the DPAChas been provided by national institutions, in particular the institutions partici-pating in the Gaia
Multilateral Agreement. This research has made use of Pythonand astropy, corner, matplotlib, numpy, pyregion, scipy and statsmodels mod-ules.
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Appendix A: Protostellar evolutionary stage L bol (L )10 M e n v ( M ) ,f M ,f = 0.1 M Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone ,f M ,f = 0.1 M Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
Class 0 protostars(M env > M / ) Class I protostars(M env < M / )Border zone
IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst IRS 7BCMM4SHOPS 383post-outburst HOPS 383 pre-outburst
Fig. A.1.
Envelope mass versus bolometric-luminosity diagram com-paring the pre- and post-outburst positions of HOPS 383 (green data)with previously reported Class 0 / I low-mass protostars detected in X-rays (blue and red data). The border zone between Class 0 ( M env > M (cid:63) /(cid:15) , where (cid:15) = . ffi ciency) and Class I( M env < M (cid:63) /(cid:15) ) protostars is defined by M env = . L bol (constant ac-cretion; dotted line) and M env = . L bol0 . (exponentially decayingaccretion; dashed line). The protostellar evolutionary tracks are shownfrom 0.01 M (cid:63), f to 0.9 M (cid:63), f , where M (cid:63), f is the final stellar mass. The risingtime along the tracks is indicated by white dots every 10 000 yr. Figure A.1 shows the envelope mass ( M env ) versusbolometric-luminosity ( L bol ) diagram, which is used as an evo-lutionary diagnostic (Saraceno et al. 1996; André et al. 2000;André et al. 2008). In this scenario, the Class 0 and I borderzone is defined by M env = . L bol (constant accretion, see An-dré & Montmerle 1994) and M env = . L bol0 . (exponentiallydecaying accretion, see Bontemps et al. 1996). Propagating themodel and flux uncertainties of HOPS 383 (Safron et al. 2015),we gather that: M env = . ± . M (cid:12) , L post-outburstbol = . + . − . L (cid:12) ,and L pre-outburstbol = . + . − . L (cid:12) (obtained from the preoutburst-to-outburst flux ratio at 24 microns of 35 ± M env is larger than M (cid:63) , con-firming its bona fide Class 0 stage. For comparison purposes, theinfrared source IRS 7B in the R Coronae Australis star-formingcore, which has a deeply embedded ( N H ∼ × cm − ) X-raycounterpart that was proposed as a Class 0 protostar candidate(Hamaguchi et al. 2005), is on the side of the Class I stage inthis plot, based on the bolometric luminosity reported in Groppiet al. (2007). The bolometric luminosity of the Class 0 candi-date CMMS4 (Kamezaki et al. 2014) is not accurate enough fora robust conclusion.We computed the protostellar evolutionary tracks assumingthe following (Bontemps et al. 1996; André et al. 2000; Andréet al. 2008): an initial envelope mass, M env (0), decays exponen-tially with a characteristic timescale, τ = yr; a mass accre-tion rate, ˙ M acc ( t ) = (cid:15) M env ( t ) /τ , where the local star-formatione ffi ciency is (cid:15) = . L bol ( t ) = L acc ( t ) + L (cid:63) ( t ), where L acc ( t ) = G ˙ M acc ( t ) M (cid:63) ( t ) / R (cid:63) ( t ) is the accretion lu-minosity, where G is the gravitational constant, and M (cid:63) (t), R (cid:63) ( t ),and L (cid:63) ( t ) are the protostellar mass, radius, and interior stellarluminosity, respectively, on the birthline of Fig. 1 in Palla & Stahler (1999) where a constant accretion rate of 10 − M (cid:12) yr − is assumed to compute the radius-mass relation (Stahler 1988;Palla & Stahler 1992). We estimate with this toy model that thefinal stellar mass of HOPS 383 will be ∼ . M (cid:12) and that less than ∼
10% of it has already been accreted onto the central protostar.
Appendix B: WISE and NEOWISE-R photometry . - m p h o t o m e t r y ( m a g n i t u d e ) Fig. B.1.
Decay of the episodic accretion of HOPS 383 in the mid-infrared. The 2010 and 2014–2018 data points are from WISE and NE-OWISE observations with the W2 filter at 4.6 µ m, respectively. Theblack and white dots are the profile-fit (Table B.1) and aperture (Ta-ble B.2) photometry for each observing epoch, respectively. The dashedvertical line is the arrival time of the first detected photon of the X-rayflare. Table B.1.
WISE and NEOWISE-R 4.6- µ m profile-fit photometry ofHOPS 383. Instrument (cid:104)
MJD (cid:105) ∆ MJD N (cid:104) w2mpro (cid:105) σ (day) (mag) (mag)WISE 55 263.89 0.99 6 10.606 0.036WISE 55 455.31 0.73 6 10.493 0.063NEOWISE-R 56 726.88 4.54 8 11.000 0.032NEOWISE-R 56 919.85 0.72 7 11.401 0.023NEOWISE-R 57 087.73 0.99 10 11.374 0.052NEOWISE-R 57 281.88 1.12 8 11.712 0.143NEOWISE-R 57 446.82 0.98 7 11.945 0.051NEOWISE-R 57 646.54 0.13 2 12.466 0.041NEOWISE-R 57 811.16 0.26 2 12.732 0.155 Notes.
The profile-fit photometry was computed from the single-exposure (L1b) source tables using the mean and standard deviation ofthe N profile-fit fluxes in the W2 band (w2mpro) obtained during eachobserving epoch. We note that ∆ MJD is the time interval between thefirst and the last frame for each epoch.
The mid-infrared flux of HOPS 383 started to rise between2004 and 2006, peaked by 2008 with a 24 µ m flux that was 35times brighter than the pre-outburst flux, and showed no largedecay between 2009 and 2012 (Safron et al. 2015). The skyposition of HOPS 383 was scanned in mid-infrared twice peryear by the Wide-field Infrared Survey Explorer (WISE) mission Article number, page 7 of 12 & A proofs: manuscript no. 38185
Table B.2.
WISE and NEOWISE-R 4.6- µ m aperture photometry of HOPS 383. Instrument (cid:104)
MJD (cid:105) ∆ MJD N MAGZP w2mcor (cid:104) w2mag (cid:105) σ (day) (mag) (mag) (mag) (mag)WISE 55 263.89 0.99 13 19.596 0.331 10.155 0.037WISE 55 455.11 1.12 15 19.692 0.331 10.226 0.037NEOWISE-R 56 726.88 4.54 10 19.669 0.322 10.616 0.037NEOWISE-R 56 918.17 4.87 16 19.642 0.322 10.888 0.038NEOWISE-R 57 087.66 1.12 14 19.645 0.326 10.917 0.037NEOWISE-R 57 281.94 1.25 16 19.641 0.326 11.137 0.038NEOWISE-R 57 446.82 0.98 7 19.641 0.326 11.274 0.038NEOWISE-R 57 649.88 8.78 17 19.644 0.326 11.599 0.038NEOWISE-R 57 814.43 7.99 15 19.644 0.326 11.708 0.038NEOWISE-R 58 013.16 1.11 16 19.643 0.326 12.557 0.042NEOWISE-R 58 170.74 6.48 12 19.643 0.326 12.307 0.042NEOWISE-R 58 377.31 1.11 13 19.643 0.326 12.565 0.045 Notes.
The aperture photometry in each observing epoch was computed from coadded single exposures of a good quality. The aperture correction(w2mcor) is taken from the single exposure (L1b) source tables. The uncertainty of the magnitude zero-point (MAGZP) is fixed to 0.037 mag. (Wright et al. 2010) in 2010 and by the Near-Earth Object Wide-field Infrared Survey Explorer (NEOWISE) reactivation mission(Mainzer et al. 2014) in 2014–2019 for which several single ex-posures were obtained at each observing epoch .The All-Sky, WISE 3-band Cryo, and NEOWISE Reacti-vation 2015–2019 data releases provide single exposure (L1b)source tables for 2010 and 2014–2018. The profile-fit photom-etry of HOPS 383 in the W2 band at 4.6 µ m is available from2010 March to 2017 February. As the full-width half-maximum(FWHM) in the W2 band is limited to 6 (cid:48)(cid:48) , HOPS 383 lieson the wing of the point-spread-function of a brighter source(2MASS J05353060-0459360), which is located at 15 (cid:48)(cid:48) . × FWHM in radius miti-gates this contamination. For each observing epoch, we selectedthe best profile-fit fluxes ( w rchi ≤
2) from the best single-exposure frames ( qual _ f rame > qi _ f act > saa _ sep > moon _ masked =
0) and we computed the mean and stan-dard deviation of these fluxes. The mean profile-fit flux decaysby 2.1 ± qual _ f rame >
0) and stacked them withthe WISE / NEOWISE coadder using simple area weighting anda pixel scale of 1 (cid:48)(cid:48) .
375 for the final image. HOPS 383 is stillvisible after 2017 February in these stacked images as a faintsource. To obtain aperture photometry, we used a standard 8 (cid:48)(cid:48) . (cid:48)(cid:48) radius background-aperture located 22 (cid:48)(cid:48) west of it, on a region free of point sourceswith lower extended emission compared to the 50 (cid:48)(cid:48) − (cid:48)(cid:48) ra-dius annulus used in the single exposure source tables. Theaperture correction (w2mcor) is taken from the single exposure(L1b) source table. The uncertainty of the magnitude zero-point(MAGZP) is fixed to the value of the WISE all-sky single ex-posures ( σ MAGZP = .
037 mag). The aperture flux decays by2.40 ± ± https: // irsa.ipac.caltech.edu / Missions / wise.html ber, which was three months before our Chandra observations(Table C.1).
Appendix C:
Chandra
ACIS-I data
Appendix C.1: Observations
The
Chandra observations are listed in Table C.1.
Table C.1.
Log of
Chandra observations with ACIS-I (PI: N.G.).
ObsID Start Time Exposure(UT) (s)18927 2017-12-13T04:00:48 37 41120882 2017-12-14T01:24:04 32 65220883 2017-12-14T17:09:31 13 814
Appendix C.2: Data reduction
We used the CIAO package (version 4.9 with the calibrationdatabase version 4.8.1; Fruscione et al. 2006) to reprocess thedatasets (Table C.1) with the tool chandra_repro . The sourcedetection was performed in each observation with the CIAO tool wavdetect with a detection threshold 1 × − and including in-formation from the point spread function (PSF). Twelve imageswere used per ObsID corresponding to three energy bands (0.5–2, 2–7, and 0.5–7 keV) and four resolution images (0 (cid:48)(cid:48) .
25, 0 (cid:48)(cid:48) . (cid:48)(cid:48) , and 2 (cid:48)(cid:48) ) that were searched for with wavdetect on variouspixel scales (1–2, 1–4, 1–8, and 1–16 pixel scales, correspond-ing to 0 (cid:48)(cid:48) . (cid:48)(cid:48) .
5, 1 (cid:48)(cid:48) –8 (cid:48)(cid:48) , and 2 (cid:48)(cid:48) –32 (cid:48)(cid:48) , respectively) to mitigate thespatial variation of the PSF on the detector (Broos et al. 2010).The obtained wavdetect source lists were then merged by pairsusing the source matching IDL program match_xy in the Toolsfor ACIS Review & Analysis (TARA) package (Broos et al.2010).The events were extracted using the ACIS Extract (AE) soft-ware package (version ae2018june14; Broos et al. 2010) fol-lowing the multi-pass validation procedure (revision 1099 of https: // cxc.cfa.harvard.edu / ciao / http: // personal.psu.edu / psb6 / TARA / http: // personal.psu.edu / psb6 / TARA / ae_users_guide.htmlArticle number, page 8 of 12icolas Grosso et al.: a powerful X-ray flare from the Class 0 protostar HOPS 383 . The AE source apertures were centered on sourcepositions with a default PSF fraction of 90% at 1.49 keV, whichwas determined from point source simulations computed by ray-tracing with MARX (Davis et al. 2012). The AE backgroundregions were built to get enough background counts in order tohave the source photometry errors within 3% of the values thatwould be obtained with no uncertainty in the background. Newpositions were determined for isolated sources with an o ff -axisposition lower than 5 (cid:48) using the mean position of the extractedevents. These positions were used to determine the astrometricposition o ff sets between the ObsIDs and a reference catalog. Weused the Gaia data release 2 (DR2) positions and position er-rors (Gaia Collaboration et al. 2016, 2018) that we propagatedwith topcat (Taylor 2019) from the J2015.5 reference epochto the J2017.948 Chandra epoch (using the same zero radial ve-locity and radial velocity error as in the Celestia 2000, Hipparcos& Tycho catalog) and added 2MASS sources without the
Gaia
DR2 best neighbor (Marrese et al. 2019); we excluded a 2MASSsource that was resolved by
Gaia and
Chandra . The new posi-tions were also determined for isolated sources with an o ff -axisposition larger than 5 (cid:48) using the PSF correlation position. Forcrowded sources (PSF fraction < ff sets between the ObsIDs. Appendix C.3: Astrometry
The final X-ray positions are mainly registered on
Gaia
DR2(with 45
Gaia and 12 2MASS sources), that is to say, on theICRS world coordinate system. The position of the X-ray coun-terpart of HOPS 383 converted to the J2000 world coordinatesystem is 05 h m s . ◦ (cid:48) (cid:48)(cid:48) .
432 with a positional errorof 0 (cid:48)(cid:48) .
099 that includes, added in quadrature, the first-ObsID ref-erence o ff set uncertainty of 0 (cid:48)(cid:48) .
029 to the AE positional error of0 (cid:48)(cid:48) .
095 (Eqs. (1)–(3) in Broos et al. 2010).The radio position of JVLA-SE and NW in Figs. 1b and cwere determined from the pixel maximums in Fig. 1 of Galván-Madrid et al. (2015) converted to grayscale, assuming a J2000world coordinate system, with positional errors estimated us-ing
FWHM / (2 √ S / N ), where S / N is the signal-to-noiseratio. For the Spitzer counterpart, we adopted a conservative po-sition error of 1 (cid:48)(cid:48) , corresponding to the matching radius used inMegeath et al. (2012).
Appendix C.4: Timing analysis
The source aperture with a PSF fraction of 96% was usedfor the timing and spectral analysis of HOPS 383. The back-ground is negligible (Appendix C.5). We estimated the countrates versus time ( (cid:99) CR ) from the arrival time ( t i ) of the events( N =
28) with a kernel estimator of the density ( (cid:98) f ; Feigelson& Jogesh Babu 2012) with a constant bandwidth ( h ): (cid:99) CR ( t , h ) = N (cid:98) f ( t , h ) / DTCOR , where
DTCOR = . (cid:98) f ( t , h ) = (1 / hN ) × (cid:80) Ni = K { ( t − t i ) / h } , where weused the Epanechnikov kernel, the inverted parabola defined as K ( y ) = .
75 (1 − y ) with − ≤ y ≤
1. We chose for h : the rule-of-thumb bandwidth, h rot , designed for unimodal distribution http: // personal.psu.edu / psb6 / TARA / procedures / http: // gea.esac.esa.int / archive / http: // / ∼ mbt / topcat / (Eq. (6.9) in Feigelson & Jogesh Babu 2012, and cyan curve inFig 2c); and the optimal bandwidth, h cv , maximizing the cross-validation (Eq. (6.10) in Feigelson & Jogesh Babu 2012, andgreen curve in Fig 2). The latter is ∼ α . This wascomputed from the pilot densities previously obtained with h rot (blue curve in Fig 2c) and h cv (red curve in Fig 2c). Appendix C.5: Spectral analysis
The source-plus-background spectrum was limited to the flaretime interval ranging from the first to the last detected eventand the 0.5–9.9 keV energy range. An energy-dependent aper-ture correction was applied by AE to the ancillary response file,which calibrates the extraction. The background spectrum wasextracted from the full exposure of the first observation in the0.1–10 keV energy range.The simultaneous spectral fitting of the background andsource-plus-background spectra was made with XSPEC (ver-sion 12.10.1) using the C-statistic applied to unbinned data.The background cumulative spectrum was modeled with theAE cplinear model, a continuous piecewise-linear functionwith ten vertices (Broos et al. 2010). The source spectrum wasmodeled with an interstellar medium absorption (Wilms et al.2000) and an isothermal collisional-radiative plasma (Smithet al. 2001) with the typical coronal abundances of pre-mainsequence stars (Güdel et al. 2007) plus an emission line withzero width at 6.4 keV ( TBabs*(vapec+gaussian) ) to whichthe background model, re-scaled to the source aperture, wasadded to fit the source-plus-background cumulative spectrum(data flow diagram in Fig. 10 of Broos et al. 2010). The modelparameters, namely the hydrogen column density, the plasmatemperature and normalization, the line normalization, and theten parameters of cplinear , were first estimated by C-statisticminimization (C-stat = ∼ = cplinear parameters since the backgroundcontribution to the source-plus-background spectrum (0.1 count)was negligible. We performed a Markov chain Monte Carlo(MCMC) (Hogg & Foreman-Mackey 2018) to compute theprobability density functions of the four physical parameters ofthis model using the Jeremy Sanders’ xspec_emcee program that used a Python implementation of the MCMC ensemblesampler with a ffi ne invariance proposed by Goodman & Weare(2010), emcee (version 2.2.1; Foreman-Mackey et al. 2013).We customized xspec_emcee to output subsidiary Markovchains for the acceptance fraction (Foreman-Mackey et al. 2013)and the following physical values: soft and hard fluxes; softand hard absorption-corrected fluxes; and line equivalent widthand counts. We used uniform priors for the physical parameters, N w =
46 walkers that were started clustered around the previousparameter estimates, and N iter = iterations.To test for convergence of the Markov chains, we followedHogg & Foreman-Mackey (2018) by using the integrated au- https: // heasarc.gsfc.nasa.gov / xanadu / xspec / https: // github.com / jeremysanders / xspec_emcee https: // emcee.readthedocs.io / en / v2.2.1 / Article number, page 9 of 12 & A proofs: manuscript no. 38185 a N o r m a li z e d c u m u l a t i v e c o un t s Fe I-XVII
XSPEC best fit with no Gaussian line (C-Statistic=677.0, dof=1309) extracted spectrum (28 counts in 11879 s)source model:TBabs(N H /cm =5.8e+23)*vapec([Z], kT/keV=10.8, EM/cm =1.1e+54)bkg region spectrum (184 counts)bkg region modelscaled bkg model (0.1 counts)source+bkg model1 2 3 4 5 6 7 8 9 10Energy (keV)0.10.00.1 R e s i d u a l s ( m o d e l - d a t a ) E x t r a c t e d c u m u l a t i v e c o un t s b N o r m a li z e d c u m u l a t i v e c o un t s Fe I-XVII
XSPEC best fit with Gaussian line (C-Statistic=671.8, dof=1308) extracted spectrum (28 counts in 11879 s)TBabs(N H /cm =5.0e+23)*vapec([Z], kT/keV=6.4, EM/cm =9.0e+53)gauss([E l =6.4 keV, sigma=0 eV], EW=1.1keV)source model: TBabs*(vapec+gauss)bkg region spectrum (184 counts)bkg region modelscaled bkg model (0.1 counts)source+bkg model1 2 3 4 5 6 7 8 9 10Energy (keV)0.10.00.1 R e s i d u a l s ( m o d e l - d a t a ) E x t r a c t e d c u m u l a t i v e c o un t s Fig. C.1.
X-ray spectrum of HOPS 383 and background determination.
Panels a and b : top panels : source-plus-background extracted cumula-tive spectrum (gold filled area) and XSPEC best-fit model (green), sourcemodel (blue), background region cumulative spectrum (orange) andmodel (red), and scaled background model (pink). The vertical dashed-dotted line is the energy of the 6.4 keV line arising from neutral or low-ionization iron.
Bottom panels : model-minus-data residuals.
Panel b :Absorbed coronal thin emission (cyan) and 6.4 keV Gaussian emissionline (silver filled area). tocorrelation time that we computed for each walker and thenaveraged these estimates for each parameter, and retained thelargest value, τ (top panel of Fig. C.2). The MCMC burn-inphase was supposed to last less than N burn-in = τ =
10 008iterations (Foreman-Mackey et al. 2013), as this was illustratedby the mean of the walkers’ acceptance fractions converging to ∼ ff before producingthe corner plots (Fig. C.3a) from which the parameter me-dian values and 90% confidence interval were determined (Ta-ble C.2). For comparison purposes, Fig. C.4 shows the source-plus-background extracted cumulative counts spectrum (gold),the model for the parameter median values (green), and a ran-dom subset of 200 MCMC models (gray).In this MCMC, the number of e ff ectively independent datapoints is N = ( N iter − N burn-in ) × N w ≈ . × and the relative See Dan Foreman-Mackey’s blog on autocorrelation time estimation(2017 October 16): https: // dfm.io / posts / autocorr / . We followed Hogg & Foreman-Mackey (2018), who are in favorof reporting median and quantiles, which are based on integrals, andagainst mode, which is not.
Fig. C.2.
Burn-in phase of the MCMC samples.
Top panel : estimatorof the integrated autocorrelation time versus the number of iterations.The dashed horizontal lines are the integrated autocorrelation time in aniteration-number unit for the model parameters after 10 iterations with46 walkers; the burn-in phase is lower than ten times the maximum ofthese values. Bottom panel : acceptance fraction versus the number ofiterations. The mean of the walkers’ acceptance fractions (dotted line)converges after the burn-in phase. accuracy is 100 √ τ/ N = . Appendix C.6: Quiescent luminosity
We used the CIAO tool aprates to compute the source countrate in the 2–8 keV energy band during our observations andexcluding the flare time interval, which led to a live time of ∼ ∼ . (cid:48)(cid:48) and corresponding to aPSF fraction of 90% at 1.49 keV. The background aperture wasa ∼ (cid:48)(cid:48) . ∼ (cid:48)(cid:48) annulus centered on the source position where theneighbor source region was excluded with an area of ∼ (cid:48)(cid:48) ,corresponding to a PSF fraction of 0.9% at 1.49 keV. The to-tal counts in source-plus-background and background apertureswere 1 (at 3.0 keV) and 304, respectively. Inside the source-plus-background aperture, we estimated a quiescent source count rateof 0.011 count ks − (0–0.057 count ks − , 90% C.I.).From a tbabs*vapec model with N H , X = . × cm − and a typical plasma quiescent temperature ( kT = . × − erg cm − s − and 2 . × erg s − in the 2–8 keV energyband, respectively. Article number, page 10 of 12icolas Grosso et al.: a powerful X-ray flare from the Class 0 protostar HOPS 383
Table C.2.
Markov chain Monte Carlo results on the flare time interval.
Parameter name Symbol Value UnitMedian 90% C.I.AbsorptionHydrogen column density N H , X cm − Plasma emissionTemperature kT T (a) EM cm − Hard-band intrinsic luminosity (a) L X , intr (2-8 keV) 1.7 0.7–7.2 10 erg s − Soft-band intrinsic luminosity (a) L X , intr (0.5-2 keV) 1.2 0.3–22.7 10 erg s − Full-band intrinsic luminosity (a) L X , intr (0.5-8 keV) 3.0 0.9–29.9 10 erg s − Neutral to low-ionization iron line emissionLine flux Flux(Fe 6.4 keV) 6.1 1.7–14.1 10 − ph cm − s − Equivalent width EW(Fe 6.4 keV) 1.1 0.2–3.1 keVTotal count number N(Fe 6.4 keV) 5.1 1.4–10.8 counts
Notes. ( a ) Assuming a distance of 420 pc (Safron et al. 2015).
Fig. C.3.
Posterior probability and covariance distributions of theflare-emission model parameters.
Panel a : hydrogen column density,plasma temperature, X-ray intrinsic luminosity in the 2 − . × with about 4 . × independent MCMC sam-ples. The dashed and dotted vertical lines in the diagonal plots show themedian value and the 90% confidence interval (C.I.) for each parame-ter. The contours in the other plots are 11.8, 39.3, 67.5, and 86.5% C.I.(corresponding to 0.5, 1, 1.5, and 2 σ levels for a 2D Gaussian) for eachpair of parameters. Panel b : alternative last row with count number fromthe emission line. N o r m a li z e d c u m u l a t i v e c o un t s MCMC with N=4.6e+07 sets ( N / =45174 independant sets) extracted spectrum (28 counts)random subset of 200 MCMC models[bkg]+TBabs*(vapec+gauss) model with median MCMC parameter valuesgauss model with median MCMC parameter value1 2 3 4 5 6 7 8 9 10Energy (keV)0.50.00.5 R e s i d u a l s ( m o d e l - d a t a ) E x t r a c t e d c u m u l a t i v e c o un t s Fig. C.4.
X-ray spectrum of HOPS 383 and Markov chain Monte Carlo(MCMC) models.
Top panel : source-plus-background extracted cumu-lative spectrum (gold) and median MCMC model (green), 6.4 keVGaussian emission line with median MCMC model (cyan), and a ran-dom subset of 200 MCMC models (gray).
Bottom panel : model-minus-data residuals.
Appendix D: SOAR Spartan data
Appendix D.1: Observations
The Spartan observations are listed in Table D.1.
Appendix D.2: Data reduction
An end-to-end pipeline, THELI (version 1.9.5; Schirmer2013), was used for the reduction of the calibration and obser-vation data (Table D.1) obtained with det3, including astrometry https: // / theli / Article number, page 11 of 12 & A proofs: manuscript no. 38185
Table D.1.
Log of SOAR observations with Spartan (proposal ID 2016B-0930, PI: N.G.; observers: K.H. and D.P.).
Filter Wavelength Width Start Time DIT (a)
NDIT (b)
Exposure (c)
FWHM (d) (microns) (microns) (UT) (s) (s) (arcsec) K Notes.
We also obtained standard infrared calibration data: dark exposures and dome flat exposures with the lamp on and o ff to account for thethermal emission from the telescope. ( a ) Detector integration time. ( b ) Number of frame exposures with the target. ( c ) Total exposure of the stackeddithered frames. ( d ) Image quality of the stacked dithered frames. with
SExtractor (Bertin & Arnouts 1996) and
SWarp (Bertinet al. 2002). We wrote a custom IDL program to automaticallydetect and correct, in the dark and flat-field corrected frames,any residual bulk images that were produced by bright stars,close to saturation in the raw frames. In each raw frame, this pro-gram identified any saturated regions centered on bright stars andlooked at these detector positions in the subsequent calibratedframes for any excess above the sky level by fitting a constantlevel or a parabola (based on an F-test). These excesses werethen subtracted or masked when too close to a source. These cal-ibrated and corrected images were registered using the 2MASScatalog and stacked using median values and a pixel scale of 0 (cid:48)(cid:48) . narrow-band filter image, the shocked molecularhydrogen emission has a higher S / N compared to the K -bandimage. We produce for validation a pure H emission image bysubtracting the K -band image scaled by a factor that we obtainby minimizing the residuals on stars in the final image (Navareteet al. 2015). Appendix E: Density model
We started from the protostellar model of HOPS 383, whichis composed of an accretion disk in an infalling envelope withbipolar cavities (Furlan et al. 2016). The fixed parameters in thismodel grid (Table 3 of Furlan et al. 2016) that was computedwith the Monte Carlo radiative transfer code, ttsre (2008 ver-sion; Whitney et al. 2003) were: – the stellar mass ( M (cid:63) = . M (cid:12) ), – the stellar e ff ective temperature ( T (cid:63) = – the disk mass ( M disk = . M (cid:12) ), – the magnetospheric truncation radius of the gas disk ( R trunc = R (cid:63) ), – the dust-disk inner radius (set to the dust sublimation radius, R sub / R (cid:63) = ( T sub / T (cid:63) ) − . = . T sub = – the scale height of the disk at the dust sublimation radius (setto the hydrostatic equilibrium solution; Eq. (35) of D’Alessioet al. 1998), – the mean molecular mass per hydrogen nucleus ( µ = . – the radial ( A = .
25) and vertical ( B = .
25) exponent indisk density law, – the fractional area of the hot spots on the star ( f spot = . – the envelope outer radius ( R env =
10 000 au), – the centrifugal radius of the envelope ( R c = R disk ), – the exponent of the cavity polynomial shape ( b cav = . – the vertical o ff set of the cavity wall ( z cav = – the stellar radius ( R (cid:63) = . R (cid:12) ), https: // gemelli.colorado.edu / ∼ bwhitney / codes / codes.html – the disk outer radius ( R disk = – the envelope (gas and dust) density at 1000 au ( ρ = . × − g cm − ) and, as ρ ∝ ˙ M env , the envelope infall rate( ˙ M env = . × − M (cid:12) yr − ), – the cavity opening angle ( θ = ◦ ), – the stellar luminosity ( L (cid:63) = L (cid:12) ), – the total (star and accretion) luminosity ( L tot = s × . L (cid:12) = . L (cid:12) , where s = .
84 is the luminosity scaling factor tomatch the observed SED flux, Eq. (3) of Furlan et al. 2016), – the inclination angle ( i = . ◦ ; corresponding in the modelgrid to cos i = .
35 with a sampling step of 0.1), and – the foreground cloud extinction ( A V , cloud =
14 mag; corre-sponding to N H , cloud = × cm − , as N H , cloud / A V , cloud = . × cm − mag − in Furlan et al. 2016).To be consistent with the luminosity scaling factor ( s ), weintroduced s (cid:48) , a scaling factor for the disk-to-star accretion rate( ˙ M disk ). Indeed, the above values of R (cid:63) (6.61 R (cid:12) ), L (cid:63) (10 L (cid:12) ),and L tot = L (cid:63) + L acc (30.2 L (cid:12) ) correspond in the model grid to˙ M disk = . × − M (cid:12) yr − , since L acc ∝ ˙ M disk , we set s (cid:48) = . M disk = . × − M (cid:12) yr − .The envelope density was given by Eqs. (1) and (2) of Whit-ney et al. (2003). We note that the bipolar cavities in Furlan et al.(2016) are free of dust (and gas); for consistency, we filled themin with gas, using a constant density ( n H = × cm − ; Whit-ney et al. 2003).The disk density was given by Eq. (3) of Whitney et al.(2003) and normalized to M disk from R trunc to R disk using a spher-ical outer boundary (Fig. 3b this work and Fig. 2b of Whitneyet al. 2003). We added accretion funnels at R trunc , which weredust-free as located inside the dust sublimation radius, assumingfor simplicity a dipole geometry for the magnetic field (Hart-mann et al. 1994), and using f spot to compute a latitude bandbetween the smallest and largest magnetic loops at the stellarsurface. The accretion-funnel density was given by Eq. (9) ofHartmann et al. (1994).The values of N H along the line-of-sight were computed bynumerically integrating these density equations. In this modelgrid, the gas-to-dust ratio is R =
100 and the protostellar dustopacity in the visible is κ ext , V = . × cm g − (Fig. 6 ofOrmel et al. 2011, (ic-sil,gra) at 0.3 Myr), leading to: N H / A V = ln 10 × (1 + R ) / (2 . µ m H κ ext , V ) = . × cm − mag − where m H is the hydrogen mass. Therefore, the observed N H , X corre-sponds to A V ∼1500 mag.