The Co-evolution of Disk and Star in Embedded Stages: The Case of the Very Low-mass Protostar
Yuki Okoda, Yoko Oya, Nami Sakai, Yoshimasa Watanabe, Jes K. Jørgensen, Ewine F.van Dishoeck, Satoshi Yamamoto
aa r X i v : . [ a s t r o - ph . S R ] S e p Draft version September 5, 2018
Typeset using L A TEX manuscript style in AASTeX61
THE CO-EVOLUTION OF DISK AND STAR IN EMBEDDED STAGES: THE CASE OF THEVERY LOW-MASS PROTOSTAR IRAS 15398 − YUKI OKODA, YOKO OYA, NAMI SAKAI, YOSHIMASA WATANABE,
3, 4
JES K. JØRGENSEN, EWINE F. VAN DISHOECK, and SATOSHI YAMAMOTO Department of Physics, The University of Tokyo, 7-3-1, Hongo, Bunkyo-ku, Tokyo 113-0033, Japan;[email protected] RIKEN Cluster for Pioneering Research, Wako, Saitama 351-0198, Japan Department of Physics, The University of Tsukuba, Tsukuba, Ibaraki 305-8577, Japan Tomonaga Center for the History of the Universe, Faculty of Pure and Applied Sciences, University of Tsukuba,Tsukuba, Ibaraki 305-8571, Japan Centre for Star and Planet Formation, Niels Bohr Institute and Natural History of Denmark, University ofCopenhagen, Øster Voldgade 5-7, DK-1350, Copenhagen K, Denmark Leiden Observatory, Leiden University, PO Box 9513, 2300 RA, Leiden, The Netherlands Max-Plank Institut f ¨ u rExtrsterrestrische Physik (MPE), Giessenbachstr.1, 85748, Garching, Germany (Received; Revised; Accepted September 5, 2018) Submitted toABSTRACTWe have observed the CCH ( N = 3 − , J = 7 / − / , F = 4 − −
2) and SO (6 − ) emissionat a 0 . ′′ − − from the systemic velocity, a velocity shift higher than 2km s − is seen for the SO emission. This high velocity component is most likely associated with theKeplerian rotation around the protostar. The protostellar mass is estimated to be 0.007 +0 . − . M ⊙ from the velocity profile of the SO emission. With this protostellar mass, the velocity structure ofthe CCH emission can be explained by the model of the infalling-rotating envelope, where the radiusof the centrifugal barrier is estimated to be 40 au from the comparison with the model. The diskmass evaluated from the dust continuum emission by assuming the dust temperature of 20 K 100K is 0.1 0.9 times the stellar mass, resulting in the Toomre Q parameter of 0.4 5. Hence, the diskstructure may be partly unstable. All these results suggest that a rotationally-supported disk can beformed in the earliest stages of the protostellar evolution. Keywords:
ISM: individual objects (IRAS 15398 − INTRODUCTIONIt is generally thought that a rotationally-supported disk is formed around a newborn star, whichcan eventually evolve into a planetary system (e.g., Bodenheimer 1995; Saigo & Tomisaka 2006;Machida et al. 2016). A thorough understanding of disk formation is therefore of fundamental impor-tance in exploring the origin of the Solar System. Such disks are indeed found in many protostellar ob-jects in the Class II and Class III stages (e.g., Williams & Cieza 2011; Dutrey et al. 2014). Althoughtheir existence was suggested for a few protostars in the younger stages (Class 0/I) (e.g., Enoch et al.2009), it has recently been established, thanks to high spatial resolution observations with interferom-eters including the Atacama Large Millimeter/Submillimeter Array (ALMA) (e.g., Tobin et al. 2012;Yen et al. 2013, 2017; Brinch & Jørgensen 2013; Murillo et al. 2013; Lindberg et al. 2014; Lee et al.2014; Ohashi et al. 2014; Oya et al. 2016, 2017; Sakai et al. 2017). These results suggest that thedisk structure is possibly formed at an earlier stage than previously thought. Although theoreticalsimulations have extensively been conducted for disk formation (e.g., Bate 1998; Hueso & Guillot2005; Inutsuka et al. 2010; Machida et al. 2011; Tsukamoto et al. 2017; Zhao et al. 2018), our under-standing is far from complete. Thus, disk formation is one of the frontiers in star-formation studies.Further through Kepler’s law it is possible to estimate protostellar masses, even when the youngstar itself cannot be observed (e.g., Lommen et al. 2008). Thus, studies of disk kinematics make itpossible to trace the amount of material that has been accreted onto the young star and thus followits build up during the embedded protostellar stages.IRAS 15398 − − H, and CH CCH, present on scales of a fewthousand au scale around the protostar (Sakai et al. 2009).Jørgensen et al. (2013) conducted sub-arcsecond resolution observations toward IRAS 15398 − CO + ( J = 4 −
3) emission at scales of 150 − + through the gas-phasereaction with H O which has sublimated from the grains due to the enhanced protostellar luminositycaused by a recent accretion burst. Bjerkeli et al. (2016b) observed the HDO (1 , − , ) emissionwith ALMA, and found that it is localized at the cavity wall in the vicinity of the protostar. Theextent of the emission is also consistent with the interpretation that a recent accretion burst has takenplace. In addition, Bjerkeli et al. (2016a) find bullets with spacings consistent with the time-scalefor relatively recent accretion bursts.Oya et al. (2014) reported the distribution of the H CO (5 , − , ) line toward this source withALMA at high angular resolution of 0 . ′′ ∼
80 au), and characterized a bipolar outflow extendingalong the northeast-southwest axis at about a 2000 au scale. The analysis of the outflow structureindicates that it is oriented almost in the plane of the sky with an inclination angle of 70 ◦ withrespect to the line of sight. This in turn suggests that the disk/envelope system is seen edge-on. Thisfeature is further verified by CO observations with the Submillimeter Array (SMA) (Bjerkeli et al.2016a). Oya et al. (2014) derived an upper limit on the protostellar mass to be 0.09 M ⊙ from thevelocity structure of the H CO emission. In the envelope, they pointed out the possibility that arotationally-supported disk structure may have already been formed around the protostar based onthe detection of high velocity components of H CO associated with the protostar whose velocity shiftis as high as ∼ − . Recently, Yen et al. (2017) also found an upper limit of 0.01 M ⊙ for theprotostellar mass, based on their observation of the C O ( J = 2 −
1) line (resolution of about 0 . ′′ − . ′′ ∼
30 au) to characterize the disk/envelopestructure of IRAS 15398 − OBSERVATIONObservations of IRAS 15398 − N = 4 − N = 3 − α , δ )= (15 . h . m . s − ◦ ′ . ′′ − at 250 GHz.The bandpass calibrator was J1517 − . ′′
04. The total on-source time was 21.61 minutes. The synthesized-beam sizes are 0 . ′′ × . ′′
15 (P.A.58 ◦ ) for the continuum image, 0 . ′′ × . ′′
29 (P.A. 60 ◦ ) for the CCH image, and 0 . ′′ × . ′′
16 (P.A. 55 ◦ )for the SO image. The rms noise levels for the continuum, the CCH emission, and the SO emissionare 0.12 mJy beam − , 4 mJy beam − , and 4 mJy beam − , respectively, for the channel width of 61kHz. RESULTS AND DISCUSSIONS3.1.
Distribution
Figure 1(a) shows the moment 0 map of the CCH ( N = 3 − , J = 7 / − / , F = 4 − − ◦ ) is clearly seen.This feature is essentially similar to those for the CCH ( N = 4 − , J = 7 / − / , F = 4 − − α , δ ) = (15 . h . m . s ± . ◦ ′ . ′′ ± . ± − at 1.2 mm is also consistent with previ-ous observations at 0.8 mm (19 mJy beam − with a beam size of about 0 . ′′ × . ′′
37; P. A. − ◦ )(Jørgensen et al. 2013) and 0.6mm (36 mJy beam − with a beam size of about 0 . ′′ × . ′′
37; P. A. − ◦ ) (Bjerkeli et al. 2016b)assuming optically thin emission from dust with the opacity law κ ∝ ν β with β ≃ ◦ ; the arrow in Figure1(b)), and prepare the intensity profiles of the CCH, SO, and 1.2 mm continuum emission along theline (Figure 1(d)). Apparently, CCH shows a double peak in its intensity profile. The intensity peakof CCH appears on both sides at a distance of about 70 au from the continuum peak. Althoughthis double-peaked feature was marginally reported in the CCH ( N = 4 − , J = 7 / − / , F =4 − −
2) emission by Jørgensen et al. (2013) and Oya et al. (2014), it is further confirmed inthe present high-resolution observation.In contrast to the CCH distribution, the SO distribution is concentrated in the vicinity of theprotostar. Figure 1(c) shows the integrated intensity map of the SO (6 − ) line in contourssuperposed on that of CCH in a color map. The peak position of the SO distribution almost coincideswith the continuum peak. This feature is clearly seen in the intensity profile along the envelopedirection, as shown in Figure 1(d), where the SO emission shows a single-peaked distribution betweenthe two intensity peaks of CCH.A similar difference between the CCH and SO distributions is reported for another Class 0 low-mass protostar, L1527. The gas distribution around the protostar of L1527, which is a prototypicalWCCC source, was explored with ALMA by Sakai et al. (2014a,b). According to their results, CCHand c-C H are distributed outside the centrifugal barrier of the infalling-rotating envelope, while SOresides at the centrifugal barrier and/or inside it. A centrifugal barrier stands for the perihelion of theballistic motion of the infalling-rotating gas, where the gas cannot fall inward under the conservationlaws of the energy and angular momentum. It has been proposed that the temperature is raisedaround the centrifugal barrier due to a weak accretion shock of the infalling gas, and SO is likelyliberated from grain mantles and enhanced there (Sakai et al. 2014a,b; Aota et al. 2015; Miura et al.2017). In contrast, CCH seems to be broken up by gas phase reactions or depleted onto grain mantlesinside the centrifugal barrier. 3.2. Kinematics
Keplerian Motion
First, we investigate the kinematic structure of the SO emission, which is well concentrated aroundthe protostar. Figure 2(a) shows the position-velocity (PV) diagram of the SO emission along theenvelope direction (P.A. 130 ◦ ) centered at the continuum peak. Red-shifted and blue-shifted com-ponents can be recognized in the southeastern and northwestern parts, respectively (Figure 2(a)).More importantly, the maximum velocity shift from the systemic velocity is as high as about 3 kms − . These components likely correspond to the high velocity components marginally detected inthe H CO emission (Oya et al. 2014). These results imply that SO traces the rotating disk struc-ture around the protostar. The PV diagram along the line perpendicular to the envelope directiondoes not show a significant velocity gradient, indicating no infall and outflow motion (Figure 2(b)).Hence, the observed rotational motion is most likely the Keplerian rotation. In Figure 2(a), the bluecontours represent the model of Keplerian rotation with a protostellar mass of 0.007 M ⊙ . It seemsto explain the velocity structure of the PV diagram observed for the SO emission reasonably well.Note that the intensity around the central position in the model is much stronger than that in theobserved PV diagram. This is most likely due to self-absorption or absorption by the foregroundgas, especially around the systemic velocity. In this study, we focus on the kinematic structure ofthe disk, and the effects of radiative transfer and self-absorption are not included in the model. Onthe other hand, the PV diagram cannot be explained by the infalling-rotating envelope, even if theradius of the centrifugal barrier is set to 1 au (almost the free-fall motion) (Figure 2(c)). The countervelocity components, which are red-shifted and blue-shifted for the northwestern and southeasternsides, respectively, appear in this model as opposed to the observed PV diagram.Thus, the protostellar mass of this source is evaluated to be only 0.007 +0 . − . M ⊙ , assuming Keplerrotation (Figure 2(d)). This very small value is consistent with the upper limits reported so far: < M ⊙ (Oya et al. 2014); < M ⊙ (Yen et al. 2017). Although IRAS 15398 − − . This is slightly differentfrom the value (5.0 − − ) reported in Oya et al. (2014) and Yen et al. (2017). This differenceis discussed later.The protostellar mass derived above is comparable to, or even smaller than, the mass of the firsthydrostatic core expected from star formation theories (0.01-0.05 M ⊙ : Penston 1969; Larson 1969;Masunaga et al. 1998; Saigo & Tomisaka 2006). Since the protostellar mass is so small, the gravityof the system may not be well represented by a central force field. In fact, the dust mass evaluatedfrom the 1.2 mm continuum data is between 0.006 and 0.001 M ⊙ for an assumed dust temperatureof 20 K and 100 K, respectively, which are not much different from the protostellar mass. Here, we employ the mass absorption coefficient of 6.8 × − cm g − at 1.2 mm (Ward-Thompson et al.2000). Although the temperature of 100 K seems too high, it is adopted as the upper limit. Hence,the rotation motion may not be exactly Keplerian, and the mass of 0.007 M ⊙ might be an apparentvalue. Nevertheless, the small protostellar mass would not change drastically, even if such an effectis considered (Mestel 1963). To explore this effect in more detail and to derive the protostellarmass definitively, we need higher sensitivity observations of SO and other molecules. Even if themotion can be approximated by the Keplerian motion, a caveat should be mentioned. With thecurrent resolution, we cannot be sure whether the materials in the beam centered at the protostarhas already accreted onto the star. In this case, the derived mass could be an upper limit to themass of the prototstar.To assess the stability of the disk, we evaluated the Toomre-Q parameter (Toomre 1964;Goldreich & Lynden-Bell 1965) by assuming Keplerian motion. The derived Q parameter is 0.4and 5 for temperatures of 20 K and 100 K, respectively. Thus, either the disk is relatively warm andof low mass, or it is in the unstable regime. Such an unstable part may be responsible for futureaccretion bursts. 3.2.2. Infalling-rotating Envelope
By using the protostellar mass estimated from the Keplerian motion, we examine whether thekinematic structure traced by the CCH emission is consistent with an infalling-rotating envelope.Figure 3(a) is the PV diagram of the CCH emission along the envelope direction, while Figure 3(b)shows that along the line perpendicular to the envelope direction. The dashed lines are the model ofthe Keplerian rotation with the protostellar mass of 0.007 M ⊙ . The Keplerian rotation model doesnot explain the velocity gradient along the line perpendicular to the envelope direction (Figure 3(b);dashed lines). The contours in Figures 3(a) and (b) are the results of the infalling-rotating envelopemodel (Oya et al. 2014) around the protostar with a protostellar mass of 0.007 M ⊙ , where the radiusof the centrifugal barrier is assumed to be at 40 au. The infalling-rotating motion seems to roughlyexplain the observed kinematic structure of CCH, although there is weak emission outside of themodel in the observed PV diagrams probably due to contributions from the outflow cavity. It shouldbe noted that, in the infalling-rotating envelope case, the systemic velocity of 5.3 km s − , whichis similar to the previous report (Yen et al. 2017), gives a better fit. Hence, the envelope and thedisk could have slightly different systemic velocities. This difference is small, but significant. It mayoriginate from the small protostellar mass, because, in this case, the disk system and the envelopecould have different centers of mass.0 3.3. Disk around a Very Low-Mass Protostar
IRAS 15398 − r .
40 au, and hence, the effect of the infalling-rotating components to the derivation ofthe protostellar mass can be ignored. Kristensen et al. (2012) and Jørgensen et al. (2013) reportedenvelope masses of 0.5 M ⊙ and 1.2 M ⊙ , respectively, and hence, the protostar will grow further.Thus, the very low mass of the protostar means that it is in its infancy. In fact, the dynamicaltimescale of the outflow of this source is reported to be 10 − yr (Oya et al. 2014; Bjerkeli et al.2016a). The present results therefore means that a rotating disk structure can be formed at a veryinfant stage of protostellar evolution.Figures 4 (a) and (b) show the comparison of the protostellar masses with the bolometric luminostiesand the disk masses, respectively. Red marks represent the result for IRAS 15398 − − M acc averaged over the protostar life is estimated to be about ∼ × − M ⊙ yr − from the protostellar mass of 0.007 M ⊙ and the above dynamical timescale. It can also beestimated by using of the following relation (Palla & Stahler 1991):˙ M acc = LR star GM , where L is the luminosity and R star the radius of the protostar. We evaluate the current accretion rate˙ M acc to be 2 . × − M ⊙ yr − by using L of 1.8 L ⊙ (Jørgensen et al. 2013) and R star of 2.5 R ⊙ (e.g.,1Palla 1999; Baraffe & Chabrier 2010). It should be noted that we roughly employ the bolometricluminosity for L and the average radius of the protostar for R star . The average and current accretionrates are almost within a range of cannonical values 10 − − − M ⊙ yr − (Hartmann et al. 1997).It should be noted that the mass loss rate due to the outflow is reported to be (3.2 − . × − M ⊙ yr − (Yıldız et al. 2015) and 7 × − M ⊙ yr − (Bjerkeli et al. 2016a), which is comparable toor smaller than the above estimates of the accretion rate.Although the very low mass of this protostar can naturally be interpreted as its infancy, thealternative possibility is that this protostar may evolve into a brown dwarf. A planetary systemcould be formed even around a brown dwarf, as predicted by theoretical models for the formationof Earth-like planets around a brown dwarf (Payne & Lodato 2007). However, this interpreta-tion may be not the case for this source, because the envelope mass around the protostar is ashigh as 0.5 − M ⊙ (Kristensen et al. 2012; Jørgensen et al. 2013). 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Parameters of the Observed LineTransition Frequency (GHz) Sµ ( D ) E u k − (K) Beam sizeCCH N = 3 − , J = 7 / − / , F = 4 − . ′′ × . ′′
29 (P.A. 60 ◦ ) N = 3 − , J = 7 / − / , F = 3 − − . ′′ × . ′′
16 (P.A. 55 ◦ ) (a) (b) ( J y / B e a m ) ・ ( k m / s ) Right Ascension(J2000) D e c li n a t i o n ( J ) color: CCH((cid:9490)=3-2, (cid:9486)=7/2-5/2, (cid:9482)=4-3 and 3-2) Contour: Continuum
310 au
CCH((cid:9490)=3-2, (cid:9486)=7/2-5/2, (cid:9482)=4-3 and 3-2)P.A. 220°
155 au -09ʼ05” -09ʼ10”15 43 02.75 02.50 02.25 02.00 01.75 smh ssss
P.A. 130° ( J y / B e a m ) ・ ( k m / s ) Continuum C o n t i n uu m ( m J y / B e a m ) L i n e s ( J y / B e a m ) ・ ( k m / s ) -2 -1 0 1 2 CCHSOBeamsize
Anglar offset[arcsec]
155 au
SENW
CCHSO (d) (c)
Contour: SO(6₇-5₆)
CCH((cid:9490)=3-2, (cid:9486)=7/2-5/2, (cid:9482)=4-3 and 3-2) m s s s P.A. 130° m s s s -34°09ʼ06” -09ʼ07” -09ʼ08” -09ʼ07” -09ʼ08” -34°09ʼ00” -34°09ʼ06” Figure 1. (a) The moment 0 map of the CCH ( N = 3 − , J = 7 / − / , F = 4 − −
2) lines. The arrowrepresents the outflow axis (P.A. 220 ◦ ). (b) The color map is a blow-up of the central part of panel (a). Theblack contours are the continuum map contour levels are 10 σ , 20 σ , and 40 σ , where σ is 0.12 mJy beam − .The white contours are the moment 0 map of the CCH ( N = 3 − , J = 7 / − / , F = 4 − − σ , 3 σ , 4 σ , 5 σ , and 6 σ , where σ is 8 mJy beam − km s − . The dashed arrowrepresents the envelope direction (P.A. 130 ◦ ). (c) The color map is a blow-up of the central part of panel(a). The contours are the moment 0 map of the SO (6 − ) line. Contour levels are 3 σ , 4 σ , where σ is 8mJy beam − km s − . The black cross shows the continuum peak position. (d) The intensity profiles of CCH( N = 3 − , J = 7 / − / , F = 4 − −
2) and SO (6 − ) along the envelope direction shown bythe dashed arrow in panel (b). Their noise levels are described by σ , where σ is 8 mJy beam − km s − , onthe bottom of this figure. The noise level of the continuum is described by 3 σ , where σ is 0.12 mJy beam − .The abscissa is the angular offset from the continuum peak. Beam size
P.A. 130°
Beam size NE P.A. 220° V L S R [ k m / s ] y J m a e b / (a) (b) Beam size
Envelope direction
NW SE
P.A. 130°
Perpendicular to the envelope direction
62 au -0.02-0.01 J y / b e a m V L S R [ k m / s ] NW SE SW -0.8” 0” 0.8” (c) (d) Counter velocity component Counter velocity component
Beam size
Envelope direction
NW SE
P.A. 130°
62 au
Envelope direction ⦿ ⦿ ⦿ -0.8” 0” 0.8” Figure 2. (a) The PV diagram of the SO (6 − ) emission along the envelope direction shown by thedashed arrow in Figure 1(b). The origin is the continuum peak position. The blue contours are the model ofKeplerian rotation around the systemic velocity of 5.5 km s − with the protostellar mass of 0.007 M ⊙ , theouter radius of 40 au, and the inclination angle of 70 ◦ (0 ◦ for a face-on configuration). (b) The PV diagram ofthe SO (6 − ) emission along the line perpendicular to the envelope direction centered at the continuumpeak position. (c,d) The color maps are the PV diagrams of the SO (6 − ) emission along the envelopedirection. The blue contours are the infalling-rotating envelope model with the protostellar mass of 0.007 M ⊙ , the outer radius of 40 au, the radius of the centrifugal barrier of 1 au, and the inclination angle of 70 ◦ .The curves are the model of Keplerian rotation with the different protostellar mass assumed. Green, blueand black show the cases of 0.010, 0.007 and 0.004 M ⊙ . (a) J y / b e a m V L S R [ k m / s ] P.A. 130° J y / b e a m V L S R [ k m / s ] Beam size
P.A. 220° (b) -2” -1” 0” 1” 2”Angular Offset
Beam size
NW SE
NE SW
Figure 3.
The PV diagram of the CCH ( N = 3 − , J = 7 / − / , F = 4 −
3) emission along the envelopedirection (a), and along the line perpendicular to the envelope direction (b). The dashed lines show themodel of Keplerian rotation with the protostellar mass of 0.007 M ⊙ . The white contours are the results ofthe infalling-rotating envelope model with a protostellar mass of 0.007 M ⊙ , the outer radius of 155 au, theradius of the centrifugal barrier of 40 au, and inclination angle of 70 ◦ . star log (cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444) (cid:9444)(cid:9489) [(cid:9489) ]
10 ⦿ -2.5 -2.0 -1.5 -1.0 -0.5 0 0.5 l og (cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444) (cid:9444) (cid:9489) [ (cid:9489) ] ⦿ d i s k l og (cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444) (cid:9444) (cid:9488) [ (cid:9488) ] b o l ⦿ -1.50-0.5-1-2-2.5-31.510.50-0.5-1 (1)(1) (2)(2) (3)(3) (4)(4) (5)(5) (6)(6)(7)(7) (8)(8)(9)(9) (10)(10) (11)(11)(12)(12)(13)(13) (14)(14)(15)(15)(16)(16)(17)(17) -2.5 -2.0 -1.5 -1.0 -0.5 0 0.5 star log (cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444)(cid:9444) (cid:9444)(cid:9489) [(cid:9489) ]
10 ⦿ (b)(a)
Figure 4. (a) Comparison between the protostellar masses and the bolometric luminosities. (b) Comparisonbetween the protostellar masses and the disk masses. They are for the sample of protostars in previous studieslisted below. Red marks with error bars represent IRAS 15398-3359 (this study). Blue and Green marksshow the protostars of T bol <
70 K (Class 0) and T bol >>