6

Thebault et al.: Planet signatures in debris discs

THE	
  STAR	  FORMATION	  NEWSLETTER:	
  No.	
  239	  10	
  Nov	  2012	
Regarding the spatial distribution of the steady state disc forthese eccentric planet cases, we see that, even for a massive µ =2 × 10−3 planet, almost no resonant structure is longer visible
in the inner or outer discs, and the Lagrange points, althoughstill present, are much less pronounced, especially in the e p =0.2 case. In addition, the outer disc gets strongly depleted while
the inner disc assumes an asymmetric shape aligned with theplanet’s orbit.

担当：麻生(東大)	
  2013/2/15	

3.2. Collisional system: embedded planet

Fig.2 shows synthetic head-on images, in scattered light, at colliPlanet	
  Signatures	
  in	
  Csteady-state,ollisionally	
  Ac3ve	  sional and dynamical
for different planet
orbits and Debis	
  Discs:	
  sca8ered	
  light	
  images	
  optical depthτ ). For our nominal
2 × 10Steve	
  case, the main
Philippe	
  Thebault,	
  Quen0n	  Kral	
  τ a= nd	  Ertel	
  structures observed in the pure N-body runs, such as for example the Lagrangian points and the 2:1 resonances in the e = 0l 惑星があってcollisionもradiaDonも考慮したときの惑星が観測できる条件	  run, are still clearly visible, but they are much less pronounced.
The other consequence of the coupled effect of collisions and ral 散乱光で見える粒子はradiaDonを受ける←size-­‐dependent←collision	
  diation pressure is that it injects small, high-β grains in regionsthat are devoid of parent bodies. These regions are the external region
beyond the outer edge
130 AU) of thebody(DyCoSS)コード	  Fig. 3. Synthetic midplane luminosity profiles for a disc seenl Dynamics	
  and	
  Collisions	  at	
  (atSteady	
  Sparenttate	
  disc, and, more importantly, the dynamically unstable region (the edge-on, in scattered light, for the standard µ = 2 × 10 andgap) close to the planet’s orbit. The global effect of collisions e = 0 case and for different values of the disc’s average opticalstep1:	
  norb個の惑星の位置Φ
毎にradiaDonを入れずにN-­‐bodyを安定させる	  and radiation pressure is thus to itend
to homogenize the surface depth (i.e., collisional activity). The shown profiles corresponddensity profile, as is clearly illustrated for runs with higher op- to a cut along the X axis when the planet passes at its maximum
qdsのように新しい粒子を入れて、
tical depthを決めて(スケーリングし)、dN∝s(i.e., collisional activity), where the differences with elongation along thisaxis (the planet’s location is marked by a
step2:	
  step1からτ
rθ
the collisionless N-body simulation get much more pronounced. black circle). Profiles are computed assuming a grey scattering
For the ep=robabilityを考えてΔt=t
0 orbit, the accumulation around the Lagrangian function.
The profile for the large parent bodies is directly deradiaDon,	
  collision	
  /nfromorbrived
orb毎に進化を追う	
  points is for example replaced by a horseshoe co-orbital ring for
the PB runs presented in Section 3.1.
higher value of τ (Fig.2h). Conversely, for lower values of τ ,4Thebault et al.: Planet
signatures in debris discs
step3:	  それぞれの位置からΦ
the spatial profile becomes much closer
to that of the parent bod0まで追った粒子の分布を全部足す	
  collisional activities in the disc (as parameterized by its average
0

0

−3

p

−3

p

p

0

ies (Fig.2g).

0

Table 1. Set-up for the nominal Despitecase runs.
denotedof collisions
radiation pressure on the purely dynamical bethis Parameters
blurring effectof collisions,for all the andhead-onby a * are explored as free parameters
haviourof perturbed
images displayed in Fig.2, the planet’s radial
location
remains discs.
theseparentrelatively easy to determine from the positionFor
of theradialgap.body runs (hereafter PBR), we let the system
PLANETevolveuntil aposition
dynamical steady state has been reached,i.e., until
For high-mass low e p cases, even the planet’s
azimutal∗−3
Planet/Star mass
µ = 2 × 10
remainseasy to spot in between the well
Lagrangianthedefined
disc’s spatial
structure does not vary between two passages∗
eccentricity
e p = 0 with the noticeable exception of thepoints,0 case (Fig.2h)
ofhighthe τperturbing
body at the same orbital phase (see TBO12).
semi-major axis
awhereAUsteady production of small collisionp = 75thefragments
For thecase oferases
a planet on a circular orbit, this happens relaPARENT BODY RUN
almost all azimutal structures in the co-orbital
tivelyregion.
quickly, typically less than a hundred planet orbits, for all
Number of test particles
NPB The= 2 ×situation
105
becomes very different if the system is seen
the planet masses considered (see Fig.1a, b and g). For the nomInitial radial extent
30
<
r
<
130
AU
(embedded
planetcase)edge-on. Fig.3 shows the midplane luminosity profile alonga−3
inal µ =While
2 × 10theand e p = 0 case, several pronounced structures45< r cut< 75
AU (inner-disc/outer-planet
radialpassingby the planet’s currentcase)position.
disare visible.gap
Themost prominent one is the gap in the middle ofInitial eccentricity
0tribution≤ e ≤ 0.01
of large parent bodies shows a pronounced
arounddisc,dropping
corresponding
to the chaotic unstable zone surroundInitial surface density
Σthe∝ planet’s
r−1
radial position (Fig.5), with thethe flux
by a
ing the effectplanet’slocation (e.g. Wisdom 1980), where only the
COLLISIONAL RUNS
factor ∼ 5 in the 50-90AU region, the coupledof colli∗−3
Average optical depthτ = 2 and× 10radiation
two stablesions
pressure on small grains
greatly co-orbital
alters this Lagrange points L4 and L5 are populated.
−4
Number of test particles
Npicture.
=2
×
105for a relatively low τ0 = 4This
Even
× 10gap
value,
the lu- to the so-called ”feeding zone” surroundcorrespondsnumSize rangea
sminosity
40scutaround the planet falls to ing
less thethanplanet’sa factor 2.
It Clearly visible are also the 1:2 and 2:1
cut ≤ s ≤ drop
orbit.
is of lessthands30% in the nominal τ0 = 2mean
× 10−3case, andfalls
Size distribution at t = 0 b dN(s)∝ s−3.5
motionresonances(hereafter MMR) at 47 and 119 AU, redown to a few percent for even more collisionallydiscs
spectively.active
The apparent
gaps at the resonances’ location is duea(τ=0.01).
Furthermore,
this
small
flux
variationisnotpeakedscut is the radiation pressure blow-out
size
0
to the fact that resonantFig.
particles
reachhigh but
eccentricitiesand τ0 = 2 × 10−3 and4. Sameas Fig.3,
with e p = 0e andbat body
the planet’s
location but diluted over a rather extended region,
when released from the parent
population
large
radialexcursions,thus
spendingashortfractionof
theirorwhenvarying
the
mass
of
the
perturbing
planet.
so that the planet’s effect on the profile is a wide and poorly debit at the exact radial resonant location. For a lower planet massl embeddedのときは、	
  • 円盤の外やgap内に粒子が入り、面密度が均質化される	  • 半径方向の光度分布は惑星の効果が見えにくくなる	  • 惑星がhigh	
  mass,	
  eccentricなときgapは見えにくくなるが、inner	
  diskがeccentricに見える	  • edge-­‐onのときは”惑星なし”と比べないと惑星があるかわからない	  l 惑星が外にあるときは、	
  • 円盤がtruncateされ、惑星が重いほど内側まで切られる	  • [rout,	
  1.2rout]より外のスロープは惑星の質量にあまり依らない	
  • すごくsteepになるのはcollisionが少ない(τ〜2×10–4)ときだけ	  • 惑星がeccentricだとringもeccentricになる	  l ring本体の非軸対称は、その摂動がringの周期より早ければ惑星は内
側にある	
  

Fig. 10. Inner ring/outer planet: head-on synthetic images in scattered light, at steady state, for four different set ups. The graphson the left hand side show the disc profile when the planet passes at apoastron, and the graphs on the right hand side correspond to periastron passages. The dotted curve shows the planetary orbit. An animated version of these graphs can be found at

MHD	
  Modeling	
  of	
  a	
  Disk-­‐Wind	
  from	  a	
  High-­‐Mass	
  Protobinary:	
  the	
  
case	
  of	
  Orion	  Source	
  I	
  Bhargav	
  Vaidya	
  and	  Ciriaco	
  Goddi	  l VLBI、SiOメーザーでOrion	
  Source	  Iのhigh-­‐mass	
  protostarからのアウトフ
ローを分解	  l モデルはMHD、disc-­‐wind、radiaDonなし	  l 中心星は10Mo、2a0=7	
  AU離れた円軌道binary(本当はa=2-­‐5	
  AUの
eccentricな軌道)	
  l 内側は圧力が高いのでsub-­‐Kepler	  l r=50a0,	
  z=150a0より外はfree	  l 物質は鉛直に運ばれてbow	
  shockを作り、それが上下に去って
collimated	
  jetが現れる	  l ﬂowは最初subsonicで、Lorentz力でsuperfast	
  speedまでなる	  l SiOメーザーの励起条件はn(H2)〜109cm–3,	
  T<2000	
  Kで、シミュレーションでは15	
  AUより外→Fig	
  1でも14	
  AUくらいのdark	
  laneが見える	  l Fig	  1のX-­‐shapeとVΦも再現〜20km/s	
  l β0が大きいとLorentz力が小さくなって遅くなる	  l β0=1、0.35G	
  @	
  r=40	
  AU	
  AU,z=20	
  AUがよさそう	
  l 今後の課題→discの影響、星風、中性粒子による拡散、化学、偏光観測	  MHD models for Orion Source I disc winds

L51

3.2 Comparison between MHD wind models and windproperties observed from SiO masers

3 R E S U LT S A N D D I S C U S S I O N

We propose that the SiO masers in Orion Source I are excited as amoderately rotating disc wind interacts with the ambient medium
At the onset of our reference run, material at the base of the disc isvia slow MHD shocks and transfer of angular momentum. Similarejected in vertical direction along the magnetic
field.
This
material
mechanisms
have been successfully applied to explain the excitation
sources provided strong evidence that a dynamicalinteraction occurred
aboutambient
500 yr agomedium.
between massive YSOs
and BN (Kaufman & Neufeld 1996). Here, we qualitatively
propagates outwards producing bow shocks
in theofSource
H2 OImasers(e.g. Goddi et al. 2011). N-body numerical simulations showed thatsteady-state properties of a single-fluid, rotating MHD
As this shock propagates outside of the domain,
the flow
collimates
the dynamicalinteraction
between a binary ofcompare
two 10 M#thestars(Source I)isanda single 10beyond
M# star (BN) maywind
lead to(launchedsimultanefrom an underlying accretion disc) with the obdue to magnetic forces until a steady equilibrium
attained,ous ejection of both stars and the BN/KL outflow as well as binary
servedphysical
anddynamical properties of the outflow molecularwhich no further
evolutionofmagnetic
field
lines
takes
place.
The
hardening, while preserving the original circumbinary disc aroundFigure 1. Velocity field (in km s−1 ) of the SiO v = 1, 2 and J = 1–0 maser
SourceI(Ballyetal.
2011;
Moeckel&Goddi2012).
Mechanicalcomponent
probedby
SiO masers in Source I.
steady-statevaluesofthedensity,poloidalandazimuthalvelocities
emission in Orion Source I as observed with the VLBA (Matthews et al.
energy conservation (and virial theorem) implies an orbit semimajor
2010). The stellar velocity is ∼6 km s−1 . The bulk of the SiO emission
Threephysicalpropertiesof the wind from the MHD modelling
of
the
wind
are
shown
in
Fig.
3,
along
with
the
magnetic
field
lines
on themodel
estimate forof theSource
kinetic energy
axisFigurea0 = 2–5 au,
is located within four ‘arms’ of an X-shaped pattern, while a ∼14 au thick
2.depending
CartoonI.
band with no SiO
the 7-mm continuum
source (Reidet al. (whiteareconsistent(inemissionred),harbours
the poloidal
velocity
versorsof the arrows),
runaway starstemperatureand the molecular gas flow
(Ballyet al. 2011; with SiO maser observations: temperature/density
2007; Goddi et al. 2011). The continuum emission traces an ionized disc,Goddi et al. 2011). For simplicity, we neglect profiles,
the eccentricity
andmorphology
and velocity field.
and MHDcriticalmagnetic
field lines idewhile the SiO contours
masers probe material
ablated fromthe surfacesurfaces.
of the disc in The assume
that the system is an equal-mass circular binary, comprised
3.1 MHD-driven wind profiles

Downloaded from http://mnrasl.oxfordjournals.org/ by guest on February 12, 2013

a wind.

Figure 2. Cartoon model for Source I.

and magnetic fields to show that radiation fromh masses >30 M# does modify the dynamics ofsources provided strong evidence that a dynamical interaction occurred about 500 yr ago between massive YSOs Source I and BN(e.g. Goddi et al. 2011). N-body numerical simulations showed thatthe dynamical interaction between a binary of two 10 M# stars(Source I) and a single 10 M# star (BN) may lead to simultaneous ejection of both stars and the BN/KL outflow as well as binary
hardening, while preserving the original circumbinary disc aroundSource I (Bally et al. 2011; Moeckel & Goddi 2012). Mechanicalenergy conservation (and virial theorem) implies an orbit semimajor
axis a = 2–5 au, depending on the estimate of the kinetic energyof the runaway stars and the molecular gas flow (Bally et al. 2011;
Goddi et al. 2011). For simplicity, we neglect the eccentricity andassume that the system is an equal-mass circular binary, comprisedof two 10 M# stars orbiting in a circle of radius a = 3.5 au.
A high-resolution study of Hα and [O ]λ6300 line profiles ina pre-main-sequence spectroscopic binary has shown evidence ofa bipolar jet as being launched by the whole binary as opposed to

Downloaded from http://mnrasl.oxfordjournals.org/ bof two 10 M# stars orbiting in a circle of radius a0 = 3.5 au.
A high-resolution study of Hα and [O I]λ6300 line profiles ina pre-main-sequence spectroscopic binary has shown evidence ofa bipolar jet as being launched by the whole binary as opposed toeach star individually (Mundt et al. 2010). Similarly, we assume
that the binary members orbit inside the inner radius of a truncatedcircumbinary disc and that the wind is launched at this radius.
The cartoon figure representing our model assumption is shown in
Fig. 2.
To study the wind launching from Source I, we perform axisymmetric ideal MHD simulations in cylindrical coordinates (r, φ, z)
2 MODEL ASSUMPTIONS
using the PLUTO code (Mignone et al. 2007). We essentially prescribe
The initial conditions adopted to model the disc wind from Oriona hydrostatic corona (density and pressure) threaded with force-free
Source I are reported in Table 1 and are derived from either obsermagnetic field. An adiabatic equation of state, P = γ γ−1 ρ γ , witheld (in km s−1) ofand/or
the SiOv =of 1,
andSource
J = I1–0
maser
vationsmodellingthe 2
system.
is believed
to be
γ = 5/3, relates the pressure with the density in the flow. For thean eccentric, hard binary with a total mass of 2 × 10 M# , based
central gravity, we treat the system of equal-mass binary members as
ource I as observedwith the VLBA (Matthews et al.
on studies of the dynamics of YSOs and molecular gas in Orionan equivalent 5 M# (reduced mass) object at their barycentre. The−1
as N-bodysimulations
et al.
2011; Goddi
elocity is ∼6BN/KL,
km sas well. The
bulk
of the(Bally
SiO
emissionaccretion disc which is treated as a boundary forms the launching
et al. 2011). In particular, proper motion measurements of radio
base of the wind (e.g. Ouyed & Pudritz 1997). Further, we chooser ‘arms’ of an X-shaped pattern, while a ∼14 au thick0
the wind-launching point to be at a distance equal to the radius of
Table 1. Parameters chosen for the reference run tossion harbours the 7-mm
continuum
source(Reid
etal.
the binary orbit, i.e. r0 = a0 , where the material is rotating with astudy MHD-driven outflow from Orion Source I.
sub-Keplerian velocity v φ = χ v 0 . Further, the disc is considered to011). The continuum emission traces an ionized disc,be denser than the hydrostatic corona and this density contrast is
Quantity
Reference value
prescribed by δ. In a realistic (thick and hot) disc, the central gravprobe material ablated from the surface of the disc in
ity is balanced radially by contributions from thermal pressure andMasses of binary (m1 , m2 )10 M# , 10 M#
Orbital separation (2a0 )7 au
centrifugal rotation. This can be ensured by consistently deriving
0 state. Temperature contours (black)
Reducedstellar(µ∗ )5densityM#
Figure3.massNumber(left), poloidal velocity(middle) andof the reference run in steadythe sub-Keplerianity(χ )azimuthal
based on thevelocity
choice of (right)
density contrast.
Wind-launching radius (r0 )3.5 auFor
example,smallervaluesofδ
implyahotter(denser)discanda
I of three MHD critical surfaces – magnetosonicfor T = 500, 1000, 2000 and 5000 K are overlaid in the left-hand panel. The middle panel shows the positions
Keplerian velocity at r0 (v 0 ) 35.7 km s−1
larger contribution of thermal pressure resulting in the disc to rotate−14 g cm−3
Volumedensity(yellow),
at r0 (ρ 0 ) Alfvén
5.0 × 10surface
surface(magenta) andwith
magnetofast
arrowsthe right-hand panel represents the velocity vectors which are
sub-Kepleriansurface
speeds to(black).
maintain White
radial balance.
Weinchose
Plasma beta at r0 (β 0 )parallel to magnetic1.0
field lines shown by solid redδ ∼ 3lines.
as a reference value resulting in the underlying disc to rotate
Density contrast (δ)
2.85
at a speed 0.8 times the Keplerian speed.
of radiation forces and magnetic fields to show that radiation fromluminous stars with masses >30 M# does modify the dynamics ofoutflows. However, since the luminosity for a 2 × 10 M# binary
would be less than that of a single 20 M# star, we adopt that modelin the ideal MHD regime (neglecting radiative forces) to numerically simulate the dynamics of the disc wind inside 100 au from
Orion Source I.

Downloaded from http://mnrasl.oxfordjournals.org/ by guest on February 12, 2013velocity is updated each time with the value at the grid cell justhoop stress which collimates the flow. The right-hand panel of Fig.
MHD models for Orion Source I disc winds
L51
above the boundary. Further, ideal MHD approximation permits
3 finally shows the azimuthal component of the wind velocity, withto set the flow velocity along the poloidal magnetic field (v p %Bp ).
values ranging from 10 km s−1 at the outer edge to 25 km s−1 close
Thus, the mass flux at the disc boundary is not fixed but consistently
to the inner wind-launching point. The transfer of angular momenderived by causal interaction between the outflowing gas and thetum from the wind to the medium may impart rotation to the latter,boundary value. Detailed description of these boundary conditions
as observed in SiO masers (Section 3.2).
are given in Vaidya et al. (2011).

Tearing	
  the	
  Veil:	
  interac3on	
  of	
  the	  Orion	  Nebula	
  with	  its	
  neutral	
  environment	
  

Paul	
  P.	
  van	
  der	
  Werf,	
  W.M.	
  Goss	  and	
  C.R.	
  O’Dell	  l 12台のEVLA、13台のVLAでHI	
  21cm線を観測	
  l 1420.4MHz連続波はfree	
  free	  opacity	
  τﬀ=1.1	
  l Fig.10のCはHuygens	
  regionで端に弧状にHI	
  opacityが大きいところがある	
  l Eはdark	
  laneでHIの吸収が見える	  l HI	
  emissionは	
  1.	
  HIの吸収と同じ速度で、吸収が弱い場所	
  26
Van der Werf et al.
2.	
  南西に伸びるBar(図で北東-­‐南西)の南東のPDR	
  3.	
  高速度成分	
  4.	
  dark	  lane	  (次のページ図のE)	
  に見られる	
  Fig. 22.— Schematic geometry of M42 and the neutral Veil, illustrating the line-of-sight depths (approximately but not exactly to scale) of
the various components. Orange indicates ionized gas, green indicates molecular gas and blue indicates atomic gas. The outflow originatingin Orion-S, which powers the kinematics of HH 202 and H i component F is schematically indicated by a set of thin lines.

Tearing the veil of the Orion Nebula15

いろいろあってconclusions	  l HII	
  regionとOrion	  Nebula	
  Clusterが中性領域に与えるradiaDve、mechanicalなfeedbackを調査	  l 複雑な構造に対してgeometryを決めた	
  l radiaDve→	
  • ONCからのnon-­‐ionizing	
  UVがM42(ionized)とOMC-­‐1(molecular)の間にOrion	
  Bar	
  PDRを作る	  • Orion	  Bar	  PDRからのHI	
  emissionを捉えて、初めてedge-­‐on	
  PDRでHIのabundanceや他のtracerとの関係を示した	
  • PDRからのHI	
  emissionを観測して、PDRの中では5-­‐10%が原子でほとんど分子とわかった	  • HIの柱密度のピークはUVで励起されるH2の振動遷移のピークと一致	  →H2がUVを吸収して原子になるtwo-­‐step	
  processを確認	  l mechanical→	  • ionized	
  ﬂowとmolecular	
  ﬂowがHII領域の中性な部分と相互作用	  • HH202がVeilとショックを形成→吸収線でHI領域に速度ずれを検出(AとB)→Fig.12のような層構造	
  A	  mapping	
  study	  of	  L1174	
  with	
  13CO	  J=2-­‐1	
  and	
  12CO	  J=3-­‐2:	
  star	  forma3on	  triggered	
  by	
  a	
  Herbig	
  Ae/Be	
  star	  Jing-­‐Hua	
  Yuan,	
  Yuefang	
  Wu,	
  Jin	
  Zeng	  Li,	
  Wentao	
  Yu	
  and	
  Mar0n	
  Miller	  l L1174領域、Kolner	  Observatorium	
  Fur	  Submillimeter	
  Astronomie	
  3m	
  l 6つのmolecular	
  cores、5-­‐31	
  Mo、サイズは0.17-­‐0.39	
  pc	
  l HD	
  200775の周りは星風で圧縮されて空洞	
  →Herbig	
  Ae/Be	
  starが星形成を誘発	  l HD	
  200775は10.7	
  Moと9.3	
  Moのbinary、15000	
  Loで主星は105	
  yr	
  l 5つのコアはHD	
  200775に掘られた空洞の周りにある	  l LTE、Tex=20	
  K、opDcally	
  thinとして柱密度とMLTEを計算→MLTE<<Mvir	
  l 6個のIRAS	
  point	
  source	
  ▲、2MASSのうちYSO	  candidateが17個＊	  A mapping study of L1174 963
A mapping study of L1174

957

Downloaded from http://mnras.oxfordjour

Figure 2. Velocity-integrated intensity map of 13 CO J = 2 − 1 emission. The velocity interval covers −2 to 7 km s−1 . The contour levels range from 4.8 to
13.1 K km s−1 with a step of 1.5σ (1σ = 0.69 K km s−1 ). The peak positions of the six molecular cores identified in this paper are labelled with plus symbols.
13
The asterisks mark the candidate YSOs from the 2MASS PSC, while the filled triangles indicate IRAS point sources in the field of view, whose corresponding Figure 8. Distribution of line centroid of CO J = 2 − 1. The overlaid contours are the same as the ones in Fig. 2.
identities are listed in Table 3. The red star represents the Herbig Be star HD 200775.
√
13

Table 1. Fitting parameters of cores.

13 CO (J = 2 − 1)

12 CO (J = 3 − 2)

CO using the equation of σ = "v/2 2 ln(2). While F = 0, we
obtain the lower limits of the required external pressure. Then, theresults are divided by the Boltzmann constant κ; these are tabulated
in Table 2. Note that all six cores need external pressures of P /k ∼are redshifted for the eastern subregions II and III. This velocity
variation is also detected with 13 CO J = 2 − 1 line emission. The
velocity centroid shows a tendency to increase with ascending rightascension (J2000) for the six cores (see Table 1). Fig. 8 presentscore	
  1—0.43	
  km/sのinfallが見られ、on	
  going	
  star	  formaDon	
  core	
  2—軽くてdiﬀuse、星形成が始まる前	
  core	
  3—南からHD	
  200775に圧縮されdeeply	  embedded、early	
  stage	  core	
  4—core	
  2と同じ	  core	
  5—densest、HD	
  200775にかなり圧縮されている、global	
  expansion	
  core	
  6—largest、most	  massive、diﬀuse、global	
  expansion	  l どのコアも線幅が広い(1.8-­‐2.4	
  km/s	  cf.	
  音速0.27	
  km/s)→core1,3はinfall、
他はturbulentな環境が残っている	
  l MLTE<<Mvirを支える外圧は105	
  Kcm–3ほどで、乱流は105	
  yrしか続かない(tﬀ〜106	  yr)のでそのうち星になるだろう	  l HD	
  20075の東西で速度勾配0.94	
  km/s/pc→星風またはcloud	
  rotaDon	  l 2MASS	
  index	
  α=[J–H]/(1.8[H-­‐Ks]-­‐0.1035)が小さいYSOは若い	  →HD	
  200775から遠い2MASS	
  sourceほど若いとわかった	
  

aw
the
am
Th
is t
A
hav
to
wit
int
obs
ger

AC

Figure 10. The 2MASS index as a function of the angle distance to HD
200775 for the central 10 YSO candidates.

YSO candidate. This indicates that the more aged YSOs, with larger

We
hel
tio
110
gra
Ba
tio
Inf

Our	
  astrochemical	
  heritage	  Paola	
  Caselli	
  and	  Cecilia	
  Ceccarelli	  l 76ページのreview	  l まずabstract	  • 地球創世から伝わるものをastrochemical視点で考える	  • 流星や彗星からの情報の他、天の川銀河でのSun-­‐like	
  star形成もわかってきた	  • pre-­‐stellar	
  core,	  protostellar	
  envelope,	
  protoplanetary	  diskと進化するにつれて化学物質も複雑になっている	  • 流星や彗星から太陽系初期の情報を述べる	
  • 星形成進化段階で観測された科学組成を述べる	  18
Caselli & Ceccarelli• 各進化段階を繋げて、どれだけのものが化学的に受け継がれたのかを述
4.2 Ice formation and evolutionべる	
  Av !3 mag

CH4

NH3

CO2

!

H2O CO

_

⌇

N2

CH3OH

CH4

NH3

D2CO

CO2

H2O

!

CO

H2CO

Av "10 mag

!

_

"

3 ! Av !10 mag

!

CO

⌇

⌇

⌇

!

⌇

UV photons

⌇

!

!

⌇!

⌇!⌇

⌇

ICE

"

_

Cosmicrays

2"Stellar UV& X-rays! !D H ! !

!! ! !!!+ ! !!!! T<16K

cosmicrays

Deuteration zone!

!! !! !!

~~~

!! !!

PURE GAS DISK

T>20K

layer
cular
16K<T<20K
mole
rm
Wa
mistryud che
o
cl
rk
Da
H3+
D3+
H2D+

o to

!! ! !!! ! ! !! ! !!! ! !!! ! !

!

~~~

!

n

te d

Dark-cloud zone

HOH
CD2

Magnetosphericaccretion

Ph

in a

!

~
~~

~~~

ISRF

m
-d o

io n

!

~
~~

Outer edge
re g!

Caselli & Ceccarelli⌇38

Fig. 6 Ice mantle evolution within a pre-stellar core, from the outer-edge, where the core
merges with the surrounding molecular cloud, to the dark-cloud zone and deuteration zone
as depicted in Fig. 5. Ice mantles become thicker and richer in complex organic molecules
Fig. 14 Schematic structure of a ”naked” protoplanetary disk, adapted from [324], [139], moving toward the center of a pre-stellar core, where star- and planet-formation takes place.
0.03

0.1

1

10

100

distance in AU[395], [138] and [37]. The various regions are labeled. The black dots with various sizes
represent the coagulated dust in the disk midplane. See text for details.

Interstellar dust grains are crucial for the chemical and physical evolutionl ものすごくいろいろあってconcluding	
  remarks	
  • 地球上の水は2×10–3	
  ME<<800	
  ME	
  in	
  prestellar	
  core,	
  1.5	  ME	  in	
  protoplanetary	  disk→太陽系の月、彗星、小惑星、Kuiper	
  Belt	  Objectが水を持っている	  • Polycyclic	
  AromaDc	
  Hydrocarbonsなどの星間Complex	  Organic	
  MoleculesがInterplanetary	
  Dusty	
  ParDclesなど太陽系の有機物形成に寄与したのだろう	  • 彗星に含まれるメタノールはhot	
  corinos(protostarなど?)より10倍少ない	  • HDO/H2O比は海で1.5×10–4、彗星で同じくらい、envelopeやhot	
  corinosでは
6×10–4から0.2、この違いはopen	
  quesDon	  • H2Oでない分子のD/H比はH2Oに比べて大きい、地球や彗星で10倍、IDPsで1%ほど大きい→pre-­‐,	  proto-­‐stellar	
  phaseと太陽系の繋がり	
  • 彗星や星形成領域のHCNとHNCでは15Nがリッチな一方、NH3やH2H+ではリッチ
44
Caselli & Ceccarelli
ではなくこれらはpre-­‐stellar	
  coreでD化されている	  Our astrochemical heritage

13

9%

103P/Hartley2"
COMET"COMETS%
OTHER"COMETS"

IOM"HOT"SPOTS"

5"
SOM"

15N/14N%x10#3%

15N/14N"x10,3"9"

IOM%HOT%SPOTS%

5%
SOM%

CHONDRITES"CHONDRITES%

EARTH"

1"

PRESTELLAR%CORES%&%PROTOSTARS%EARTH%

SUN"
1"
2"

3"

4"

5"

D/H"x10,4"
Fig. 4 15 N/14 N versus D/H in comets, chondrites, hot spots in the IOM of meteorites and
IDPs, SOM in meteorites, Earth and Sun.

in comets applies to the organic material in carbonaceous chondrites and IDPs.

HCN%
HNC%

1%

NH3%
N2H+%

SOLAR%NEBULA%
10#5%

10#4%

10#3%

10#2%

10#1%

1%

D/H%
Fig. 17 15 N/14 N versus D/H in comets, chondrites, hot spots in the IOM of meteorites
and IDPs, SOM in meteorites, Earth, Solar Nebula and pre-stellar cores and protostellar
envelopes.