Star Forma)on Newsle0er#274 16-20本目
麻生 有佑

Dense gas in the Galac)c central molecular zone is warm and heated by turbulence
Adam Ginsburg et al. (2015)
Ø Central molecular zoneは極限的な物理状態での星形成を調べられる最も近い場所である。
Ø あらゆる物理を測れるガス温度がダスト温度と異なるという問題が永らくある。密度が足りないとdecoupleして合わない可能性があるが、銀河円盤では合って、中心では合わない。
Ø 加熱方法も光子、X線、乱流、宇宙線など議論が分かれている。

Ø –0.4°<l<1.6°, APEX, p-H2CO (J=3-2)@~214 GHzなど, Δv=0.8 km/s, binary dila)on(マスク)。
Ø 温度マップを作るために32,1-22,0/30,3-20,2の比をとった(比が高いと高温)。
A&A proofs: manuscr
Ø dendrogram clamp ﬁndでS/Nが良い領域を抜き出す。

Ø RADEXでモデルグリッド; n=102.5-7 cm–3で20個、dV/
dR=5 km/s/pc, N(H2CO)=1011-15.1 cm–2で30個、T=10-350 Kで50個で計算したものをスプライン補間して100x101x250。
Ø 観測側はH2CO 30,3-20,2の強度、速度、線幅、
32,1-22,0/30,3-20,2比、32,2-22,1/32,1-22,0比、CH3OH 42,2-31,2の強度で記述されるGaussiansと、Clump ﬁndの結果の2種類を考えた。
Fig. 6. The derived temperature4vs the measured
ratio
5 cm
–3 p-H CO S (3Ø Clumpごとに得られた(ベストフィットの)ライン比と温度の関係を得て、n~10
symbols represent leaves (compact と10clumps) andthe small symbols repre
The temperature includes constraints from the assumed constant H CO
color coded by signal-to-noise in the ratio R , with gray S/N < 5, blu
で温度マップを作った。
The black lines show the modeled temperature as a function of R for
2

⌫

2,1

2

1

1

abundance XH2 CO = 1.2 ⇥ 10 9 . While the measured values appear con
permit any density in the range 104 cm 3 < n < 105 cm 3 ; this plot mDense gas in the Galac)c central molecular zone is warm and heated by turbulence
Adam Ginsburg et al. (2015)の続き
Ø AV>10に興味があるのでUV加熱は無視し、星や超新星も局所的であり、X線は密度が3桁足りない(Ao+2013)。
Ø 宇宙線と乱流の加熱がダストと輝線の冷却と平衡とする。
Ø 宇宙線は深くまで浸透するので加熱は一様とするとζCR<10–14/s
Ø 乱流だけから予想される温度と観測された温度はおよそ合う。
Ø 予想より低温の点は線幅がbulk mo)onか、冷却されやすい高密度領域。
Ø Herschel HiGalのダスト温度よりガス温度の方が高いà平衡にな
く、UV加熱を無視したことと合う。
Ø CMZは高密度な割に星形成トレーサーが検出されず、乱流が重力収縮を抑えていると言われている。今回の観測でも乱流が卓越しているとわかった(cs=0.5-0.9 km/s)。CMZのように高密度であればジーンズ長は長く、乱流によるsub-Jeans-massfragmenta)onの検出も期待される。

A&A proofs: manuscri

Ginsburg et al: APEX C(a)Fig. 10. Plots of the p-H2 CO-measured temperature against the pred
5.1. While no trend is obvious in the data, most of the high signal-to-noline. There are few points with TG,measured > TG,turb , so turbulence gen
to heat the gas to the observed temperatures. The regions with predic
explanation: either their heating rate has been overestimated or their co
described in Figure 12; the blue symbols are compact ‘clump’ sources aindicate regions with lower limits on the measured temperature, i.e. any
TG > 150 K. The gray dashed line shows the TG,turb = TG,measured rela1level are outlined. (b) The same as (a) for the dendrogram-extract
fg > 3 are excluded. Regions with TG,turb > 180 K are indicated withthe ratio R1 , with gray S/N < 5, blue 5 < S/N < 25, green 25 < S/N <in blue and gray are ignored, the agreement between the predicted tem(a)Fig. 12. (a) The fitted temperature of region-averaged spectra as desc
temperature. As in Figure 9, the blue symbols are compact ‘clump’ sour

which point the dust and gas strongly decouple (Laibe, Gonzalez &Maddison 2012). The second parameter involved is the dust-to-gas
ratio ϵ. For ϵ of order unity, dust affects the gas as much as gas affects
dust due to aerodynamic drag. Although the dust backreaction due todust inertia is negligible in most astrophysical systems, it becomes
critical in regions of discs where dust concentrates, such as particles
traps.
In this Letter, we present a mechanism where grain growth andfragmentation result in dust distributions that produce apparent gaps
when the dust backreaction is taken into consideration. We provide
the details on our numerical simulations and synthetic observables
in Section 2, discuss the origin of the apparent gap in Section 3, andconclude in Section 4.

potential of the star and planet (see Fouchet et al. 2007). The discis initially set up as in Fouchet, Gonzalez & Maddison (2010) witha constant surface density $(r) = 19.67 kg m−2 with radial extension from 4 to 120 au, a vertically isothermal temperature profile
T(r) = 15 (r/rp )−1 K, and is free to evolve. Grains have an initially
uniform size s0 = 10 µm and are allowed to grow and fragment.
In Gonzalez et al. (2015), we presented a suite of simulations thatvary the fragmentation threshold: Vfrag = 10, 15, 20, 25 m s−1 and+∞ (i.e. growth only), evolved for 100 000 yr or ∼400 planetary
orbits, and discussed our choice of parameters and their influenceon the results. In particular, our disc temperature is colder thanwhat observations suggest (e.g. Williams & Best 2014). Changing its value would affect gap formation (only moderately√ for ahigh-mass planet), as well as fragmentation (Vrel ∝ cs ∝ T ). In
the mid-plane, St is independent of T, leaving dust dynamics unchanged. In this Letter, we focus on the Vfrag = 15 m s−1 case, forwhich we see a second apparent gap.
Fig. 1 shows the gas and dust volume densities ρ g and ρ d , aswell as the grain size s in the disc mid-plane at the end of thesimulation. The planet at 40 au has opened a gap in the gas phaseand much of the disc interior to the gap has been accreted on toss low-density area.
relIn the dusts phase,ss thetheSPHstar, leaving a central
gap is more prominent, with sharper edges. Some material remains
in the inner disc and grains have accumulated at the dust trap located at the gap outer edge. These properties are commonly seen
in studiesJ of planet gaps in gas+dust discs (see e.g. Paardekooper& Mellema 2004; Fouchet et al. 2010). An 5
unexpected feature isfrag ring in the dust disc at r ∼ 90 au visible
the second high-density
in the density map (centre panel of Fig. 1) and in the radial profile of the mid-plane volume density (Fig. 2). Large grains (mmto cm in size) are abundant in the high-density regions (right-handpanel of Fig. 1 and bottom panel of Fig. 2), where relative velocities between grains are lowest and grain growth is more efficient.
The dust population is best characterized in Fig. 3, showing the
distribution of grains in the meridian plane and their radial size

ALMA images of discs: are all gaps carved by planets?
J.-F. Gonzalez et al. (2015)

We study the evolution of gas and dust in protoplanetary discs using
our 3D, two-phase, smoothed particle hydrodynamics (SPH) code,
described in Barrière-Fouchet et al. (2005). We self-consistently
treat the aerodynamical drag between gas and dust, SPH
includingthe backreaction of dust on gas. We follow the standard SPH
frag
implementation of artificial viscosity (Monaghan 1989) and use
αSPH = 0.1 and βSPH = 0.5, corresponding to a uniform Shakura &o parameter*∼ 10−2 (see Fouchet et al. 2007,
Sunyaev (1973)
αSS o–1for a discussion). Arena & Gonzalez (2013) showed that the expected properties of turbulence are naturally reproduced by the
SPH formalism. Grain growth is treated as described in Laibe et al.
(2008), via the prescription developed by Stepinski & Valageas(1997) based√ on the determination of the relative velocity of grains,
Vrel ∝ cs αSS St/(1 + St), as a function of local properties of thegas phase (e.g. its sound speed cs ) and grain size. We implement
grain fragmentation by defining a velocity threshold, Vfrag , to which
we compare Vrel . Collisions occurring at speeds below the thresh-

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Ø HL Tauで見つかったようなギャップは全て惑星が作ったのか？dust pile-upを代案として
シミュレート。
Ø ダストの成長はストークス数とg/d比で決まる。HL Tauでは数十AUでSt~1となり、そこではガスとdecoupleした粒子が最も速くmigrateする。g/d~1となると、空気力学
2 M E T H O D S A N D R E S U LT S
(aerodynamics)的にダストの慣性がガスに影響を与える。
2.1 Hydrodynamic simulations
Ø 3D, two-phase (gas & dust), SPH, α =0.1, β =0.5, α =0.01, V ∝c √(α St)/(1+St)
Ø VrelとV を比べて、小さいとgrain growth、大きいとsha0ering (サイズが小さくなる)。
Ø Mdisc=0.01 M , M =1M , g/d=0.01, 40 AUに5 M の惑星, 面密度は19.67 kg m–2で一定、温度は15(r/4 AU) K、ダスト粒子は初期に10 μm, V =15 m/s, 10 yrまで。
Ø ギャップの外縁60 AUにトラップされるのは普通。加えて、90 AUにも溜まっている。
Ø 60 AUのダストは成長しきっていて、90 AUのダストは成長中で内向きに移動している。

ALMA images
discs:
are all
gaps carved by planets?
L38 ofJ.-F.
Gonzalez
et al.
J.-F. Gonzalez et al. (2015)の続き
Ø ALMAで擬似観測。d=140 AU, θ=0.1”, face-on。
Ø 20 AU以内のinner disk, 40 AUのギャップ, 60
AUのリング, 75 AUの狭いギャップ, 90 AUのリ
ングは全て観測できる。

Figure 4. Simulated ALMA observations of a disc viewed nearly face-on at
d = 140 pc and δ = −23◦ for an integration time of 1 h and angular resolutionof 0.1 arcsec. From left to right: λ = 350 µm, 850 µm, and 1.3 mm. The
scale on each image is in arcsec, with the beam size represented at its
bottom-left corner, and the colour bar gives the flux in mJy beam−1 . The
planet’s position is marked with a cross.

Ø 2つのリングの原因が違うということが大事。
Ø 惑星が軌道の外に圧力の極大を作り、ダストが集まり成長したのが内側のリング。
Ø VfragよりもV
relが小さいことで成長し、St=1に近いところに溜まったのが外側のリング。
Figure 2. Azimuthally
averaged radial profiles of the gas and dust volume
densities (top), dust-to-gas ratio (middle), and mean grain size (bottom) inthe mid-plane after 100 000 yr. The vertical dashed line shows the planet
orbital radius, the vertical dash–dotted lines show the locations of pressure
maxima (see Fig. 6) and the horizontal dotted line shows the initial dust-togas ratio.

Ø Vfragが大きいと外側のリングはより内側にでき、内側のリングだけになる。
Ø Vfragが小さいとより外側になるが、時間がかかるのでできない。
Ø Vfragはシリケートや氷粒子では1-10 m/s、porous aggregatesだと数10 m/s。
Ø 外側のギャップも観測量としては軽い惑星と変わらないので2つ惑星があった場合も調べていて(Cuello in prepara)on)、惑星2つの場合は外側のギャップが今回のものよりダ
Figure 5. Radial cut of the brightness profile in the ALMA synthetic images
ストでのみ深い(惑星を軽くするとガスで浅くなる)ので区別できるかもしれない。
for λ = 350 µm, 850 µm, and 1.3 mm. The vertical dashed line shows theplanet orbital radius and the vertical dash–dotted lines show the locations ofpressure maxima (see Fig. 6).

then passed to the CASA simulator for ALMA to obtain synthetic
images for a given observing configuration (wavelength, angularresolution, integration time). For more details about this procedure,see Gonzalez et al. (2012).

sumoto 2012, Bate et al. 2014) show thas performing simulations in which the sink particles are
tween gas accretion onto the protosta
significantly more massive than individual stars. However,Possible smoking-gun
evidenceforini)almass
segrega)onin
re-virialized
post-gasaccording to the high resolution hydrodynamic calculations
driven outflow
results in star formatio≈ 30% in the neighborhood of proto-st
Machida & Maatsumoto (2012) and Bate et al. (2014),
expulsion globularby
clusters
sub-pc scale. This would imply the overal
and direct observation (Lada & Lada 2003) only small fracHosein Haghi et al.tion(2015)of gas convert into stars.
clump (sub-pc to pc scale) would be les
ing its innermost regions. Such proto-sta
Based on the hydrodynamical simulations without feedare the only available ones that can track
back, Kruijssen et al. (2012) also claimed that the SFE isØ 銀画面のほとんどの星が埋もれた星団でできたかもしれない。
and outflow (driven by radiation and mag
higher in the regions where stars have formed, and henceØ 高密度領域ではIMFは観測的にも不変↔︎M
Fのスロープαと中心集中度c=log(r)に massiv
t/rcmoredividual proto-stars. Other
the long-term evolution of star clusters is only weakly afthat intend to form star clusters, either
相関 (De Marchirela)on)。
fectedby gas removal. They argued that this is because
or include only partial feedback. The SFof the low gas fractions within the pre-cluster cloud core,lations can never be reliable. The Dalecaused by the accretion of gas The
on to
the ofstars.
In this case 5eﬀect
gas expulsion
Ø ガスの動きは計算せず球対称重力ポテンシャルとビリアル平衡として、時間変化させ
lations that go up to globular cluster m
it however needs to be explained how the outer parts of theTable 1. The initial and the final parametersof theN-bodyruns after
t=100 Myr
evolution.
Column1 gives out theparticle representing a sub-cluster itself,
gas cloudareexpelled
without
heatingandpushing
る。ガス質量はτ
=0.6Myrからτ
で減少(gasexpulsion)。
the model name, in which the first twoDdigit numbers after ‘rh’ denote
initial half-mass radius in parsec,
GEthe
ticles themselves are of pc scale. Of courinner
gas ofcontentsof theeﬃciency
pre-cluster
core.
It also is
the last two digit numbers denote thepercentage
the star formation(sfe), andcloudthe degreeof mass5
Ø星団メンバーのN体はNBODY6を使う。銀河中心から8.5
kpcを220
segregation is written after the letter’S’. The initial
mass of allmodels is M = 10HIIM⊙ regions. Column 2ifgives
all overkm/sの円軌道。
the clump (the SFE can, in pra
diﬃcult
to explain
ultracompact
gasthecan siminitial 3D half-mass radius. Columns 3, 4 and 5 contain the adopted star formation eﬃciency, gas expulsion
the /router regions), gas expulsion should
ply beaccreted
ontostars
despitetheirluminosity.
In pcでrØ 潮汐外場は天の川銀河のもの。Half
mass半径は1,5,7,10
timescaleand initial degree of segregation,
respectively.
Column6 is thethe initial
fillingfactor, wherert0 =54.3 pc
h tが>0.1, <0.1を試す。
for all models. Note that the gas massisnotincludedinthecalculation
of
the
tidal
radius.
The
final
slope
of
the
impact on the cluster. This is the init
the short-lived
UCHIIm phase, the gas is heated significantly
5 Mmass
0.85 inside the tidal
radiusis presented
massfunctionafter 100 Myr ofM=10
evolution in the
0.50 ! M ⊙ ! Kroupa
Ø
N=170000,
canonical
IMF(α=1.3と2.3),
T=0-100
4, rangeothat weMyrplausibly adopt in our computa(upbeing
to 10
K following
or more)in theimmediatesurroundingin column 7, the initial canonical value
2.3. The
columns
give thefinal mass, the3D half-mass of theradiusthe simulated star clusters
aftersegrega)onは、S=1でfull100
Myr of evolution,
the finalconcentrationc = segrega)on
logto), as well of
We assume that the gas expulsion
t /rcoutflow
stars(within0.05 pc),
whichthen leadsthe
10 (ranØ ofprimordialmass(重いほど中心近くにある)as the two-body relaxation time of the star clusters after 100 Myr of evolution.
bedded phase at a certain time, τD , w
residual gas on larger scales.
f
0.6 Myr and corresponds to the UCHI
in α2
theMr<rt rsimulations
by
Model
rh0
ε
τGE However,S(rh /rt )0
c
Trh
hf
[pc]
[Myr]
[0.50 − 0.85]M⊙
[M⊙ ]
[pc]
[Gyr]
time, it is assumed that the mass of gDale, Ercolano, & Bonnell (2012) including the influ(1)
(2)
(3)
(4)
(5) photoionizing(6)(7)(10) stars(11)
nentially on the gas-expulsion characteri
enceofradiation (8)from(9)O-type
it isover which the natal gas is removed fromrh10-S0-sfe25
10
0.25 shown
1
0.0
36741
22.3
1.38 and13.14 agrees
that 0.18
the global2.16
SFE is between
20-30%
rh1-S0.9-sfe33
1
0.33
0.1
0.9
0.02
2.27
67392
4.4
2.40
1.81
time-varying total mass in gas is given b
is0.02
observed instar-forming
and embedded
rh1-S0.9-sfe25
10.25 with
0.1 what
0.9
2.21
65016 regions
5.2
2.10
1.50
rh3-S0.9-sfe25
30.25 clusters
0.1
0.9 which0.05 suggests
1.85that the
34930
12.9
1.55
4.40
star formation eﬃciency
Mgas (t) = Mgas (0) e−(t−τD )/τGE .
rh5-S0.9-sfe25
5
0.25
1
0.9
0.09
1.96
54274
18.6
1.61
9.36
can
at
most
be
ε
≈
30%
(see,
e.g.,
Lada
&
Lada
2003),
rh7-S0.9-sfe25
7
0.25
1
0.9
0.13
1.45
41716
21.6
1.34
11.02
rh10-S0.9-sfe33
10
0.33 although1
0.9 on0.18
22.6
1.195
the scale 1.08
of whole31486
molecular
clouds11.81
the star
The impact of the early gas expulsi
rh10-S0.9-sfe25
10
0.25
1
0.9
0.18
0.81
20060
21.5
1.178
10.08
formationeﬃciency
is2.28at most66038
a few(Murray
namical evolution of star clusters has b
rh10-S0.9-NoGE
10
1.0
–
0.9
0.18
16.2 percent
1.70
8.84
rh1-S0-NoGE
1
1.0
–
0.0
0.02
2.29
73827
1.9
2.65
10.10
2011). High-resolution mapping in star-forming regions has
Baumgardt & Kroupa (2007) using a lar
rh3-S0-NoGE
3
1.0
–
0.0
0.05
2.27
74379
4.3
2.44
12.40
uncoveredthat
outflows,
even
from
low-mass
proto-stars,
tegration. They showed that the survival
rh5-S0-NoGE
5
1.0
–
0.0
0.09
2.30
72668
7.0
2.42
8.01
rh10-S0-NoGE
10
1.0– very0.0 powerful
0.18
2.22
71039 clearing
12.7
1.84
are
in destroying
andcloud7.50
material
erties of star clusters are strongly influe
rh20-S0-NoGE
20
1.0
–
0.0
0.36
2.26
55029
22.8
1.67
6.30
away
from
the
regionsub-psof the expulsion of the residual-gas that
rh30-S0-NoGE
30
1.0
–
0.0
0.54 star-forming
2.34
10864 on35.0
0.86 scales
3.40 (e.g.,
Bally et al. 2011, 2012, McGroarty et al. 2004, Frank et al.
depending on the parameter values fo

The eﬀect of gas expulsion

9

Possible smoking-gun evidence for ini)al mass segrega)on in re-virialized post-gas
expulsion globular clusters
Hosein Haghi et al. (2015)の続き
Ø 初めにmass segrega)onがないとMFはほぼ時間変化しない。
Ø あると、)dally underﬁlling (rh/rt<0.1)だとIMRがα=1.3の範囲(m<0.5 Mo)のみ時間変化。
Ø mass seg.とtidally ﬁllingだと0.5<m0.85 Moでもベキが浅くなる。
Ø mas seg.とgas expulsionによって、cが大きいほどαが急という観測の傾向を再現。
Ø αの振動は外側の速度分散と関連。外の軽い星は中に、中の重い星は外に(revirialize)。
The eﬀect of gas expulsion 9Ø 本当は銀河系も時間変化するので形成時期で潮汐場は異なる。
10

Haghi et al.

2.5

Slope α [ 0.3 MO• < m < 0.8 MO• ]Figure 4. The same as Fig. 3 but for a tidally-filling cluster with rh0 =10 pc. The slope of the
2 MF in the mass range m = 0.50M⊙
and m = 0.85M⊙ significantly changes away from the canonical value of 2.30. The MF shows a nonlinear trend for m ! 0.30M⊙ , below
which it becomes much shallower than the canonical IMF. An explanation for this specific form1.5is represented in Sec. 5
10.5

similar mass in the local
universe (e.g., Franx et al. 2008;
0Williams et al. 2010; Mosleh et al. 2011; Law et al. 2012;
Our simulations show that only in the case where a cluster-0.5
Mosleh, Williams & Franx2013). I.e., the fluctuation in theends up with a very large rh /rt , will the low mass stars leave
pre-MW potential grew
-1 with time as it began assembling
the cluster. Note that all clusters in Table 1 which have aduring the first few Gyrand it is only during the past 8-10
-1.5
significantly altered MF due to gas expulsion end up with
Gyr that the MW finally settled down in the form we know
final 3D half-mass radii of about ≃21 pc. This is compatible-2
it today. So, when the GCs
were relatively
young,
their envi0.5
1
1.5
2
2.5
Figure 5. The time evolution
thetypicalMF-slope,
for mainse- globular clusters such as
withofthe
size αof2 ,remotehaloronment would have been much more dramatic.
Some,
but
Concentration
c=log10(rt/rc)
quence stars within the tidal
radius in
the mass
range14,
0.50M! are still larger than the
Palomar
4 andPalomarbut ⊙they
not all, of these GCs would have formed in a strong loM ! 0.85M⊙ . The initialof stars ininthe
all MPPmodelsis with a median half-mass
sizenumberof the clusters
plot
cal tidal field such that they lost a significant fraction of
5
N=170000 (corresponding to
M = around
10 M⊙ )<
withKroupaFigure 7. The slope α of the MF at the end of the N-body inradiusrh an>=input
8.9 ±5.9 pc.
stellar population
asrh0
a =10
resultexpulsion
their
Figure 4. The same as Fig. 3 buttheirfor a tidally-filling
cluster with
pc. ofThetheslopeof the MF ofin themass range m = 0.50M⊙
canonical IMF with the slope of α = 2.3. The initial half-mass
4.4

Implications for the very young MW→赤い線は観測のeye-ball ﬁt。■はgas expulsionなし。
●はmas seg.なし。緑▲は両方なし。○が今回。

on
ti
of
m
T
se
co
(2

(t
0.
F
a
d
ca
th
se

The Early ALMA View of the FU Ori Outburst System
A.S. Hales et al. (2015)
Ø 一部の若い星は通常の1000倍の質量降着率を示すことがあり、100年で0.01 Moを中心星へ降着させる(FU Orionis or FU Ori objects)。
Ø 考えられる理由は熱的不安定、重力+磁気回転不安定、disk fragmenta)on、惑星伴星。
Ø 1-10 AUとそれより外の円盤やエンベロープの関係を調べる必要がある。
Ø FU Oriは450 pcにあるバイナリー(FU OriとS)で、バイナリーのFU Oriはそれなりの数がある。
Ø FU Ori, ALMA Cycle 0, Δv~0.1 km/s, 12CO (3-2), HCN (4-3), HCO+ (4-3), θ=0.63”x0.51”
Ø 連続波は北の星にピークがあるが、singleよりdouble Gaussianでよくフィットできた。
Ø 12COは北と南を覆い、HCO+は南にピークがあり、HCNは北にピークがある。
Ø 12COで最も速度勾配が大きい東西にPV図を切るとケプラー回転。

⬅︎連続波(左上)、12CO(右上)HCO+(左下)、HCN(右下)The Early ALMA View of the FU Ori Outburst System
A.S. Hales et al. (2015)の続き
Ø 連続波は南北とも円盤を持つことを示し、半径は45 AU以下。円盤はバイナリー間隔
225 AUの0.2-0.5倍にtruncateされるという予想と合う(Artymowicz & Lubow 1996)。
Ø 直径で3000 AUを超えるエンベロープ成分は検出されなかった。
Ø 円盤のダスト質量は北で2x10–4 Mo, 南で8x10–5 Moだが、光学的に厚い可能性あり。
Ø 中心星質量は過去の研究から北が0.3 Mo, 南が1.2 Moなので、同じ温度を仮定すると重い星の方に軽い円盤があることになる。
Ø 高密度トレーサーのHCO+が南のみで検出されたので、北の円盤は南の円盤より低質量であると示唆される。
Ø 北のみで検出されたHCNは円盤の上空でできるので、北の円盤の方が高温と示唆される。àHCO+の検出と話は合う。

The Evolu)on of Inner Disk Gas in Transi)on Disks
K. Hoadley et al. (2015)
Ø 惑星ができるr<10 AUは光学的に厚いので調べづらいので、遷移円盤の穴に残ったガスで調べる。
Ø H2は他分子の104倍存在し、5000 Kまで生き残り、UVからもselfshieldされる。しかし双極子がないので、一番強いFUV (912-1700Å)を見る。星からのLyαで励起される。
Ø Hubble Space TelescopeのCosmic Origins SpectrographとSpace Telescope Imaging
Spectrograph (TW Hyaのみ)で撮られたFUV H2データを使う。普通の円盤と遷移円盤。
Ø 速度分解能はCOSでΔv=17 km/s, SITSでΔv=12 km/s。
The Evolution of Gas in Transition Disks

6

Ø モデルの仮定: ケプラー回転、円盤表面の浅く暖かい層、Lyα吸収はLTE。
Hoadley et al. 2015
Table 1
Ø 面密度はexponen)al tail付きの標準的なモデル。r/z=const.なので2D？
Stellar Parameters
Table 3
Parametric Values Explored in Modeling Framework

Parameter

Valuesz/r(2, 3, 4, 5, 6, 7) × Hp

Table 4
Minimum χ2 Statistics for Each Progression Fit

Units
Target[0,1]Progression[0,2][v′ ,J ′ ] Target
[1,4]
[1,7]

Spect.
Type

M⋆
(M⊙ )

d
(pc)

Av

id
(◦ )

Age
(Myr)

vsini
(km s−1 )ref.aAA Tau (2011)
5.37
6.48
5.52
1.25
AA Tau (2013)
1.78
5.29
4.24
1.62
AA TauK7
0.8
140 0.5 75
6.4 ± 0.2
11.4
2,8,9,11,12,15,17,25
BP Tau
2.82
51.75
5.28
2.97
BP TauK7
0.73
140 0.5 30
5.9 ± 0.3
7.8
4,8,9,12,15,17,18,30
q-1.0, -0.5, -0.25, -0.1, -0.05,
CS Cha
4.56
5.14
4.19
2.62
0.0, +0.05, +0.1, +0.25, +0.5
DF Tau A
2.69
13.30
7.21
7.37
CS ChaK6
1.05
160 0.8 60
6.4 ± 0.1...
5,6,15,19,22
DM Tau
6.12
19.55
7.68
37.95 A
DF
Tau
M2
0.19
140 0.6 85
6.3 ± 0.5
16.1
9,11,12,16,17,18
T1AU
500, 1000, 1500, 2000, 2500,
K
GM Aur
3.84
6.72
1.47
1.74
DM 35.81
Tau
M1.5
0.5
140 0.0 35
6.6 ± 0.2
10.0
3,11,12,15,17,25
3000, 3500, 4000, 4500, 5000
HN Tau A
41.71
63.16
13.52
LkCa15111.03103.3014.14
GM151.65
Aur
K5.5
1.20
140 0.1 55
6.9 ± 0.2
12.4
3,8,9,11,12,15,17,25
Hoadley et al. 2015
rc
0.1, 0.5, 1.0, 3.0, 5.0, 7.5, 10.0, 20.0
AU
RECX 11
2.40
9.48
1.09
0.93
HN TauA
K5
0.85
140 0.5 40
1.9 ± 0.9
52.8
7,9,11,18,23
RECX 15
42.45
90.01
13.98
63.32
LkCa15K3
0.85
140 0.6 49
6.4 ± 0.3
12.5
3,8,10,12,15,17,18
SU Aur
25.73
39.31
13.24
21.07
M H2
5×10−10 , 10−10 , 5×10−11 ,M⊙
Table 3
Table 4 2.64
−12 , 10−12 ,TW Hya3.29
3.63
2.1511
RECX
K4
0.80
97
0.0 704.0 ± 1.5...
14,15,20,21
10−11 , 5×10
rametric Values Explored in Modeling
Framework
Minimum χ2 Statistics
for Each
Progression
Fit
−13
−13
−14
−14
UX
Tau
A
104.69124.2313.14
123.165×10, 10, 5×10, 10
RECX 15M2
0.40
97
0.0 606.0 ± 1.0...
15,20,21,32
V4046 Sgr
12.82
13.09
5.93
2.86
SU Aur
G1
2.30
140 0.9 62
2.5 ± 0.9
65.0
1,4,9,11,18
Values were Values
chosen to reproduce theUnits
desired H2 features
meter
Progression [v′ 2,J ′ ]All model-to-data reduced-χ statistics for simultaneous(Herczeg et al. 2004; France et al. 2012b). The only parameters Target
TW[vHya
K6
0.60
54
0.0
4
10.0 ± 6.0
6.0
3,13,16,24,28,29,31
[0,1]
[0,2]
[1,4]
[1,7]
′
′
emissionline
fitting,transitioning
fromexcited
state
,J
].
r
(2,
3,
4,
5,
6,
7)
×
H
without aforementioned constraints
were z/r, γ, q and rc because
p
2
Tau
K2
1.30
140 0.2 35
6.1 ± 0.3
25.4
3,8,11,17,19
All χ5.37
statistics 6.48
are calculated
vobs =UX[-250,
250] A
literature estimates of these values were not known. γ and AA
q Tau (2011)
5.52between
1.25
2km/s.1.78
The largest5.29
source of errors
come
V4046
Sgr
K5
1.75
83
0.0 346.9 ± 0.1
14.2(+13.7) 25,27,28,29
0.0,
0.5, 0.75,
1.5, 1.75, 1.99
were0.25,
constrainedby1.25,the power-law
functionality role they play AA
in Tau (2013)4.24 in the
1.62χ statistics
the linear estimationof the
FUV backgroundcontinuum
the models, and rc was estimated around <r[1,7] > calculated by BP Tau from 2.82
51.75
5.28
2.97
a
attempt-1.0,-0.5,-0.25, -0.1, -0.05,
Akeson et al. (2002); (2) Andrews & Williams (2007); (3) Andrews et al. (2011);
Franceet al.
(2012b).
CS Cha beneath
4.56the emission
5.14line. Because
4.19 the models
2.62 do not (1)the backgroundbeneath each emis0.0, +0.05, +0.1, +0.25, +0.5
DF Tau A to find
2.69
13.30 continuum
7.21 levels
7.37(4)
Bouvier(1990); (5) Espaillat et al. (2007a); (6) Espaillat et al. (2011); (7) France et al.
of the FUV
the appropriate energy for photo-excitation:
DM Tau sion line,
6.12 extraction
19.55
7.68continuum
37.95 had to be done
(2012b);with lower signal-to-noise
more uncer-(8) Furlan et al. (2011); (9) Gullbring et al. (1998); (10) Hartmann et al. (1987);
500, 1000, 1500, 2000, 2500,
K
GM Aur manually.
3.84 Targets6.72
1.47 2 1.74 haveU
3
in the continuum
the χ 35.81
statistics(11)becomelarge
λLyα3000, 3500, 4000, 4500,
5000g[B:v′ ,J ′ ]Hartmann& Stauﬀer (1989); (12) Hartmann et al. (1998); (13) Herczeg & Hillenbrand
HN Tau A tainty
41.71
63.16 flux, so
13.52
errors in thecontinuum14.14
dominatethe fitting. Only the
σH2 =
Alu ,
(7)LkCa15 as the111.03103.30151.65(2008); (14) Ingleby et al. (2011); (15) Ingleby et al. (2013); (16) Johns-Krull & Valenti (2001);
g[X:v,J] AU
show
fits for all
targets (with χ2 <0.1, 0.5, 1.0, 3.0, 5.0, 7.5, 8πc10.0, 20.0
RECX 11 [1,4] progression
2.40
9.48 decent 1.09
0.93
(17)
Kenyon et al. (1994); (18) Kraus & Hillenbrand (2009); (19) Lawson et al. (1996); (20)
we focus on
the relative
results 63.32
of the [1,4]progression
RECX 15 15), so
42.45
90.01
13.98
−10 ,is10the−10 , rest
−11 ,the remainderof the21.07
Discussionsection. et al. (2001); (21) Lawson et al. (2004); (22) Luhman (2004); (23) McJunkin et al.
where5×10
λLyαwavelengthof the stellarSU Aur emission
Lawson25.73lines for39.31
13.24
5×10frame
M⊙
H2
−11
−12
−12
TW Hya
2.64
3.29
3.63
2.15
Lyα line
needed10 profile, 5×10, 10 to ,excite the warm H2 in ground(2013); (24) Pontoppidan et al. (2008); (25) Quast et al. (2000); (26) Ricci et al. (2010);−13 , 10−13 , 5×10−14 , 10−14
104.69124.2313.14 123.16
5×10
state
[X
: v, J] up to energy level [B : v ′ , J ′ ], and Alu UX
is Tau A(27) Rodriguez et al. (2010); (28) Rosenfeld et al. (2012); (29) Rosenfeld et al. (2013); (30)V4046 Sgr
13.09
5.93
2.86
4 12.82
COS − DM Tau
γ0.0, 0.25, 0.5, 0.75, 1.25, 1.5, 1.75, 1.99

probability
that theH2 desiredin population[X : v, J] will bes were thechosento reproduce
H2 features

Synthesized H disk model

Simon et al. (2000); (31) Webb et al. (1999); (32) Woitke et al. (2013).

of the observed FUV H2 fluorescent emission originates.

from further out in the disk (the peak of the radiation located at
rpeak = 2.50 AU) than the 2011 radiation distribution (rpeak = 0.75
AU), a consequence of the inner disk “shadowing” produced by the
extra absorber on the AA Tau sightline (Bouvier et al. 2013).

progression
exploredof
in thisstudy.
focus
analysis DisksThe Evolu)on
Inner
DiskWeGas
inour
Transi)onon the [1,4] radiation distributions for all targets in our
K. Hoadleyal.and(2015)の続き
surveyto define etinnerouter radial disk boundaries,
Tau epochs. The 2013 emission lines are narrower withwhich describe where the bulk (90%) of the emitting H2slightly larger peak fluxes than the same H2 emission
atmosphere resides. We define the 90% emitting regionlines observed in 2011 (Schneider et al. 2015). This sugØ最も大丈夫そうな[v’,J’]=[1,4]へのベストフィットから90%のフラックスが入るrを決める。
as follows:inとroutinnermost
gests that lessflux is contributed fromthe10
Hoadley et al. 2015
! F (H ,r)2
disk. The modeling results for the [1,4] progression are≤ 0.95 for r > rin
Ftot (H2 )shown in Figure 8. Each AA Tau epoch was modeled
FH2 ,obs = F (H2 ,r)
≤
0.95
for
r
<
r
out
independently, and the models reproduce the same rest
Ftot (H2 )wavelength emission lines. Figure 8 also shows the radi12
Hoadley et al. 2015
We use rin and rout to evaluate the evolutionary beation distributions of [1,4] fluorescence for each epoch inØ 例えば、AA
Tauは2013年には2011年より内側の寄与
1.0haviorof the H2 radiation.
Figure7 presents a schematic
the AA Tau disk. The 2011 emission includes a large conof how
the inner and outer radial
0.8 boundaries encapsulateが減っている(過去の研究でも報告されている)。
tribution from material inside 1 AU (rin,2011 = 0.08±0.01
90% of the total H2 flux produced in the disk atmosphere.
AU; rpeak,2011 = 0.75 AU; rout,2011 = 4.17±2.04 AU),0.6
àextrapotential
absorber,pseudo-transi)ondisk
We analyze
evolutionary
characteristics ofthewhile the 2013 [1,4] emission “appears” to have shifted
molecular disk atmosphere by 0.4comparing the FUV H2
outward in the disk (rin,2013 = 0.15±0.02 AU; rpeak,2013
radiation distributions to other0.2dust and molecular disk
=radial
2.50
AU;whichrout,2013= 7.59±2.75
AU).of OurmodelsFigure 7. The inner and outerboundaries
defineFigure8. We present a comparison
the [1,4] progression
ob- indiØ Primordialからtransi)onに進化すると放射が外へ移動
where 90% of the total radiation is arising from the disk. The blackobservables.
served with HST -COS for AA Tau during the 2011 and 2013
catethatof theradiusof Thedetectable
Hlabel0.0line represents the normalized radial
distribution
the [1,4]inner
proepochs.
On the left:
column under the 2011
represent the
2 fluorescence
Primordial
gressionto the total amount of flux produced by the [1,4]
3 observed fluorescent emission line profiles cascading from the [1,4]
−200
−100
0
100 emission200
for SU Aur. The green
verticalthelines show
the radial
excited state, withoverlaid
best-fit modeled
emission lines
in blue.
している。
Velocity (km sprogression
)
from
[1,4]
progression
has
moved
outwardradiallyinboundaries that encapsulate 90% of the total emission. For r , we
The column to the right, labeled 2013, shows the observed [1,4]4.3. Case Study: Model Figure
Robustness
using
AA
Tau
start
at
the
outermostradiusandintegrate
inwardtosmallerradii
fluorescentemissionlines,
withmodeledemissionlines
overlaid
inFigure11. Wepresent as
the relation
betweenthe estimated
r
the
disk
the
“extraabsorber”
moved into our field of
10. A comparison of observed [1,4] progression line prothe disk
until4095%
of the total
progression
flux isfrom
accounted
orange.radiation
The 2013 observed emission line profiles appear narrower
and r [1,4]quantities,
determined
the best-fit modeled
files of targets withinclination angles between
30 and. TheØWeダスト進化の指標13-31μmのλF冪指数が大きいfor.
r priis definedby starting
innermostradiusthan theirdistributions
for atall thetargets.
The innerradial boundary(r 2011) is counterparts (Schneider et al. 2015). On the right:
two broadestline profiles, BP framework
Tau and HN
Tau Likewise,A, are intheλexplore how robust ourmodeling
is
at
view
in
the
AA
Tau
disk.
Schneider
et al. (2015), using
definedinTau
our andmodelstheinner
progression
out until in the thedefined as theradius offlux
H radiation
disk comparison
that marks of the total [1,4] progression flux radiating from themordial phase. The two narrowest line profiles,DM
UX and integratingwhereis atleast 90%for.
of thetotalradiationin Tauthe in the 2011 and 2013 observations. The 2013 modTau A, are in the transition phase.
95% of the total emission flux
accounted
The
resultingan-is accounted
disk offorAA
anindependent
modeling
to estimate
the raidentifying
where the fluorescing
H2 outresidesin PPDs.
disk. Likewise,
thediskouterboundary(r) is defined
nulus between Wer and r outerrepresents
thering oftheradial
majorityels predict
thattechniquethe observed H fluorescenceemission originates(transisionに近い)とrも大きい。
as the outermost
radiusof H radiation that encompasses
90% ofinもrof the observed FUV H fluorescentemissionoriginates.
from further
out in the disk (the peak of the radiation located at
the total
amountorigins
of radiation accounted
for in fluorescence
the rinner disk.
The
=
2.50
AU)
than
the
2011
radiation
distribution
(r
=
0.75
dial
of
H
in
the
AA
Tau
disk,
come
compare two epochs of HSTthe-COS
data
on
AA
Tau
(2011
Transitional
2modeled r and r ,blue diamonds with error bars represent each
degree of Lyα penetration into the disk by clearing
a consequenceof the inner disk “shadowing” produced by the
progression
explored
inand
thistheH
study.
We focus
our analysis
ßダストが蒸発したり沈殿したりするとwarm
blackができ
dashedline represents
a linear fit toAU),
the data.
The
H
from
the
inner
disk
atmosphere;ortheevolution
of
2
extra
absorber
on
the
AA
Tau
sightline(Bouvieretal.
2013).
totheranka correlation
similarconclusion:
observed 2013 H2 emission
and 2013), where the 2013 theobservations
occuron during
a distributions
Spearmancoeﬃcient])
theThis[1,4]evolving
radiationforall targets
in([ρ,
ourn] = [0.70, 5.5 × 10 the
disk surface temperature distribution.
two radial quantities suggest a strong increasing trend
surveyinner between
and outerdiskboundaries,
radiationto
structureis also observablein the tolinedefine
profilesbetweenthem,radial
indicating
that
theAU
whole emitting
regionis moving compared
Tau epochs.
The 2013 emissiontolines2011.
are narrower
with“dimming”
event
from
X-ray
near-IR
wavelengths.
within∼1
isreducedAdditionる領域も後退する。
outward.
where thebulk (90%) of the emitting Hof the [1,4] progression, as seen in which
Figure describe
10. As theslightly larger peak fluxes than the same H emission
atmosphereresides.
We
definethe90%
emitting
regionPPDs
in
oursurveyevolvefromprimordial
totransition
ally, Schneider et al.
find(Schneider
thatet theouterBased on the duration of disks,
thethedimming,
Bouvier
et al.
lines (2015)
observed in 2011
al. 2015).
This sug-radial
majority of the observed Hasemission
follows:migrates
gests that less flux is contributed from the innermostrHya!
to larger
radii.
1.0 extent of the H
fluorescenceluminosity
doesn’tchange(2013) TW
suggestan
obscuration
at
r
>
8
AU;
this
hypoth2
disk. The modeling results for the [1,4] progression arer≤ 0.95 for r > r
LkCa15
We compare estimates of r and r
to investigate
12
et al. 2015
shown Hoadleyin Figure
8. Each AA Tau epoch was modeled
F
=
r
the
idea
that
the
radialdistributionsoffluorescing
H
V4046
SGR
significantly
betweenepochs,andwhichisreproduce
a result
consistentesis is strengthenedby themigrate
gas-to-dust
ratio (NH /Av )≤ 0.95 for r < rindependently,
the models
the same
rest
0.1
outward in the disks as PPDs evolve. Figure 11r
UX Tau A
wavelength
emission lines. Figure
8 alsoshows the radi1.0
presents a comparisonof r FUV
and r , H
which
shows
the
withintheerrorsonourmodeledrestimatesof theof the absorber
and the evolutionoftheemisSU Aur
out
2
We
use
r
and
r
to
evaluatetheevolutionary
beation distributions of [1,4] fluorescence for each epoch inannulus of H emission extending further out into the
GM AurhavioroftheH
radiation.
Figure7presentsaschematic
disk
as
the
inner
disk
radiusmoves
outward.
A
line
canthe AA Tau disk. The 2011 emission includes a large conAA
Tau epochs.
sion (Schneider
et al. 2015).
We utilize theof line
pro0.8 boundaries
DM Tauhowthetheinnerinnerand outer radial
encapsulate
be fit to represent the relationship between
andtribution from material inside 1 AU (r= 0.08±0.01
CS Cha between AA Tau
90%of the total H flux produced
inThethe disk“extra
atmosphere.
outerHST
radiating-COS
disk radii forour survey
targets:
file changes
FUV
observingabsorber”
obscuresinner
diskAU),H2 flu10
AU; r= 0.75 AU; ther= 4.17±2.04
RECX−15
We analyze potential evolutionary
characteristics of the
0.6
while
the
2013
[1,4]
emission“appears”
tohave shifted
log
(r
(H
))
=
0.79
log
(r
(H
))
+
1.39,
(13)
epochsRECX−11
to determine how those changes relatemolecular
to radial
inFUV
theH 2013
HST-COS
spectrum,makingdisk atmosphereorescenceby comparing theoutwardin thedisk (r FUV= 0.15±0.02
AU; r
HN Tauradiation
to 1other dust and molecular disk
where both log (r (H )) and log(r (Hdistributions)) are=disk
2.50 AU;r=deficit
7.59±2.75 ofAU).inner
Our models
indi-emis0.4
H2 radiation
distributions
in
the
disk.
AA
Tau
appear
as
a
with
a
disk
observables.
DF Tauin units of AU, and the coeﬃcients[1.39 ± 0.22,
cate that the inner radius of detectable H fluorescence
BP Tau0.79 ± 0.21] between
are computed froma χ
minimizationfrom the3 [1,4] progression has moved
outwardin
ThereAAarenoticeable
diﬀerences
theobservedsion
eﬀectively,apseudo-transition
disk.
Ourradiallymodeling
−1
0
1
2
0.2
4.3.
Study:
Tau(χ=0.896) of a linear function between
logCase(r (H)) Model Robustness using AAn Tau
the disk as the “extra absorber” moved into our field ofand logof(r the(H )).2011
The Spearman
correlation
co- robust our
FUV H2 emission
line profiles
andrank
2013
AA
framework
wasis atcapable
theetchange
in emisWe
explorehowmodeling framework
view in of
the identifying
AA Tau disk. Schneideral. (2015), using
Figure 12.
with an observable dust
0.01
0.10
1.00eﬃcient between
10.00 r and r indicates a statistically sig0.0 Comparison of r and r
AA Tau − 2011
1446.12 Å

5

4

4

3

3

2

2

1

1

BP TauHN Tau A
DM TauUX Tau A

10

0.005

Normalized Flux

Flux (10−14 ergs cm−2 s−1Å−1)F(H2,r)/FTOTAL(H2)0−200 −100

0.010

AA Tau − 2011
AA Tau − 2013

15

0

0
100 200−200 −100

1489.57 Å

0

100 200

1489.57 Å

5

4

4

3

3

2

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1

1

0
−200 −10000100 200−200 −100

1504.76 Å

0

100 200

FH2 (r) [10−13 ergs cm−2 s−1 ]0.015

AA Tau − 2013
1446.12 Å

5

10

5

1504.76 Å

r OUT (H2) [AU]

8

0.000

0.1

6

6

4

4

2

1.0

10.0

2

0
−200 −100rDISK (AU)100100 200−200 −100

Velocity (km/s)

0

0

0.1

100 200

1.0
rDISK (AU)

Velocity (km/s)

10.0

The Evolution of Gas in Transition Disks

11

0.030
0.025

0.1

−1

r IN (H2) [AU]
in

◦

0.020
0.015

in

◦

out

out

0.010

in

F(H2,r) / F TOTAL (H2)

2

out

out
2

in

AA TauBP TauDF TauHN TauRECX 11RECX 15

2

peak

in

2

out

−3

0.005
0.000

2

0.1

1.0

10.0

0.025

CS ChaDM Tau
GM AurSU AurUX Tau A
V4046 SGR
LkCa15
TW Hya

peak

0.020
0.015

The Evolution of Gas in Transition Disks

15

0.010

2

2

0.005

cavityinout

H2 ,obs

in,[1,4]2out,[1,4]

in

out

in

2

out

10

10

in

2

in

2

10

out

2

2

2
min

10

10

out

in

1.0

rDISK (AU)outsion line profiles between the 2011 and 2013 AA Tau
observations and found that the bulk of the 2013 AA
Tau [1,4] radiation in the disk originated at larger radii
than the 2011 H2 fluorescence. We expect our models
are therefore capable of distinguishing between H2 fluorescence10
evolution in diﬀering disk types.

in,2011

5. DISCUSSION

peak,2011

2

2

13−31

out,2011

We created 2D radiative transfer models to simulate
observed HST -COS and -STIS FUV H2 emission lines
to understand where the majority of the radiation arises
in PPDs.
We use the best-fit model results to define thein,2013inner and outer radii of warm H2 (rin , rout ) and examine
if and how the molecular distributions change as PPDs
out,2013
evolve. We compare rin and rout to other dust and molecular tracers that help describe the evolutionary state of
the PPDs. Table 1 provides a detailed list of inner2disk
observables for each target, including dust cavity radius
1 inner disk CO radius (rin,CO ). We also look(rcavity ) andat where the theoretical snow lines in the disks exist andhow these radii relate to the H2 disk emission.

5.1. Radiation Distribution of Modeled H2 Fluorescent
Emission
Figure 9 presents the normalized radial distributions ofwarm H2 transitioning from excited state [1,4] for all targets. We modeled 6 primordial disks (AA Tau, BP Tau,DF Tau A, HN Tau A, RECX-11, and RECX-15) and8 transition disks (CS Cha, DM Tau, GM Aur, LkCa
15, SU Aur, TW Hya, UX Tau A, and V4046 Sgr) tocompare the radial distribution of warm H2 in the disk
atmospheres as the dust disk evolves. The H2 radial distributions of the diﬀerent dust disk stages appear qualitatively diﬀerent. The primordial disk population (topplot
in Figure 9) generally starts radiating significantly
peak,2013
in the very inner disk (r ! 0.05 AU), and the radiation
only extends out to a few AU, consistent with the simple estimates of the average H2 emitting radius presented
by France et al. (2012b). The generalized transition disk
radiation behavior (bottom plot) starts further out inthe disk (r ∼ 0.1 AU) and extend significantly further
out into the disk (r ∼ 10 AU). These diﬀerent behaviors
suggest structural changes in any of all of the following:
the spatial distributions of warm H2 in populations [v,J];

2

in

10.0

Figure 9. The normalized modeled radiation field distribution of H2 fluorescence cascading from the [v′ ,J ′ ] = [1,4] energy level for all
targets. Each radiation distribution was calculated from the minimum reduced-χ2 model parameters that best reproduce the observed H2emission lines. The top plot represents radiation distributions for all primordial disk targets, and the bottom plot shows the distributions
for transition disk objects. The two disk evolution types appear to show an evolving H2 FUV radiation field; primordial disks generally
radiating more inward in the disk, with the bulk of the radiation occurring within r ! 1 AU, and transition disk H2 radiation starting atlarger radii (r ∼ 0.1 AU) extending to larger radii (r ∼ 10 AU).

Normalized Flux2

r OUT (H2) [AU]

out

0.1

in

BP TauHN Tau A
DM TauUX Tau A

2

2

10

0.000

F (H2 ,r)Ftot (H2 )F (H2 ,r)Ftot (H2 )r OUT (H2) [AU]in,COr IN (H2) [AU]

2

out

in
out
independent
modeling technique to estimate the raidentifying where the fluorescing
H2 resides
in(Furlan
PPDs.et We
evolution diagnostic
n13−31
al. 2009). an
In thetop plot:nificant correlation between the variables (ρ = 0.70) with
−200
−100
0
100target point
200in
rDISK (AU)triangleerror
each
0.1 Tau disk, comedial originsof H2 fluorescence in the AA
compare two epochs of Each
HSTblue-COS
datawithon
AA bars
Taurepresents(2011(km s )the best-fit linear cora small probability that the sample is randomized (n =our survey. The black dashed Velocity
line representsr (H ) [AU]to a similar conclusion: the observed 2013
H2 emission
and 2013), where the 2013 observations occur during a
−3
−1

IN

2