Star Formation NewsletterNov 14, 2014
Shin Koyamatsu

1. Kompaneets model fitting of the Orion-Eridanus superbubble
Andy Pon, Doug Johnstone, John Bally and Carl Heiles	

概要• OB Associationからの風や超新星によりISMにSuperbubbleと呼ばれる大きな穴が空く• 最も近い大質量星形成領域であるオリオン分子雲には20°×45°の大きく偏平したSuperbubbleがある• Orion-Eridanus superbubbleをKompaneetsモデルでフィット• SuperbubbleのEridanus側が太陽から遠い方向にあるモデルの方が合うBisnovatyi-Kogan, Blinnikov & Silich 1989; Mac Low et al.
1989;
Tenorio-Tagle, Rozyczka & Bodenheimer
1990;
2.2 Orion–Eridanus Superbubble
Tomisaka 1992, 1998; Stil et al. 2009), Kompaneets models
In this paper, all references to north or souhave been shown to provide an excellent first-order solution
1本目つづきcreasing or decreasing declination and referen
for the shape of a bubble expanding in an exponential
KompaneetsModelwest refer to increasing or decreasing right asc
atmosphere (e.g.
Mac Low et al. 1989). Kompaneets modotherwise specified. References to above or b
are the current, de facto
choice for analytic models of• els
純粋なexponential
atmosphereへのショック波の伝播のsemi-analytic解
will refer to directions along the major axis
superbubbles. For this paper, we will use the continuous• injection
天の川銀河の円盤中のsuperbubbleの進化に適用できる(W4superbubble,
Eridanus superbubble
either towards the lowe
variant described in detail by Basu et al. (1999)
tion of the exponential atmosphere or the hOphiuchus
andwill refer superbubble,
to such a modellocal
simplybubble)
as a Kompaneets
portion, respectively. Thus, the top of the b• model.
exponential atmosphere内を膨張するbubbleの形状のよい1次近似解
the portion of the bubble farthest from the sThe three key assumptions of the Kompaneets model• arekey(1)assumptionrounded by lower density regions of the expo
the pressure within the bubble is spatially uniform,
sphere, and the bottom of the bubble will be(2)1.thebubble expands in the direction of the surface normal
bubble内の圧力は空間的に一定
the bubble closest to the source and surroun
at all locations, and (3) the expansion speed of the bubble is
2. bubbleはすべての場所で表面に垂直な方向に拡がる
density portions of the exponential atmosphgiven
by Hugoniot jump conditions for a strong shock, wherewe will use upper and lower to refer to direct
the3.exteriorpressure is negligible compared to the interior
膨張のはやさは強いショックに対するHugoniot
jump条件
major axis of the bubble. While ‘above’ or ‘topressure. The Kompaneets model predicts that the location
will generally refer to larger absolute magnitudof the wall of the superbubble is given by!
"#$latitude, we emphasize that we intend these1 z/2Hy2−z/Hto directions along the major axis of the supr(z, y) = 2H arccos e
1−+e,(1)24H 2
not a direction along a line of constant Gala
since the Orion–Eridanus superbubble’s majo
where r is the perpendicular distance from the central axis
to make an angle of approximately 30◦ to theof the superbubble to the bubble wall, z is the distance along
Galactic plane in the plane of the sky.
the central axis from the driving source, H is the scaleheight
The Orion star-forming region is the neaof the exponential atmosphere, and y is a parameter withforming region to the Sun that is actively
dimensions of length which varies from 0 to 2H and essenand is located 400 pc away in the constella
tially specifies the evolutionary stage of the bubble, with(Hirota et al. 2007; Menten et al. 2007; Sa
more evolved bubbles having a larger value of y.
2007). It is approximately 130 pc distant fromAt early times in the evolution of a Kompaneets model,plane and contains both young subgroups ofthe superbubble remains relatively spherical and the y pamolecular clouds in which stars are actively
rameter is roughly equal to the radius of the superbubble. After a bubble expands to a radial size approximately
Bally 2008).

6

1本目つづきA. Pon et al.

Kompaneets Fitting of Ori–Eri Superbubble

3!"#$!$%&"'&"()*$+,,-$!"#$%$

Figure 1. Orion–Eridanus superbubble as seen in Hα. Labels forthe various major components of the bubble have been added tothe image from Di Cicco & Walker (2009). Arcs A, B, and C arecollectively referred to as the Eridanus filaments. The locations ofadditional small molecular clouds have been indicated, althoughmany of these clouds are not obviously visible in the image, suchthat the arrows have been placed to point towards the known
locations of the clouds.

45◦ elongated superbubble structure, with enhanced Hαflux within the confines of the bubble (e.g. Pickering1890; Barnard 1894; Elliott & Meaburn 1970; Elliott 1973;
Isobe 1973; Reynolds, Roesler & Scherb 1974; Sivan 1974;
Johnson 1978; Reynolds & Ogden 1979; Boumis et al. 2001;
Haﬀner et al. 2003). The east side of the superbubble isdominated by a bright crescent of Hα emission, which is referred to as Barnard’s Loop (Pickering 1890; Barnard 1894).
The west side of the superbubble is composed of a hook

Figure 4. Schematic diagram of models T and A. Model T isshown in the top panel and model A is shown in the bottom panel.
In each of these schematic diagrams, the superbubble major axis
is inclined roughly 30◦ into the page, in order to match the oﬀset
betweenthe2. superbubble
and thenormalbytothethe
Galactic
FigureDiﬀuse X-rayaxis
emission
observed
ROSAT
satelplane
Thus,Thethe topmajor
axisshowsof thethebubblelite apparent(SnowdeninetFig.
al. 1.
1997).
panel
0.25 inkeV,◦modelis stillfrom shows
the normal
to keV,the Galactic
R12 Abandand 35the inclined
bottom panel
the 0.75
R45 band.
plane.
The contours are integrated intensities of Hα from the WHAM
survey (Haﬀner et al. 2003) and are logarithmically spaced witheach contour corresponding to a factor of 2 increase in intensity.
The 1.
lowestcontour corresponds
to an integrated intensity of 10 R.
Table
Kompaneetsmodel properties
The locations of Arc A, Arc B, and Barnard’s Loop are labelled.

eter, rathe
bubble, asparameters
bubble in t
the inclina
sky, and th
cation is coand the faa distanceThe b
the bottom
model A (fgiven in thof this modA slig
reasonable
Because thend and nmuch widebubble oriefits, howev
In model A
the closest
300 pc dist
bubble is 2
T. For a dfrom this m4

DISCU

Model T dogy of theArcs A andmajor axis
the bubble
driving sou2. Mid-J CO observations of Perseus B1-East 5: evidence for turbulent
dissipation via low-velocity shocks
Andy	  Pon,	  Doug	  Johnstone,	  Michael	  J.	  Kaufman,	  Paola	  Caselli,	  and	  Ren	  é 	  Plume	  	  概要• 巨大分子雲は超音速乱流を含んでいるが、MHDシミュレーションによるとすぐに消失する• 乱流の散逸は分子雲内に温かいガスを作る可能性がある• Perseus B1-East 5領域のCO (J=5–4, 6–5)をHerschelで観測• 過去の低回転遷移の測定量も合わせてPDRモデルをフィット• B1-E5では6–5の積分強度が異常に強い → 温かいガス成分• 体積フィリングファクタ0.15%の低速度ショック領域での乱流散逸によるショック加熱• 乱流エネルギーの散逸率は3.5×1032 erg s-1

2本目つづきMid-J CO Observations of Perseus B1-E5

7Mid-J CO Observations of Perseus B1-E5

Figure 5. Spectra of various CO lines observed towards B1-E5.
The black, green, orange, and red lines correspond to the J = 1→ 0, 3 → 2, 5 → 4, and 6 → 5 transitions of 12 CO, while thedark blue and light blue lines correspond to the J = 1 → 0 and
2 → 1 transitions of 13 CO. The J = 5 → 4 and 6 → 5 intensities
have been scaled up by factors of 7 and 50, respectively, while the12 CO J = 1 → 0 intensities have been scaled down by a factor of
1.5.

978 being found towards infrared dark clouds (Fontani et al.
2012). The 12 CO J = 1 → 0 data indicate that the depletion
factor for B1-E5 is approximately 1.5, suggesting that very
little CO has frozen out of the gas phase in B1-E5. Since
the Kaufman et al. (1999) PDR code uses a steady statechemistry, PDR models that do not include freeze out areselected.
The 13 CO observations can also be used to calculate
the CO depletion factor. Under the assumption that the12CO J = 1 → 0 line is highly optically thick, the excitation
temperature of the gas is given by⎡ ⎛⎞⎤−1TuTex = Tu ⎣ln ⎝+ 1⎠⎦ ,(2)uTR + exp T T/T( u bg )−1

Figure 6. Integrated intensities of various 12 CO lines from the Kaufman et al. (1999) PDR models are shown as the lines, while theobserved integrated
ex intensities are shown as symbols. Black, blue, and red lines correspond to densities of 102 , 103 , and 104 cm−3 ,respectively, while dashed, solid, and dotted lines correspond to ISRFs of 0.31, 3.1, and 31 Habing, respectively. Triangles denote
integrated
intensities of the 8 km s−1 component while squares are used for the 3 km s−1 component.
Rbg The error bars show three times
the uncertainties in the integrated intensities. The column density for the PDR models has been set to 7 × 1021 cm−2 , to match thetotal column density of B1-E5.
uwhere T is the excitation temperature of the emitting gas,
T is the observed peak intensity of the line, T is the
2.7 K background temperature, and T is the upper energy
level of the molecular line when expressed in units of kelvin.
traces lower density gas than the 8 km s component. This
vided by a model with a density of 10 cm and column
Since
3 forand8 skmcomponent,s−1 12 CO
J = 1lines
are having lower inis consistentwith→the 30kms component
densityof 4 both× 10 cmthe, while
the 8 kmtegratedintensities
in C O andNH lines (Sadavoy et al.
thebest-fitting
model ithas aisdensity
just under
10 bothcm
flattopped,likelythatlinesareoptically
thick,2012; Sadavoy et al., in preparation).
and a column density of 1 × 10 cm .
suchthatsincethe
temperature
forthe thecompoFor
3 km s twocomponent,
the Kaufman et al. (1999)
In general,largerexcitation
ISRFs imply a larger
incommodels with a density
ing energy flux, the integrated intensities of the CO lines−1 and ISRF of 10 cm and 3.1 Habing,nentswithwouldandK,to with3 cm
kmands 0.31 component
or 10Habing, also reasonably reproduce theincrease
increasingbe
G , 11in orderfor 16a cloudmain- theline ratios of the CO J = 1 → 0 and 5 → 4 lines, withtain an equilibrium between heating and cooling terms. The
being
thecoolercomponent.
respect to the 3 → 2 line. The best-fitting meudon model
CO SLED also becomes flatter with increasing G , as the13latter eﬀect
the 3opticallykm s ratios also
has a anddensity of 10 cm , buttemperatureof CO emittingthat
gas increases.
Assuming
theThisCO lines toarethin
the best-fitting kosma-τ model has a slightly higher density
is not necessarily trivially expected, as a larger ISRF willof 10 lines
cm . The
CO ratios
are roughly consistent withcause
and othermolecules, totemperature
only form at larger optithatCO,theexcitation
for theseis thatderivedthe meudon models. All three codes
cal depths, where12
the gas need not be significantly warmer
13 underpredict the COfromCOJ ISRFs.
= 1 Higher→ 0densities
line, also
the column
J = 6 → 5 todensity
3 → 2 ratio, ofwith theCOkosma-τ model being athanwherethe
CO forms
in lowerfactor of 17 low, a 6σ discrepancy.
produce larger intensities and flatter SLEDs, as the gas dencan
bederivedusing
sity approaches the large critical densities of the higher J
The ratio of the CO J = 1 → 0 to 3 → 2 lines is muchlines.
⎛⎞−1 than the 3 km s
closer to unity for the 8 km s component
component, which makes higher density models (10 cm
Based upon the integrated
intensities, the best-fitting
28πW
ν ofkunity
gl ISRF
1 of 10for the Kaufman1et al. 1999 and meudon codes, 10 cm
models withan ordera density⎝al. have⎠ better for the 8 km s) and meudon) code)*fit this ratio*−for the kosma-τcmN, =for both the Kaufman et(1999)3hν
c Ahdo a reasonable
gu expcomponent thanhνthe 3 km s component. The meudon 10codes. These models
job at reproducing
theexp−1−1
4.2

22−2−1−3−1−1

52218−3

3−2−1

34

4.10Kaufman et al. models−3−3120

The Kaufman et al. (1999) PDR models are run for semiinfinite, constant density slabs of material illuminated on
one side. The models use a microturbulent doppler linewidthof 1.5 km s −1 , which would correspond to an FWHM of
2.5 km s−1 that is just slightly larger than the observed
linewidths of B1-E5. In comparing the PDR model output
with the observations of B1-E5, it is assumed that the PDR
emission fills the beam. While the Kaufman et al. (1999)
models handle the chemistry of most common molecules
quite well, they do not self-consistently calculate the selfshielding of 13 CO. Therefore, only the 12 CO data will becompared to the Kaufman et al. (1999) PDR models. PDR
models with an ISRF of log(G0 ) = -0.5, 0.5, and 1.5 (G0= 0.31, 3.1, and 31 Habing) are selected, where the average−33−1

4.2−3−31312

12−1−145

3−3−3−1−34−1

12kTandintegrated intensities of lower J lines, both 13 COex COlines for the meudon code, but underpredict the integrated
intensity of the 12 CO J = 6 → 5 line. A lower G0 of 0.31and a higher density of 104 cm−3 also produces reasonable−hνet al. (1999) code.
integrated intensities from the Kaufman
lSince the 3 km s−1 component haskTlower
ex integrated intensities, these models suggest that the 3 km s−1 component×)1)
1 − exp**g expQ)

kTbg reproduces the integrated intensity
cm−3 model reasonably
ratios of the 12 CO J = 5 → 4 and 1 → 0 lines with the 3→ 2 line, but again slightly underpredicts the ratio of the 6→ 5 to 3 → 2 lines. In Fig. 10, the error bar for the 6 → 5−Epointl that crosses the 104 cm−3 model line is that of the 3−1−1kTkm
exs component, not the 8 km s component. The 8 km
−1s component error range for the ratio of the 6 → 5 to 3 →*,(3)c 0000 RAS, MNRAS 000, 000–000⃝

where W is the integrated intensity of the line, A is the Einstein A coeﬃcient for the transition, h is Planck’s constant,

3. A Spitzer View of The Giant Molecular Cloud Mon OB1 East/NGC 2264
Rapson,	  Valerie	  A.,	  Pipher,	  Judith	  L.,	  Gutermuth,	  Robert	  A.,	  Megeath,	  S.	  Thomas,	  Allen,	  Thomas	  S.,	  	  Myers,	  Philip	  C.,	  Allen,	  Lori	  E.	  	  

概要• 巨大分子雲Mon OB1 Eastの3.6, 4.5, 5.8, 8.0, 24μm画像をSpitzerで観測• 星周円盤 and/or 星周エンベロープをもつYSOを赤外超過により同定・分類、また分子雲の物質に対する分布を調査• YSOの局所面密度とダスト減光から見積もられる分子ガスの柱密度とのあいだにべきの相関関係を発見• 内部の星形成領域NGC 2264の内と外で異なるべき、円盤の割合• Mon OB1 cloudでは非一様な同時でない星形成が起きている30:00.0
30:00.0
9:00:00.0

9:00:00.0

42:00.0

6:40:00.0

Right ascension

01632

10:00:00.0

Declination

10:00:00.0
30:00.0

Declination

30:00.011:00:00.0

11:00:00.0

3本目つづき42:00.0

6:40:00.0

Right ascension

Fig. 1.— Left:0 Spitzer 4.5µm invertedgrayscale48image of Mon
OB1 East.
NGC 2264
is outlinedby the red
111
127
32 95
16 80
64 12780 14395 159111
6448
squares which encase the S Mon region (top square) and Cone region (bottom square) of Sung et al. (2008).
All Spitzer data outside of the two red boxes represent Mon OB1 cloud. The position of the two reference
regions used to estimate the number of field stars towards Mon OB1 East are outlined by the magenta boxes
to the east and west of the cloud. We used 2MASS data for all sources within the magenta boxes to estimate
the number of field stars. Right: Same Spitzer 4.5 µm inverted grayscale image of Mon OB1 East but withClass 0/I (red), Class II (green) and transition disk (blue) sources overlaid.

1431593本目つづき– 26 –– 27 –2Fig. 9.— YSO mass surface density vs. molecular gas mass column density of all Spitzer identified Class Fig. 10.— Histograms of ⌃Y SO /⌃gas for NGC 2264 (top left), Mon OB1 cloud (top right), and all of Mon
0/I (red triangles), Class II (green squares) and transition disk sources (blue diamonds) for NGC 2264 (top OB1 East (bottom), separated by evolutionary class. Histograms for Class II YSOs are in black, normalizedleft), Mon OB1 cloud (top right), and all of Mon OB1 East (bottom). Overplotted are polygons marking the to the peak bin. Class I YSOs are in red, and are multiplied by a factor of 3.7, the median ratio of Class II
fit to, and estimated spread in, data of the Mon R2 molecular cloud (black; N=2.67 ± 0.02), and Ophiuchus to Class I sources in nearby clusters (Gutermuth et al. 2009) and then normalized to the peak bin. Green ismolecular cloud (gray; N=1.87 ± 0.03) from Gutermuth et al. (2011). Hatched regions show where AV < 1 the Class II histogram, normalized to the peak bin and divided by 30. This is the upper limit to the expected
frequency of edge on Class II sources that could be misclassified as Class 0/I objects.(Gutermuth et al. 2009,
and the gas column density cannot be reliably determined from the near-IR extinction.
2011). Vertical dashed lines mark the boundary between young protostar rich regions (left), embedded cloud
populations (center) and regions where gas has been dispersed (right; Gutermuth et al. 2011)