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)

The Astrophysical Journal, 837:151 (7pp), 2017 March 10

O’Dell, Kollatschny, & Ferland

Figure 1. This Hα 0.66×0.55 pc(353″×295″) image of the Orion Nebula is from the 0 2 pixel–1 MUSE data. The samples employed in generating the data in
Table 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,
as discussed in Section 3. Those proplyds ionized by q 2 Ori A are shown in enlargements of 2″×2″ , except for 244.3–439.8, which is a sample of 4″×4″. The
location of the ﬁrst of the Spitzer Space Telescope observations discussed in Section 4 is shown south of q 2 Ori A and further deﬁned in Figure 4.

shifted to near the intersection. However, our adopted value forthe 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

F([O III])/F(Hβ) and F(He I)/F(Hα) we see that all samples are
consistent with ionization by q1 Ori C, with the caveat of theremarks in Section 2.3 about the high Tstar values for Samples C
and F. Taken at face value, F([O III])/F(Hβ) and F(He I)/F(Hα)
indicate that Samples C and F have spatial distances about
equal to rsky , while Samples G, A, and B demand progressively
larger corrections to the true distance. In the left and rightpanels it appears that Sample D could also be associated with
q1 Ori C, but in the middle panel we see that no plausible
correction for distance would make the line ratio compatible

ratios it lies too far to the right of Tstar for q 2 Ori A , which rules
out q 2 Ori A .
The identiﬁcation of the inner ﬁve samples as being ionizedby q1 Ori C falls within the accepted model of this region of
the Orion Nebula, i.e., that it is a concave irregular surface. The
sample apparently ionized by q 2 Ori A is consistent with theBright Bar and the regions southeast from there being ionizedby the nearest hot star. Sample E, although lying outside theBright Bar, is probably ionized by q1 Ori C, as discussed
earlier in this section.
3. Identifying the Ionizing Stars Using Observationsof the Proplyds
There is a completely independent method of identifying thedominant source of ionization using high-resolution images of
the Orion Nebula proplyds. These are young stars withcircumstellar material in or projected on H II regions. Underfavorable orientations and locations the interior protoplanetarydisks are rendered visible in silhouette. More relevant here,
these disks are surrounded by much larger envelopes of gas anddust. A nearby hot star will form an ionization front in this gas

4. Using Spitzer Space Telescope Infrared
Spectra Line Ratios
Published Spitzer Space Telescope observations of theinfrared [Ne III] 15.6 μm and [Ne II] 12.8 μm lines are also
useful for determining the ionizing star. These are particularly
useful because they both arise in the same He+ + H+ zone andare unaffected by emission from the Heo + H+ zone and anychanges of abundances, as is the case for our optical line ratios.
In the study of Rubin et al. (2011) observations were made at
11 locations to the southeast of the central Orion Nebula asshown in Figure 4. These samples all lie outside the Bright Bar,
with the innermost being close to q 2 Ori A . The progression ofincreasing distances is I4, I3, I2, I1, M1, M2, M3, M4, V1, V2,
and V3. They all lie in a smooth region between the Bright Bar
and the boundary of the Extended Orion Nebula (EON).
Sample V2 falls near the edge of the sharp southeast boundaryof the EON, and Sample V3 falls on this boundary. Thesesamples do not agree with the results from the others because
Sample V3 is certainly an ionization front viewed edge-on andSample V2 is probably a region that is rising toward thatfeature (Rubin et al. 2011), and thus our Cloudy models do notprovide a guide to the conditions (Section 2.3). They are not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 .
5. Discussion
The results of Sections 2 and 3 show similar but not identicalresults. Some of the differences may be explained by themethodology.
The F([O III])/F(Hβ) and F(He I)/F(Hα) line ratios areindicators of the conditions in the He+ + H+ zone, as [O III]and He I emission arises only there, as does most of the Hαemission. The gas density decreases (O’Dell 2001) withincreasing distance from the MIF; therefore, most of theemission used in the F([O III])/F(Hβ) and F(He I)/F(Hα) line
ratio method occurs near the onset of the He+ + H+ zone, witha lesser Hα 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.
This means that the ratios using these two ions will giveslightly different results. In the Heo + H+ zone the [N II]
emission is well constrained, but the fact that the Hα emission
comes mostly from the He+ + H+ zone means that theF([N II])/F(Hα) ratio is intrinsically not as good a diagnostic.
In contrast, the proplyds give an indication of the ionizationof objects located somewhere along the line of sight from theforeground Veil and the MIF. The Orion Nebula Cluster iscentered near the Trapezium and monotonically decreases in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)

embedded 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.

below). This rules out the H II region scenario proposed by
López-Sepulcre et al. (2013) for VLA 12.
Although the position of the knot VLA 12C at 6 and 3 cm
appears somewhat (∼1 5) displaced to the northeast of theposition of HOPS 108 (as obtained from 24 and 8 μm data;
Megeath et al. 2012), at 1.3 and 0.7 cm the radio emission
appears to be closer to the HOPS 108 IR position (seeFigure 1). This shift in the position could indicate either an
opacity gradient or that the radio emission observed at shorter
wavelengths originates in a different object than the emission
observed at longer wavelengths. This last suggestion is
conﬁrmed by our data of subarcsecond angular resolution atlong wavelengths (Figure 4), which reveal two radio sources
spatially separated by ∼2″. We identify the northeasternmost ofthese two radio sources with the VLA 12C knot and the othersource, whose position coincides with the IR source, with theradio counterpart of the HOPS 108 protostar. Recent ALMA
observations with an angular resolution of 0 12 (J. Tobin et al.
2017b, in preparation) reveal a compact source at 870 μm
whose position coincides within <0 1 with the cm position(Table 3) conﬁrming the nature of HOPS 108 as a protostar andits association with the cm radio source. Furthermore, thespectral index of HOPS 108 obtained from the centimeter datathat separate its emission from that of VLA 12C shows that
HOPS 108 has a positive spectral index (α=0.7±0.3;
Table 3 and Figure 3), thus it is associated with thermal (free–free) emission, in contrast to the knots of VLA 12, whose
emission is non-thermal. From the empirical correlation
between cm ﬂux density and bolometric luminosity for
protostellar objects, (Sn d 2 mJy kpc2) = 0.008 (L bol L :)0.6(Anglada et al. 2015), and using our measured ﬂux density at
3 cm (Table 3), we estimate a bolometric luminosity
Lbol;7 Le for HOPS 108. This value of the bolometricluminosity is close to the lowest end of the range considered by
Furlan et al. (2016) in the modeling of this object, being muchsmaller (Table
than 3 andearlyFigure
estimates(Lbol;1000
Le;as Crimier
regime3), conﬁrming
its nature
a YSO.
et al.
et studied
al. 2013)
that wereSo
far,2009;
MIPS López-Sepulcre
2297 has not beenin greatdetail,based
but itonisseparating the emission of HOPS 108 from that of the VLA
12 knots conﬁrms HOPS 108 as a low-luminosity object.
The radio source VLA 11, which was associated with FIR 3(HOPS 370) and reported as unresolved by Reipurth et al.
(1999), is clearly resolved as a very elongated bipolar radio
jet along a PA;5° in our A-conﬁguration data at 5 cm
(Figure 4). The radio source is also resolved in the 3.6 cm
image obtained by combining archive data from 1994 to 2000.
The spectral index is α=0.27±0.10 (Table 3, Figure 3),
which is a reasonable value for a thermal free–free radio jet(Anglada 1996). HOPS 370 is known to be associated with anoutﬂow extending along a PA similar to that of the radio jet(González-García et al. 2016; see Figure 2, left). The
morphology of VLA 11 appears to change in the 0.7 cm image(Figure 1), which shows weak extended emission in a direction
perpendicular to the jet, suggesting a noticeable dust contribution at this wavelength. To prevent dust contamination, the0.7 cm data point has been excluded in the ﬁts to determine thespectral index of this and the remaining thermal sources.
In the lower-resolution images at 6 and 3 cm (Figure 1),
VLA 11 appears to be blended with the radio emission of MIPS
2301 (Megeath et al. 2012, an IR source also known as MIR
22, Nielbock et al. 2003) located ∼3″ south of VLA 11. MIPS
2301 is well separated, but only marginally detected, in thehigher-resolution 3.6 cm image shown in Figure 4 and in the
A-conﬁguration images at 5 cm obtained with the robust
weighting parameter set to values >1 (not shown in Figure 4).18
With the data currently available we obtain a positive spectral
index, α=0.30±0.23 (Table 3, Figure 3), consistent withthermal free–free emission from a YSO. More sensitive
observations are required for a better characterization of thenature of the radio emission of this object.
The source MIPS 2297 (Megeath et al. 2012, also known as
MIR 23; Nielbock et al. 2003) is marginally detected at 1.3 cm
and 0.7 cm (Figure 1), but clearly detected in the 0.9 cm
observations of J. Tobin et al. (2017a, in preparation). It shows
ALMA
observationsof higher
angular resolution(J. Tobina positive
spectral index(α=1.5±0.5)
in the centimeter
et al. 2017b, in preparation). From the empirical correlation

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)

Table 1
Parameters of the Observed Linesa
MoleculeCS
NH2CHOc
HCOOCH3c
SiOc
CH3CHOc
SO

TransitionJ=5–4
12 0,12 –110,11
205,16 –195,15; AJ=6–5
141,14 –131,13 ; A
JN=6 7 –56

Frequency (GHz)244.9355565
247.390719
249.0474280
260.5180090
260.5440195
261.8437210

Eu (K)
35.3
78.1
141.6
43.8
96.3
47.6

Sμ2 (D2)b
19
157
50
58
82
16

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

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)
46
0
×

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

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°.

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
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,
and the line maps were obtained after subtracting the continuum
component directly from the visibilities. Synthesized beam sizes

(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 arealso point-like at the resolution of the observations. In addition
to NH2CHO and HCOOCH3, we tentatively detected thet-HCOOH emission concentrated near the protostar.
The distribution of SiO (J=6–5) is different from those of
the above molecular species. It has a slight extension toward
the northeastern direction from the continuum peak, as shownin Figure 4. The extension is signiﬁcant, considering thesynthesized beam of this observation (~0. 5). The size of the
SiO distribution deconvolved by the synthesized beam is
0. 70 ´ 0. 43 (P.A. = 35°).

reproduce the high-velocity-shift part of the PV diagrams of CS
and SO by the Keplerian disk model (Figures 11 and 14). A
radius of the emitting region of the maximum velocity-shiftcomponent (∼6 km s−1) in the disk is estimated to be as small
as 4 au. We also show the results of the Keplerian disk model
combined with the infalling–rotating envelope model in
Figure 15. In addition, the results of the Keplerian disk modelare overlaid on those of the infalling–rotating envelope modelin Figures 12 and 13 for reference.
In L483, the kinematic structure of the gas traced by CS
around the protostar is similar to those of H2CO in L1527 andH2CS in IRAS 16293–2422 A, which are explained by a
combination of the infalling–rotating envelope component andthe (possible) Keplerian disk component inside the centrifugal
barrier (Sakai et al. 2014b; Oya et al. 2016). Hence, the CS
distribution in L483 is different from those in L1527 and TMC–
1A (Sakai et al. 2014b, 2016; Oya et al. 2015), where CS
resides only in the infalling–rotating envelope. Such a distribution of CS in L483 seems to resemble the IRAS 16293–2422 A
case. In IRAS 16293–2422 A, the emission of the normal
isotopic species of CS seems to come from the disk, as well as
the infalling–rotating envelope (Y. Oya et al. 2017, in
preparation), although the C34S emission traces the envelope
outside the centrifugal barrier (Favre et al. 2014). The differenceof the behavior of CS would originate from the higher
bolometric luminosity of L483 (13 L :; Shirley et al. 2000)