A 1000 AU Scale Molecular Outflow Driven by a Protostar with an age of <4000 Years
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A 1000 AU Scale Molecular Outflow Driven bya Protostar with an age of ∼ < Ray S. FURUYA, Yoshimi KITAMURA, and Hiroko SHINNAGA Institute of Liberal Arts and Sciences, Tokushima University, Minami Jousanjima-Machi 1-1, Tokushima, Tokushima 770-8502, Japan Institute of Space and Astronautical Science, Japan Aerospace Exploration Agency, Yoshinodai 3-1-1, Chuo-ku, Sagamihara, Kanagawa252-5210, Japan Chile Observatory, National Astronomical Observatory of Japan, Current address: Department of Physics, Faculty of Science,Kogoshima University, Korimoto 1-21-35, Kagoshima, Kogoshima 890-0065, Japan (Received; Revised; Accepted for publication in ApJ on December 12, 2018)
ABSTRACTTo shed light on the early phase of a low-mass protostar formation process, we conducted interfero-metric observations towards a protostar GF 9-2 using the CARMA and SMA. The observations havebeen carried out in the CO J = 3 − µ m with a spatial resolution of ≈
400 AU. All the continuum images detecteda single point-like source with a beam-deconvolved effective radius of 250 ±
80 AU at the center of thepreviously known 1.1 – 4.5 M (cid:12) molecular cloud core. A compact emission is detected towards theobject at the Spitzer
MIPS and IRAC bands as well as the four bands at the
WISE . Our spectro-scopic imaging of the CO line revealed that the continuum source is driving a 1000 AU scale molecularoutflow, including a pair of lobes where a collimated “higher” velocity ( ∼
10 km s − with respect tothe velocity of the cloud) red lobe exists inside a poorly collimated “lower” velocity ( ∼ − ) redlobe. These lobes are rather young (dynamical time scales of ∼
500 – 2000 yrs) and the least powerful(momentum rates of ∼ − − − M (cid:12) km s − yr − ) ones so far detected. A protostellar massof M ∗ ∼ < . M (cid:12) was estimated using an upper limit of the protostellar age of τ ∗ ∼ < (4 ± × yrs and an inferred non-spherical steady mass accretion rate of ∼ × − M (cid:12) yr − . Together withresults from an SED analysis, we discuss that the outflow system is driven by a protostar whose surfacetemperature of ∼ ,
000 K, and that the natal cloud core is being dispersed by the outflow.
Keywords:
ISM: evolution — ISM: jets and outflows — ISM: individual (GF 9–2, L 1082C,WISE J205129.83+601838) — submillimeter: ISM — stars: formation — stars: proto-stars INTRODUCTION1.1.
Background
Because of the prominent progress in observational and theoretical studies over the past decades, the early phasesof the formation of an isolated low-mass star are now certainly well understood. Nevertheless, our knowledge of itsearliest phase is obviously limited in the observational studies. The theoretical studies in the late 1960s and the 1970sindicated that, once a cloud core lost support against the collapse due to its self-gravity, the core collapses isothermallyin a runaway fashion, which is followed by an accretion process onto a protostellar core. The first kernel of a mass
Corresponding author: Ray S. [email protected]@[email protected] a r X i v : . [ a s t r o - ph . GA ] D ec Furuya et al. produced at the center of the collapsing cloud core is believed to evolve towards a hydrostatic object which becomesopaque to the dust continuum emission, and thus the cooling by the dust emission becomes inefficient. Such anadiabatic hydrostatic core is referred to as a first hydrostatic core, a first core (Larson 1969). The thermal evolutionof first cores has been extensively studied by e.g., Masunaga et al. (1998) who performed 1D radiative hydrodynamiccalculations to track their evolution. Once the central temperature is elevated to about 2000 K, corresponding to thevolume mass density of ρ ∼ − g cm − , molecular hydrogen commences being dissociated into atomic hydrogen.The dissociation uses the thermal energy which originates from the released gravitational energy due to the accretion.After the dissociation is completed, a second hydrostatic core forms. The second cores are believed to correspond tothe protostars observed as “class 0” objects (Andr´e et al. 1993). Although feasibility of detecting the first cores havebeen long discussed (e.g., Boss & Yorke 1995; Omukai 2007; Tomisaka & Tomida 2011; Commer¸con et al. 2012; Tomidaet al. 2013), it is an enormously difficult task for observers to identify the first cores simply because they are veryshort-lived objects.Despite such difficulty, the growing numbers of the first core candidates have been reported (e.g., Belloche et al.2006; Enoch et al. 2010; Chen et al. 2010; Dunham et al. 2011; Pineda et al. 2011; Tsitali et al. 2013; Pezzuto et al.2012; Hirano & Liu 2014; Friesen et al. 2014; Maureira et al. 2017) based on unbiased surveys which observed thethermal dust continuum emission at submillimeter (submm) to infrared (IR) bands with space telescopes togetherwith bolometer cameras on ground-based millimeter (mm) to submm telescopes. The theoretical studies predictedthat the duration of the first core phase should continue at most a few times 10 yrs which is an order of magnitudeshorter than the lifetime of the class 0 objects (e.g., Andr´e et al. 2000). Therefore only a handful of the first core objects should exist in nearby molecular clouds when we consider the number of the class 0 sources identified so farand the relative duration between the first and second core phases (Masunaga et al. 1998; Masunaga & Inutsuka2000). Therefore the number of the candidates is too many to be consistent with this statistical argument. This raisesa question that some of the candidates may have been misidentified and/or the duration of the first cores may belonger than the theoretical predictions. We therefore believe that it is still required to identify protostars at its earlyevolutionary stage, and study their physical properties.1.2. Previous Observations
In this context, we performed a detailed study of the natal cloud core harboring an extremely young low-massprotostar GF 9-2 at a distance ( d ) of 200 pc (Wiesemeyer 1997, see a summary in Poidevin & Bastien 2006) locatedin the GF 9 filament (Schneider & Elmegreen 1979). The filament was studied by near-infrared (NIR) extinction(Ciardi et al. 1998), and optical and NIR absorption polarization (Poidevin & Bastien 2006) observations. Poidevin& Bastien (2006) showed that there exist well aligned pc-scale magnetic fields whose overall direction appears tobe almost perpendicular to the filament, and discussed that the magnetic fields must have regulated the formationand evolution of the filament. In the filament there are almost equally spaced seven dense cloud cores traced by theNH (1,1) lines (Furuya et al. 2008, hereafter Paper II). Subsequently, Furuya et al. (2014) [hereafter Paper IV] studiedthe physical properties of the tenuous ambient gas surrounding the dense cloud core GF 9-2 by observing the J =1–0transitions of CO, CO and C O molecules, which covered ∼ one-fifth of the whole filament. We found that thefilament around the core is supported by turbulent and magnetic pressures against self-gravity, and argued that thecore has formed through gravitational collapse triggered by the local decay of the supporting force(s).The cloud core GF 9-2 has been studied in various molecular lines by many authors. Using the NIR extinction maptogether with the CO (1–0) and CS (2–1) line data, Ciardi et al. (1998, 2000) identified an ∼ M (cid:12) molecular clump of the “GF 9-Core” with a size of ∼ × protocore of the “Southwestern Condensation” (Furuya et al. 2014), and the IRAS point source PSC 20503+6006 (see the leftpanel of Figure 10 in Ciardi et al. 2000). The dense cloud core is cross-identified as L1082 C (e.g., Bontemps etal. 1996; Caselli et al. 2002, and references therein). Caselli et al. (2002) estimated that the core has a mass of0.35 ± M (cid:12) over a radius ∼ .
14 pc through the N H + (1–0) line observations using the FCRAO 14 m telescope( θ HPBW = 54 (cid:48)(cid:48) ). Figure 2 in Caselli et al. (2002) and Figure 10 in Ciardi et al. (2000) clearly show that the IRASpoint source is located at the ∼ (cid:48)(cid:48) (corresponding to 0.08 pc at d = 200 pc) south-south-west position from the corecenter. Bontemps et al. (1996) observed the dense cloud core L1082 C, i.e., GF 9-2 in the CO (2–1) line using CaltechSubmillimeter Observatory (CSO) 10.4 m telescope ( θ HPBW = 30 (cid:48)(cid:48) ) as an outflow survey towards low-mass embeddedyoung stellar objects (YSOs). However, no molecular outflow was detected with the upper limit of 1.5 × − M (cid:12) km s − yr − in outflow momentum rate. In contrast to the negative detection of an outflow, Furuya et al. (2003) ∼ < O maser emission at 22 GHz towards the core center using the Nobeyama 45 m telescope. Althoughthe beam size of the maser observations was 75 (cid:48)(cid:48) (corresponding to 0.072 pc at d = 200 pc), the presence of the masersstrongly suggests that star formation activity has already commenced within a radius of 0.036 pc centered on the cloudcore. Subsequently, we carried out higher resolution observations: the N H + (1–0) and H CO + (1–0) lines usingNobeyama 45 m telescope ( θ HPBW (cid:39) (cid:48)(cid:48) ) and Caltech Owens Valley Radio Observatory (OVRO) millimeter (mm)array (synthesized beam size θ syn (cid:39) (cid:48)(cid:48) ) to study the cloud core (Furuya et al. 2006, hereafter Paper I), HCO + (3–2)line using CSO 10.4 m telescope (25 (cid:48)(cid:48) ) to detect gas infall in the core (Furuya et al. 2009, hereafter Paper III), and CO(1–0) and (3–2) lines using Nobeyama 45 m (17 (cid:48)(cid:48) ) and CSO 10.4 m telescopes (22 (cid:48)(cid:48) ), respectively, to assess the negativedetection of the outflow (Paper I). Here we scaled the clump mass and size estimated by Ciardi et al. (1998), the cloudcore mass and size by Caselli et al. (2002), and the outflow momentum rate by Bontemps et al. (1996) using d = 200pc (Wiesemeyer 1997) because Bontemps et al. (1996); Ciardi et al. (1998, 2000); Caselli et al. (2002) adopted d = 440pc. Hereafter, we adopt the distance to the object of d = 200 pc.In Paper I, we argued that the central object(s) deeply embedded in the GF 9-2 core has (have) not generated anextensive molecular outflow, as also suggested by Bontemps et al. (1996). The absence of an extensive outflow shouldprovide us with a rare opportunity to investigate the physical properties of the natal core free from the disturbance bythe outflow. In Paper III, we reported that “blueskewed asymmetry profiles” in the optically thick HCO + (1–0), (3–2),and HCN (1–0) lines were detected within a radius of r (cid:39) (cid:48)(cid:48) (corresponding to 0.03 pc) suggestive of the presence oflarge-scale gas infall. The estimated infall velocity was shown to have reasonable consistency with the predictions fromthe runaway collapse model in the accretion phase (Larson 1969; Penston 1969; Hunter 1977; Whitworth & Summers1985, hereafter the LPH solution), and an infall rate of ˙ M inf (cid:39) . × − M (cid:12) yr − was deduced (Paper III). On theother hand, analyzing the N H + (1–0) and H CO + (1–0) line images that were produced by combining the visibilitydata taken with the OVRO mm-array and the 45 m telescope, we revealed that a density profile of ρ ( r ) ∝ r − holdsover the annulus of 0.003 ∼ < r/ pc ∼ < H + emission ( (cid:39) (cid:48)(cid:48) ;Paper I), we set an inner radius for the region showing the ρ ( r ) ∝ r − profile to be 0.006 pc. This inner radius gives anupper limit of the protostar’s age of t protostar ∼ < × yrs (Paper I). All the results strongly suggest that the GF 9-2core has undergone gravitational collapse from the initially unstable state, and has just formed a protostar(s) in thepast ∼ < × yrs in its center. In order to assess whether or not the protostar has launched a compact molecularoutflow(s) and to search for its driving source(s), we performed interferometric observations of the central region ofthe core in the CO (3–2) line and continuum emission. OBSERVATIONS AND DATA RETRIEVALTo resolve the issues summarized in §
1, we carried out mm- and submm- wavelength interferometric observationswith the Combined Array for Research in Millimeter-wave Astronomy (CARMA) ( § ( § µ m and 1.1 mm, respectively, providedus with 2.0 and 1.8 times higher angular resolution images than the CARMA did at 3.3 mm, while the CARMA3.3 mm observations gave us 3.0 and 2.3 times larger fields of view (FoV) and could detect 5.4 and 4.0 times largerspatial structures than the SMA 840 µ m and 1.1 mm observations, respectively. In addition, we used archival Spitzer
Space Telescope and Wide-field Infrared Survey Explorer data. to verify the presence of mid-infrared to far-infraredemission towards the protostar ( § CARMA Observations
The aperture synthesis observations tuned at 91.18 GHz, which is the middle frequency of between the N H + (1–0)and HCO + (1–0) emission, were carried out using the CARMA with C, D, and E configurations (project code: c0296).The visibility data used to produce a 3 mm continuum emission image were obtained by 5, 4, and 5 partial tracks Support for CARMA construction was derived from the Gordon and Betty Moore Foundation, the Kenneth T. and Eileen L. NorrisFoundation, the James S. McDonnell Foundation, the Associates of the California Institute of Technology, the University of Chicago, thestates of California, Illinois, and Maryland, and the National Science Foundation. CARMA development and operations were supportedby the National Science Foundation under a cooperative agreement, and by the CARMA partner universities. The Submillimeter Array is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute ofAstronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica. This work is partially based on the data taken with
Spitzer
Space Telescope, which is operated by the Jet Propulsion Laboratory,California Institute of Technology under a contract with NASA. This publication makes use of data products from the Wide-field Infrared Survey Explorer, which is a joint project of the Universityof California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics andSpace Administration.
Furuya et al. performed in 2009 May, April, and June, respectively. Since our primary goal of the observations is to image thesemolecular lines by combining with the single-dish data taken with the Nobeyama 45 m telescope (Paper I), we carriedout nineteen-field mosaic observations. However, in this paper, we concentrate on discussing the continuum emission,and the line data will be published in another paper.We used J1849+670 as phase and gain calibrators, and 3C454.3, 3C345, J1642+689, and J1751+096 as passbandcalibrators. The flux densities of J1849+670 were determined from observations of Mars and MWC 349. The uncer-tainties of our flux calibration were estimated to be approximately 10%. The visibility data were calibrated and editedusing the MIRIAD package. For the continuum data, we merged the visibilities in both of the sidebands to improvethe image sensitivity: the representative frequency of the continuum image was set to be the center frequency of thedual bands. The image construction was done using the MIRIAD package.2.2.
SMA Observations
The aperture synthesis observations were made using the SMA at the wavelengths ( λ ) of 1.1 mm and 840 µ m on 2010August 10 and 11, respectively (project code: 2010A-S56). The observations were in the Compact-North configurationusing the eight and six antennas at λ =1.1 mm and 840 µ m, respectively. During the observations, the optical depth ofthe terrestrial atmosphere at 225 GHz ( λ = 1.3 mm) measured with the CSO monitor was fairly stable ( τ ∼ . τ ∼ . CO(3–2), N H + (3–2) and HCO + (4–3) at the 840 µ m band as well as HCO + (3–2) at the 1.1 mm band. Note that theN H + and HCO + lines are typical high-density gas tracers. We configured the correlator so that we can utilize themaximum bandwidth of 4 GHz in each polarization for the continuum detection.For phase and amplitude calibrators, we observed 3C 418 and BLLac, whose angular distances to the source are9.2 and 21 degrees, respectively, every 5 minutes. The calibration of the passband response function was done byobserving 3C 279, 3C 454.3, 3C 345, and 3C 84 at the beginning and ending of the observations. The scale factors forconverting into the absolute flux densities were determined by observations towards Neptune and Uranus. All thevisibility data were calibrated and edited using the MIR package. Image construction was subsequently done with theMIRIAD package. For the continuum emission data, we produced an image at each sideband by concatenating thedual polarization data. 2.3. Spitzer Space Telescope Archive Data
We retrieved infrared (IR) images towards the 3 mm continuum source (Paper I) from the
Spitzer Science Center data archive. These images are taken at wavelengths of 3.6, 4.5, 5.8, and 8.0 µ m taken with the Infrared Array Camera (IRAC; Fazio et al. 2004) (data ID µ m with the Multiband Imaging Photometer for Spitzer(MIPS; Rieke et al. 2004) (data ID Spitzer
Space Telescope. The point spread function(PSF) sizes of the Spitzer IRAC and MIPS images ranges between 1 . (cid:48)(cid:48) µ m band, which are comparable tothose in our interferometric observations and 18 (cid:48)(cid:48) at the 70 µ m bands. These fully calibrated data were used withoutspatial smoothing, and 1 σ noise levels over emission-free regions are 2 . × − , 2 . × − , 9 . × − , 0 . × − ,0 . × − , and 0 . × − MJy sr − for the 3.6, 4.5, 5.8, 8.0, 24 and 70 µ m band data, respectively.2.4. WISE Archive Data
We retrieved infrared (IR) images towards the 3 mm continuum source from the
Wide-field Infrared Survey Explorer (WISE; Wright et al. 2010) data archive. These images are taken at wavelengths of 3.4, 4.6, 11.6, and 22.1 µ m. ThePSF sizes of the WISE images are ∼ (cid:48)(cid:48) at the 3.4, 4.6 and 11.6 µ m bands, whereas ∼ (cid:48)(cid:48) at the 22 µ m band, whichare ∼ Spitzer images. RESULTS AND ANALYSIS3.1.
Continuum Emission at 3.3 mm, 1.1 mm, and 850 µ m Bands In this subsection we describe results obtained from the mm and submm continuum emission maps (Figure 1), whichare followed by an analysis of a continuum spectrum (Figure 2).3.1.1.
Maps and Flux Measurements ∼ < § µ m ones by SMA (Figures 1b - e) were produced in each sideband.At all the bands, we detected single point-like sources whose peak positions agree with each other within the spatialresolutions at R. A. = 20 h m . s , Decl = 60 ◦ (cid:48) . (cid:48)(cid:48)
23 in J2000. In the CARMA 3.3 mm image the object showsa weak elongated structure to the east. No other significant emission was detected over the FoVs (Table 1) at all thebands. This would exclude the “protobinary” hypothesis which we discussed in § JMFIT in AIPS package and
IMFIT in CASAto perform beam-deconvolution. Table 2 summarizes our flux density ( S ν ) measurement at each frequency with anassumption that the morphology of the object can be approximated by a 2D elliptical Gaussian. A comparison ofthe synthesized beam sizes (Table 1) and the beam-deconvolved source sizes (see the Θ maj × Θ min values in Table 2)indicates that our observations barely resolved the emission at all the bands, i.e., they are detected as slightly extendedsources.After assessing the beam-deconvolved sizes, we measured the area of the emanating region in the plane-of-sky ( A s )to calculate its effective radius ( R eff ) by A s = π (Θ maj × Θ min ) ≡ πR . For further analysis, we use the mean R eff of 250 ±
80 AU calculated from the SMA bands of 1.1 mm and 850 µ m because the synthesized beam size ( θ syn ) of ∼ (cid:48)(cid:48) at SMA observations is better than those at 3 mm.In order to compare the CARMA 3 mm measurement with the previous one at 3 mm using the OVRO mm-array( θ syn ∼ (cid:48)(cid:48) ; Paper I), we checked consistency between the photometric results from the above Gaussian-fitting methodand those measured over the region enclosed by the 3 σ level contours using the CARMA and SMA data. This isbecause we adopted the latter method in Paper I. We found a reasonable consistency between the two methods,suggesting that the 2D elliptical Gaussian approximation adopted in the beam-deconvolution processes is reasonable.3.1.2. Uncertainty in the Flux Measurements
Based on the minimum spatial frequencies of our observations (Table 2) and the discussion in Appendix of Wilner& Welch (1994), we estimated that our SMA observations missed approximately 70% of the flux densities for theextended components of the envelope with respect to the expected zero-spacing flux densities assuming a model witha power-law radial density profile (Paper I).Table 2 presents the mm- and submm continuum flux densities where we included the previous OVRO measurement(Paper I). We verified that the 3 mm flux difference between the OVRO and CARMA measurement is real. Wequalitatively argue that the difference would be caused in the process of synthesis imaging due to the various differences,e.g., those in spatial-frequency ( D λ ) ranges (3.8 ≤ D λ /kλ ≤
67 for OVRO vs. 1.93 ≤ D λ /kλ ≤ u, v ) coverage, their weighting functions (the OVRO data were imaged with natural weighting, whilethe CARMA data with “Robustness of +2”), correlator bandwidths (4096 MHz for OVRO vs. 938 MHz for CARMA),usage of multi-frequency synthesis method in CARMA, and the beam sizes adopted for the photometry (5 . (cid:48)(cid:48) × . (cid:48)(cid:48) . (cid:48)(cid:48) × . (cid:48)(cid:48)
84 for CARMA). Among these possible causes, we argue that the difference in the D λ rangesmay be the most dominant one. However, it is not trivial to explain quantitatively the photometric difference betweenthe OVRO and CARMA results.3.1.3. Continuum Spectrum over the Millimeter and Submillimeter Bands
Figure 2 presents an interferometric mm- and submm continuum spectrum including the OVRO measurements.Compared to the 1.1 mm and 850 µ m flux densities, those at 3.3 mm drop almost one order of magnitude (Table 2;see also Figure 2). Assuming the power-law spectrum with S ν = S ν α , we obtained the best-fit spectral index of α = 2 . ± .
3. If we assume that the observed continuum emission at mm and submm regime is fully attributedto optically-thin thermal dust emission represented by a single-temperature graybody emission and that the dustmass absorption coefficient can be written by κ ν ∝ ν β , the best-fit α value obtained between the wavelength range of λ = 3 mm – 850 µ m leads to the β -index of 0 . ± .
3. The inferred β is comparable to those measured at the wavelengthrange towards circumstellar disks around class II sources ( β ∼ < β ∼ β -indexis smaller than those expected for an envelope harboring a protostar. This yields that there might be forming anextremely-dense compact optically-thick region around the protostar or there could still exist a remnant of a firsthydrostatic core (e.g., Larson 1969)(described in § Furuya et al.
It is theoretically shown that small dust grains can grow in the innermost densest part of infalling envelopes, e.g.,Hirashita & Omukai (2009); Ormel et al. (2011). We therefore estimate mass of the GF 9-2 envelope to infer a meanvolume density ( n ) to see such a possibility. Assuming that the observed continuum emission is attributed to opticallythin thermal dust emission, the total mass of the circumstellar materials, M csm , can be calculated by M csm = S ν d κ ν B ν ( T d ) where κ ν is the dust mass absorption coefficient at frequency ν , and B ν ( T d ) the Plank function with dust temperatureof T d . In these calculations, we adopted the usual assumption that κ ν has a form of κ ( ν/ν ) β where κ is a referencevalue at a reference frequency, ν . We used the κ = 0.1 cm g − at ν = 1.2 THz with β = 1.8 (Planck Collaboration2011b). Notice that the κ and β values lead to κ
231 GHz of 0.007 cm g − which falls between the two values used inPaper I. Moreover we assumed that the dust in the region is well-coupled with the gas. Hence we considered that theexcitation temperature ( T ex ) of the N H + (1–0) lines (Paper I) represents the gas ( T gas ) and dust ( T d ) temperatures,i.e., T ex = T gas = T d . We adopted the mean excitation temperature of (cid:104) T ex (cid:105) = 22 . ± . σ level contour of the 840 µ m continuum emission in Figure 1e (Appendix A). We calculated M csm values for the individual bands (Table 2); the mean value of ∼ × − M (cid:12) for the higher resolution SMAresults and the R eff =250 ±
80 AU leads to the order of n is 10 cm − . This may be too low for dust grains to coagulate,hence it is unlikely that dust grains over the entire envelope has already grow as in more evolved objects.3.2. Continuum Emission in the Infrared Image Data
Spitzer data
In Figure 3, we present the
Spitzer
IRAC and MIPS images centered on the peak position of the SMA 357 GHzsource (Figure 1e). Assuming that the absolute positions in the IRAC and MIPS images agree with each other within0 . (cid:48)(cid:48)
36, which is the pixel size common to all the
Spitzer images, these images were “registered” to the 3.6 µ m one by“reprojecting” using astropy package. Subsequently these images were compared to the SMA images, whose positionalaccuracies are estimated to be approximately ∼
10% of the synthesized beam sizes of 0 . (cid:48)(cid:48) ∼ (cid:48)(cid:48) (corresponding to ∼ µ mbands with the detection threshold of the 3 σ level in each image.The flux density at each band was computed from the image in unit of MJy sr − with aperture photometry usingtask apphot in IRAF package with a standard manner; we integrated the emission inside an optimized circle centeredon the source after subtracting the background. Notice that the 70 µ m flux density is considered as upper limitsbecause the HPBWs of the aperture encompassing the source is significantly larger than that of interest for us. Thephotometry results are summarized in Table 4. 3.2.2. WISE data
Towards the 3 mm continuum source, we clearly detected a point-like source in all the four
WISE bands. The objectis identified as
WISE J205129.83+601838 in the WISE All-Sky Release Source Catalog from which we obtained itsVega magnitudes. These magnitudes were converted into flux densities in unit of Jy (Table 4) with an equation of F ν = F ν, × ( − m Vega / . where F ν, is a zero magnitude flux density (Wright et al. 2010) and m Vega calibrated WISEVega magnitudes of the source. The resultant fluxes are shown in Table 4.3.3. CO (3–2) Line Emission
Next we present the results from analysis of the SMA CO (3–2) line data (Figure 4) and from re-analysis ofthe previously published CO (3–2) line data taken with the Caltech Submillimeter Observatory (CSO) § § Velocity Channel Maps The Caltech Submillimeter Observatory was operated by the California Institute of Technology under the grant from the US NationalScience Foundation (AST 05-40882). ∼ < CO (3–2) emission taken with the SMA. The blueshifted emissionbetween v LSR = − . − . − is mainly seen towards the southwest of the continuum source, although theemission is much weaker than the redshifted one. Similarly the redshifted emission is also seen towards the southwest,and is detected over a wider velocity range between v LSR = − . . − . In the velocity channels where theintense emission is detected, e.g., the panels between v LSR = − . . − , the negative contours are seenalong the northeast–southwest direction. We argue that such an artifact is caused by the limited ( u, v ) coverage inour SMA observations. 3.3.2. Spectral Line Profile
Averaging the CO emission along the velocity channel (Figure 4) over the region enclosed by the 3 σ level contourof the total integrated intensity map (Figure 5), we produced an interferometric CO (3–2) spectrum (Figure 6).Here we selected only the bright emission seen in the central peak of Figure 5 to extract the CO emission towardsthe continuum source. The specific intensity scale in unit of Jy beam − of the spectrum was converted to brightnesstemperature ( T b ) in K using the relation of T b = c k B ν s (cid:82) Ω s I ν dΩ where k B is the Boltzmann constant, c the lightvelocity, and Ω s the source solid angle. The CO (3–2) spectrum taken with SMA shows a noticeable line profile. Ithas well-defined high velocity wing emission, especially on the redshifted side.Figure 7 compares the single-dish and interferometric CO (3–2) spectra in the T b scale taken with the CSO 10.4 mtelescope (Paper I) and the SMA, respectively. The SMA spectrum was produced by averaging the CO emission insidethe CSO beam of Ω b = π θ ( θ HPBW = 22 (cid:48)(cid:48) ). On the other hand, the previously published CSO data werereprocessed by revising the velocity ranges for the emission-free channels to determine the baseline of each spectrum(Figure 8). Our re-analysis is motivated by the fact that we did not know the presence of the high-velocity wingemission (Figure 6) when we had reduced the data for Paper I. This re-analysis allowed us to detect the redshifted tailemission up to v LSR = 3.6 km s − and at a single channel centered on v LSR = 9.2 km s − . Notice that the high-velocitywing emission is detected only towards the center position (Figure 8).Using the spectra shown in Figure 7, we computed the integrated intensities of the redshifted emission to be (cid:82) T b d v = 1 . ± . − for the CSO spectrum and 0 . ± .
09 K km s − for the SMA one over a velocityrange of 0 . ≤ v LSR /km s − ≤ . § ∼ § θ HPBW = 22 (cid:48)(cid:48) beam detected the redshifted wing emission onlytoward the center position (Figure 8).The most straightforward interpretation for the high-velocity wing emission is that the protostar of interest is drivinga molecular outflow. This is simply because the velocity difference between the systemic velocity ( v sys ) and the terminalvelocity of the redshifted emission ( v t , red ) of | v t , red − v sys | = 9.5 km s − is comparable with typical velocities of theoutflow lobes observed in low-mass class 0 sources (an order of 1 – 10 km s − ; e.g., Bachiller 1996; Tobin et al. 2016),and is larger than a typical velocity of the envelope gas motions identified as rotation and/or infall (an order of 0.1 – 1km s − ; e.g., Walker et al. 1986; Mardones et al. 1997; Momose et al. 1998). However, the velocity difference on theblueshifted side is moderate: | v t , blue − v sys | = 3.3 km s − . In this paper, we consider that the blue one is attributedto an outflow system rather than rotation.Another clear feature seen in the SMA CO (3–2) spectrum is that no-emission is seen around the v sys . We considerboth or either of the following two causes. One is that the bulk emission near v sys is extended, thus it was resolvedout by the interferometer observations. Here the largest detectable angular size scale in our SMA observations wasestimated to be 27 (cid:48)(cid:48) (Table 1), corresponding to 0.026 pc at d = 200 pc. Therefore, the extended ( ∼ > . v sys is optically thick, producing aself-absorption trough in the spectrum.3.3.3. Spatial and Velocity Structures of the Outflow Gas
Next we examine the spatial and velocity structures of the outflow gas unveiled by our SMA CO (3–2) linespectroscopic imaging. Figure 9 presents integrated intensity maps of the compact outflow lobes overlaid on the
Furuya et al.
357 GHz continuum image which has the highest angular resolution among our interferometric continuum images.In order to produce the lobe maps, we defined two LSR-velocity ranges by considering all the features seen in thechannel maps (Figure 4) and the spectrum (Figure 6): the blueshifted emission by − . ≤ v LSR / km s − ≤ − . . ≤ v LSR / km s − ≤ +7 .
0. After defining the LSR-velocity ranges, we calculated acharacteristic velocity of each lobe ( v ch ), which is an intensity weighted mean LSR-velocity, in order to compare withother first core candidates in the literature. This is because the majority of the previous studies which address outflowproperties driven by first core candidates and very low luminosity objects preferred to use v ch to represent the flowvelocity rather than using the terminal velocity. We therefore calculated v ch , blue = − . − in the LSR-velocityfor the blue lobe, and v ch , red = +2 . − for the red lobe.The redshifted lobe is seen to the southwest of the continuum emission, whereas the blueshifted one spreads towardsboth southwest and northeast. This spatial-velocity structure seen in Figure 9, especially the redshifted one, can berecognized in the early CO (3–2) single-dish map taken with the CSO 10.4 m telescope (see Figure 8d in Paper I).However, a caution must be used to compare Figure 8d in Paper I and Figure 9 because the former was obtained fora velocity range of − . ≤ v LSR / km s − ≤ +0 .
6, whereas the latter +0 . ≤ v LSR / km s − ≤ +7 .
0. Namely, because ofthe small overlap between the two velocity ranges defined in the CSO and SMA data, we argue that the Figure 8d inPaper I represents mostly the ambient gas, hence we do not consider Figure 8d in Paper I for the discussion about theoutflows.Comparing the spatial extent of the blue and red lobes with that of the 357 GHz continuum emission (see Figure 9),we suggest that the “circum”stellar envelope is, highly likely, being evacuated by the outflow. Such dispersal of theparental gas by outflows has been observed at the core-scale at 0.1 to 0.01 pc with single-dish telescopes (e.g., Moninet al. 1996; Yoshida et al. 2010) and at the envelope-scale with interferometers (e.g., Momose et al. 1996; Velusamy& Langer 1998; Arce & Sargent 2006). This is discussed in § . ∼ < v LSR / km s − ∼ < . ∼ (cid:48)(cid:48) (correspondingto ∼ . ∼ < v LSR / km s − ∼ < +3 .
2) of Figure 4. It is also interesting that the former component shows a fan-shapedstructure opened to the southwest (Figure 9b), which is well-recognized especially in the velocity channel panels of0 . ∼ < v LSR / km s − ∼ < +2 . . ∼ < v LSR / km s − ∼ < +6 . . ≤ v LSR / km s − ≤ +5 . . < v LSR / km s − ≤ +7 .
0. These velocity ranges are shown bythe horizontal color bars in Figure 6. Note that the “red IV” is composed of the main emission associated with thecontinuum source and the isolated emission seen to the southwest (see Figure 10a). For the “red IV”, we considerthat our observations were not sensitive enough to detect a structure possibly connecting the mainbody and thesouthwestern island.Last, we point out an observational fact that the boundary velocity between the “red IV” and “red HV” lobes isclose to the LSR-velocity of the H O maser emission at 22 GHz (Furuya et al. 2003)(see Figure 6). The co-existence ofthe prominent CO wing emission and the maser transition which is widely believed to be excited in a shocked region(e.g., Hollenbach et al. 1993, 2013; Furuya et al. 2001), strongly suggests that the masers are collisionally excited inthe shocked region between the red-IV lobe with the ambient medium (see Figure 6).In summary, our SMA observations clearly detected a prominent high-velocity outflow in the CO (3–2) emissiontowards the central object in the core. Because no other continuum sources are detected, the continuum source mustbe driving the outflow. The compact outflowing gas has the following features: (i) both the blue and red lobes arebright at the southwest of the continuum source, (ii) bipolarity between the blue- and red lobes cannot be uniquelyidentified with respect to the continuum source position, (iii) the blue lobe exhibits an elongated structure along thenortheast-southwest direction, and (iv) the red lobe can be resolved into the dual structures of the “red IV” lobe andthe “red HV” one. The “red IV” lobe has a wide opening angle, whereas the “red HV” lobe is more collimated thanthe “red IV”: the former apparently traces the central axis of the latter.3.3.4.
Physical Properties of the Compact Outflow Lobes ∼ < θ ), flow velocity, dynamical time scale( τ dyn ), mass ( M lobe ), mass loss rate ( ˙ M outflow ), and momentum rate ( F outflow ). These parameters are summarized inTable 3.For the red lobes, we calculated the outflow properties for the “red IV” and “red HV” ones (see Figure 10). Wespatially divided the blue lobe into the northeastern and southwestern ones (see Table 3). The two blue lobes areseparated by the line of P.A.= − ◦ which passes the peak of the continuum emission. Notice that the two blue lobesare defined over the common LSR-velocity range ( § θ would be ∆ θ ∼ ◦ .The dynamical time scale τ dyn is estimated from the ratio between the maximum extent of the lobe and the flowvelocity which is given by v flow = | v t − v sys | ( § i is available from observations, we assumed an outflow inclination angle ( i ) of 45 ◦ . Here we defined i bythe angle between the outflow axis and the line of sight. In order to see the uncertainty caused by the unknown i values, we present Figure 11 which shows how outflow velocity and τ dyn change with outflow inclination angles. Weobtained the τ dyn range from 500 to 2000 yrs for the “red HV” and “red IV” lobes by correcting for i = 45 ◦ . Here wedid not use v ch calculated in § τ dyn estimates because the use of v ch overestimates the τ dyn values, whichpropagate to the other physical properties.In § v LSR = 9 . − component in the central panel of Figure 8). Assuming that the 9.2 km s − component represents line-of-sight velocity of a redshifted lobe, we can have another inference of τ dyn : the beam sizefor the single-dish CO (3–2) observations and the flow velocity given by v flow = | v LSR − v sys | / cos(45 ◦ ) (cid:39)
17 km s − yields an upper limit of τ dyn ∼ < l/v flow = 900 yrs where l ∼ < ( θ HPBW / / sin(45 ◦ ) = 3100 AU. The upper limit does notcontradict with the range of τ dyn (cid:39)
500 – 2000 yrs estimated from the interferometric images.For calculating M lobe we assumed that the CO wing emission is optically thin and its excitation is in LTE. Wekeep using the CO/H abundance ratio of 10 − (Dickman 1978) for a comparison with previous papers. We adoptedthe range of the excitation temperature to be between 7.3 K and 22.6 K. Here, the lower limit is given by adding thetemperature of the cosmic microwave background emission to the peak brightness temperature of the CO spectrum(Figure 6), whereas the upper limit is the gas temperature described in § T ex of thegas over the regions enclosed by the 3 σ contours of the blue- and redshifted lobes to be 20.4 ± ± T ex of the N H + (1–0) lines, which generally probe a static dense core gas, traces the temperatures of the outflowinggas where the CO transitions are excited, it is reasonable to adopt such a T ex range rather than using a single value.By integrating the lobe emission shown in Figure 9, we estimated M lobe to be of the order of 10 − − − M (cid:12) . Thelobe masses would be underestimated for the following two reasons. One is that the lobe emission was assumed to beoptically thin. The other is the uncertainty in defining the boundary velocities between the outflowing gas and theambient quiescent gas due to the effect of the resolve-out and/or the self-absorption ( § τ dyn and M lobe values. Notice that the majority of the red lobe mass isattributed to the “red IV” one.After obtaining M lobe , the mass-loss, momentum rates, and mechanical luminosity are estimated to be ˙ M outlow ∼ − − − M (cid:12) yr − , ˙ F outlow ∼ − − − M (cid:12) km s − yr − , and ˙ L outlow ∼ − − − L (cid:12) , respectively (Table3). The “red IV” lobe is one order of magnitude more powerful than the blue lobes (see F outflow in Table 3). In § F outflow values suggests that the “red HV” lobe is too powerless to drive the “red IV”lobe.We argue that the outflow lobes in the GF 9-2 core is clearly one of the smallest, the least massive, and the leastpowerful outflows when we compare with other sources compiled in e.g., Beuther et al. (2002); Takahashi et al. (2008);Takahashi & Ho (2012); Bally (2016). DISCUSSION4.1.
Excluding “Class 0 proto-brown dwarf” and Wide Binary Scenarios Furuya et al.
Considering an upper limit of the luminosity reported in the previous work ( L bol < .
3; Wiesemeyer 1997), Palau etal. (2014) categorized this object as a Very Low Luminosity Objects (VeLLOs; see Kauffmann et al. 2005; di Francescoet al. 2007, for definitions) more specifically, a “class 0 proto-brown dwarf”. However, this interpretation may beunlikely because of the “envelope infall rate” with the order of 10 − M (cid:12) yr − ( § M acc ∼ − − − M (cid:12) yr − for VeLLOs discussed in e.g., Pineda et al. (2011). It should benoticed that this comparison should be made among infall rates measured in the core-scale at 0.1–0.01 pc, not in theenvelope-disk scale at ∼ < AU, because majority of the previous works used data obtained with the core-scale.In the previous work we pointed out that there is a positional offset of 6 (cid:48)(cid:48) between the 3 mm source and the peak ofthe N H + (1–0) emission (see Figure 15 in Paper I), and proposed that the offset can be interpreted as the presence ofa binary system. Clearly such a wide binary scenario is rejected by the higher resolution SMA observations. However,we do have neither positive evidence to support the presence of a closer binary whose separation is smaller than ourbeam sizes ( ∼ <
400 AU) nor negative one to reject such a possibility. Therefore, we keep our assumption that theoutflow is driven by a single object to make our discussion as simple as possible.4.2.
Age of the Outflow Driving Source
It should be noticed that the range of dynamical time scale of the outflow lobes (500 – 2000 yrs; § – 10 yrs; e.g., Bachiller 1996; Bontemps et al. 1996;Arce & Sargent 2006; Curtis et al. 2010; Velusamy et al. 2014). Contrary to the red lobes, the blue lobes seem to be afew times older than the red ones, but their τ dyn values are still comparable to those of the youngest class 0 sources.Figure 12 demonstrates the compactness of the outflow system by a comparison of the size scales between the0.1 pc-scale molecular cloud core and the 1000 AU-scale circumstellar structure, i.e., envelope. Here the cloud core isrepresented by the H CO + (1–0) line whereas the envelope by the 350 µ m continuum emission (Paper I). The totalextent of the outflow is about 1 /
40 of that of the cloud core and is about 1 / µ m envelope, suggestive of on-goingdisk-mediated accretion. This orthogonality is not affected by the absolute position accuracy of the 350 µ m image( ∼ < (cid:48)(cid:48) ; Paper I). Hence, if an edge-on disk exists, its elongation should be almost parallel to that of the envelope.In addition to the outflow dynamical time scale, we can obtain another constraint on the source age from the size ofthe free-fall region centered on a forming protostar. Despite the presence of the compact outflow, in the radial columndensity profile, N ( r ) (Figure 11 in Paper I), we identified that the best-fit power-law profile with an index of − ∼ <
600 AU and a free-fall profile with an index of − / r ∼ >
600 AU. Because the gasmotions over the core is well described by the extended Larson-Penston’s solution for t >
0, i.e., the LPH solution(Papers I and III), the absence of an N ( r ) ∝ r − / profile, i.e., ρ ( r ) ∝ r − / in volume mass density, in the r ∼ >
600 AUregion indicates that free-fall region around the protostar has not yet expanded up to the radius of ∼
600 AU, yieldingto an upper limit of the central free-fall region ( r ff ∼ <
600 AU). Recall that we conservatively adopted r ff ∼ < (cid:48)(cid:48) separation-binary interpretation ( § r ff ∼ <
600 AU and T = 20 K ( § r ff ∼ < T = 10 K used in Paper I, the upper limit of the elapsed time since the first kernel of amass has formed is updated as, τ ∗ = r ff (2 ∼ c iso (1) ∼ < (4 ± × yrs (cid:18) T
20 K (cid:19) − / (cid:16) r ff
600 AU (cid:17) . (2)The revised upper limit is consistent with the dynamical time scales of the red and blue lobes (see Table 3). Last, the r ff ∼ <
600 AU does not contradict with the envelope inner radius (corresponding to R diskmax in Table 5; described in § “Stellar Mass” of the Outflow Driving Source Next we attempt to estimate the stellar mass using the mass infall rate in the envelope and the stellar age. Consideringthe two facts that the gas is globally ( r ∼ < (cid:48)(cid:48) , i.e., r ∼ < ρ ( r ) ∝ r − profile is identified for r ∼ >
600 AU ( § M sphinf = 2 . × − M (cid:12) yr − ) is considered to be valid in the region of 600 ∼ < r/ [AU] ∼ < ∼ < first core candidate.However, this spherical infall rate would not hold for the inner r ∼ <
600 AU region. This is because the compactoutflow has already launched, suggestive of non-spherical disk-mediated accretion onto the forming star due to thedispersal of the envelope by the outflow ( § M acc = f ˙ M sphinf for r ∼ <
600 AU where an efficiency coefficient f is written as (4 π − Ω outflow ) / (4 π ) with a total solid angle ofthe blue and red lobes viewed from the central star (Ω outflow ). Assuming the shapes of the bipolar outflow lobes areconical, we calculated the total solid angle of the lobes as Ω outflow = 2 · π [1 − cos( θ/ ∼ . θ ∼ ◦ which is the widest opening angle of the “red IV” lobe (Table 3). We obtained f ∼ .
6, leading to the non-sphericalsteady accretion rate of, ˙ M acc = f ˙ M sphinf ∼ × − (cid:18) f . (cid:19) M (cid:12) y r − . (3)Then the object may have acquired a stellar mass of, M ∗ = ˙ M acc τ ∗ = f ˙ M sphinf τ ∗ ∼ < . M (cid:12) (cid:18) f . (cid:19) (cid:32) ˙ M sphinf . × − M (cid:12) yr − (cid:33) (cid:18) τ ∗ × yrs (cid:19) , (4)since the formation of the central point source. Here the upper limit in Eq.(4) is due to the upper limit in the τ ∗ estimate [Eq.(2)]. Notice that the M ∗ value does not change significantly even if one consider the range of the outflowdynamical time scales, τ dyn , as the age of the object.We must keep in mind that the current accretion/outflow rates may much differ from the time-averaged rates duringthe formation process of the possible disk-outflow system or a possibility that an accretion rate from the envelopeonto a putative disk should differ from an accretion rate from the disk onto the central object. However, because suchmore rigorous discussion is beyond the scope of this paper, we assume that the accretion rates onto the disk and theforming star are equal to each other and steady for further discussion.4.4. Spectral Energy Distribution of the Outflow Driving Source
In order to obtain further constraints on the nature of the source, we produced Figure 13 where spectral energydistribution (SED) of the outflow driving source between the mm bands and Spitzer bands is presented. The modelSEDs were obtained by using a sedfitter tool by Robitaille et al. (2006) together with large sets of SED modelsprovided by Robitaille (2017). As emphasized by the author, we have to keep in mind the limitations of the tool andthe model sets because they provide simplified SED model sets to search for a first-order-of-approximation pictureof a YSO of interests among the various combinations of physical parameters characterizing a YSO. Each model setin Robitaille (2017) allows us to search for the best-fit solution in a wider range of the model parameters than theoriginal one (Robitaille et al. 2006). Each set consists of a protostar or a pre-main sequence (PMS) star surroundedby an accretion disk and a rotationally infalling envelope with a bipolar outflow, and also considers the scatteringand reprocessing of the stellar radiation by dust. Moreover, some model sets include a bipolar cavity dispersed by anoutflow. In each model set, model SEDs were computed at 10 different viewing angles by convolving a frequency-filterfunction and an aperture size at each measurement; we used the default ones for the Spitzer measurements, and addedthose for our CSO, SMA, CARMA, and OVRO observations.To fit the photometry results (Tables 2 and 4), we set a distance range of 200 ±
20 pc, and used an extinction curveby Whitney et al. (2003, 2004), allowing to estimate the foreground A v as well (see the A v value in each panel ofFigure 13). We dealt with the 70 µ m flux density as an upper limit because the HPBW of the aperture encompassingthe source is much larger than those imaged by SMA, CARMA and OVRO interferometers. Notice that the numberof the free parameters in the model sets (see Table 2 in Robitaille 2017) ranges between 2 and 12, whereas we havefifteen data points ( n data = 15), including the upper limits. A caution must be used for the WISE 22 µ m and Spitzer24 µ m data whose aperture sizes are comparable to that for the 350 µ m (Table 4). This is because we did not considerthem as upper limits because the 22 and 24 µ m emission most likely comes from the central compact component(s)rather than from the extended envelope traced by the 70 and 350 µ m emission. Indeed, we found that all the modelsgave us unrealistic SED fits when we set the 22 and 24 µ m fluxes as upper limits.As mentioned above, we attempted all the model sets summarized in Table 2 of Robitaille (2017), and selected themost reasonable models from each model set based on the criterion of χ − χ < n data . In practice, we rejected2 Furuya et al. the 16 out of 18 model sets in Table 2 of Robitaille (2017) because these model sets failed to reproduce the observedSEDs, whereas the remaining two model sets of spu-smi and spubhmi gave reasonable fits. Here the first characterof s in the model-set names denotes that they commonly consider emission from a central s tar, the second one of p does p assive disk, the third u does so-called Ulrich-like envelope (an envelope with radial power-law dependency of ∼ / R c , 1/2 inside; Ulrich 1976). The fourth b in spubhmi indicates that this modelset considers a b ipolar-outflow cavity, whereas not in spu-smi . The fifth h indicates that the spubhmi model set leavesradius of the inner hole produced by the outflow as a free-parameter, whereas the fifth s means that the spu-smi adopted a dust-sublimation radius as an inner hole radius. The last two characters of mi indicates that the both modelsets commonly considered the ambient medium (m) together with the interstellar dusts (i) , as described earlier.On the basis of the inferred physical parameters, we rejected spu-smi because the stellar radius suggested by themodel sets ( R ∗ = 0 . R (cid:12) ) is too small and the stellar surface temperature ( T ∗ ∼ . × K) is too high for aprotostar ( § spubhmi yields the decent SED fit (Figure 13) with the reasonable physical parameters(Table 5). The model gave plausible values of R ∗ ∼ R (cid:12) , and T ∗ ∼ × K as well as those characterizing theenvelope and disk. The inferred inclination angle of i ∼ ◦ from the two models gives a constraint in our analysis in § § θ ∼ ◦ may not contradictwith our estimate in Table 3 when we consider the uncertainty of our measurement (∆ θ ∼ ◦ ; § R disk value does not affect the estimate of the protostar’s age based on the radial density profile ( § T ∗ ∼ ∼
200 AU radius disk with an infalling envelope. Moreover, we argue that the outflow system whoselobe axis is almost parallel to the plane of sky may not have completed dispersing its natal cloud core.4.5.
Evolutionary Stage of the Outflow Driving Source: Constraints from the Outflow Properties
The SED analysis suggested that the outflow is viewed almost from its side (see e.g., Figure 7 in Tomisaka & Tomida2011). In this subsection, we discuss this geometry along with the outflow properties ( § Terminology of the Outflow Velocity and Inclination Angle Dependency on the Flow Velocity and Dynamical Timescale
Before addressing the outflow geometries, we must clarify the terminology of the outflow velocities when we compareto those in theoretical studies. Most of the theoretical studies refer to “high velocity” as those having v flow of theorder of 10 − km s − and “low velocity” as 1 km s − ∼ < v out ∼ <
10 km s − (e.g., Bachiller 1996). Note that thesetheoretical velocities are the 3D one. In contrast, observers argue outflow velocities projected onto the line of sight.Furthermore there is no widely accepted unique definition of the velocity range for “high” and “low” velocities even ifwe limit to discussing outflows traced by low- J CO lines. Considering these limitations and the uncertainties in ourobservations, we do not directly compare observed outflow velocities and those in numerical simulations. Therefore,we keep using the words of the IV and HV defined in § v flow / cos i and the “real” dynamical time scale of τ dyn cos i/ sin i for the four lobes strongly depends on i . In addition, the opening angles of the lobes are rather large of θ ∼ ◦ − ◦ with the uncertainties of ∆ θ ∼ ◦ (Table 3). We therefore discuss two representative cases of the pole-on outflow at i = 20 ◦ and the edge-on one at i = 70 ◦ , as done in e.g., Belloche et al. (2006).4.5.2. “Pole-on” vs. “Edge-on” Outflows Based on the Outflow Age and the Elapsed Time of Mass AccretionA Pole-on Outflow (e.g., i = 20 ◦ ) — Figure 11 suggests that the 3D-velocities of the pole-on lobes range from 3 to 9km s − , leading to the corrected ages of τ dyn cos i/ sin i ∼ − yrs. In the pole-on geometry the large spatial-overlapping between the blue and red lobes together with the absence of the redshifted CO emission to the northeastof the continuum source is explained to some extent. This is because the outflow gas moving away from us along theoutflow axis should be obscured by the circumstellar materials.However, there are two major caveats in the pole-on outflow hypothesis. If the pole-on is the case, all the lobesshould show more roundish morphology to form a symmetric roundish lobe pairs, but the “red HV” and blue lobesare elongated. It is also impossible to explain the southwest island (see Figure 9b) along with the pole-on hypothesis.The other major caveat is that the i -corrected τ dyn values of the order of 10 − yrs would be larger than the sourceage of τ ∗ ∼ < × yrs in Eq.(2). We therefore conclude that a pole-on geometry is unlikely. ∼ < An Edge-on Outflow (e.g., i = ◦ ) — If we are observing an edge-on outflow, the 3D-velocities of the lobes would beapproximately 10 – 30 km s − (Figure 11), leading to the lobe ages of τ dyn cos i/ sin i ∼ < × yrs. It is well-knownthat protostars, i.e., class 0 sources, show highly-collimated high-velocity outflows whose dynamical time scale ranges10 – 10 yrs (e.g., Gueth & Guilloteau 1999; Santangelo et al. 2015). The inferred age range of the GF 9-2 outflowagrees with the range, and is considered to represent elapsed time of the ongoing accretion process ( § θ ∼ ◦ − ◦ (Table 3).However, the edge-on hypothesis does not explain why the morphology and velocity structure of the blue and redlobes are not symmetric; we should detect redshifted lobe(s) to the northeast of the continuum source. Such asymmetrymay be explained by “realistic” MHD simulations which includes turbulence (e.g., Matsumoto & Hanawa 2011).Keeping all the strong and weak points in mind, it is possible that the central object has experienced its secondcollapse τ dyn cos i/ sin i ∼ < × yrs ago. In this interpretation, the poorly collimated “red IV” lobe (Figure 10)would represent a remnant outflowing gas which is the fossil of a lobe that had been driven by the object when it hadbeen at the first core stage, and the “red HV” lobe (Figure 10) is considered as a fresh lobe currently being driven bythe protostar. If this is the case, the dynamical time scale of the “red IV” lobe ( τ dyn cos i/ sin i ∼
500 yrs for i ∼ ◦ ;Figure 11) may be interpreted as the duration of the first core stage. In this interpretation, the small β of 0.4 ± § Summary of the Subsection
Considering the geometry of the compact outflow lobes in conjunction with the τ dyn and τ ∗ estimates, we suggestthat the outflow geometry is edge-on which reconciles with the results from the SED analysis. SUMMARYInterferometric observations of the CO (3–2) line and the continuum emission at 3.3 mm, 1.1 mm, and 850 µ mbands were carried out towards the deeply embedded protostar at the center of the dense molecular cloud core GF 9-2using CARMA and SMA. Furthermore we analyzed the Spitzer and
W ISE satellite images and the single-dish CO(3–2) spectra previously taken with CSO. The main findings of this research are summarized as follows.1. At the center of the cloud core, we detected a compact continuum source which is considered to be representinga circumstellar envelope at 1.1 mm and 850 µ m bands. The beam-deconvolved effective radius, R eff , of thecontinuum source was measured to be 250 ±
80 AU at the 1.1 mm and 850 µ m bands where we attained the linearresolution of ∼
400 AU at the distance of the source (200 pc).2. Towards the position of the submm source, an infrared source, designated as WISE J205129.83+601838, is clearlydetected at wavelengths between 70 µ m and 3.4 µ m in the Spitzer MIPS and IRAC images as well as WISE ones.This detection clearly indicates that the object is at a protostar phase. Our SED analysis using the sedfitter tool (Robitaille 2017) suggested its surface temperature of ∼ ∼ R (cid:12) .3. Our spectroscopic imaging of the CO (3–2) line clearly detected a 1000-AU scale molecular outflow systemdriven by the continuum source. The system exhibits collimated- and poorly collimated redshifted outflowlobes with the line-of-sight velocities of ∼
10 km s − and ∼ − , respectively. The blueshifted exhibits acollimated lobe extending both towards the northeast and southwest of the central source. Our analysis showedthat the outflow lobes are one of the youngest (dynamical time scales of ∼
500 – 2000 yrs) and the least powerful(momentum rates of ∼ − − − M (cid:12) km s − yr − ) ones so far detected. A comparison between the continuumand outflow maps suggests that the innermost part ( r ∼ < independently driven. Based on the outflow morphology, velocity structure of the lobes, compar-isons with outflow models and the results from the SED analysis, we concluded that the outflow axis is not farfrom parallel to the plane of the sky, i.e., the edge-on geometry.4 Furuya et al.
5. Although the outflow properties agree with those measured in the VeLLOs and the “Class 0 proto-brown dwarfs”,we excluded such interpretations on the basis of the large spherical infall rate of the order of 10 − M (cid:12) yr − .We also excluded the wide-binary interpretation based on the previous lower resolution data because only aunresolved ( ∼ <
400 AU) continuum emission was detected by the interferometric observations.6. We argued that it has not passed τ ∗ ∼ < (4 ± × years since the protostar formed at the center of the cloudcore. This is because a radial volume density profile with a form of ρ ( r ) ∝ r − / , which proves the presence ofa freely falling gas towards the central object, was not identified outside the r (cid:39)
600 AU region in radius whichhas a consistency with the disk radius inferred from the SED analysis. The upper limit of the protostellar agesuggests that the total mass accreted onto the central object is M ∗ ∼ < . M (cid:12) .Given the uniqueness of the source properties, follow-up high-resolution and high-sensitivity observations alongwith simulation studies are required towards a more complete understanding of the physics in a low-mass protostarformation process.The authors sincerely acknowledge the anonymous referee whose comments significantly helped to improve quality ofthis paper, especially critical comments on the Spitzer data. R. S. F. gratefully acknowledges John M. Carpenter andAndrea Isella for their generous help and discussion at the CARMA observations and the data reduction process. R. S.F. also sincerely thanks Masahiro N. Machida, Tomoyuki Hanawa and Shu-ichiro Inutsuka for fruitful discussion, andTakeshi Inagaki for data analysis with Python. This work was partially supported by the JSPS Institutional Programfor Young Researcher Overseas Visits ( Wakate Haken ) at Subaru Telescope of National Astronomical Observatoryof Japan (NAOJ) for R. S. F. and H. S. and the AWA support program at Tokushima University for R. S. F. Dataanalyses in this work were partly carried out on the computer system operated by Subaru Telescope and that byAstronomy Data Center of NAOJ. The authors gratefully acknowledge all the staff at CARMA, SMA, CSO and theSpitzer Science Center, and the MIR software group at CfA, the AIPS software group at NRAO and the GILDASsoftware group at IRAM.
Facilities:
SMA, CARMA, OVRO mm-array, Spitzer Science Telescope, CSO 10.4 m telescopeAPPENDIX A. EXCITATION TEMPERATURE MAP OF THE N H + (1–0) EMISSIONAs reported in Paper I, we combined the visibility data of the N H + line taken with the OVRO mm-array andthe single-dish Nobeyama 45 m telescope (see § H + transition (see § T ex and optical depth of the lines withrelatively high accuracy. In addition, such a combined image allows us to analyze the spatial structure of the gas withan angular resolution that can be achieved by interferometers and the analysis is free from the “missing flux” problem.In § H + (1–0) emission over a circular region with a radiusof R eff = 250 ±
80 AU centered at the submm source [ (cid:104) T ex (cid:105) = 22 . ± . T ex map (Figure 14) which was used to produce the column density map of N H + shown in Figure 9b of Paper I. Inaddition, we measured the mean T ex of 20.4 ± ± σ level contours of theoutflow lobes (see Figure 14b). REFERENCES Andr´e, P., Ward-Thompson, D., & Barsony, M. 1993, ApJ,406, 122Andr´e, P., Ward-Thompson, D., & Barsony, M. 2000,Protostars and Planets IV, 59 Andr´e, P., Ward-Thompson, D., & Motte, F. 1996, A&A,314, 625Andrews, S. M., & Williams, J. P. 2007, ApJ, 659, 705Andrews, S. M., & Williams, J. P. 2007, ApJ, 671, 1800 ∼ < Arce, H. G., & Sargent, A. I. 2006, ApJ, 646, 1070Bachiller, R. 1996, ARA&A, 34, 111Banerjee, R., & Pudritz, R. E. 2006, ApJ, 641, 949Bastien, P., Jenness, T., & Molnar, J. 2005, AstronomicalPolarimetry: Current Status and Future Directions, 343,69Bate, M. R., Tricco, T. S., & Price, D. J. 2014, MNRAS,437, 77Beckwith, S. V. W., Henning, T., & Nakagawa, Y. 2000,Protostars and Planets IV, 533Belloche, A., Parise, B., van der Tak, F. F. S., et al. 2006,A&A, 454, L51Bally, J. 2016, ARA&A, 54, 491Beuther, H., Schilke, P., Sridharan, T. K., et al. 2002,A&A, 383, 892Bonnor, W. B. 1956, MNRAS, 116, 351Bontemps, S., Andre, P., Terebey, S., & Cabrit, S. 1996,A&A, 311, 858Boss, A. P., & Yorke, H. W. 1995, ApJL, 439, L55Caselli, P., Benson, P. J., Myers, P. C., & Tafalla, M. 2002,ApJ, 572, 238Chen, X., Arce, H. G., Zhang, Q., et al. 2010, ApJ, 715,1344Ciardi, D. R., Woodward, C. E., Clemens, D. P., Harker,D. E., & Rudy, R. J. 1998, AJ, 116, 349Ciardi, D. R., Woodward, C. E., Clemens, D. P., Harker,D. E., & Rudy, R. J. 2000, AJ, 120, 393Commer¸con, B., Launhardt, R., Dullemond, C., & Henning,T. 2012, A&A, 545, A98Crutcher, R. M., Nutter, D. J., Ward-Thompson, D., &Kirk, J. M. 2004, ApJ, 600, 279Curtis, E. I., Richer, J. S., Swift, J. J., & Williams, J. P.2010, MNRAS, 408, 1516Dickman, R. L. 1978, ApJS, 37, 407di Francesco, J., Evans, N. J., II, Caselli, P., et al. 2007,Protostars and Planets V, 17Dunham, M. M., Crapsi, A., Evans, N. J., II, et al. 2008,ApJS, 179, 249-282Dunham, M. M., Chen, X., Arce, H. G., et al. 2011, ApJ,742, 1Dunham, M. M., & Vorobyov, E. I. 2012, ApJ, 747, 52Dunham, M. M., Stutz, A. M., Allen, L. E., et al. 2014,Protostars and Planets VI, 195Ebert, R. 1955, ZA, 37, 217Emerson, J. P. 1988, NATO Advanced Science Institutes(ASI) Series C, 241, 21Enoch, M. L., Lee, J.-E., Harvey, P., Dunham, M. M., &Schnee, S. 2010, ApJL, 722, L33Fazio, G. G., Hora, J. L., Allen, L. E., et al. 2004, ApJS,154, 10 Friberg, P., Bastien, P., Berry, D., et al. 2016, Proc. SPIE,9914, 991403Friesen, R. K., Di Francesco, J., Bourke, T. L., et al. 2014,ApJ, 797, 27Furuya, R. S., Kitamura, Y., Wootten, H. A., Claussen,M. J., & Kawabe, R. 2001, ApJL, 559, L143Furuya, R. S., Kitamura, Y., Wootten, A., Claussen, M. J.,& Kawabe, R. 2003, ApJS, 144, 71Furuya, R. S., Kitamura, Y., & Shinnaga, H. 2006, ApJ,653, 1369 (Paper I)Furuya, R. S., Kitamura, Y., & Shinnaga, H. 2008, PASJ,60, 421(Paper II)Furuya, R. S., Kitamura, Y., & Shinnaga, H. 2009, ApJ,692, 96 (Paper III)Furuya, R. S., Kitamura, Y., & Shinnaga, H. 2014, ApJ,793, 94 (Paper IV)Gueth, F., & Guilloteau, S. 1999, A&A, 343, 571Hollenbach, D., Elitzur, M., & McKee, C. F. 1993,Astrophysical Masers, 412, 159Hollenbach, D., Elitzur, M., & McKee, C. F. 2013, ApJ,773, 70Hirano, N., & Liu, F.-c. 2014, ApJ, 789, 50Hirashita, H., & Omukai, K. 2009, MNRAS, 399, 1795Hennebelle, P., & Teyssier, R. 2008, A&A, 477, 25Hunter, C. 1977, ApJ, 218, 834Kauffmann, J., Bertoldi, F., Evans, N. J., II, & C2DCollaboration 2005, Astronomische Nachrichten, 326, 878Kenyon, S. J., Hartmann, L. W., Strom, K. M., & Strom,S. E. 1990, AJ, 99, 869Kitamura, Y., Momose, M., Yokogawa, S., et al. 2002, ApJ,581, 357Kudoh, T., & Shibata, K. 1995, ApJL, 452, L41Larson, R. B. 1969, MNRAS, 145, 271Looney, L. W., Mundy, L. G., & Welch, W. J. 2000, ApJ,529, 477Machida, M. N., Inutsuka, S.-i., & Matsumoto, T. 2010,ApJ, 724, 1006Machida, M. N. 2014, ApJL, 796, LL17Mardones, D., Myers, P. C., Tafalla, M., et al. 1997, ApJ,489, 719Masunaga, H., Miyama, S. M., & Inutsuka, S.-i. 1998, ApJ,495, 346Masunaga, H., & Inutsuka, S.-i. 2000, ApJ, 531, 350Matsumoto, T., & Hanawa, T. 2011, ApJ, 728, 47Maureira, M. J., Arce, H. G., Dunham, M. M., et al. 2017,ApJ, 838, 60Mellon, R. R., & Li, Z.-Y. 2008, ApJ, 681, 1356Momose, M., Ohashi, N., Kawabe, R., Nakano, T., &Hayashi, M. 1998, ApJ, 504, 314 Furuya et al.
Momose, M., Ohashi, N., Kawabe, R., Hayashi, M., &Nakano, T. 1996, ApJ, 470, 1001Monin, J.-L., Pudritz, R. E., & Lazareff, B. 1996, A&A,305, 572Natta, A., Palla, F., Preite-Martinez, A., & Panagia, N.1981, A&A, 99, 289Ogino, S., Tomisaka, K., & Nakamura, F. 1999, PASJ, 51,637Omukai, K. 2007, PASJ, 59, 589Ormel, C. W., Min, M., Tielens, A. G. G. M., Dominik, C.,& Paszun, D. 2011, A&A, 532, A43Ossenkopf, V., & Henning, T. 1994, A&A, 291, 943Ouyed, R., & Pudritz, R. E. 1997, ApJ, 484, 794Palau, A., Zapata, L. A., Rodr´ıguez, L. F., et al. 2014,MNRAS, 444, 833Penston, M. V. 1969, MNRAS, 144, 425Pezzuto, S., Elia, D., Schisano, E., et al. 2012, A&A, 547,A54Pineda, J. E., Arce, H. G., Schnee, S., et al. 2011, ApJ, 743,201Poidevin, F., & Bastien, P. 2006, ApJ, 650, 945Planck Collaboration, Abergel, A., Ade, P. A. R., et al.2011, A&A, 536, A25Preibisch, T., Ossenkopf, V., Yorke, H. W., & Henning, T.1993, A&A, 279, 577Price, D. J., Tricco, T. S., & Bate, M. R. 2012, MNRAS,423, L45Pudritz, R. E., & Norman, C. A. 1986, ApJ, 301, 571Richer, J. S., Shepherd, D. S., Cabrit, S., Bachiller, R., &Churchwell, E. 2000, Protostars and Planets IV, 867Ricci, L., Testi, L., Natta, A., et al. 2010, A&A, 512, A15Ricci, L., Testi, L., Natta, A., & Brooks, K. J. 2010, A&A,521, A66Rieke, G. H., Young, E. T., Engelbracht, C. W., et al. 2004,ApJS, 154, 25Robitaille, T. P., Whitney, B. A., Indebetouw, R., Wood,K., & Denzmore, P. 2006, ApJS, 167, 256Robitaille, T. P. 2017, A&A, 600, A11Saigo, K., & Tomisaka, K. 2006, ApJ, 645, 381Saigo, K., Tomisaka, K., & Matsumoto, T. 2008, ApJ, 674,997-1014Schneider, S., & Elmegreen, B. G. 1979, ApJS, 41, 87Shu, F. H. 1977, ApJ, 214, 488Simon, M., Dutrey, A., & Guilloteau, S. 2000, ApJ, 545,1034Stahler, S. W., Shu, F. H., & Taam, R. E. 1980, ApJ, 242,226 Santangelo, G., Murillo, N. M., Nisini, B., et al. 2015,A&A, 581, A91Shibata, K., & Uchida, Y. 1985, PASJ, 37, 31Stahler, S. W., Shu, F. H., & Taam, R. E. 1980, ApJ, 241,637Takahashi, S., Saito, M., Ohashi, N., et al. 2008, ApJ, 688,344-361Takahashi, S., & Ho, P. T. P. 2012, ApJL, 745, L10Tobin, J. J., Hartmann, L., Chiang, H.-F., et al. 2011, ApJ,740, 45Tobin, J. J., Stutz, A. M., Megeath, S. T., et al. 2015, ApJ,798, 128Tobin, J. J., Stutz, A. M., Manoj, P., et al. 2016, ApJ, 831,36Tomida, K., Machida, M. N., Saigo, K., Tomisaka, K., &Matsumoto, T. 2010, ApJL, 725, L239-L244Tomida, K., Tomisaka, K., Matsumoto, T., et al. 2013,ApJ, 763, 6Tomida, K., Okuzumi, S., & Machida, M. N. 2015, ApJ,801, 117Tomisaka, K. 1998, ApJL, 502, L163Tomisaka, K. 2002, ApJ, 575, 306Tomisaka, K., & Tomida, K. 2011, PASJ, 63, 1151Tomisaka, K. 2014, ApJ, 785, 24Tsitali, A. E., Belloche, A., Commer¸con, B., & Menten,K. M. 2013, A&A, 557, A98Ulrich, R. K. 1976, ApJ, 210, 377Yoshida, A., Kitamura, Y., Shimajiri, Y., & Kawabe, R.2010, ApJ, 718, 1019Vaytet, N., & Haugbølle, T. 2017, A&A, 598, A116Velusamy, T., & Langer, W. D. 1998, Nature, 392, 685Velusamy, T., Langer, W. D., & Thompson, T. 2014, ApJ,783, 6Walker, C. K., Lada, C. J., Young, E. T., Maloney, P. R., &Wilking, B. A. 1986, ApJL, 309, L47Wiesemeyer, H. 1997, Ph.D. dissertation. Univ. Bonn,Bonn, Germany, 1997Wilner, D. J., & Welch, W. J. 1994, ApJ, 427, 898Winkler, K.-H. A., & Newman, M. J. 1980, ApJ, 236, 201Whitney, B. A., Wood, K., Bjorkman, J. E., & Wolff, M. J.2003, ApJ, 591, 1049Whitney, B. A., Indebetouw, R., Bjorkman, J. E., & Wood,K. 2004, ApJ, 617, 1177Whitworth, A., & Summers, D. 1985, MNRAS, 214, 1Wright, E. L., Eisenhardt, P. R. M., Mainzer, A. K., et al.2010, AJ, 140, 1868 ∼ < Figure 1.
Interferometric continuum emission maps towards the center of the GF 9-2 low-mass star forming core at frequenciesof (a) 91 GHz ( λ = 3.3 mm) taken with the CARMA, (b) 268 GHz (1.12 mm), (c) 280 GHz (1.07 mm), (d) 345 GHz (870 µ m),and (e) 357 GHz (840 µ m) taken with the SMA. All the contours, except for the central thick one in (a), are drawn with the 3 σ intervals, starting from the 3 σ levels, where the 1 σ levels mean the RMS noise levels of the images. The single thick contour inthe panel (a) represents the 5 σ level. The horizontal bar in (a) indicates the linear size scale of 0.01 pc assuming the distanceto the source to be 200 pc. See Table 1 for the image sensitivity and the synthesized beam sizes which are indicated by thehatched ellipses at the bottom-right corners of the five panels. Figure 2.
Plot of the mm and submm continuum spectrum of the compact source embedded in the GF 9-2 cloud core. Theflux densities presented in this panel are summarized in Table 2. The broken-line indicates the best-fit power law spectrum witha form of S ν ∝ ν α where the best-fit spectral index is α = 2 . ± . § Figure 3.
Continuum emission images in unit of MJy sr − towards the center of the GF 9-2 core at wavelength of (a) 70 µ m,(b) 24 µ m, (c) 8.0 µ m, (d) 5.8 µ m, (e) 4.5 µ m, and (f) 3.6 µ m taken with the Spitzer
MIPS filters [(a) and (b)] and IRAC filters[(c)–(f)]. The center positions of all the images are reprojected to the position of the 840 µ m (357 GHz) source ( § Spitzer images are described in § Furuya et al.
Figure 4.
Velocity channel maps of the CO (3–2) emission towards the submm continuum source (Figure 1) in the GF 9-2cloud core. The stars in these panels indicate the position of the continuum source. Each channel map is averaged over a 0.4km s − bin whose central velocity in unit of km s − is shown at the top-left corner. The size of the each panel is 24 (cid:48)(cid:48) × (cid:48)(cid:48) ,corresponding to 4800 AU × d = 200. All the contours are the 3 σ intervals starting from the 3 σ level where 1 σ corresponds to 122 mJy beam − . The systemic velocity of the cloud core is v LSR = − .
48 km s − (Paper I). ∼ < Figure 5.
Total integrated intensity map of the CO (3–2) emission produced from the velocity channel maps (Figure 4).The emission is integrated over − . ≤ v LSR / km s − ≤ +7 . σ intervals starting from the +3 σ level where the 1 σ RMS noise level is 0.62 Jy beam − km s − . The synthesized beamsize is indicated by the hatched ellipse at the bottom-left corner. The horizontal bar at the top-right corner indicates the linearsize scale of 0.01 pc. This map is used for producing the SMA CO spectrum shown in Figure 6
Figure 6. CO (3–2) emission spectrum towards the center of the GF 9-2 cloud core observed with the SMA in brightnesstemperature ( T b ) scale in units of K (see § σ contour of a total integrated intensity map(Figure 5) in order to fully detect high-velocity tails. The solid angle of the 3 σ region is Ω s = 6 . × − sr, i.e., an area A of 27.3 arcsec , corresponding to an effective radius of R eff = (cid:112) A/π = 2 . (cid:48)(cid:48)
9, i.e., 580 AU at d of 200 pc. The green and blackvertical dashed-lines at v LSR = − .
48 km s − and +5 . − , respectively, show the systemic velocity of the cloud (Paper I)and the velocity of the H O maser emission at 22 GHz (Furuya et al. 2003). The horizontal blue and red thick bars under thespectrum indicate the velocity ranges for producing the outflow lobe maps shown in Figure 9, and the orange and magenta thinbars under the red one indicate those for Figure 10. Furuya et al.
Figure 7.
Comparison of the CO (3–2) line spectra taken with the single-dish CSO 10.4 m telescope (thick histogram) andwith the SMA interferometer (yellow-shaded thin histogram). Both the spectra are shown in brightness temperature scale inunits of K by averaging over the CSO beam area with θ HPBW = 22 (cid:48)(cid:48) centered on the submm continuum source. The CSOspectrum towards the center (see Figure 8) was Hanning-smoothed to increase signal-to-noise (S/N) ratio. The horizontaldashed line indicates the the 3 σ noise level for the CSO spectrum. The RMS 1 σ noise levels are 60 mK with a velocity resolutionof 1.6 km s − for the CSO spectrum and 16 mK with a resolution of 0.4 km s − for the SMA one. The vertical green dashedline indicates the systemic velocity of the cloud (Paper I). The horizontal blue and red bars under the spectrum are the same asthe thick ones in Figure 6, and indicate the velocity ranges for producing the outflow lobe maps shown in Figure 9. Notice thatthe telescope pointing accuracy in the CSO observations was better than 5 (cid:48)(cid:48) (Paper I), and that the absolute flux calibrationsof both the CSO and SMA observations have uncertainties of ∼ Table 1.
Summary of the Interferometric Observations
Spatial Frequency Synthesized BeamArray Emission f cent ∆ f a λ b FoV c Range LAS d θ maj × θ min P.A. Sensitivity(GHz) (MHz) (mm) (arcsec) ( kλ ) (arcsec) (arcsec) (deg) (mJy beam − )CARMA Continuum 91.181 938 3.29 88 1.93 – 101.7 107 3.89 × −
84 0.32SMA Continuum 267.755 3960 1.12 38 10.3 – 102.7 20 2.63 × × × × CO(3–2) 345.796 88.4 0.867 30 7.7 – 132.3 27 2.13 × e a Total bandwidth per polarization. b Wavelength. c Field-of-view. d Largest detectable angular size scale. e Sensitivity per velocity channel width of 0.4 km s − . ∼ < Figure 8.
Nine spectra of the CO (3–2) line in the main-beam brightness temperature scale in units of K observed with a21 (cid:48)(cid:48) grid centered at the mm-continuum source of the GF 9-2 cloud core center. The data were previously taken with the CSO10.4 m telescope (Paper I; θ HPBW = 22 (cid:48)(cid:48) ), and are re-analyzed in this study ( § α, ∆ δ ) in units of arcsecond. The 1 σ RMS noise level of each spectrum in a velocity resolution of 0.4 km s − is shown in units of mK below the parenthesis. Figure 9.
Integrated intensity maps of the outflow components observed by the CO (3–2) emission (blue and red contours)overlaid on the 357 GHz continuum emission image (black contours with grayscale map) (Figure 1d). The map size is 21 (cid:48)(cid:48) × (cid:48)(cid:48) .The color contours in the two panels show (a) the blue lobe of − . ≤ v LSR / km s − ≤ − .
2, and (b) the red lobe of +0 . ≤ v LSR / km s − ≤ +7 .
0. See also the blue and red thick bars under the spectrum in Figure 6 for the velocity ranges. All thecontours are plotted with the 3 σ intervals starting from the 3 σ levels. The RMS noise levels of the blue and red maps are 10.4and 60 mJy beam − km s − , respectively. The ellipse in the box at the bottom left corner of each panel shows the size of thesynthesized beam of the CO (3–2) observations. The linear size scale of 0.01 pc is shown by the bar in the panel (a). Furuya et al.
Figure 10.
Integrated intensity maps of (a) the redshifted intermediate-velocity (“red IV”) outflow lobe and (b) the redshiftedhigh-velocity (“red HV”) one identified in the redshifted outflow lobe (Figure 9b). The map size is 18 (cid:48)(cid:48) × (cid:48)(cid:48) . The “red IV”lobe map is obtained by integrating the emission over +0 . ≤ v LSR / km s − ≤ +5 . . 0. See also the orange and magenta thin bars under the spectrum in Figure 6 for the velocity ranges. Allthe contours are plotted with the 3 σ intervals starting from the +3 σ level. The RMS noise levels of the “red IV” and “red HV”maps are 334 and 60 mJy beam − km s − , respectively. The star marks the peak position of the submm continuum emissionlocated at R. A. = 20 h m . s , Decl = 60 ◦ (cid:48) . (cid:48)(cid:48) 23 in J2000 ( § Figure 11. Plots of the outflow velocity ( v flow / cos i ; the left panel) and dynamical time scale ( τ dyn cos i/ sin i ; the right panel)of the outflow lobes as a function of the inclination angle ( i ). Here i is defined by the angle measured from the line-of-sight(l.o.s.; i = 0 ◦ ), hence a lobe with i = 90 ◦ is parallel to the sky plane. We refer to an outflow whose axis has i = 0 ◦ as apole-on outflow, whereas i = 90 ◦ as an edge-on outflow ( § § v flow / cos i and τ dyn cos i/ sin i curves in these plots pass the values in Table 3 at i = 45 ◦ . Note that the cyanand blue curves in the right panel agree with each other. ∼ < Figure 12. Comparison among the dense molecular cloud core traced by the H CO + (1–0) emission (green contour), thecircumstellar envelope observed by the 350 µ m continuum emission (black contours) and the newly detected compact molecularoutflow (the blue- and red contours; see Figure 9). The H CO + line and 350 µ m continuum emission maps are taken fromPaper I. All the contours are plotted by the 3 σ intervals starting from the 3 σ levels. The ellipses in the bottom-left and bottom-right corners are the beam size of the SHARCII bolometer for the 350 µ m imaging and the synthesized beam size for the SMA CO (3–2) line observations, respectively. The dashed circle at the center indicates the field-of-view of the SMA CO (3–2)observations (see Table 1). Figure 13. Model fitting to the GF 9-2 protostar performed by sedfitter tool using different YSO SED models in Robitaille(2017). The solid black line shows the best-fit model. See Tables 2 and 4 for the flux densities, Table 5 for the inferred parametersand, § Furuya et al. Figure 14. (a)(left) Excitation temperature ( T ex ) map of the N H + (1–0) line (color), where the 350 µ m (white contours; seePaper I) and 357 GHz (black contours; see Figure 1e in this paper) continuum emission maps are overlaid. The color bar on theright-hand side shows the temperature scale in K. The T ex map was obtained from our analysis of the N H + hyperfine structurein the combined image of the OVRO mm-array and single-dish Nobeyama 45 m telescope data, which is free from missing fluxdensities. Note that the T ex map was obtained in Paper I, but was not presented there. The contour parameters are the sameas in Figure 12 for the 350 µ m emission and in Figure 1e for the 357 GHz emission. The ellipses in the boxes at the bottomleft and right corners show the synthesized beam sizes of the 357 GHz continuum emission image ( θ maj × θ min = 2 . (cid:48)(cid:48) × . (cid:48)(cid:48) ◦ ; Table 1) and the N H + T ex map (4 . (cid:48)(cid:48) × . (cid:48)(cid:48) − ◦ ; Paper I), respectively. See Appendix A for details.(b)(right) Magnification of the central region of the plot shown in panel (a). The cyan and blue contours, respectively, presentthe 3 σ level ones of the blue and red outflow lobes shown in Figure 9. ∼ < Table 2. Summary of the Interferometric Photometry of the Continuum Emission Frequency Θ maj × Θ mina P. A. b R effc I peak ν d S ν e M csmh 2D Gaussian f σ g (MHz) (arcsec × arcsec) (degree) (AU) (mJy beam − ) (mJy) (mJy) ( M (cid:12) )89964 i × ± 15 680 ± 150 1.06 ± ± ± ± × ± 15 550 ± 450 1.62 ± ± ± j ± × ± ± 50 25.2 ± ± ± ± × ± ± 50 29.6 ± ± ± 12 0.022 ± × ± ± 110 28.6 ± ± ± ± × ± 12 260 ± 90 19.7 ± ± ± ± Note —See § a Beam-deconvolved source size obtained by task JMFIT in AIPS package with an assumption that the intensity dis-tribution of the source is approximated by a 2D elliptical Gaussian whose FWHMs along the major and minor axesare Θ maj and Θ min , respectively. b Position angle. c Effective source radius in AU. R eff is calculated from π (Θ maj × Θ min ) d = πR where d is the distance to thesource in pc. The uncertainty is estimated from the difference caused by the elliptical and circular approximations. d Peak intensity obtained from the 2D elliptical Gaussian fitting. The error in the fitting is calculated from the RMSnoise level of each image (Table 1). e Total flux density. f S ν obtained from the 2D elliptical Gaussian fitting by deconvolving the synthesized beam (see Table 1.). g S ν obtained by integrating the emission over the area enclosed by the 3 σ level contour. h “Circumstellar” mass estimated from the S ν values in the column 6. i Continuum image was taken from Paper I. All the values are obtained from the re-analysis with the same methodas that of the other band data. j Except for the weak emission elongated to the east (see Figure 1a). Furuya et al. T a b l e . P h y s i c a l P r o p e r t i e s o f t h e O u t fl o w L o b e s l / s i n i a P . A . b θ c v f l o w / c o s i d τ d y n c o s i s i n i e M l o b e f ˙ M o u t f l o w s i n i c o s i g F o u t f l o w s i n i c o s i h L o u t f l o w c o s i i L o b e . K . K . K . K . K . K . K . K ( AU )( d e g )( d e g )( k m s − )( y r s )( − M (cid:12) )( − M (cid:12) y r − )( − M (cid:12) y r − y r − )( − L (cid:12) ) B l u e : N o rt h e a s t ∼− ∼ . . . . . . × − × − B l u e : S o u t h w e s t ∼ ∼ . . . . . × − × − R e d : I V ∼ ∼ . . R e d : HV ∼ ∼ . . . . × − a O u t fl o w l o b e l e n g t h c a l c u l a t e d f r o m t h e m e a s u r e d l o b e l e n g t h s ee n i n F i g u r e s nd . H e n ce w ec o rr ec t e d f o rt h e un k n o w n i n c li n a t i o n a n g l e o f t h e o u t fl o w a x i s ( i ) b y a ss u m i n g i = ◦ . b O u t fl o w l o b e p o s i t i o n a n g l e m e a s u r e db y e y e . T h e un ce rt a i n t y i s t y p i c a ll y ◦ . c O u t fl o w l o b e o p e n i n ga n g l e m e a s u r e db y e y e . T h e un ce rt a i n t y i s t y p i c a ll y ◦ . d O u t fl o w v e l o c i t y g i v e nb y v f l o w = | v t − v s y s | / c o s i . e D y n a m i c a l t i m e s c a l e g i v e nb y τ d y n = l / v f l o w . f O u t fl o w l o b e m a ss o b t a i n e db y t h e C O i n t e g r a t e d i n t e n s i t y o v e rt h e v e l o c i t y r a n g e d e fi n e d i n § . . . T h e l e f t a nd r i g h t c o l u m n v a l u e s c o rr e s p o nd t o t h e m a ss e s i n t h ec a s e s o f T e x = . K a nd . K , r e s p ec t i v e l y . g O u t fl o w m a ss l o ss r a t ee s t i m a t e db y ˙ M o u t f l o w = M l o b e / τ d y n . h O u t fl o w m o m e n t u m r a t ee s t i m a t e db y F o u t f l o w = M l o b e v f l o w / τ d y n . i O u t fl o w m ec h a n i c a ll u m i n o s i t y e s t i m a t e db y L o u t f l o w = M l o b e v f l o w . ∼ < Table 4. Summary of the Continuum Photometry withCameras Wavelength Camera Aperture S ν ( µ m) (arcsec) (mJy)350 SHARC-II 8.4 ∼ < a 70 MIPS 18.0 ∼ < a 24 MIPS 6.0 184.0 c . ± . d . ± . d c c . ± . d c c . ± . × − Note —See § a Considered to be an upper limit because the aperture is signifi-cantly larger than the source size. b Taken from Paper I. c Photometric uncertainties in Spitzer images are typically 10%. d Photometry at the WISE bands are taken from those for WISE J205129.83+601838 in the WISE All-Sky Release Source Cat-alog. Table 5. Summary of the SED model fits Property Symbol Unit Results a Stellar radius R ∗ R (cid:12) T ∗ K 3400Disk mass [dust] M dustdisk M (cid:12) . × − Disk outer radius c R diskmax AU 50Disk flaring power β disk · · · p · · · − . h AU 3.7Envelope density [dust] ρ env0 g cm − . × − Cavity density [dust] ρ cav0 g cm − . × − Cavity opening angle θ deg 20Cavity power c · · · i deg 65 Note —Results from SED fitting using YSO models in Robitaille(2017). See § a See Figure 13. The best-fit model is g5QAQXBF 07 in the modelset of spubhmi . bb