DISK DISPERSAL AROUND YOUNG STARS
THE STAR FORMATION NEWSLETTER

4Parameterizing
the viscosity
by ! ! "cs H (Pringle and Rees 197
An electronic publication dedicated to early stellar evolution
and molecularclouds

Shakura and Sunyaev 1973), where scale height H and sound speed
are given by H ! rcs /(GM" /r )1/2 and cs ! (kT /mH )1/2 , we find ! $ r ,No. 237 — 20 Sept 2012 Editor: Bo Reipurth (reipurth@ifa.hawaii.edu)
the temperature in the disk follows a T $ r #1/2 power law. Given that thviscous timescale t! ! r 2 /! , it follows that t! $ r :The dispersal of protoplanetary disks around binary stars
Richard Alexandert! ! 105 yr" #1r0.01
10 AU

" # "

#Accepted by ApJ Letters http://arxiv.org/abs/1209.0779
The viscous dispersal timescale is plotted in Fig. 1 for two values of(10#3 and 10#2 ) thought to best represent the effective viscosity in pr
アイディア：連星系での光蒸発を考える。中心星付近で潮汐に
toplanetary disks (see, e.g., Hartmann et al. 1998). Assuming that diangular momentum is constant and that the disk mass resides mostly
よってinner cavityがあく。→内側がoptically thinになるので、
the outer disk radius rd , it follows that rd (t ) $ Md (t )#2 . Therefore,
円盤は内側から一気に晴れ上がる。→連星系での円盤寿命は光蒸
the disk accretes matter onto the central star, it preserves its angular m
発で決まっている。
mentum by expanding to greater size. For times greater than the initi
timescale t! 0 ,計算：連星系周りでの1D円盤進化計算(粘性円盤+光蒸発+潮tt #1/2
andr d ! rd 0
Md (t ) ! Md 0汐)。
t! 0t! 0
F IG . 1.— Surface density evolution of the standard models with a = 1AU.
特に２つの場合を考える：
The upper panel shows the strong wind case (typical of X-ray photoevaporaF IG . 2.— Disk lifetimes tdisk plotted as a function of binary separationtion), with Σ(R) plotted at 0.1, 0.5, 0.7, 0.85, 0.9, 0.95 & 1.0t
disk . (Herethe curves
are normalised
to the lifetime
1AU. disk
The
whereare allthe
initialradius andmassatoaf=theat t! 0 . Ardthe
d 0clarity,
0 and a.MFor
1, weak wind (EUVに対応)
disk lifetime tdisk = 1.68Myr). In this case the majority of the disk is removed
single-star models, plotted for reference, have the same viscous time-scales
by the photoevaporative wind, and the disk is cleared rapidlythe
once disk
the surface
gr
ws,
visc
us
timescaletheuterregi
ns,
wherethe maoooo
asthe
the standard
model
set. The upperin
panel
shows
the strongwindmodels;Rc=1.4AU, Mdot̲thick<Mdot̲thin
density drops below ! 1g cm−2 . The lower panel shows the weak wind caselower panel the weak wind models. The sharp increase in disk lifetimes
lothenger.
Therefore, viscous accretion may account for th(typical of EUV photoevaporation), with Σ(R) plotted at the resides,
same fractionsgetsoffora < 0.25R
c is caused by the “switch” to direct photoevaporation of thethe disk lifetime (tdisk = 13.74Myr). Here accretion dominates the evolution,
2, strong wind (X-rayに対応)
disk once the binary has cleared a sufficiently large inner cavity.
and the wind only clears the disk once it has undergone substantial viscous
evolution.
Rc=5.0AU, Mdot̲thick=Mdot̲thin" #the binary. The high mass loss rate also has a strong impact on
the radial profile of the disk: the surface density never reaches
結果は、1の場合で連星間距離(a)を変えた時に円盤寿命に大the power-law profile characteristic of accretion- or decretiondominated disks. Eventually the wind prevails completely,
きく影響。
and the disk is then rapidly cleared from the inside out. This
final clearing retains the “two-time-scale” behavior characteristic of photoevaporative disk dispersal (Clarke et al. 2001;
Alexander et al. 2006b), but with the important caveat that ac(参考)円盤の寿命
cretion and angular momentum transport play only a minor
右図(Hollenbach et roleal.)は円盤の散逸タイムスケールを表す。
in the disk’s evolution.
By contrast, in the weak wind case the initial accretion rate
内側では降着、外側では光蒸発が支配的。
exceeds the wind rate by a factor ! 100, and the disk’s evolution is governed primarily by the viscosity. At R ! 10AU thedisk is close to the Σ ∝ R−3/2 power-law of a decretion disk
and, as in the single-star case, only once the disk has undergone substantial viscous evolution does the photoevaporative
wind trigger disk clearing. The inner edge of the disk remains
at a fixed position throughout, and although the final clearing
again shows the characteristic two-time-scale behaviour, here

" #the majority of the disk has been accreted on to the binary,with only a small fraction (7%) removed by photoevaporation.
Further insight can be gained by looking at Figure 2, whichshows how the disk lifetime varies as a function of binary
separation a in the different models. The most prominent features are a sharp increase in the disk lifetime at small separations (a " 0.25Rc), and a progressive decrease in disk lifetimes at large separations (a ! 5–10AU). The increase in disk
lifetimes at small a occurs because the photoevaporative wind
has a characteristic radius: when the binary clears a cavity inthe disk that is larger than the critical radius Rc , the wind isdriven by direct irradiation of the inner disk edge. This increases the efficiency of the wind, and clears the disk more
rapidly. The increase in tdisk for small a is much larger whenaccretion on to the binary is suppressed, and the critical valueof a also increases for a lower mass ratio q. This effect is alsomore pronounced in the weak wind models. This is primarily due to the more prominent role of accretion in the weakwind models, but is also partly because direct irradiation in-Orion Revisited - I. The massive cluster inFigure 1. Timescales for disk dispersal: t! is the viscous timescale for "front of the Orion Nebula Cluster
10#2 and 10#3 ; tSE is the stellar encounter (tidal stripping) timescale for Trap
zium cluster conditions (see sections II.B and III); tws and tws% are stellar wiJo ̃ao Alves and Herv ́e Bouy
stripping timescales for wind and disk parameters summarized in section I
Accepted by A&A http://arxiv.org/abs/1209.3787 tc (evap) is the photoevaporation timescale by the central star (strong wind casand tE (evap) is the photoevaporation timescale for an external star (Trapeziuconditions) for the conditions summarized in sections II.D and III.

目的：オリオンA 分子雲のForeground Stellar Populationを調べた(CFHT, CAHA, Spitzerのデータ)。
結果：B型星からM型星について
特に空間分布について、NGC1980のあたりにピーク
この研究の結果から推察した NGC1980 の特性 : 年齢4-5Myr, 星の数2000個(ONCの中でも多い)→この分布は、従来ONCとL1641Nのポピュレーションだと思っていたものとオーバーラップする。
以上より、従来ONCと思っていたものは実は複数のpopulationの組み合わせだった
1)トラペチウムなどの最も若い星、2)foreground, 主にNGC1980, 3) ﬁeld starthe left-, central and the right-hand panels respectively.

On the evolution of irradiated turbulent
clouds: A comparative study between
modes of triggered star-formation
S. Anathpindika and H.C. Bhatt
Accepted by MNRAS http://arxiv.org/abs/1208.6114

分子雲中のガスは常にionizing radiationに晒されている。小さい分子雲はionizing radiationで壊れてしまうが、ある程度大きい分子雲はionizing radiationから生き残って新たな星形成領域となりうる。
On the evolution of irradiated turbulent clouds: A comparat→これを調べるために、ionizing radiationも考慮したSPHシ
Figure 10. Left-panel: A rendered column-density plot showing the mid-planeofmass
ミュレーションを行った。
few protostellar objects. The protostars, represented by the sink-particles, have be
1
M(ρtNo on
IR the cloud surface, though another potential site of star-formation
shockedlayer
radiationの強さは典型的なB型星からB型星クラスター
N =4.10 sstron
N density=5.10 s plot showing the mid-plane of the irradiated cloud in case 4. As incolumn
N =7.10 sまでパラメータを振った。
layer on the irradiated surface of the cloud (t ∼ 0.1Myr ≡ 0.33t ). Shocked fila
LyC
LyC
LyC

48 -1
49 -1
51 -1

ff

are a(201the f
5 DISCUSSION
0.1
sugga relStar-formation according to Krumholz & McKee (2008), is un- duce
likely to commence unless gas density is higher than a threshold. smal
The proposition is attractive, for such a threshold will essentially cloud
reduce the problem of star-formation to simply an investigation ults
into the processes likely to assist the assembly of a sufficiently in Fi
0.01
dense volume of gas. However, it is common knowledge that thedynamical state of gas is also inextricably interwoven with its tem- ing t
1525perature,
where for convenience,
we include turbulence as justan- cludi
log10(n/[cm-3])
other thermal component. It is therefore difficult to define a density few
12thresholdfor the formation of stars.
Gravity+turbulence51only-1
posin
density to spawn stars.

pn

結果、radiationであぶられる所に密度の高いﬁlament構造ができて、そこで星形成が活性化された。

更に確率分布関数を計算。ionizing radiationは星の分布関数に大きく影響すると思われる。

NLyC=7.10 s
N =5.1049 s-1

LyC
48 -1
NLyC=4.10s
There are two crucial issues before us : (i) theprocess,usually the treferred to as a trigger, that assembles pockets of dense gas, and (ii) ment
The nearby population of M-dwarfs with WISE:
A search
forthesewarm
11
theprocesses
that maintainpockets at the right temperature filam
so they can spawn stars. Clearly, this latter question is beyond the tuall
circumstellar dust
10.5
remit of the present work, however, by maintaining isothermality even
Henning Avenhaus1, Hans Martin Schmid1 and Michael R.
Meyer1
withinthe cloud, at least the first half of the problem can be ex- the e10

Accepted by A&A http://arxiv.org/abs/1209.0678

log10(E(k)/[ergs g-1])

11.5

9.5

WISEのデータを使ってデブリディスクサーベイを行った。
9
従来太陽型星がよく調べられてきたが、M型矮星についてはあまり調べられて来なかったのでM型矮星につ8.5いて調べた。サンプルは103個のM型矮星。
8
結論として、MIRでの赤外線釣果はひとつも見られなかった。(従来の結果と比べると22μmではSpitzerの0.5
0.6
0.7
0.8
0.9
1log (k)ものよりも検出限界は10-15%良くなっている。)
Figure 11. Upper-panel : Figure showing the probability distribution func→この研究からのM型星周りのデブリディスクを持つ確率は0+1.3%
tion (PDF) for gas density within the cloud at an advanced stage of its evol(年齢の影響?>30Myr)
ution in each of the 4 test cases. A log-normal distribution fits all the PDFs
10

reasonably well, though they also have a prominent power-law tail. We notethat the width of the log-normal distribution is narrower for the irradiatedclouds and the narrowest for the choice of the strongest flux. The best-fitting
log-normal distributions have also been shown for comparative purposes.
On the structure of molecular clouds
Lower-panel : The power-spectrum for each of the 4 test cases has been
The spectra show remarkable similarity to the Kolmogorov specJavier Ballesteros-Paredes1, Paola D’Alessio1 and Leeplotted.
W. Hartmann2
trum. We note that maximum power is cascaded down to smaller spatial
scales in case 4.

Accepted by MNRAS http://arxiv.org/abs/1209.3346

case shows that a progressively smaller mass of the cloud is reLarson scaling relation は観測の解釈の間違いだ！という論文
tained with increasing strengths of the incident flux. The observed＊scaling relation : ρ R^-1
M R^2column density
N=constant−4
rate of photoablation in cases 3 and 4 is a few times ∼ 10 M#yr−1 , which is consistent with that calculated using Eqn. (4) above.
Plotted in the lower-panel is the timescale on which gas within thecloud overhauls a fiducial density threshold, ρthresh ∼ 10−18 g cm−3 .
Thus, with relatively weaker shocks resulting from damped turbulence in case 1, it takes considerably longer, more than 70 % of thefree-fall time, to overhaul a fiducial density threshold of 10−18 gferented,
cloudformthatgrew
any i5.1

Thethe e
ity w
plott
to beevide
4 spa
in ca
has dends
distras ca
the ihowesteepWe show that the inter-cloud Larson scaling relation between mean volume density and size ρ ∝ R−1, which in turn implies that mass M ∝ R2, or that the column density N is constant, is an artifactof the observational methods used. Specifically, setting the column density threshold near or above
the peak of the column density probability distribution function N-pdf (N ∼ 1021 cm−2) produces
the Larson scaling as long as the N-pdf decreases rapidly at higher column densities. We argue thatthe physical reasons behind local clouds to have this behavior are that (1) this peak column density
is near the value required to shield CO from photodissociation in the solar neighborhood, and (2)gas at higher column densities is rare because it is susceptible to gravitational collapse into muchsmaller structures in specific small regions of the cloud. Similarly, we also use previous results toshow that if instead a threshold is set for the volume density, the density will appear to be constant,implying thus that M ∝ R3. Thus, the Larson scaling relation does not provide much information on
the structure of molecular clouds, and does not imply either that clouds are in Virial equilibrium, or
have a universal structure. We also show that the slope of the M − R curve for a single cloud, whichtransitions from near-to-flat values for large radii to α = 2 as a limiting case for small radii, depends
on the properties of the N-pdf.
A Disk-Wind Model for the Near-Infrared Excess Emission in Protostars
Alissa Bans1 and Arieh K ̈onigl1
Accepted by The Astrophysical Journal http://arxiv.org/abs/1207.1508
4Bans and KöniglFig. 1.— Schematic representation of a centrifugally driven disk wind from the vicinity of a low-luminosity protostar. The gaseous wind
uplifts dust from the disk surface beyond the sublimation radius rsub but remains dust-free closer to the protostar. The innermost radius
r0min of the wind-launching zone may (but need not) coincide with the radius rm (Equation (1)) where the disk is truncated by the stellar
magnetosphere. The outer radius r0max of this zone is possibly limited to ∼ 2 rsub by dust absorption of stellar FUV photons in the wind.
Most of the NIR excess emission is produced by reprocessing of the stellar radiation by the dusty wind.

Protostellar
systems,
ranging fromlow-luminosity
T Tauriand Herbigstars
to highTableclassical1 lists theparametersof the Ae
threerepresentative
In the radial
self-similarity
formulation,
any pointinluminosity
Be vector
stars, exhibit
near-infraredexcesstheir(1993b)
spectrathat
thatweisutilize
dominatedwind(NIR)
solutions
frominSafierin ourthe wind isspecified byHerbig
the radiusR, whosea cylincalculations.
drical coordinates{r, inφ z}
by a bumpthearemonochromatic luminosity witha peak near 3 microns. The bump can be
If the wind transports away all the excess angular moR = {r0 ξ(χ),
r0 χ} . emission component(2)
approximatedby a φ,thermalof tem-ofperature∼ 1500K thatof theorder rateof thementum
the accretedgas, andthe ismass
accretion
Here χ ≡ sublimation
z/r0 is the dimensionlessverticalcoordinateoftemperature of interstellar dust grains.
In therim”with rathroughthecurrently
disk, Ṁin ,popularremains “puffed-up
nearly constanta point alonga streamline
whose
footprint intersects
thedius,propagatesthen the ratioof Ṁoutthe
to opticallyṀin can bethin
approximatedscenario,the bump
represents
stellar radiationthatthroughinner
disk at a radius r0 (with the subscript ‘0’ denoting thebyln{(r/r)/[2(λ−1)]}(seeKönigl&
Salmeron
0max
0min
thephysicalsurrounding
accretionand is absorbed and reemitted by the dust that resides
justequatorialregion
plane).ofAllquantities
can bedisk
repre2011). According to this estimate, Ṁout /Ṁin should lie
sented as beyond
productstheof dusta powerlaw in r0 anda function
sublimation
radiusrsub. However,thismodel
cannotaccount
forthestrongest
in the range ∼ 1 − 5 % for the models listed in Table 1.
of χ. In particular,
the mass indensity
be written
bumps measured
thesecansources,and itaspredicts
a pronounced
secondarybounce in determinations
theThisis consistent withthe observational! "−3/2of thisInratio
low-mass
protostars.
However, using ourinterferometricthisinpaper
we present
an alternativer0visibility curve that is not observed.
η(χ) ,
ρ = ρ1best-fit valuesof Ṁradiation
HAe
that weout for theinterpretation,
attributes
the bump(3)to reemissionof stellar
bytwo
dust
thatstars
is uplifted
rwhich1model in this paper (which are ! 4 × 10−8 M" yr−1 ; seedisk atbythe
a centrifugally
drivenwind. Winds
areexpression
a leading would
candidate
the acFigureof3),thisthetype
above
implyformass
where ρ1 fromis thethedensitydisk surface at
the fiducretion
rates
that
areinexcess
of
those
that
havebeencial disk radius
r,
which
we
henceforth
taketobe
1AU,
origin 1of the strong outflows associated with protostars, and there is observational evidence fordisk
inferredobservationally
for
these
twosources(whichareand η(χ) isobtained
from
thesolutionof
theMHDwind
winds originating on scales ∼ rsub. Using a newly constructed
Monte
Carlo
radiative
transfercode
equations. Using Equations (2.1), (2.7c), and (5.2) in" 10−7 M" yr−1 ; see below). The discrepancy could befocusing
on low-luminosity
Blandfordand& Payne(1982),and normalizingsources,
by values
reduced if we have overestimated Ṁout (for example, beappropriate to our modeled HAe stars, we can express
cause of our neglect of other likely contributions to thethe fiducial
3 density as
NIR excess emission, as discussed in Section 3, or because
#$!
"−1/2
we have overestimated the ratio r0max /r0min in EquaṀoutM∗tion(4)). Note in this connection that a mass outflow
ρ1 =3.5 × 10−15
2 M"
5 × 10−8 M" yr−1rate ∼ 1 × 10−8 M" yr−1 was inferred in MWC 275 from"−1!
observations of Herbig-Haro knots in the bipolar outflow
ψ0 ln (r0max /r0min )emanating from this source (Wassell et al. 2006). It isg cm−3 ,×0.03
2.5
also conceivable that the mass accretion rate in these(4)intermediate-mass sources undergoes a measurable de-we show that this model can account for the NIR excess emission even in bright Herbig Ae stars
such as AB Auriga and MWC 275, and that it successfully reproduces the basic features of thevisibilities measured in these protostars. We argue that a robust dusty outflow in these sources could
be self-limiting — through shielding of the stellar FUV photons — to a relatively narrow launchingregion between rsub and ∼ 2rsub. We also suggest that the NIR and scattered-light variability
exhibited by a source like MWC 275 can be attributed in this picture to the uplifting of dust clouds
from the disk.