星形成ゼミ
2017/6/23
SFN #292 #36-#40
長谷川哲夫（国立天文台）

1

292 Which Stars Are Ionizing the Orion Nebula?
R. O’Dell, W. Kollatschny, and G. J. Ferland#36 C.
ApJ 837, 151 (2017)
問題：オリオン星雲の励起星は θ1 Ori C だけか？
方法：
1. Cloudyモデルと比較• MUSEによる可視域輝線強度比• Spitzerによる中間赤外輝線強度比
2. Proplydsの形（励起星の方向に明るいリム）
結論：• 星雲の内域は θ1 Ori C が主な励起源• Bright Bar 外側は星雲の南東縁までθ2 Ori A が主な励起源

Ferland 2006), which is the ratio of the density of photons that
can ionize hydrogen to the local hydrogen density (nH ). We canapproximate the local hydrogen density as the electron density(n e ) since hydrogen is essentially completely ionized and the»10% contribution from He I is small.
We can calculate the photon density (r LyC ) from the total
ionizing photon luminosity of the ionizing star (Q LyC ) and thetrue distance to the emitting layer (r ), and this leads to therelation

Ionization Parameter

U=

Q LyC
4p r 2 c

´

1
,
ne

or near edge
can be devel
we use thethat it will b
We calcu
F(He I,667.8
variety of cTable 1. Th
ionizing sta
used the s
Badnell et aof the line r
ratios are areddening b
compared w(1 )where c distanceを採用→Uの上限値
is the velocity of light.
• r はprojected• Ne は red [SII] doublet ratioから推定
2.2. Calculations

In this section we present a series of new ionization
simulations of the Orion H II region. We use version C17.00
of Cloudy, the spectral simulation code last described byFerland et al. (2013). The geometry is that used by Baldwin
et al. (1991). It is assumed that the physical thicknessof the H+
The AstrophysicalJournal, 837:151 (7pp), 2017 March 10
O’Dell, Kollatschny, & Ferland
θ1 Ori C (O6)を励起星として
ioniza;onparameterを計算した場合
layer that constitutes the H II region is much thinner than the
The Astrophysical Journal, 837:151 (7pp), 2017 March 10
O’Dell, Kollatschny, & Ferland
separation to the ionizing stars. The geometry is plane parallel.
The layer is assumed to be in hydrostatic equilibrium with theH+ layer held back against the molecular cloud by the
combination of gas pressure in the X-ray-emitting hot gas
surrounding the star cluster and the absorbed outward
momentum in the starlight.
This model was originally applied to central regions of thenebula, close to the Trapezium, where we are most likely
viewing the H+ layer nearly face-on. However, the bright
Huygens region is most likely bowl shaped, as described byFerland et al. (2013) and Wen & O’Dell (1995). The observed
surfacebrightness
ofeachthelayer
is affectedlineby
Figure 2. Ionization parameter vs. ionizing star temperature
is shown
in all three panels,
depictingextinction-insensitive
ratios.two
Dashedthings.
vertical linesFirst,
indicate
the temperatures of the two candidate ionizing stars. The dotted lines indicate the expected line ratio using Cloudy for varying values of the distance-dependentasthedistance
of
the
star
increases,
the
ﬂux
of
ionizing
ionization parameter assuming that the ionizing star is q Ori C . The seven samples from Figure 1 are shown. Those probably ionized by q Ori C are in blue andcircled. Arrows indicate the direction of the correction
to the ionization
parameter sets
for the fact
it was calculatedusing the minimumfalls
possibleas
distancefromthe star,r −2. Thisphotons,which
thethat surface
brightness,i.e., the distance in the plane of the sky. The length of the arrow only suggests the magnitude of that correction. Blue circles and lettering indicate samples most likely
ionized by q Ori C .
causes the surface brightness to decrease with increasing
radius.
At the same time, the “tilt” of the H+ layer to our line of
2 Ori A , which rules
Tstar forratiosit lies too far
to the 837:151
right of(7pp),
4. Using Spitzer Space Telescope
Infrared& Ferland
The
AstrophysicalJournal,2017q March
10
O’Dell, Kollatschny,
2
OriA.A (O9.5 V)を励起星として
ioniza;onparameterを計算した場合
sight increases(see Figure
3 of Wen
O’DellSpectra&Line
Ratios 1995), whichoutθq 2 Ori
The identiﬁcation of the inner ﬁve samples
as being
causestheionizedsurface Publishedbrightness
increasea factorof of
SpitzertoSpace
Telescopebyobservations
theby q1 Ori C falls within the accepted model of thisof
-1,regionμm and[Ne II] 12.8relative
μm lines arealso
infrared[Ne III] 15.6
where
θ
is
the
viewinganglemeasuredto thecos(q)
the Orion Nebula, i.e., that it is a concave irregular surface. The
for determining the ionizing star. These are particularly
towiththethelayer.useful
As longas our line of sight passes throughsample apparently ionized by q 2 Ori A normal
is consistentuseful because they both arise in the same He+ + H+ zone andBright Bar and the regions southeast from
there
being
ionizedthe entire layer, theemission-line
Heo + H+ zonedo
are relative
unaffected byemission from theintensities
andnotany
by the nearest hot star. Sample E, although lying outside thechangessurface
of abundances,
as is the case for our optical line ratios.
1change,
only
their
total
brightness.
Bright Bar, is probably ionized by q Ori C, as discussed
In the study of Rubin et al. (2011) observations were made atThis simple geometry
breakswhenviewingangle
Figure 1. This Hα 0.66×0.55 pc(353″×295″) image of the Orion Nebula is from the 0 2 pixel MUSE data. The samples employed in generating the data in earlier in this section.
11 locations
to thedown
southeastof the thecentralOrion Nebula
asTable 1 are shown as lettered rectangles the size of the sample. The arrows indicate the direction(s) of the orientation of the bright edges of some of the many proplyds,
shown
in Figurelie outside the BrightBar,
approaches edge-on,
since
we4. These
thensamples
see allsuccessive
regions
as discussed in Section 3. Those proplyds ionized by q Ori A are shown in enlargements of 2″×2″ , except for 244.3–439.8, which is a sample of 4″×4″. The
2
with the ofinnermostto q Ori Athe
. The progressionoflocation of the ﬁrst of the Spitzer Space Telescope observations discussed in Section 4 is shown south of q Ori A and further deﬁned in Figure 4.
projected
on the plane
the beingsky,close
resolvingionization
3. Identifying the Ionizing Stars Using
Observations
increasing distances is I4, I3, I2, I1, M1, M2, M3, M4, V1, V2,of the Proplyds structure. This occurs within the Bright Bar (Sellgren et al. 1990)
and V3. They all lie in a smooth region between the Bright Bar
shifted to near the intersection. However, our adopted value forF([O III])/F(Hβ) and F(He I)/F(Hα) we see that all samples are1

By compexpected va
that star. Th
the equation

where rsky iionizing star
red [S II] rat
q 2 Ori A . T
once in Figu
q1 Ori C ansame line ra
We do nTherefore, aan upper lim
derived poi
ratio and Tst
larger differ
greater dista
If the obs
projection althen that sam
two points (Clie to the righin a very ﬂata smaller diwould be lar

1

1

–1

2

2

with ionization by q Ori C, with the caveat of the
採用した距離：388(+-8)consistentpc (Kounkel+2017)

the distance is already smaller than many modern determinations,as summarized in O’Dell & Henney (2008).
It is more likely that the Cloudy model that we adopt is
beginning to break down in Samples C and F as a result of thelarge viewing angle. Our model assumes that we are observing
along a line of sight that goes through the entire H+ layer. Whenthe MIF is highly tilted, we resolve the ionization structure on
the plane of the sky and thus sample along layers within thecolumn. This compromises a comparison of line-of-sight models

1

There is a completely independent method of identifying theand the boundary of the Extended Orion Nebula (EON).

remarks in Section 2.3 about the high Tstar values for Samples C dominant source of ionization using high-resolution images ofSample V2 falls near the edge of the sharp southeast boundary
and F. Taken at face value, F([O III])/F(Hβ) and F(He I)/F(Hα) the Orion Nebula proplyds. These are young stars withof the EON, and Sample V3 falls on this boundary. These
indicate that Samples C and F have spatial distances about Figure
circumstellar
in ornow
projected
on parameters
H II regions.
Under using distancesamplesdoq 2nottheprobably
resultsionized
frombythebecauseOri agree
A . The withq 2 others
Ori A are
3. Same as material
Figure 2, except
the ionization
are calculated
fromsamplesin redequal to rsky , while Samples G, A, and B demand progressively andcircled. orientations and locations the interior protoplanetaryfavorable
Sample V3 is certainly an ionization front viewed edge-on andlarger corrections to the true distance. In the left and right disks are rendered visible in silhouette. More relevant here,
Sample V2 is probably a region that is rising toward thatpanels it appears that Sample D could also be associated withthese disks are surrounded by much larger envelopes of gas andfeature (Rubin et al. 2011), and thus our Cloudy models do not
q1 Ori C, but in the middle panel we see that no plausible
dust. A nearby hot star will form an ionization front in this gas
provide a guide to the conditions (Section 2.3). They are notcorrection for distance would make the line ratio compatible

2

2

h 10

directions. Thistion as a ratios
observer. Thisimilarities andng

the position of
the sky of the
an compare theame line ratio
quare of r /rsky ,mes larger thans four times tooparation in thef r with respectrsky /r ) and the×tan Θ. Whator negative canues with respect
From the ratio

Figure 3. Same as Figure 2, except now the ionization parameters are calculated using distance from q 2 Ori A . The samples probably ionized by q 2 Ori A are in red
and circled.

Figure 4. Spitzer samples (after Figure 1 of Rubin et al. 2011); see text for
details. The east boundary of the EON appears at the lower left (southeast), andthe Trapezium stars lie in the center of the bright northwest feature.

star, since Sample I4 is closest to q 2 Ori A and our Sample D
in the plane of the sky and the I3, I2, and I1 samples progress
farther away. Their locations in Figure 5 indicate a successionof corrections for projection effects, and at large distances the
samples are close to Tstar for q 2 Ori A .

Figure 5. Like Figures 2 and 3, except now we use the ratio of the [Ne III]15.6 μm and [Ne II] 12.8 μm lines. The symbols are samples from the study of
Rubin et al. (2011) that begin close to q 2 Ori A and proceed to the southeast
with a maximum distance of 12.08 from q1 Ori C . Colors indicate which star
was assumed to be the ionizing source, blue for q1 Ori C and red for q 2 Ori A .
Samples V2 and V3 fall at and near the edge of the sharp southeast boundary of
the EON and are not used in the discussion in Section 4.

O’Dell, Kollatschny,& Ferlanda lesserHα contribution from the narrower Heo + H+ zone.

The distributions of doubly ionized oxygen (producing the[O III] emission) and He I ions (producing the 667.8 nm
emission) are slightly different within the He+ + H+ zone.
Positionsof Samples
with similar
Respectbuttonot
Their
Ionizing Stars
The resultsof Sections
2 and 3 show
identical
This means that the ratios using these two ions will give
results. Some of the differences may be explained by thedifferent results. In the Heo + H+ zone the [N II]
rsky (pc)Θ slightly Z(pc)SampleIonizing Starr /rskymethodology.
emission is well constrained, but the fact that the Hα emission
and
F(He
I
)/F(Hα)line
ratiosare
The F([O III])/F(Hβ)
q1 Ori
C
0.054
1.77
A
+ + H+ zone,
mostly from the He+ + H+ zone means that theas [O III] −60°comes −0.094
indicatorsof theconditionsin the He
q1 OriC only there,0.028
B He I emission−0.134 ratio is intrinsically not as good a diagnostic.
arisesas does most4.37of the Hα −78°F([N II])/F(Hα)
andOri densityCemission.
Theq1gasdecreases
2001) with 28° In contrast,
C
0.189(O’Dell 1.134
0.101the proplyds give an indication of the ionization
located somewhere along the line of sight from theincreasing
distance
therefore, mostq 2 Orifrom
A the MIF;0.060
D
2.81 of the −72°of objects−0.180Veil and the MIF. The Orion Nebula Cluster isand F(He I)/F(Hα)
line 60°foregroundemissionused qin1 Ori
the F([OC III])/F(Hβ)
E
0.254
1.78
0.444
ratio method occurs
near the onset of the He+ + H+ zone, withcentered near the Trapezium and monotonically decreases in
1
5. DiscussionTable 2FG

q Ori Cq1 Ori C

0.179
0.139

1.134
1.53

28°

555°

0.096
0.196

The region outside of the Bright Bar changes at about theposition of the HH203 and HH204 shocks. To the northeastof these shocks F([N II])/F(Hα) is high and F([O III])/F(Hβ) is
low. This is reversed to the southwest of these shocks andagrees with our Sample D and Sample E values. This transitionis illustrated well in an F([N II])/F(Hα) image, such as Figure
24 of O’Dell et al. (2015).
The southwest high-ionization region is marked by numerous crenellations lying immediately southeast of the Bright Bar.
These features have sharp boundaries and are likely to be
shocks driven by a series of outﬂows from the Orion-S starformation region (O’Dell et al. 2015).

3

292 Star Formation Under the Outflow: The Discovery of a Non-thermal Jet from#37 OMC-2 FIR 3 and Its Relationship to the Deeply Embedded FIR 4 Protostar
Mayra Osorio, Ana K. Díaz-Rodríguez, Guillem Anglada, S. Thomas Megeath, Luis F. Rodríguez, John J. Tobin, Amelia
M. Stutz, Elise Furlan, William J. Fischer, P. Manoj, José F. Gómez, Beatriz González-García, Thomas Stanke,
Dan M. Watson, Laurent Loinard, Roland Vavrek, and Carlos Carrasco-González
ApJ 840, 36 (2017)
The Astrophysical Journal, 840:36 (12pp), 2017 May 1
Osorio et al.
問題：FIR 4の正体は？ FIR 3の双極流との関係は？
方法：
1. VLAによる新たな多波長観測（2014-2015年）• 6, 10, 22, 44 GHz
2. VLAアーカイブデータとの比較で固有運動を検出• 1994, 1998, 2000年の8.5 GHz観測
結果と議論：
1. 領域内に10の電波源、うち7つはYSO（Table 2）
2. FIR 3に付随する熱的電波源（ｆ−ｆ）から南西に~100 km/
s で移動する3つの非熱的電波源（VLA 12N, 12C,
12S）：ジェットのショックで加速された電子によるシンク
ロトロン放射か3. FIR 4はジェットとは別の熱的電波源として検出され、固有運動も異なる
4. この結果はIRS 3からのジェットと分子雲の相互作用によりFIR 4の形成がトリガーされたという説（Shimajiri+2008）を支持するように見える
5. FIR 4に付随する電波源の固有運動が~30 km/sと高速すぎる点が上の説と矛盾する。しかし移動してゆくVLA
12とのブレンドやFIR 4からのジェットにより見かけ上大
きな固有運動が測定された可能性もあるので、高分解能の観測を続ける必要がある。

Figure 1. VLA C-conﬁguration maps (epoch 2014) of the region around the HOPS 108 protostar. Contour levels are −3, 3, 6, 9, 12, 15, 20, 30, 50, 70, 100, 150, and
200 times the rms noise of each map. The robust weighting parameter (Briggs 1995) was set to 0. The synthesized beam is shown in the bottom left corner of eachmap. The 10.4 μm positions (Nielbock et al. 2003) are indicated by red plus signs whose sizes correspond to the positional uncertainty. For HOPS 108, which is notreported at 10.4 μm, the 24 μm position (Megeath et al. 2012) is given. Top left: C band (5 cm); rms=11 μJy beam−1; HPBW=4 75×2 46, P.A.=46°. Topright: X band (3 cm); rms=9 μJy beam−1; HPBW=2 68×1 73, P.A.=36°. Bottom left: K band (1.3 cm); rms=10 μJy beam−1; HPBW=0 84×0 71,
P.A.=5°. Bottom right: Q band (0.7 cm) map; rms=15 μJy beam−1; HPBW=0 46×0 36, P.A.=−20°.

4embedded within one of these cores. A bolometric luminosity
Lbol;1000 Le for FIR 4 was obtained by Crimier et al. (2009),

uncertainties in the amount of external heating and theinclination of the protostar, Furlan et al. (2014) found that

The Astrophysical Journal, 840:36 (12pp), 2017 May 1

Osorio et al.
Table 2
Positions of the Radio Sources

SourceVLA 11MIPS 2301
MIPS 2297
VLA 12N
HOPS 64
VLA 12CHOPS 108
VLA 12S
VLA 16
VLA 15

AlternativeNames
FIR 3, HOPS 370, MIPS 2302, MIR 21
MIR 22
MIR 23
MIPS 2293, MIR 24, FIR 4d/e
FIR 4, MIPS 2289

FIR 4jPositiona
NatureClass I YSOYSOClass II YSO
Jet knotClass I YSO
Jet knotClass 0 YSO
Jet knotYSO
YSO

Referencesα(J2000)

δ(J2000)

1
2
2
3
1
3
1
3
3
3

05 35 27.6337±0.0007
05 35 27.639±0.008
05 35 27.4711±0.0020
05 35 27.385±0.011
05 35 26.9960±0.0023
05 35 27.111±0.012
05 35 27.086±0.004
05 35 26.9389±0.0017
05 35 26.824±0.003
05 35 26.4091±0.0011−05 09 34.368±0.011−05 09 37.09±0.12−05 09 44.06±0.03−05 09 50.68±0.16−05 09 54.06±0.03−05 09 58.57±0.18−05 09 59.95±0.06−05 10 02.876±0.025−05 10 05.64±0.04−05 10 05.951±0.016

Note.
a
Positions derived from elliptical Gaussian ﬁts in the A-conﬁguration C-band map of epoch 2015, except for MIPS 2301 and MIPS 2297, which are not detected inthis map, and whose positions were derived from the 1994–2000 X-band and the 2014 Q-band maps, respectively. Uncertainties correspond to those of absolute
positions, and are calculated by adding in quadrature a systematic error of 0 01 (Dzib et al. 2017) to the formal error of the ﬁt, 0.5 θ/SNR (Reid et al. 1988). The
systematic error accounts for uncertainties introduced by the phase calibration process and the error of the ﬁt for those due to the source size and signal-to-noise ratio.
References. (1) Furlan et al. (2016), (2) Megeath et al. (2012), (3) This paper.
The Astrophysical Journal, 840:36 (12pp), 2017 May 1
Osorio et al.

Herschel
63scenario
micron3 mm
below). This rulesout the H II[OI]
regionproposed byseparating the emission ALMAof HOPS 108
from that of the VLA
電波スペクトル
12 knots conﬁrms HOPS 108 as a low-luminosity object.
López-Sepulcre et al. (2013) for VLA 12.
Osorio et al.
The radio source VLA 11, which was associated with FIR 3The Astrophysical Journal, 840:36 (12pp), 2017 May 1
Although the position of the knot VLA 12C at 6 and 3 cm
(HOPS 370) and reported as unresolved by Reipurth et al.
appears somewhat (∼1 5) displaced to the northeast of the(1999), is clearly resolved as a very elongated bipolar radio
position of HOPS 108 (as obtained from 24 and 8 μm data;
jet along a PA;5° in our A-conﬁguration data at 5 cm
Megeath et al. 2012), at 1.3 and 0.7 cm the radio emission(Figure 4). The radio source is also resolved in the 3.6 cm
appears to be closer to the HOPS 108 IR position (seeimage obtained by combining archive data from 1994 to 2000.
Figure 1). This shift in the position could indicate either an
The spectral index is α=0.27±0.10 (Table 3, Figure 3),
opacity gradient or that the radio emission observed at shorter
which is a reasonable value for a thermal free–free radio jet
wavelengths originates in a different object than the emission(Anglada 1996). HOPS 370 is known to be associated with anobserved at longer wavelengths. This last suggestion isoutﬂow extending along a PA similar to that of the radio jet
conﬁrmed by our data of subarcsecond angular resolution at(González-García et al. 2016; see Figure 2, left). Thelong wavelengths (Figure 4), which reveal two radio sourcesmorphology of VLA 11 appears to change in the 0.7 cm image
spatially separated by ∼2″. We identify the northeasternmost of(Figure 1), which shows weak extended emission in a direction
these two radio sources with the VLA 12C knot and the otherperpendicular to the jet, suggesting a noticeable dust contribsource, whose position coincides with the IR source, with theution at this wavelength. To prevent dust contamination, theradio counterpart of the HOPS 108 protostar. Recent ALMA
0.7 cm data point has been excluded in the ﬁts to determine theobservations with an angular resolution of 0 12 (J. Tobin et al.
spectral index of this and the remaining thermal sources.
2017b, in preparation) reveal a compact source at 870 μm
In the lower-resolution images at 6 and 3 cm (Figure 1),
whose position coincides within <0 1 with the cm position
VLA 11 appears to be blended with the radio emission of MIPS
(Table 3) conﬁrming the nature of HOPS 108 as a protostar and2301 (Megeath et al. 2012, an IR source also known as MIR
its association with the cm radio source. Furthermore, the22, Nielbock et al. 2003) located ∼3″ south of VLA 11. MIPS
spectral index of HOPS 108 obtained from the centimeter data
2301 is well separated, but only marginally detected, in thethat separate its emission from that of VLA 12C shows thathigher-resolution 3.6 cm image shown in Figure 4 and in the
HOPS 108 has a positive spectral index (α=0.7±0.3;
A-conﬁguration images at 5 cm obtained with the robustFigure 3. Spectra of the detected VLA sources. Error bars are 1σ. The upper limits are represented by arrows where the central symbol is at 3σ. The dashed lines
Table 3 and Figure 3), thus it is associated with thermal (free–least-squared ﬁts. To avoid possible dust contamination, the 0.7 cm data point has been excluded in the ﬁtting for the positive spectral index sources. The
weighting parameter set to values >1 (not shown in Figure 4).18represent
upper limits have been taken into account only when they constrain the ﬁt.
free) emission, in contrast to the knots of VLA 12, whose
With
the
data
currently
availableweobtainapositive
spectralemission is non-thermal. From the empirical correlation
data taken over ∼6 years into the same map, which is not optimalnoise ratio, is 2±2 km s (Table 3). Since this velocity is
index, α=0.30±0.23 (Table 3, Figure 3), consistent withto measure accurate proper motions, because of motions and obtained from the absolute positions of the source, its small value
between cm ﬂux density and bolometric luminosity forindicates that systematic errors are effectively very small and thatchanges in morphology of the sources over this period. Also, thethermalet al.
free–free
emissiona YSO.
sensitive
Figure2. Left panel:
image ofI] 63 μm line2016) overlaid
on the 3from
cm emission
VLA Moremap (contours,
seearchival data were observed in the X band (8.4–8.5 GHz), while
their possible residual effects on the proper motion determination
protostellar
objects,(Snthed 2[O mJy
kpc2) =emission
0.008 ((color
L bol scale,
L :)0.6González-García
Ibetter] 63 μmcharacterization
emission peaks closeofto thethethe new data have been obtained in the C band (4–8 GHz)
Figure 1), with proper motions indicatedby arrows. Both images are
tracing
the jet ejectedfrom HOPSare
370.required
Note that forthe [Oare at a level of ∼2 km s or smaller.
observationsa(Angladaetal.
2015),
and
using
our
measured
ﬂux
densityat(Table1).
Althoughthefrequency
ranges
of
the
two
bands
are
position of HOPS 108, suggesting that this protostar has been formed in a region where a strong shock interaction between the jet and the ambient medium is taking
natureoftheradio
emissionofthisobject.
almost contiguous, the differences in frequency could result in
3cm(Table3),
weestimateabolometricluminosity
place. The plus signs indicate the positions of IR sources. Right panel: ALMA image of the dust emission at 3 mm (color scale, Kainulainen et al. 2017) overlaid onapparent motions due to variations of opacity with frequency,
3.2.1. The HOPS 370 Radio Jet
source2297 to(Megeath
et easternal. 2012,
asparticularly in extended sources. Despite all these factors that
5are characterized bythe
3 cmHOPS(black contours)
and 0.7value
cm (blueofcontours,
see Figure 1). Note thatThethe radio
jet atMIPS
3 cm appears
delineate theedge alsoof theknown
dust clump
As discussed above, the knots of VLA 12
LbolVLA
;7maps
Le atfor108. Thisthe bolometriccontribute to the uncertainty in our measurements, the relatively
negative spectral indices, indicating non-thermal emission (seeMIR
23;
Nielbock
et
al.
2003)
is
marginallydetectedat
1.3
cm
traced by the ALMA image.
high velocity of the jet and the amount of time (∼18 years) thatluminosity is close to the lowest end of the range considered by
Figure 3), but their overall morphology with three aligned knots
and 0.7 cm (Figure 1), but clearly detected in the 0.9 cmelapsed between the two maps make the proper motions clearly suggests that they pertain to a radio jet. Radio jets associated with
Furlan et al. (2016) in the modeling of this object, being much
Our results show that, while VLA 11 and VLA 15
YSOs are characterized by thermal (free–free) emission (e.g.,observations of J. Tobin et al. (2017a, in preparation). It showsdetectable.
static within the uncertainties, the VLA 12N, VLA 12C
Anglada 1996), but in a few cases, non-thermal emission has beensmaller (Table
than 3 andearlyFigure
estimates(Lbol;1000
Le;as Crimier
regime3), conﬁrming
its nature
a YSO.
ALMA
observationsof higher
angular resolution(J. Tobinremain
and VLA 12S knots present measurable proper motions offound in the jet lobes at relatively large distances from the jet core,a positive
spectral index(α=1.5±0.5)
in the centimeterroughly
in
the
south
direction(Table3and∼40–100kmsshows a positive spectral index and thermal emission
et
al.
2009;
López-Sepulcre
et
al.
2013)
that
were
based
on
So far, MIPS 2297 has not been studied in great detail, but it is
et al. 2017b, in preparation). From the empirical correlationFigure 4; see also the left panel in Figure 2). HOPS 108 shows a which(Carrasco-González et al. 2013; Anglada et al. 2015; Rodríguez−1−1−1The Astrophysical Journal, 840:36 (12pp), 2017 May 1

Osorio et al.

Figure 4. Proper motions of VLA sources. First epoch maps (X band; left panels) were obtained by combining archive data from 1994–2000 (Table 1). Second epoch
maps (C band; right panels) were obtained from the 2015 data. Different weightings have been used to better emphasize the emission of each source. Synthesized
beams of ∼0 95 and rms of ∼10 μJy beam−1 were obtained for the ﬁrst epoch maps, while synthesized beams of ∼0 24–0 46, and rms of 3–8 μJy beam−1 wereobtained for the second epoch maps. Contour levels are −4, −3, 3, 4, 5, 6, 7, 8, 10, 12, 15, 20, 35, 60, 100, and 150 times the rms noise of each map. Positions of thesources derived from elliptical Gaussian ﬁts to the ﬁrst epoch maps are marked with blue plus signs, whose sizes represent their positional errors. The arrows indicate
the displacement between the two epochs.

6292 L483: Warm Carbon-chain Chemistry Source Harboring Hot Corino Activity
Oya, Nami Sakai, Yoshimasa Watanabe, Aya E. Higuchi, Tomoya Hirota, Ana López-Sepulcre, Takeshi Sakai, Yuri
#38 Yoko
Aikawa, Cecilia Ceccarelli, Bertrand Lefloch, Emmanuel Caux, Charlotte Vastel, Claudine Kahane, and Satoshi
Yamamoto
ApJ 837, 174 (2017)
背景：• 小質量星形成領域の２つの化学的タイプ• Hot corino chemistry: 飽和したcomplex organic molecules (COM) 代表例 IRAS 16293-2422• Warm carbon-chain chemistry (WCCC): 不飽和分子（炭素鎖分子など） 代表例 L1527→両者の中間的な天体を、特に100auスケールの分解能で研究することは重要• L483暗黒星雲（距離200pc）中のIRAS 18148-0440: Class 0 原始星＋回転するエンベロープ＋双極分子流、WCCC的？
観測：ALMA Cycle 2 Band 6 (Table 1)、分解能~0.5” (100au)
結果：
1. 100auスケールで分子分布に明確な差：CCHは中心100auに穴、SOやHNCOは中心のみ、CSは中心もエンベロープも双極流壁も2. WCCCとhot corinoの共存を空間的に分解した（初）：中心でCOMs (NH2CHO, HCOOCH3)を検出
3. CSで運動を調べると、回転しながら収縮するエンベロープ（中心質量0.1-0.2Msun）＋遠心力バリア@30-200au＋（おそらく）ケプラー
的円盤、の構造があると思われる。上記の化学的特徴の変化は遠心力バリアを付近で起きている。
4. 内側の円盤でもCSやSOが検出されるのは、L1527やTMC-1Aに比べてL483の光度が高く円盤の温度が高いからだろう。
5. Astrophysical
SiOの不思議な分布：分子流により遠心力バリア付近に発生した衝撃波が見えている？（謎）
TheJournal, 837:174 (15pp), 2017 March 10
Oya et al.
Table 1
Parameters of the Observed Linesa
MoleculeCS
NH2CHOc
HCOOCH3c
SiOc
CH3CHOc
SO
CCHd
t-HCOOHc(CH3)2Oc
HNCOc

TransitionJ=5–4
12 0,12 –110,11
205,16 –195,15; AJ=6–5
141,14 –131,13 ; A
JN=6 7 –56
N=3–2, J=7/2–5/2,
F=4–3 and F=3–2
12 0,12 –110,11
135,8–134,9 ; EE
12 0,12 –110,11

Frequency (GHz)244.9355565
247.390719
249.0474280
260.5180090
260.5440195
261.8437210
262.0042600
262.0064820
262.103481
262.393513
263.7486250

Eu (K)
35.3
78.1
141.6
43.8
96.3
47.6
25.1
82.8
118.0
82.3

Sμ2 (D2)b
19
157
50
58
82
16
2.3
1.7
30
143
30

Aij (s−1)−42.98×10
1.10×10−3
1.46×10−3
9.12×10−4
6.25×10−4
2.28×10−4
5.32×10−5
5.12×10−5
2.03×10−4
7.18×10−5
2.56×10−4

Synthesized Beam
0 51×0 46 (P.A. −177°. 24)
0 56×0 49 (P.A. 16°. 44)
0 52×0 45 (P.A. −177°. 18)
0 46×0 42 (P.A. −177°. 78)
0 46×0 42 (P.A. −177°. 81)
0 46×0 42 (P.A. 3°. 09)
0 98×0 92 (P.A. −78°. 06)
0 98×0 92 (P.A. −78°. 06)
0 46×0 42 (P.A. 2°. 98)
0 46×0 41 (P.A. 1°. 69)
0 48×0 41 (P.A. 3°. 87)

Notes.
aTaken from CDMS (Müller et al. 2005) and JPL (Pickett et al. 1998).
b
Nuclear spin degeneracy is not included.
cFour channels binded. The spectral proﬁles of the HCOOCH3 and t-HCOOH line (Figure 10) is obtained with binding 16 channels.
dAn outer taper of 1″ is applied.

The L483 dark cloud is located in the Aquila Rift (d=200 pc;

7

Figure 3. Integrated intensity maps of (a) SO (JN=6 7 –56 ), (b) HNCO (12 0,12 –110,11), (c) NH2CHO (12 0,12 –110,11), and (d) HCOOCH3 (205,16 –195,15 ; A). The contours
represent the 1.2 mm continuum map, where the contour levels are the same as those in Figure 2.

The Astrophysical Journal, 837:174 (15pp), 2017 March 10

Oya et al.

CCH

and SiO were observed with the B and 6 receiver in thefrequency range from 245 to 264 GHz (Table 1). Thirty-fourantennas were used in the observation, where the baseline length
ranged from 18.5 to 644 m. The ﬁeld center of the observation
was (a2000 , d2000 )=(18h 17m 29.s 910, -04 39¢39. 60). The
primary beam size (FWHM) is 23 03. The total on-source timewas 25.44 minutes, where a typical system temperature was
A backend
The Astrophysical Journal, 837:174 (15pp),60–120
2017K. March
10correlator was tuned to a resolution of
61.030 kHz and a bandwidth of each chunk of 58.5892 MHz.
The resolution corresponds to the velocity resolution of
0.073 km s−1 at 250 GHz. J1733-1304 was used for the
bandpass calibration and for the phase calibration every 7
minutes. An absolute ﬂux density scale was derived from Titan.
The data calibration was performed in the antenna-basedmanner, and uncertainties are less than 9%.
Images were obtained by using the CLEAN algorithm, where
the Brigg’s weighting with the robustness parameter of 0.5 was
employed. We applied self-calibration for better imaging. A
continuum image was prepared by averaging line-free channels,
Figure 4. Integrated intensity map of SiO (J=6–5). The white contours
and the line maps were obtained after subtracting the continuum
represent the 1.2 mm continuum map, where the contour levels are the same as
those in Figure 2.
component directly from the visibilities. Synthesized beam sizes

CS

SiO

3

Figure 2. Integrated intensity maps of (a) CCH (N=3–2, J=7/2–5/2, F=4–3 and 3–2) and (b) CS (J=5–4). The black contours represent the 1.2 mm
continuum map, where the contour levels are 10σ, 20σ, 40σ, 80σ, and 160σ, where the rms level is 0.13 mJy beam−1. The outﬂow axis is along the red arrow with a P.
A. of 105° in panel (b). The PV diagram in Figure 8 is prepared along the red arrow with a P.A. of 15° in panel (b), which is centered at the position with an offset of
4″ to the southeast from the continuum peak along a P.A. of 105°.

SO

HNCO

NH2CHOHCOOCH3

Figure 3. Integrated intensity maps of (a) SO (JN=6 7 –56 ), (b) HNCO (12 0,12 –110,11), (c) NH2CHO (12 0,12 –110,11), and (d) HCOOCH3 (205,16 –195,15 ; A). The contours
represent
the 1.2 mm Journal,
continuum837:174
map, where
the 2017
contourlevelsThe Astrophysical(15pp),
March
10are the same as those in Figure 2.
Oya et al.

and SiO wereHNCO
observed with the B and 6 receiver in thefrequency range from 245 to 264 GHz (Table 1). Thirty-fourantennas were used in the observation, where the baseline length
ranged from 18.5 to 644 m. The ﬁeld center of the observation
was (a2000 , d2000 )=(18h 17m 29.s 910, -04 39¢39. 60). The
primary beam size (FWHM) is 23 03. The total on-source timewas 25.44 minutes, where a typical system temperature was
60–120 K. A backend correlator was tuned to a resolution of
61.030 kHz and a bandwidth of each chunk of 58.5892 MHz.
The resolution corresponds to the velocity resolution of
0.073 km s−1 at 250 GHz. J1733-1304 was used for the
bandpass calibration and for the phase calibration every 7
minutes. An absolute ﬂux density scale was derived from Titan.
The data calibration was performed in the antenna-basedmanner, and uncertainties are less than 9%.
Images were obtained by using the CLEAN algorithm, where
the Brigg’s weighting with the robustness parameter of 0.5 was
employed. We applied self-calibration for better imaging. A
continuum image was prepared by averaging line-free channels,
Figure 4. Integrated intensity map of SiO (J=6–5). The white contours
and the line maps were obtained after subtracting the continuum
represent
1.2 mmcontinuum
the contour levels are the same as
Figure 5.theMoment1 mapsof (a) map,
SO (JwhereN=6 7 –56 ) and (b) HNCO (12 0,12 –110,11). The black contours represent the 1.2 mm continuum map, where the contour levels
thoseFigurecomponent
directly
fromin thevisibilities.
sizes
are theinsame
as 2.
those in Figure 2. The PV diagrams in Figures 7, 9, 11, and 15 are preparedalong theblack arrowpanel(a) (P.A. = Synthesized15°) and along beam
the direction

Figure 5. Moment 1 maps of (a) SO (JN=6 7 –56 ) and (b) HNCO (12 0,12 –110,11). The black contours represeare the same as those in Figure 2. The PV diagrams in Figures 7, 9, 11, and 15 are prepared along the black
perpendicular to it centered at the continuum peak.

SO

Figure 6. Schematic illustration of the disk/envelope system in L483. The
midplane of the disk/envelope is extended along a P.A. of 15°, and its westernside faces the observer. The line emission on the western side of the protostar
seems to be missing, according to the observation (see text).
8

perpendicular to it centered at the continuum peak.

3
protostellar sources with ALMA (#2013.1.01102.S; P.I. N.
Sakai), and the above lines are all detected in the other sources

for the spectral lines are listed in Table 1. An rms noise level forthe continuum is 0.13 mJy beam−1, while those for CCH, CS,
SO, HNCO, NH CHO, HCOOCH , and SiO maps are derived

protostellar source
Sakai), and the ab(B335 and NGC
setup (Imai et al.
their detections arefor each of these
Section 1, this sou(Hirota et al. 2009of COMs suchcharacteristic of
Yamamoto 2013)
CS, SO, HNCO, N
Figures 2–4.
The CCH (Nemission is exten
taper of 1″ is appliof the image (Figmolecule CCH ar
scale conﬁrms th
distribution has a

The Astrophysical Journal, 837:174 (15pp), 2017 March 10

Oya et al.

The Astrophysical Journal, 837:174 (15pp), 2017 March 10

Oya et al.

Figure 9. PV diagrams of SO (JN=6 7 –56 ; panels (a) and (b)) and HNCO (12 0,12 –110,11; panels (c) and (d)) along the disk/envelope direction (P.A. = 15°; panels (a)and (c)) indicated in Figure 5, and the direction perpendicular to it (P.A. = 105°; panels (b) and (d)). The contour levels are every 5σ and 3σ, where the rms levels are
6.1 and 6.5 mJy beam−1 for SO and HNCO, respectively.

whose velocity shift from the systemic velocity is larger than
5 km s−1. However, the redshifted component is brighter thanOya
et al.
the blueshifted component.
Thisfeature is also seen in its PV
diagrams (Figures 9(c) and (d)). This implies that the HNCO
compact high-velocity-shift component
concentrated
distribution
is asymmetric towardin the vicinity of the protostar. Thisasymmetrygas inthe
the southwestern part of the
protostar,
whose maximum
velocity implies
shift lessfrom
also point-like at thethe
resolutionof the observations.
In addition
system (Figure 6).
as high as about
6 km s−1.
(5.5 kms−1) theis disk/envelope
HCOOCH3,velocity
we tentatively
detected
to NH2CHO and systemict-HCOOH emission concentrated
the perpendicular
protostar.
Along thenear
line
to the disk/envelope direction
The distribution of
SiO (J=6–5)
is different
from those extendedof
NH2CHO(P.A.
= 105°),
a component
over 4.4.
a 20″
scaleandisHCOOCH3
the above molecular species. It has a slight extension towardThe
mostcomponent
notable result inthis study is detections of theobserved(Figurethislooks
the northeastern direction
from thecontinuum7(b)).
peak, asAlthoughshown
and aHCOOCH
COMs,
NH2CHOin the velocity
structure,
it likely
tracespart of3. These two species are
in Figure 4. The complicatedextension is signiﬁcant,considering
the concentrated
around the protostar, as shown in the moment 0
The sizethe maps
synthesized beam ofthe
thisoutﬂow
observation(~0. 5).
cavity.
Figure8 ofshowsthe (Figures
PV diagram
the line proﬁles of NH2CHO
3(c), (d)).along
The spectral
SiO distribution deconvolved
the outﬂow
synthesizedlobe
beamonis theline acrossbythesideof the position are dominated by
the protostar
and southeastern
HCOOCH3 toward
0. 70 ´ 0. 43 (P.A. = 35°).
redshifted(Figure10), which is similar to theprotostar, which is indicated by ana arrowin component
Figure 2(b).
The
HNCO spectrum. Since the redshifted component is enhanced
elliptic
featureofthePV
diagramcharacteristicoftheoutﬂow
for all the HNCO, NH2CHO, and HCOOCH3 spectra
5cavity wall is clearly observed. Althoughbothredshifted
consistently,
theirthedetections
are secure, i.e., they do notcorrespond
The redshifted line
and blueshifted components with respect totootherthespecies/transitions.
systemic−1
proﬁleimpliesanasymmetric
distributionof those molecules invelocity (5.5 km s ) are seen, the center velocity of the elliptic

the vicinity of the protostar, as mentioned in Section 4.3.
However, the origin of the asymmetry is puzzling and is left for
future high spatial resolution observations.

Figure 7. PV diagrams of CS (J=5–4) along the disk/envelope direction (P.A. = 15°; panels (a) and (c)) indicated in Figure 5, and the direction perpendicular to it(P.A. = 105°; panels (b) and (d)). Panels (c) and (d) are the blow-ups of the central parts of panels (a) and (b). The contour levels are every 5σ, where the rms level isdashed lines represent
systemic2017
velocity(5.5 km10s−1). The rectangle in each panel represents the spatial and velocity resolutions.
7.6
mJy
beam−1. The white Journal,
The
Astrophysical
837:174the(15pp),
March(P.A. 160°) and is slightly more extended than the 1.2 mm dust
continuum. This slightly extended component will be discussed
in Section 4. On the other hand, the SO (JN=6 7–56 ), HNCO
(12 0,12 –110,11), NH2CHO (12 0,12 –110,11), and HCOOCH3(205,16 –195,15; A ) distributions are highly concentrated to thecontinuum peak (Figure 3). The sizes of the distributions
deconvolved by the synthesized beam are 0. 56 ´ 0. 39 (P.A.
92°) and 0. 26 ´ 0. 16 (P.A. 42°) for SO and HNCO,
respectively. The HNCO distribution is almost point-like.
Similarly, the distributions of NH2CHO and HCOOCH3 are

Figure 8. PV diagrams of CS (J=5–4) along the line across the outﬂow
indicated in Figure 2(b). The contour levels are every 5σ, where the rms level is
7.6 mJy beam−1.

4. Velocity Structure
4.1. Geometrical Conﬁguration of the Disk/Envelope System

5. Analysis with the Infalling–Rotating Envelope Model
HNCOで回転が見えることはまれ
In order to understand the chemical differentiation observed
Redshift側が強いのはおそらく円盤の質量分布の非対称性のせい
above in terms of the physical structure around the protostar,

we analyze the kinematic structure of the disk/envelope
system. In L1527, TMC–1A, and IRAS 16293–2422 A, thekinematic structures of the infalling–rotating envelopes are
successfully explained by a simple ballistic model (Sakai et al.
2014b, 2016, Oya et al. 2015, 2016). Hence, we apply the samemodel to the CS data. However, we cannot reproduce theoverall velocity structure with the infalling–rotating envelopemodel. As explained below, we need to consider the twophysical components: the infalling–rotating envelope and thecentrally concentrated component.
The infalling–rotating envelope model employed in thisstudy is essentially the same as that reported by Oya et al.
(2014). Since we are interested in the velocity structure of the

前ページFig 2(b)のPA=15dの線に沿ったカット
CSで双極流のcavity wallが見えているfeature is slightly redshifted. This elliptic feature is consistent
with the conﬁguration illustrated in Figure 6. It is quite similar
to that observed in the nearly edge-on outﬂow system of IRAS
15398–3359 (Oya et al. 2014), where both the redshifted andblueshifted components can be seen in each outﬂow lobe.
Hence, it is most likely that the outﬂow of L483 blows nearlyon the plane of the sky at least in the vicinity of the protostar, incontrast to previous reports (e.g., Fuller et al. 1995). Hence, the79The Astrophysical Journal, 837:174 (15pp), 2017 March 10

Oya et al.

The Astrophysical Journal, 837:174 (15pp), 2017 March 10

Oya et al.

回転・収縮envelope

Figure 14. PV diagrams of SO (JN=6 7 –56 ), where the color maps are the same as those in Figures 9(a) and (b). The blue contours represent the results of the
Keplerian model, where M=0.15 M: and i=80°, and the emission is simply assumed to be only inside the centrifugal barrier (100 au). The contour levels are every
20% from 3% of each peak intensity.

Kepler円盤

Figure 11. PV diagrams of CS (J=5–4), where the color maps are the same as in panels (c) and (d) of Figure 7. The black contours in panels (a) and (b) represent theresults of the infalling–rotating envelope models, where M=0.15 M: , rCB=100 au, and i=80°. The blue contours in panels (c) and (d) represent the results of the
Keplerian model with the above M and i values, where the emission is simply assumed to come from the inside of the centrifugal barrier. In panels (a)–(d), the intrinsic
line width is assumed to be 0.2 km s−1, and the model image is convolved with the synthesized beam. The contour levels are every 20% from 5% of each peakintensity.

The Astrophysical Journal, 837:174 (15pp), 2017 March 10

Oya et al.

reproduce the high-velocity-shift part of the PV diagrams of CS
the (possible) Keplerian disk component inside the centrifugal
Figure 15. PV diagrams of CS (J=5–4) along (a) the disk/envelope direction (P.A. = 15°) and (b) the direction perpendicular to it (P.A. = 105°), where the color
maps are the same as in panels (c) and (d) of Figure 7. The black contours represent the results of an infalling–rotating envelope model combined with the Keplerian
and SO by the Keplerian disk model (Figures 11 and 14). A
barrier (Sakai et al. 2014b; Oya et al. 2016). Hence, the CSmodel,
where the physical parameters are the same as those for the models for Figure 11. The contour levels are every 20% from 3% of each peak intensity.
radius of the emitting region of the maximum velocity-shift
distribution in L483 isTable
different2 from those in L1527 and TMC–
Densities
Abundancesof the2016;
Molecules
towardCSthe Protostar Positionis estimatedto be asandsmallFractional
component (∼6 km s−1) in the diskColumn
1A (Sakai
et al. 2014b,
Oya et Observed
al. 2015), whereas 4 au. We also show the results of the Keplerian disk model
resides only in the infalling–rotating envelope. Such a distribu−2−10size can be regarded
6. SiO Emission
centrifugal
barrierRelative(Sakai et toal. H2017),
thiscombined with the infalling–rotating envelope modelin Density
tion of(10CS14incm
L483)seems to resemble the IRAS 16293–2422 A
Fractional
Abundance)
Column2 (10as the upper limit for the radius of the centrifugal barrier.
Figure 15. In addition, the results of the Keplerian disk modelcase. In IRAS 16293–2422 A, the emission of the normal
The distribution of SiO extends to the northeastern direction
Dust
Temperature
70 K envelope model
100
K species of CS seems130
K from the disk, as well
70 as
K
K to derive the radius
130ofKthe centrifugal
Although it is100
difﬁcult
areoverlaid
on those of the infalling–rotatingisotopic
to comefrom the continuum peak, as mentioned in Section 3. The P.A.
barrier from the SO emission because of the contamination by
in Figures 12 and 13 for reference.
the infalling–rotating envelope (Y. Oya et al. 2017, inof the direction of this extension is ∼35° (Figure 4). Figure 16
CSIn L483, the kinematic structure of
5.1±0.2
6.2±0.2
7.5±0.2
5.3±0.2
9.6±0.3
15.2±0.5
34
the disk component in L483, these sizes will directly givethe gas traced by CS
preparation), although the C S emission traces the envelope
shows the velocity channel maps of SiO. The blueshifted
HNCOthe protostar is similar to those9.6±1.1
14.1±1.6
10.0±1.1 rough estimates
17.7±2.028.6±3.2
for the size of the centrifugal
barrier. They arearoundof H2CO in L1527 and 11.5±1.3
outside the centrifugal barrier(Favre et al. 2014). The differencecomponent of SiO (v LSR < 5.2 km s−1) is offset from theconsistent withthe estimate from the3.9±1.9analysis of the PV
in IRAS 16293–2422 A, which
are explained by a 1.5±0.7
Hof the behavior of CS 1.9±0.9
would originate from the 1.3±0.6higher1.2±0.6
2.4±1.1
NH
protostar by 0 5 (∼100 au), while the weak redshifted
2CS
2CHO
combinationenvelope component and 69.9±36.0
diagrams of107±55
CS.
2000)
bolometric luminosity of80.0±41.2
L483 (13 L :; Shirley et al. 73±38
component appears at the protostar position. The position of
HCOOCH3of the infalling–rotating69.8±35.9
162±84
9

12

t-HCOOH

2.0±1.9

2.4±2.2

Tentative Detection
2.9±2.7

2.1±1.9

3.6±3.4

5.8±5.5

CH3CHO(CH3)2O<0.62<116.1<0.75<119.6Upper Limit<0.95<135.0<0.65<121.6<1.2<183.8<1.9<273.6Note. Derived under the assumption of local thermodynamic equilibrium at the temperature of 100 K. N(H2) is derived from the 1.2 mm dust continuum emission (seetext). Errors are derived from three times the rms noise of the integrated intensity.

source has been recognized as an intermediate source inprevious studies (Sakai et al. 2009a), the present observationcomparable to those reported for the hot corinos, IRAS
16293–2422 (6 ´ 10-10 and 9 ´ 10-9; Jaber et al. 2014) and10