Revealing a Centrally Condensed Structure in OMC-3/MMS 3 with ALMA High Resolution Observations
aa r X i v : . [ a s t r o - ph . GA ] F e b Draft version February 17, 2021
Typeset using L A TEX twocolumn style in AASTeX63
Revealing a Centrally Condensed Structure in OMC-3/MMS 3with ALMA High Resolution Observations
Kaho Morii,
1, 2,3
Satoko Takahashi,
4, 5,6 and Masahiro N. Machida Department of Astronomy, Graduate School of Science, The University of Tokyo, 7-3-1 Hongo, Bunkyo-ku, Tokyo 113-0033, Japanemail: [email protected] Division of Science, National Astronomical Observatory of Japan, 2-21-1 Osawa, Mitaka, Tokyo 181-8588, Japan Department of Earth and Planetary Sciences, Faculty of Sciences, Kyushu University, Fukuoka 819-0395, Japan Joint ALMA Observatory, Alonso de C´ordova 3107, Vitacura, Santiago, Chile NAOJ Chile, National Astronomical Observatory of Japan, Alonso de C´ordova 3788, Office 61B, Vitacura, Santiago, Chile, 7630492 Department of Astronomical Science, School of Physical Sciences, The Graduate University for Advanced Studies, SOKENDAI, Mitaka,Tokyo 181-8588, Japan (Received; Revised; Accepted)
Submitted to ApJABSTRACTUsing the Atacama Large Millimeter/submillimeter Array (ALMA), we investigated a peculiar mil-limeter source MMS 3 located in the Orion Molecular Cloud 3 (OMC-3) region in the 1.3 mm con-tinuum, CO ( J =2–1), SiO ( J =5–4), C O ( J =2–1), N D + ( J =3–2), and DCN ( J =3–2) emissions.With the ALMA high angular resolution ( ∼ . ′′ . ′′ × . ′′
32 (P.A.=0.22 ◦ ). The peak position coincideswith the locations of previously reported Spitzer /IRAC and X-ray sources within their positionaluncertainties. We also detected an envelope with a diameter of ∼ ◦ ) in the C O( J =2–1) emission. Moreover, a bipolar outflow was detected in the CO ( J =2–1) emission for the firsttime. The outflow elongates roughly perpendicular to the long axis of the envelope detected in theC O ( J =2–1) emission. Compact high-velocity CO gas in the (red-shifted) velocity range of 22–30km s − , presumably tracing a jet, was detected near the 1.3 mm continuum peak. A compact and faintred-shifted SiO emission was marginally detected on the CO outflow lobe. The physical quantities ofthe outflow in MMS 3 are relatively smaller than those in other sources in the OMC-3 region. Thecentrally condensed object associated with the near-infrared and X-ray sources, the flattened envelope,and the faint outflow indicate that MMS 3 harbors a low mass protostar with an age of ∼ yr. Keywords: stars: formation – stars: protostars – ISM: individual objects (OMC-3, MMS 3) – astro-chemistry – radio continuum: stars INTRODUCTIONIt is crucially important to identify how and whenstars form in their natal clouds to understand the wholepicture of star formation. However, it is difficult to iden-tify newborn stars in gravitationally collapsing cloudcores because the growth timescale of protostars is veryshort. In a gravitationally collapsing cloud, the firsthydrostatic core (or the first Larson core, hereafterthe first core, Larson 1969; Masunaga & Inutsuka 2000)forms before protostar formation. After protostar for-mation, the first core (remnant) remains around theprotostar and becomes a rotationally supported disk(Bate 1998; Machida et al. 2010). The first core rem- nant or rotationally supported disk drives an outflowthat determines the star formation efficiency (Tomisaka2002; Machida & Matsumoto 2012). Thus, the first coreand its remnant play a key role in the early star for-mation phase (Bate 1998; Masunaga & Inutsuka 2000;Machida et al. 2010). However, the timescale for thefirst core is as short as ∼ –10 yr (Larson 1969;Masunaga & Inutsuka 2000; Saigo & Tomisaka 2006).The protostar phase, during which parcels of gas con-tinue to accrete onto the protostar or circumstellar diskfrom the infalling envelope, lasts for ∼ yr. Thus, thedetection rate for the first core is 1/100–1/1000 ((10 –10 yr)/10 yr). Although the detection rate for veryyoung protostars is also low, these objects provide usefulinformation for understanding the star formation pro-cess.The Orion Molecular Cloud 3 (OMC-3) star-formingregion is located d = 392 pc from the Sun (Tobin et al.2020) and is one of the best sites for investigating theearly phase of star formation. In OMC-3, there are bothprestellar and protostellar sources MMS 1–10, whichare identified in Chini et al. (1997). Takahashi et al.(2013) reported that cores are formed by fragmenta-tion of filaments. As seen in Takahashi et al. (2013),the main filament in northern OMC-3 extends from thenorthwest to the southeast (P.A. = 135 ◦ ). The pro-tostellar sources in this region contain Class 0 and Iprotostars (e.g., Chini et al. 1997; Nielbock et al. 2003;Takahashi et al. 2008; Furlan et al. 2016). In our previ-ous studies, we reported the detailed properties of theClass 0 protostellar sources MMS 5 and MMS 6 and as-sociated compact and collimated outflows, which wereobserved by the Submillimeter Array (SMA) and ALMA(Takahashi et al. 2008, 2009; Takahashi & Ho 2012;Takahashi et al. 2012, 2013, 2019; Matsushita et al.2019).In this paper, we focus on another millimeter source,MMS 3, that is identified as SMM 4 in Takahashi et al.(2013), located around the center of the OMC-3 fila-ment (for details, see Figure 3 of Takahashi et al. 2013).The average number density of MMS 3 estimated fromthe SMA 0.85 mm dust continuum observation is n =(4 ± × cm − , which is one order of magnitude lowerthan the other very young protostellar sources in OMC-3 (Takahashi et al. 2013). The peak flux of MMS 3 is4–20 times weaker than the other protostellar sourcesin this region. In addition, to date, neither outflow norjet has been observed in MMS 3 (Reipurth et al. 1999;Aso et al. 2000; Williams et al. 2003; Takahashi et al.2008; Tanabe et al. 2019; Feddersen et al. 2020). More-over, no 9 mm continuum emission has been de-tected around MMS 3 in the recent VLA observations(Tobin et al. 2020). Thus, it appears that MMS 3is in the prestellar phase before gravitational contrac-tion has begun. Nevertheless, X-ray emission was de-tected by the Chandra X − Ray Observatory (TKH10, Tsuboi et al. 2001). Furthermore, an infrared sourcewas also detected in MMS 3 in both the
Spitzer ,with wavelengths longer than the IRAC 4.5 µ m bands,and the Herschel/P ACS/SP IRE bands (HOPS 91in Furlan et al. 2016, Megeath 2442 in Megeath et al.2012). Thus, although we could not detect a centralcondensation in MMS 3 in the SMA observation, a pro-tostar has already been born within MMS 3. In thisstudy, we investigate this peculiar object MMS 3 with ALMA observations to elucidate the early phase of starformation.This paper is structured as follows. In §
2, we ex-plain the observation method and data reduction pro-cess. The results are shown in §
3. We discuss the evolu-tionary stage of MMS 3 in § § OBSERVATIONSThe observations were made in Cycle 3(2015.1.00341.S, PI. S. Takahashi), 2016 June 30, July12, 17, and 19 (Atacama Compact 7m-array, hereafterthe ACA), 2016 January 29 (12 m array in the compactconfiguration C36-1, hereafter ALMA low resolution),and 2016 September 18 and 19 (12 m array in theextended configuration C36-4, hereafter ALMA highresolution) using the 1.3 mm band (Band 6). The phasecenter was set at R.A. (J2000) = 5 h m . s − ◦ ′ . ′′ λ and 365 k λ (ACA), 8.6 k λ and 240 k λ (ALMA lowresolution), and 9.0 k λ and 2400 k λ (ALMA high res-olution). These were insensitive to structures extendedmore than 21 ′′ for the ACA, and 13 ′′ for the ALMAlow and high resolution at the 10% level of the totalflux density (ALMA Cycle 3 Technical handbook). Ourspectral setups include CO ( J = 2–1), SiO ( J = 5–4),C O ( J = 2–1), N D + ( J = 3–2), and DCN ( J = 3–2)emissions. These were observed with velocity resolu-tions of 0.37 km s − , 0.39 km s − , 0.048 km s − , 0.046km s − , and 0.39 km s − , respectively. Line-free chan-nels corresponding to effective bandwidths of 695 MHz(ACA), 800 MHz (ALMA low resolution), and 782 MHz(ALMA high resolution) were allocated for imaging thecontinuum emission. Calibration of the raw visibilitydata was performed by the ALMA observatory with thestandard calibration method using the Common As-tronomy Software Application (CASA; McMullin et al.2007) versions 4.6.0 for ACA and ALMA low resolutiondatasets and 4.7.0 for ALMA high resolution datasets.After calibration, clean images were made using theCASA task “clean”. Natural weightings were used forthe final images. Velocity widths of 0.1 km s − (C Oand N D + ), 0.5 km s − (DCN), and 1.0 km s − (CO andSiO) were used to produce the image cubes. Data com-bination from the two different arrays were performed inthe image base for the 1.3 mm continuum, C O, N D + ,and DCN using the CASA task “feather”, and in the uv base for CO using the CASA task “concat”. An imageproduced from the ALMA low resolution datasets wasonly presented for the SiO emission . No primary beamcorrection has been applied for the presented images inthis paper, while primary beam corrected CO and C Oflux values were used for deriving physical parameters.The resulting synthesized beam sizes and 1 σ rms noiselevels are listed in Table 2. RESULTS3.1. Dust Continuum EmissionFigure 1 shows the 1.3 mm dust continuum emis-sion. In Figure 1 a , we can confirm a filamentary struc-ture of the dust distribution in the east–west direc-tion. Using two-dimensional (2D) Gaussian fitting, thesize of the filamentary structure is estimated to be(20 ′′ ± . ′′ × (7 . ′′ ± . ′′ ∼ × ∼ ◦ . The peak position of the dust con-tinuum emission is R.A. = 05 h m . s
185 and Dec = − ◦ ′ . ′′ ′′ from the peak po-sition obtained from the previous SMA observations inthe 850 µ m band (Takahashi et al. 2013). We detecteda more extended structure in the ACA observations(Figure 1(a)). This is because the maximum recover-able size of the ACA observations is 1.6 times largerthan that in previous SMA observations. The ACAobservations also achieved ∼ Theseimproved sensitivities allow the detection of more ex-tended and faint emissions in the ALMA images. Withinthe extended structure, for the first time, we detecteda compact and centrally condensed millimeter sourcewith the ALMA observations (Figures 1(b) and (c)).The peak position of the 1.3 mm continuum emissionin Figure 1(c) is R.A. = 05 h m s .
929 and Dec = − ◦ ′ . ′′ µ m, Furlan et al. 2016) andX-ray (Tsuboi et al. 2001) sources within their posi-tional errors. Thus, both the infrared and X-ray sourcesare associated with the centrally condensed structure.Using 2D Gaussian fitting, the compact 1.3 mm con-tinuum structure in Figure 1(b) was measured to be(1 ′′ . ± ′′ . × (1 ′′ . ± ′′ . ×
430 au. The size of the very compact No data combination was performed since there was no totalflux difference in the image between the ACA and ALMA lowresolution data sets. In order to directly compare the sensitivity between the ALMA1.3 mm data and the SMA 0.85 mm data, the beam surface areadifference and the spectral index between the two wavelengths( β =1.6 from Johnstone & Bally 1999) were taken into account. source detected with the ALMA high resolution (Fig-ure 1(c)) is (0 ′′ . ± ′′ . × (0 ′′ . ± ′′ . ∼ ×
130 au. The de-convolved size of MMS 3 reported by Tobin et al. (2020)is 0 ′′ . × ′′ .
20. Thus, our result is in agreement withTobin et al. (2020) within a factor of 2.Using the ratio between the major and minor axes ofthe 1.3 mm continuum emission, we can estimate the in-clination angle of the disk i with respect to the planeof the sky. The angle calculated from the ALMA highresolution image (Figure 1c) is i = 44 ◦ . On the otherhand, the inclination angle derived from Tobin et al.(2020) is i = 40 ◦ . Thus, our results are consistent withTobin et al. (2020) within ∼ ◦ .Using 2D Gaussian fitting, we also obtained the totalflux and peak flux for each image in Figure 1, summa-rized in Table 3. A structure comparable to the beamsize was detected in the ALMA high resolution (Figure 1and Table 3). Assuming optically thin dust thermalemission and an isothermal dust temperature, we esti-mate the dust mass with the total 1.3 mm flux, F . ,using M dust = F . d κ . B . ( T dust ) , (1)where d , κ . , and B . ( T dust ) are the distanceto the source (392 pc, Tobin et al. 2020), the absorp-tion coefficient for the dust per unit mass, and thePlanck function as a function of the dust temperature T dust , respectively. We adopt κ . = 0 . g − fromthe dust coagulation model of the MRN (Mathis et al.1977) distribution with thin ice mantles computed byOssenkopf & Henning (1994). Assuming a dust temper-ature of T dust = 20 and 40 K, we obtain the dust mass,column density, and number density values shown in Ta-ble 4. The gas masses estimated from Figure 1( c ) are(1.6 ± × − M ⊙ for T dust = 20 K and (7.0 ± × − M ⊙ for T dust = 40 K, in which a gas-to-dust massratio of 100 is assumed.3.2. Dense Gas TracersFigure 2 shows the integrated intensity maps for C O( J = 2–1) and N D + ( J = 3–2) emissions overlaid withthe 1.3 mm dust thermal emission. We detected a cen-trally condensed C O ( J = 2–1) emission structurearound the 1.3 mm continuum emission, presumablytracing the dense gas envelope around MMS 3. The In this paper, we define the term ‘envelope’ as the circumstellarmaterial surrounding the 1.3 mm compact continuum emissiontraced by the dense gas tracers such as C O and DCN. We usethe term ‘core’ to represent a large self-gravitating object, whichcorresponds to a molecular cloud core, as used in previous studies.
Table 1.
Observation Parameters
Parameters ACA ALMA low resolution ALMA high resolutionObserving date (YYYY-MM-DD) 2016-06-30 -07-12, -17 and -19 2016-01-29 2016-09-18 and -19Number of antennas 10 40 48Primary beam size (arcsec) 46 27 27PWV (mm) – 2.6 ∼ ∼ a ∼ ∼ ∼
21 and ∼ b J0607-0834 and J0542-0913 J0541-0541 J0607-0834Total continuum bandwidth; USB+LSB (MHz) 695 800 782Projected baseline ranges (k λ ) 5.5–365 8.6–240 9.0–2400Maximum recoverable size (arcsec) c
21 13 13Total on-source time (minutes) 16 4 16 a Antenna-based phase differences for the bandpass calibrators. b The phase calibrator was observed every 8 minutes. c Our observations were insensitive to emissions more extended than this size scale at the 10% level of the total flux density (ALMA Cycle 3Technical Handbook).
Table 2.
Image Parameters
Image Synthesized beam size (P.A.) RMS noise level Velocity width Figurearcsec × arcsec (deg) km s − ACA Continuum 8.2 × − ] – 1(a)ACA + ALMA low resolution continuum a × − ] – 1(b) 2 3ALMA high resolution continuum 0.22 × − ] – 1(c)ACA + ALMA low resolution C O( J =2–1) a × − ] 0.3 4ACA + ALMA low resolution C O( J =2–1) a × − km s − ] 4.4 2 3 5ACA + ALMA low resolution N D + ( J =3–2) a × − km s − ] 4.4 2ACA + ALMA low resolution DCN ( J =3–2) a × − km s − ] 9.5 3ALMA low resolution SiO( J =5–4) 1.9 × − km s − ] 3.0 5ACA + ALMA low resolution CO ( J =2–1) b × − km s − ] c
16, 22, 8.0 c a Data combination was done in the image base using the CASA task “feather”. b Data combination was done in the UV visibility base using the CASA task “concat”. c The different rms noise levels are due to the fact that the integrated intensity maps used different velocity ranges. A detailed description is given inthe caption of Figure 5
Table 3.
The 1.3 mm Continuum Source Parameters
Data Total flux (mJy) Peak flux (mJy beam − ) Deconvolved size (arcsec × arcsec, deg)ACA 310 ±
12 64 ± ± . × . ± . , ± . ± ± ± × ± ± ± ± ± × ± ± . (a) (b) (c)
100 au
Figure 1.
Dust continuum emission (color and black contours) obtained in the 1.3 mm band from (a) ACA 7 m array, (b)ACA + ALMA low resolution, and (c) ALMA high resolution. The black contours show 10, 12, and 14 σ (1 σ = 4.4 mJy beam − )for panel (a), 10, 15, 20, 25, and 30 σ (1 σ = 0.46 mJy beam − ) for panel (b), and 5, 7, 9, and 11 σ (1 σ = 0.093 mJy beam − )for panel (c). The white ellipse in the bottom left corner represents the synthesized beam size in each image. The black line inthe bottom right corner of each panel indicates the spatial scale. The cross in panel (a) corresponds to the positions of MMS3 identified by the SMA 850 µ m continuum emission (Takahashi et al. 2013). The circles and crosses in panels (b) and (c)correspond to the position of the infrared (HOPS 91, Furlan et al. 2016)) and X-ray sources (Tsuboi et al. 2001), respectively.The sizes of the circles and crosses indicate their positional uncertainties. Table 4.
Physical Parameters of Centrally Condensed Structure
Data Temperature (K) Gas mass (M ⊙ ) a Column density (10 cm − ) b Number density (cm − ) b ACA 20 1.1 ± ± ± × ACA 40 0.47 ± ± ± × ACA + ALMA low resolution 20 0.11 ± ± ± × ACA + ALMA low resolution 40 (4.7 ± × − ± ± × ALMA high resolution 20 (1.6 ± × − ± ± × ALMA high resolution 40 (7.0 ± × − ± ± × a The errors for the mass are obtained by propagating only the errors for the measured flux densities. b The errors for the column density and number density are obtained by propagating the errors for the measured flux densities anddeconvolved sizes.
MMS 2MMS 3
Figure 2. C O ( J = 2–1) and N D + ( J = 3–2) emissions overlaid on the 1.3 mm continuum emission. The color shows theintegrated intensity map for C O. The gray and white contours are the integrated intensity for C O and N D + , respectively.The contours for C O start at 7 σ with an interval of 3 σ (1 σ = 29 mJy beam − km s − ). The contours for N D + correspondto 3, 4, 5, and 6 σ (1 σ = 34 mJy beam − km s − ). These emissions are integrated over the range v LSR = 9–13.4 km s − . Theblack contours correspond to the 1.3 mm continuum emission (ACA + ALMA low resolution), in which the contours start at10 σ with an interval of 5 σ (1 σ = 0.46 mJy beam − ). The ellipse in the bottom left corner is the beam size. The spatial scaleis indicated in the bottom right corner. The ALMA primary beam size is denoted by a black dashed circle. long-axis length of the C O emission estimated by 2DGaussian fitting is ∼ ◦ .The N D + ( J = 3–2) emission distribution aroundMMS 3 extends from the northwest to the southeastand is elongated to the east, which is consistent withlarge-scale filaments in the OMC-3 region observed inthe millimeter and submillimeter continuum emission(Chini et al. 1997; Johnstone & Bally 1999). The ve-locity width of the N D + emission is relatively narrow( δv FWHM ∼ − ).We detected a DCN ( J = 3–2) line in the spectralwindow of SiO ( J = 5–4) and produced the image pre-sented in Figure 3. The DCN emission shows a centrallycondensed structure and seems to roughly overlap withthe C O ( J = 2–1) emission. The long-axis length ofthe DCN emission estimated by 2D Gaussian fitting is ∼ O emission. The local peak of the DCN emission is shifted ∼ ′′ . O emissions. A tem-perature in the range of 10 to 80 K is required for themain formation path of DCN molecules (Turner 2001;Salinas et al. 2017). Hence it is natural that DCN tracesthe dense gas envelope except for the region very closeto the protostar where the gas temperature should beas high as T &
100 K.Figure 4 shows the channel maps for C O ( J = 2–1).In the velocity range of v LSR = 9.7 to 10.3 km s − , theC O emission shows extended complex structures. Wecan see two emission peaks located at ∼ ′′ southeastand ∼ ′′ southwest with respect to the 1.3 mm contin-uum peak. Around a velocity of v LSR = 10.6 km s − ,the C O emission is elongated from the southwest tothe northeast across the 1.3 mm continuum emissionpeak. The C O emission is distributed in the north
Figure 3. C O ( J = 2–1) and DCN ( J = 3–2) emissions overlaid on the 1.3 mm continuum emission. The gray and blackcontours are the same as in Figure 2. The white contours are the integrated intensity of DCN. The contours correspond to 3, 6,9, and 12 σ (1 σ = 23 mJy beam − km s − ). These emissions are integrated over the range v LSR = 6.2–15.7 km s − . The spatialscale is indicated in the bottom right corner. The ellipse in the bottom left corner is the beam size. The ALMA primary beamsize is denoted by a black dashed circle. part toward the 1.3 mm continuum peak around a ve-locity of v LSR = 10.9 km s − . The extended C O emis-sion in the velocity range v LSR = 10 . . − showsa butterfly-like structure particularly in the northwestdirection with respect to the 1.3 mm continuum peak.The distribution is presumably related to the outflowcavity (blue-shifted gas) observed in the CO emission(see § v LSR = 11.2 km s − . In the range of v LSR = 11.5 to 12.1 kms − , the emission is condensed around the 1.3 mm con-tinuum peak and shows an elongated structure along thenortheast to the southwest. The elongation direction isalmost perpendicular to the OMC-3 filament traced byN D + (Figure 2). A compact C O emission associatedwith the 1.3 mm peak remains in the velocity range of v LSR = 12.1 to 12.7 km s − .In this study, we adopted a systemic velocity for MMS3 of 11.2 km s − , derived from the optically thin C O( J = 2–1) and N D + ( J = 3–2) emissions. This velocity corresponds to the central velocity where we see a dipin these two lines. The dip is confirmed over all of thearea of both lines because the ambient gas associatedwith the large-scale structure is resolved out due to thefiltering effect of the interferometric observations. Thesame value of the systemic velocity was confirmed in theoptically thin H CO + ( J = 1–0) and N H + ( J = 1–0)emissions in single-dish observations (Ikeda et al. 2007;Tatematsu et al. 2008).Assuming optically thin C O emission and local ther-modynamic equilibrium (LTE), the envelope columndensity ( N H ) is derived as N H = X − O (cid:18) h π Sµ (cid:19) (cid:18) kT ex hB + 13 (cid:19) exp (cid:18) E u kT ex (cid:19) Z T B dv, (2)where X C O , S , µ , T ex , E u , and T B are the C Oto H abundance ratio, the line strength, the relevantdipole moment, the excitation temperature, the energyof upper level above ground, and the brightness temper-ature in units of K, respectively (Cabrit & Bertout 1992;Mangum & Shirley 2015; Feddersen et al. 2020). Here,we adopted X C O of 1 . × − (Frerking et al. 1982), Figure 4.
Channel maps for C O ( J = 2–1). In each panel, contour levels of 3, 5, 10, 15, 20, 25 σ (1 σ = 16 mJy beam − )are plotted. The white symbol ‘+’ corresponds to the peak position of the 1.3 mm continuum emission. The central velocity isshown in the upper left corner of each panel. The ellipse in the bottom left corner of each panel is the beam size. The spatialscale and the ALMA primary beam (white dashed circle) are denoted in the bottom left panel. Sµ of 0.02 Debye , and E u /k of 15.8 K. We assumedthat the excitation temperature is T ex =20 K. Then, us-ing the estimated N H , the envelope mass M H can beestimated as M H = µ H m H Ω d N H , (3)where Ω is the total solid angle, µ H = 2 . d is thesource distance. The estimated envelope column densityand mass using the flux more than 7 σ are N H = 1 . × cm − and M H = 0 .
89 M ⊙ , respectively.3.3. Outflow and Jet Tracers3.3.1. Morphology
Figure 5 presents the integrated intensity maps forCO ( J = 2–1) and SiO ( J = 5–4). The CO emissionwas detected in the velocity range of v LSR = − − with greater than 3 σ in the channel maps.The detected emission consists of two components. Thefirst is in the low- and mid-velocity ranges of − − ≤ v LSR − v sys ≤
22 km s − , which traces theoutflow cavity in both the blue- and red-shifted compo-nents. The projected size of the blue- and red-shiftedcomponent is 5800 au and 4600 au, respectively. Theposition angle of the CO outflow cavity is ∼ ◦ , andperpendicular to the long axis of the dense envelope ob-served in the C O emission. The second component isthe high-velocity component, which is detected in thevelocity range of 22 to 30 km s − with respect to thesystem velocity (i.e., red-shifted). The detected emis-sion shows a compact structure associated with the 1.3mm continuum peak, which may correspond to a re-cently ejected jet. In our observations, no blue-shiftedhigh-velocity component was detected.In addition to the CO emission, a compact SiO ( J =5–4) emission was marginally detected with the 5 σ emis-sion peak in the integrated intensity map ( v LSR − v sys ≃ − ). The compact SiO emission was located1800 au east from the 1.3 mm continuum emission peakposition and gravitationally unbound. . The SiO emis-sion is located in the outflow region traced by CO, whileits velocity is not high. Thus, it is expected that theSiO emission originates from outgoing material, such asa low-velocity outflow. Alternatively, this could be at- The infall v inf and Keplerian v Kep velocity are described as v inf = (2 GM/r ) / and v kep = ( GM/r ) / , respectively. As-suming a protostellar mass of 0.1 M ⊙ and the distance from theprotostar of 1800 au, the velocities are estimated to be an orderof 0.1 km s − considering the inclination angle of 45 ◦ . The veloc-ity of the detected SiO emission exceeds both the Keplerian andinfall velocity. tributed to an outflow-envelope interaction. More ob-servations are necessary to identify the origin of the SiOemission. 3.3.2. Outflow Parameters
To investigate the outflow properties, we derive out-flow physical parameters from equations (2) and (3) as-suming LTE and an optically thin condition, in whichX C O in equation (2) is replaced by X CO . Here, weadopted the CO to H abundance ratio X CO of 10 − (Frerking et al. 1982) and E u /k of 16.59 K. We alsoassumed that the excitation temperature is T ex =20K (Aso et al. 2000). We only use the pixels above3 σ in a given channel. To obtain the total outflowmass, the blue-shifted and red-shifted CO emissionswere integrated separately. The blue-shifted M flow , b and red-shifted M flow , r outflow mass are estimated to be M flow , b = 9 . × − M ⊙ and M flow , r = 1 . × − M ⊙ ,respectively.The intrinsic outflow velocity ( v flow ) and outflowlength ( l flow ) can be calculated using the observed out-flow velocity ( v obs ) and outflow length ( l obs ) as v flow = v obs / cos i and l flow = l obs / sin i , where i is the inclinationangle of the disk.We adopt an inclination angle of the disk i = 45 ◦ .Considering the inclination angle, the blue-shifted andred-shifted maximum outflow velocities ( v max , b and v max , r ) are estimated to be v max , b = 25 km s − and v max , r = 35 km s − , respectively.The outflow momentum P flow and kinetic energy E flow are defined as P flow = M flow v flow and E flow = M flow v /
2. The momentum and energy of the blue-shifted outflow are P flow , b = 8 . × − M ⊙ km s − and E flow , b = 1 . × erg, respectively, while those of thered-shifted outflow are P flow , r = 1 . × − M ⊙ km s − and E flow , r = 3 . × erg, respectively.The outflow time derivative quantities such as themass loss rate, ˙ M flow = M flow t dyn , (4)the outflow momentum flux, F flow = P flow t dyn , (5)and the outflow kinetic luminosity, L kin = E flow t dyn , (6)are often used to evaluate the outflow activity, where t dyn is the outflow dynamical timescale estimated fromthe outflow length l flow and the maximum velocity v max as t dyn = l flow /v max . The dynamical timescales for the0 Figure 5.
Integrated intensity map for C O ( J = 2–1), CO ( J = 2–1), and SiO ( J = 5–4) emissions. The grayscale andgray contours correspond to C O integrated from 9 to 13.4 km s − . The contour levels are 10, 13, and 16 σ (1 σ = 29 mJybeam − km s − ). The blue, red, and white contours are the integrated intensity of the CO emission, in which the emission isintegrated from v LSR = − − (blue), from v LSR =11.2 to 33.2 km s − (red), and from v LSR =33.2 to 41.2 km s − (white), respectively. The contour levels correspond to 5, 7, 10, 15, 20, 25, and 30 (1 σ = 150, 200, and 28 mJy beam − km s − ).The SiO emission integrated from 11.2 to 15.2 km s − is indicated by the black contours, in which the contour levels are 4 σ and 5 σ (1 σ = 2.9 mJy beam − km s − ). The spatial scale is indicated in the bottom right corner. The black filled ellipse in thebottom left corner is the beam size. The ALMA primary beam size is denoted with a black dashed circle. blue-shifted and red-shifted outflow are estimated to be t dyn , b = 1800 yr and t dyn , r = 1000 yr, respectively. Themass loss rate, outflow momentum flux and kinetic lu-minosity for the blue-shifted component can be esti-mated to be ˙ M flow , b = 5 . × − M ⊙ yr − , F flow , b =4 . × − M ⊙ km yr − , and L kin , b = 6 . × − L ⊙ ,respectively. Those derived for the red-shifted com-ponent are ˙ M flow , r = 1 . × − M ⊙ yr − , F flow , r =9 . × − M ⊙ km yr − and L kin , r = 2 . × − L ⊙ , re-spectively.To understand the evolutionary stage and the sta-tus of the outflow in MMS 3, in Figure 6, wecompare the outflow size, mass, force (or momen-tum flux), and bolometric luminosity of MMS 3 withthose of other outflows associated with Class 0 andClass I sources in the OMC-2/3 (Takahashi et al. 2008;Takahashi & Ho 2012; Furlan et al. 2016; Tanabe et al.2019; Feddersen et al. 2020) and Taurus star-forming re- gions (Hogerheijde et al. 1998). In the figure, comparedto other developed outflow associated with most of theintermediate-mass protostellar sources in the OMC-2/3region, the outflow detected in MMS 3 is less massiveand smaller. Rather, the outflow in MMS 3 shares a sim-ilar nature to low-mass protostellar outflows detected inthe Taurus star-forming region (c.f. Hogerheijde et al.1998).In our observations, we detected an outflow inMMS 3, which was not identified in the previousmolecular outflow survey in this region (Aso et al.2000; Williams et al. 2003; Takahashi et al. 2008;Tanabe et al. 2019; Feddersen et al. 2020). This impliesthat the outflow in MMS 3 is compact and significantlyfaint, as it would not be detected in a survey type ob-servation because of the low sensitivity and low angularresolution. It should be noted that the angular resolu-tion of this observation ( ∼ ′′ ) is better than those of1past studies (e.g., ∼ ′′ for Takahashi et al. 2008 andTanabe et al. 2019 and ∼ ′′ for Feddersen et al. 2020).Also, the sensitivity of this study is about seven timesbetter than in previous studies. EVOLUTIONARY STAGE OF MMS 34.1.
Observational Characteristics of MMS 3
The target of our study, MMS 3, has been recog-nized as a mysterious object in past studies because var-ious observations of MMS 3 have exhibited both prestel-lar and protostellar natures. The peak (sub)millimeterflux for MMS 3 is 3–20 times weaker than other proto-stellar sources observed in the OMC-3 region, such asMMS 2 (SMM 3), MMS 5 (SMM 6), MMS 6 (SMM7), and MMS 7 (SMM 11), while it is comparable toprestellar sources such as MMS 4 (SMM 5) and SMM12 (see Figure 3 of Takahashi et al. 2013). The num-ber density of hydrogen molecule of MMS 3 within D ∼ n ∼ (4 ± × cm − , which is at leastone order of magnitude lower than those of the proto-stellar sources in this region. Thus, it has been con-sidered that the relatively low density (or low conden-sation) of MMS 3 is consistent with neither outflownor jet detection (Reipurth et al. 1999; Takahashi et al.2008; Tanabe et al. 2019; Feddersen et al. 2020) becauseMMS 3 seems to be in the prestellar phase. Neverthe-less, both near-infrared and X-ray sources were detectedwithin MMS 3, which is strong evidence of the existenceof a protostar (Tsuboi et al. 2001; Furlan et al. 2016).The spectral energy distribution measured in the near-to far-infrared wavelengths suggests that MMS 3 is aClass 0 source. These observational characteristics havemotivated us to study the evolutionary stage of MMS3 through direct imaging with high angular resolutionand high sensitivity observations.In this study, the first detection of a centrally con-densed compact structure with a size of ∼
150 au, whichplausibly traces a disk (Figure 1 c ), was made by ourALMA dust continuum observations with an angularresolution of θ ∼ ′′ .
2. The N D + emission was also de-tected and traces a large-scale filament in the OMC-3region. On the other hand, the C O emission traces aflattened envelope with a size of ∼ D + . The spatial distribution of the C O emissionis very different from that of the N D + emission thattraces the filament. The spatial discrepancy betweenC O and N D + is considered to be realized when par-ent molecules such as H D + are destroyed by proton ex-change with CO after the gas temperature exceeds thesublimation temperature of the CO molecules ( T &
20 K;Jørgensen et al. 2004; Crapsi et al. 2005; Salinas et al. 2017). Thus, it is expected that a mature and warmenvelope with a temperature of &
20 K already exists inMMS 3. We also detected a CO bipolar outflow (Fig-ure 5), in which cavity-like structures seem to be inter-acting with the C O flattened envelope. The directionof the outflow is perpendicular to the long axis of theenvelope. The compact disk, flattened envelope, andmolecular outflow detected in our ALMA observationsindicate that MMS 3 is not a prestellar source, but aprotostellar source, consistent with the detection of in-frared and X-ray sources in MMS 3.We also investigated the outflow properties of MMS 3.The comparisons of the outflow physical properties be-tween MMS 3 and other sources in the OMC-2/3 regionare presented in Figure 6. The masses and sizes of theoutflows associated with MMS 3 and MMS 6 are signifi-cantly smaller than those detected by the survey obser-vations in the OMC-2/3 region (Takahashi et al. 2008;Tanabe et al. 2019; Feddersen et al. 2020). This indi-cates that the outflows detected in MMS 3 and MMS 6are very young. The outflow dynamical timescale, whichis an indicator for identifying the protostellar age, es-timated for MMS 3 is 1000-1800 yr. In addition, thedynamical timescale for MMS 6 (Takahashi & Ho 2012;Takahashi et al. 2019) is as short as . σ detection) of a very com-pact SiO emission (Figure 5) could be attributed to ajet in MMS 3. Matsushita et al. (2019) showed thatthe intensity ratio between CO and SiO emission is ∼ Figure 6.
Outflow size vs. outflow mass (left), outflow size vs. outflow force (middle), and bolometric luminosity vs. outflowforce (right). The red filled diamonds denote the outflow associated with OMC-3/MMS 3. The blue filled squares denote theoutflow associated with OMC-3/MMS 6 (Takahashi & Ho 2012). The squares and circles represent the outflows observed inthe OMC-2/3 (Feddersen et al. 2020) and the Taurus (Hogerheijde et al. 1998) star-forming regions, respectively. Class 0 andI objects are shown by filled and open symbols, respectively. The bolometric luminosity of the OMC-2/3 objects is taken fromFurlan et al. (2016). Outflow parameters in OMC-2/3 estimated by Feddersen et al. (2020) plotted here are consistent withprevious several studies within factor of ∼ at their peak. Assuming a similar intensity ratio forMMS 3, the SiO emission can be detected at ∼ σ signallevel. This implies that SiO detection should be limitedaround the peak position in the CO emission. Thus,the marginal SiO detection for MMS 3 may be naturalconsidering the limited sensitivity of our observations.It should be noted that, alternatively, the SiO emissionis possibly related to the interaction between the out-flow and the envelope because the velocity of the SiOemission is not very high, as described in § Possible Evolutionary Scenarios
In this subsection, we discuss the evolutionary sta-tus of MMS 3. As described above, in addition to afaint bipolar outflow, we detected a very compact high-velocity component around the continuum peak position(Figure5). Additionally, the number density of MMS 3 iscomparably low compared with other prestellar sourcesfound in the OMC-3 region. These characteristics ofMMS 3 resemble the very early phase of star forma-tion just after protostar formation. In the collapsing(prestellar) cloud, the first core drives a wide-angle bipo-lar outflow before protostar formation (Machida et al.2008; Tomida et al. 2013). Then, the central region ofthe first core collapses to form a protostar that drives acollimated jet. Thus, a small-sized jet is enclosed by awide-angle large-scale outflow just after protostar forma-tion, as seen in Figure 5. At this epoch, the remnant ofthe first core remains around the protostar with a sizeof ∼ ∼ t dyn ∼ § ∼ L ⊙ , which is comparable to the bolometric lumi-nosity of MMS 3. Furthermore, a protostellar source atthis epoch would be observed to be similar to a prestellarsource because the protostar has just been born, and itssurrounding environment should not be affected by theprotostar. Thus, the very early star formation phasescenario seems to explain the observational propertiesof MMS 3 well. However, the velocity of the outflow( & − ) observed in MMS 3 is larger than the the-oretical prediction. Theoretical studies have shown thatthe outflow velocity just after protostar formation is lim-ited to <
10 km s − (e.g., Machida et al. 2008). Besides,as shown in § O and DCN distributions im-ply that the central protostar is surrounded by a warmenvelope with
T >
20 K. In addition, the peak shift ofDCN (1 . ′′ ∼
520 au) suggests the pres-ence of high-temperature ( >
100 K) gas in the vicinityof the central star. These results seem to contradict thevery early phase scenario.An alternative scenario to explain the observed char-acteristics of MMS 3 is the low-mass star forma-tion scenario. It is considered that intermediate-massstars in the OMC-3 region preferentially form withhigh accretion rates, because of the high mass ejec-tion rate, which is proportional to the mass accre-tion rate (Takahashi et al. 2008; Takahashi & Ho 2012;Matsushita et al. 2019). On the other hand, the outflow3parameters for MMS 3 are similar to those for low-massstars observed in the low-mass star-forming region (Fi-ure 6). Thus, the mass accretion rate onto the protostarin MMS 3 is expected to be low. The mass accretion rateis determined by the condition of the prestellar phase.As described above, the number density within the cen-tral 3000 au region is at least one order of magnitudesmaller in MMS 3 than in other intermediate-mass Class0 and I sources in the OMC-3 region (Takahashi et al.2013). The relatively low mean number density corre-sponds to the less massive core within which a low-massstar with a relatively low mass accretion rate forms, in-stead of forming an intermediate-mass star with a rela-tively high mass accretion rate. Matsushita et al. (2017)showed that a low mass accretion rate is realized, anda low-mass star forms when the natal cloud core is lessdense (for details, see also Machida & Hosokawa 2020) . Furthermore, the bolometric luminosity of MMS 3is L bol =4.2 or 3.6 L ⊙ , which is four to ten times lowerthan that of other protostellar sources in the OMC-3region (Furlan et al. 2016; Tobin et al. 2020). Duringthe main accretion phase, the bolometric luminosity isroughly proportional to the mass accretion rate. Thus,the low bolometric luminosity also means that the massaccretion rate is low . Therefore, we expect that, inMMS 3, a young protostar is growing with a low massaccretion rate. As a result, a low-mass star would formin MMS 3, while intermediate-mass stars form in otherbright millimeter sources such as MMS 5 and MMS 6 inthis region.In summary, compared to other protostellar sourcesin the OMC-3 region, MMS 3 is peculiar because thebolometric luminosity, number density, and outflow pa-rameters, such as outflow force, are as low as those nor-mally observed in low-mass protostellar sources. Also,no strong SiO emission was detected, in contrast toother Class 0 sources in OMC-3 (Matsushita et al. 2019;Hsu et al. 2020, Takahashi et al. in prep.). To explainthe characteristics of MMS 3, we proposed two possiblescenarios: (1) a very early star formation phase and (2)low-mass star formation with a low accretion rate. For Matsushita et al. (2017) pointed out that the accretion rate de-pends on the concentration of the star-forming cloud, and ahigher accretion rate is realized in a cloud with an initially highercentral density (or higher central condensation). The relatively low luminosity of MMS 3 may be explained byepisodic accretion. The time varying accretion (or episodic ac-cretion) temporally changes the bolometric luminosity. The bolo-metric luminosity becomes low in the low-accretion (or quiescent)phase. It should be noted that although the outflow intermit-tently appears with episodic accretion, we could not confirm anysign of time variability in the MMS 3 outflow (e.g., Figure 6 rightpanel). scenario (1), the protostar in MMS 3 is younger thanthe protostars in other protostellar sources and is in thevery early phase just after protostar formation. How-ever, the outflow velocity of >
10 km s − cannot sup-port this scenario. For scenario (2), a low-mass staris currently forming with a low mass accretion rate inMMS 3, while intermediate-mass stars with relativelyhigh mass accretion rates are forming in other sources.At the moment, the observational signatures do not con-tradict scenario (2). The difference in observationalcharacteristics between MMS 3 and other bright sourcesmay be attributed to the initial condition of star for-mation and the surrounding environment of star form-ing cores, suggested by the previous SMA observations(Takahashi et al. 2013). SUMMARYWe investigated the protostellar source MMS 3 in theOMC-3 region with 1.3 mm continuum, CO ( J = 2–1),C O ( J = 2–1), SiO ( J = 5–4), N D + ( J = 3–2), andDCN ( J = 3–2) emissions and obtained the followingresults. • With sub-arcsecond angular resolution, we de-tected compact 1.3 mm continuum sources witha size of ∼
150 au and a mass of ∼ − M ⊙ forthe first time, presumably tracing a compact dustydisk. The peak position of the 1.3 mm continuumemission corresponds to both the infrared (HOPS91) and X-ray sources. • A flattened envelope with a size of ∼ O ( J = 2–1) emission. A faintand compact bipolar outflow was also detected inthe CO ( J =2–1) emission for the first time. Theoutflow direction is roughly perpendicular to themajor axis of the flattened envelope and disk de-tected in the C O and 1.3 mm continuum emis-sions, respectively. The outflow maximum gas ve-locity is ∼
42 km s − , assuming a disk inclinationangle of i = 45 ◦ . The SiO ( J = 5–4) emission ismarginally detected at almost the same positionas the CO outflow. • Two possible scenarios were proposed to explainthe evolutionary stage and status of a protostarin MMS 3. One is the low-mass Class 0 stagescenario, while the other is the very early phasescenario. Although many observational character-istics can be explained by both scenarios, the highoutflow velocity cannot support the latter. Com-paring the outflow properties in MMS 3 with thosein different star-forming regions, it is expected4 that a low-mass star is forming, for which the ac-cretion rate in MMS 3 is expected to be lower thanthose of protostellar sources in the OMC-3 region.ACKNOWLEDGMENTSWe thank the anonymous referee for providing usvery helpful comments and suggestions. This paperuses the following ALMA data: ADS/JAO. ALMA No.2015.1.00341.S. ALMA is a partnership of ESO (rep-resenting its member states), NSF (USA) and NINS(Japan), together with NRC (Canada),
M OST andASIAA (Taiwan), and KASI (Republic of Korea), incooperation with the Republic of Chile. The JointALMA Observatory is operated by ESO, AUI/NRAO,and NAOJ. K. Morii is very grateful for support fromthe SOKENDAI and NAOJ Chile Observatory (cur-rently, NAOJ ALMA Project) while visiting a coau-thor S. Takahashi through the SOKENDAI summerstudent program 2018. This work was supported byJSPS KAKENHI grants JP17KK0096, JP17K05387,and JP17H06360. REFERENCES
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