星形成ゼミ
2021/04/30
SFN 338 #34-#39

長谷川哲夫

1

to determine which is younger and which is older. Sz102 and Merin28 are the most evolved sources and show signs of having

338 Star formation in ‘the Brick’: ALMA reveals an active proto-cluster in the Galactic centre cloud#34 G0.253+0.016
2021 MNRAS.503…77W

2

fragmentation, suggesting that the large-scale turbulence and strong magnetic field in the cloud do not dominate on these scales,
and thatThestar Milkyformationon thescale ofindividual
protostars
is similar
thatResultsin Galactic
environments.
Way
Project:
ProbingStar
Formation
withtoFirston disc
Yellowballs
from DR2338#352021 ApJ…911…28W

The Milky Way Project: Probing Star Formation with First Results on Yellowballs from
DR2

Grace Wolf-Chase, C. R. Kerton, Kathryn Devine, Anupa Pouydal, Johanna Mori, Leonardo Trujillo, Aurora
Cossairt, Sarah Schoultz, Tharindu Jayasinghe, Matthew Povich F Yellowballs (YBs) were first discovered during
the Milky Way Project citizen-science initiative (MWP; Simpson et al. 2012). MWP users noticed compact, yellow regions inSpitzer Space Telescope mid-infrared (MIR) images of the Milky Way plane and asked professional astronomers to explain these"yellow balls." Follow-up work by Kerton et al. (2015) determined that YBs likely trace compact photo-dissociation regions
associated with massive and intermediate-mass star formation. YBs were included as target objects in a version of the Milky
Way Project launched in 2016 (Jayasinghe et al. 2016), which produced a listing of over 6000 YB locations. We have measured
distances, cross-match associations, physical properties, and MIR colors of 500 YBs within a pilot region covering the l= 30 - 40
degrees, b= +/- 1 degree region of the Galactic plane. We find 20-30high-mass star formation capable of becoming expanding
H II regions that produce MIR bubbles. A majority of YBs represent intermediate-mass star-forming regions whose placementin evolutionary diagrams suggest they are still actively accreting, and may be precursors to optically-revealed Herbig Ae/Be
nebulae. Many of these intermediate-mass YBs were missed by surveys of massive star-formation tracers and thus this catalog
provides information for many new sites of star formation. Future work will expand this pilot region analysis to the entire YB
catalog.

12

6

The Astrophysical Journal, 799:153 (9pp), 2015 February 1

Kerton et al.

Figure 4. MWP GLIMPSE/MIPSGAL images, 4.5 µm (blue), 8 µm (green), and 24 µm (red), showing yellowballs with (in white) and without (in yellow) RMS
counterparts. The top panel shows part of W 33, with a cluster of yellowballs located between two large bubbles. The lower panel shows two isolated yellowballs andtwo yellowballs associated with a larger H ii region (see the text for details). The Galactic coordinates of the image center and the angular size of the image are shownat the top of each frame.

8Figure 1. YB spatial distribution. Heliocentric distances and Galactic longitudes are plotted for all 493 YBs in the pilot region
with CO associations (diamonds). The average statistical uncertainty in the distance is ⇠ 0.5 kpc. Lines show the spiral arm
traces along this line of sight from Reid et al. (2016): Per = Perseus, SgN and SgF = Sagittarius near and far sections, ScN and
ScF

9

16

Wolf-Chase et al.

First Results on Yellowballs from DR2

17

Figure 5. Radar-plot visualizations for five distance-independent quantities for di↵erent classes of objects. Values plotted are
medians, and the maxima are indicated for each radial line. Di↵erent classes are connected via di↵erent colored pentagons.
Left: Protostellar clumps and H II regions identified in the Hi-GAL survey (EMS17) are compared with YB–Hi-Gal matcheslacking counterparts in other surveys. Right: Four YB–Hi-GAL subsets are compared. The RMS and WISE-K subsets arealmost entirely high-mass star-forming regions, while the WISE-Q and YB–Hi-GAL only subsets are a mix of intermediate-mass
regions and young high-mass regions.

extend the KWA15 analysis to the larger DR2 sample and examine YB colors using additional IR wavelengths. The
histogram of the log(F12 /F8 ) results for DR2, similar to that shown in KWA15, is shown in Figure 6.
KWA15 found that the average value of their sample (-0.19) was more negative than the average color sample of
WISE H II regions (-0.09) (Anderson et al. 2012) and that the YBs in their sample without RMS counterparts had aneven more negative average of -0.23. The colors of our new YB sample yields an even more negative mean log(F12 /F8 )value of 0.43 ± 0.15. Size measurements (based on MWP user radii) confirm that our new YB sample indeed contains
a larger number of very compact sources. Whereas the mean (± standard deviation) diameter of the RMS-matched
YBs with distances from KWA15 is 0.94 pc ± 0.71 pc, the mean diameter of all YBs with good distance estimatesin our current sample is 0.75 pc ± 0.53 pc. The median diameters are smaller (0.77 pc and 0.63 pc, respectively),as relatively few YBs are significantly larger than 1 pc. The mean diameter of the subset of YBs with counterpartsin catalogs of H II regions and MYSOs (CORNISH, RMS, WISE-C,G,K) is 0.96 pc ± 0.62 pc and the mean (±standard deviation) log(F12/F8) color of these objects is -0.28 ± 0.31; however, the mean diameter of YBs lacking
such counterparts is 0.65 ± 0.45 pc and the mean log(F12/F8) color of this subset is -0.47 ± 0.24. Many compact YBs
were missed in the KWA15 sample, which was drawn from images at lower zoom levels, where citizen scientists would
have been more likely to notice large and bright YBs. Additionally, only YBs with RMS counterparts had associateddistances in the KWA15 sample; our current sample indicates that YBs with counterparts in catalogs of H II regions
and MYSOs are biased to physically larger objects.
In addition to measurements at 8 and 12 µm, we use longer MIR wavelengths at 24 and 70 µm to examine the
YB population in our DR2 sample. We show a MIR color-color plot in Figure 7 for the 219 YBs in our pilot region
with reliable 8-µm and 24-µm photometry as well as counterparts in the Hi-GAL survey (EMS17). The color-color
plot indicates sources with cross-matches in the RMS and WISE catalogs. The average H II region colors and H II
region - PNe cut-o↵ colors established by Anderson et al. (2012) are also indicated in Figure 7. WISE sources wereseparated into categories by Anderson et al. (2014): “K” for known, “G” for group, “C” for candidate, “Q” for radio
quiet. YBs with counterparts in the RMS catalog and/or WISE K, G, and C sources are, unsurprisingly, grouped
near the expected colors for H II regions. We find a significant number of sources with WISE Q (n = 44) or no

Figure 4. Bolometric Luminosity – Clump Mass plots for various subsets of the YB–Hi-Gal Sources. Red lines show Molinari
et al. (2008) clump evolutionary tracks. Small squares, along the vertical portion of each track, correspond to 5 ⇥ 104 year
intervals, and isochrones are the red short dashed lines. The blue upper diagonal line represents a 90th percentile lower limit for
Herschel -defined H II regions, and the blue lower diagonal line is a 90th percentile upper limit for the pre-stellar objects in the
EMS17 sample of Hi-GAL clumps(see text for details). For sources with good SED fits (black), mass and luminosity error bars
are shown (most are smaller than the symbol used). Sources with poor SED fits (grey) do not have error bars shown for mass.

10

338 Formation and evolution of protostellar accretion discs. I. Angular-momentum budget, gravitational#36 self-regulation, and numerical convergence2021 MNRAS.502.4911X
Formationand evolution of protostellar accretion discs. I. Angular-momentum budget,
gravitational self-regulation, and numerical convergence

Wenrui Xu, Matthew W. Kunz F We investigate the formation and early evolution of a protostellar disc from a magnetizedpre-stellar core using non-ideal magnetohydrodynamic (MHD) simulations including ambipolar diﬀusion and Ohmic dissipation.
The dynamical contraction of the pre-stellar core ultimately leads to the formation of a first hydrostatic core, after ambipolar
diﬀusion decouples the magnetic field from the predominantly neutral gas. The hydrostatic core accumulates angular momentum
from the infalling material, evolving into a rotationally supported torus; this ‘first hydrostatic torus’ then forms an accreting
protostar and a rotationally supported disc. The disc spreads out by gravitational instability, reaching ⇠30 au in diameter at⇠3 kyr after protostar formation. The total mass and angular momentum of the protostar-disc system are determined mainly
by accretion of gas from an infalling pseudo-disc, which has low specific angular momentum because of magnetic braking; theirremoval from the protostar-disc system by outflow and disc magnetic braking are negligible, in part because the magnetic field
is poorly coupled there. The redistribution of angular momentum within the protostar-disc system is facilitated mainly bygravitational instability; this allows formation of relatively large discs even when the specific angular momentum of infallingmaterial is low. We argue that such discs should remain marginally unstable as they grow (with Toomre Q ⇠ 1-2), an ideathat is broadly consistent with recent observational estimates for Class 0/I discs. We discuss the numerical convergence of ourresults, and show that properly treating the inner boundary condition is crucial for achieving convergence at an acceptablecomputational cost.

The Perseus ALMA Chemistry Survey (PEACHES). I. The Complex Organic Moleculesin Perseus Embedded Protostars

Yao-Lun Yang, Nami Sakai, Yichen Zhang, Nadia M. Murillo, Ziwei E. Zhang, Aya E. Higuchi, Shaoshan
Zeng, Ana López-Sepulcre, Satoshi Yamamoto, Bertrand Lefloch, Mathilde Bouvier, Cecilia Ceccarelli, Tomoya
Hirota, Muneaki Imai, Yoko Oya, Takeshi Sakai, Yoshimasa Watanabe F To date, about two dozen low-massembedded protostars exhibit rich spectra with complex organic molecule (COM) lines. These protostars seem to possess
diﬀerent enrichment in COMs. However, the statistics of COM abundance in low-mass protostars are limited by the scarcity ofobservations. This study introduces the Perseus ALMA Chemistry Survey (PEACHES), which aims at unbiasedly characterizing
the chemistry of COMs toward the embedded (Class 0/I) protostars in the Perseus molecular cloud. Among the 50 embedded
protostars surveyed, 58protostars have CH3 OH, CH3 OCHO, and N-bearing COMs, respectively. The detectability of COMs
depends on neither the averaged continuum brightness temperature, a proxy of the H2 column density, nor the bolometricluminosity and the bolometric temperature. For the protostars with detected COMs, CH3 OH has a tight correlation with
11
CH3 CN, spanning more than two orders of magnitude in column densities normalized by the continuum brightness temperature,suggesting a chemical relation between CH3 OH and CH3 CN and large chemical diversity among the PEACHES samples at the6

W. Xu and M. W. Kunz

Figure 1. Azimuthally averaged profiles at di↵erent epochs from our fiducial 3D simulation. The total mass enclosed within 10 au (including
the accretor mass), M10 , is provided for reference. Each panel displays the density (colour), magnetic-field lines (black streamlines), andvelocity (white quivers). Each row captures a di↵erent phase of evolution; from top to bottom: the flattening of the pre-stellar core
along magnetic-field lines and the advection-di↵usion of the magnetic field into an hourglass shape, the formation of the first hydrostatic
core/torus on ⇠au scales, and the spreading of the rotationally supported disc through gravitational instability.

⇡0.2 G, a value similar to that found in previous calculations

122.1 Initial conditions

Pre-stellar flattening

The first hydrostatic

Gravitational instability

midplane v [km s 1 ][g cm 2 ]midplane nn [cm 3 ]We employ the same physical initial conditions for all simandspherical,
collapsecore/torus
and diskandspreading
Protostellar disc formation
evolution 7
ulations. The pre-stellar core is initiallywith thenumber density of neutrals given by
1015
1015
1015
4 W. Xu⇢ andnM.
W.
Kunz
1024.46kyr
1026.15kyr1027.00 kyr
0nn (r) = max, n1 .
(1) kyr
1013
1013
1013
1025.55
1026.31 kyr1027.41 kyr
2
1 + (r/r0 )1025.93 kyr1026.64 kyr
kyr(cf. figure 8a of Kunz & Mouschovias 2010).
small
di↵usion 1027.84
and promotes
1011 Starting the sim1011 grains also increases ambipolar1011
4
3
We choose
initialcentraldensityn0 time= 10
cm gasandulationat a anlower
densitya↵ordsenoughfor thetodisc 9formation (Zhao et al. 2018a).
In Appendix B we show
109
10
109
characteristiccore size r0 = lines
0.1 pc,representativeof a typflatten
along magnetic-field
beforethe core contracts
the abundance of di↵erent charged species and the resulting
7
10
107 0
107 0 model.1
ical NH3 core
(Jijinathe
et magnetical. 1999).
The
backgroundden3
dynamically.
Becausefield
evolvesunderneardi↵usivities
100
101
102
10non-ideal
10
101
102 for our
103 dust
10
10
102
103
3
sity n1 = 500
cm the, attained
incollapse
the initialat radiiflux-freezing
duringdynamicalof astatesupercritical
4
104
104
r 0.436pc,atprovidesa numericalphysicallyrepcore,starting
a lower initial
densityfloor
also10thata↵ordsthe initially3
33
resents thedensityfieldof theambient
molecular
gas consistently
surrounding
2.4 10Computational
domain and10spatial
resolution
uniform
magnetic
time
to adjustand10evolve
2
2
2
the
core;
this
gas
is
excludedfromthecalculationoftheself1010
10
with the density profile.
We usea spherical-polar grid withan outer radial bound1
gravitational potential in order to avoid 10
unphysicalbehavior
101
101
18
ary
at
2
au
(1.27
pc,
or
12.7r
)
and
a fiducial inner ra00
at large r (see Appendix A). Adopting a 10
mean
mass per neu100
100
dial boundary
at rin2 = 1 3 au. The
radial(r) 2 grid is3 log2.2
ofnstate
tralEquation
particle m= 2.33mp (accounting for100molecularhydro101
102
103
100
101
10
10
100
101
10
10
uniform
withafiducial
resolution
of
⇡14
cells
per
factorof 2
gen with 20% He by number), the total self-gravitating
mass
101
101
101
For physical simplicity and computational
expediency,
we(dr/r ⇡ 0.05). To reduce computational cost, we only simuvr
enclosed within r is then '1.43 M . The initial
temperature0model the thermal0 evolution of the system10using
the100 boundary condition at thev following
late100✓ 2 [0, ⇡/2] with a reflectingof the gas is set to T0 = 10 K, giving an initial isothermalbarotropic equation of state (EoS)
for thetemperature:1
midplane;for our 3D simulations,
2 [0, ⇡]
10 km
10 1
10 1 we only simulate
sound speedh cs0 ⌘ (kB T0 /mn )i1/2 ' 0.188
s 1.
with2 a periodic boundary condition.
The polar (✓) grid is
2/3 initially uniform,The
core is threaded by anvertical mag10 2
10
10 2
T (n
,
(2)
n ) = T0 1 + (nn /nadia )non-uniform
with
the
grid
spacingdecreasing
towards the
netic field with strength B0 = 25 µG. In 3the innermost flux
3
3
10
10
10
0
1
3
1 the polarthecell103size 10
is0 ⇡1/310ofcell
102
10midplane;
100
101 midplane
102
102
103 size,tubes of
the =
core,
ratio
to magwherenadia
2 ⇥the1011
cm of3 self-gravitating
is value of10 nn massat 10which
the
and
the
fiducial
resolution
is
28
cells
in
2D
and
24
cells
in 3D,
netictransitions
flux is '1.5
times thecentral value
a forlower
colgassmoothly
fromcriticalan isothermal
EoS (atresulting in a midplane d✓ ⇡ 0.03 rad (⇠2 deg). In 3D, the
1/2
lapse, (3/2)(63G)(Mouschovias& Spitzer1976).
The
densities)
to a polytropic
EoS with index= 5/3
mimickazimuthal ( ) grid is uniform with a fiducial resolution of 16
coreadiabatic
has supercriticalratio outTheto abreak-point
cylindricaling
evolution mass-to-flux(at higher densities).
cells for 2 [0, ⇡].
radius
Rsc is' chosen
1.27r0 . such that equation (2) accurately reproof
the EoS
To test the numerical convergence of our simulations, in
For ther dependence
r0 , the coreof
is theset central
into uniform
rotation withduces
core temperature
oninitial
denSection
6 we compare simulations covering a range of di↵erent
angularvelocity⌦0 =radiative
0.2cs0 /r0non-ideal' 1.22 ⇥MHD
10 14 simulation
rad s 1 . Thissity
obtainedin theofrin and resolutions. In particular, we vary rin between 0.25 au
corresponds
to a ratioand(2010,
thermalprotostarformation
bybetween
Kunz &rotational
Mouschovias
seeenergy
theirand 2 au and increase the ✓ resolution by up to a factor of 4.
13 (Barof '0.011,
again
with
a typicalnNH
core
figures
4a and4b),consistentat least up
to densities
cm 3 ,
n &3 10ranco& point
Goodman1998). temperature
To avoid
excessively
highspecificatwhich
the central
becomes
K andaveraged quantities at di↵erent epochs from our fiducial 3D simulation. Each column plots three
Figure 2.
Radialprofiles&100
of azimuthally
epochscorresponding
to a row in
Fig. 1. To facilitate comparison, earlier epochs are also plotted in later panels using thin lines. 13
angularmomentum
the to
outer
part
ofEoS
ourfordomainapproaches
7/5. Weinplan
refinethisfuture(wherework.
2.5 Boundary conditions
r/r0 can be at most ⇠10), we set the angular velocity to[kyr 1 ]

103

r 3/2

102

101

Keplerian out to 10au

100 0
10

101

102

r [au]
Figure 5. Rotation profile of the disc at the end of the fiducial 3D
simulation.

mass flux [g s 1 ]1021

Figure 3. Snapshots of midplane density in the fiducial 3D simulation. The first four panels correspond to the last four epochs
shown in Fig. 1. The initially axisymmetric first hydrostatic torus
develops spiral arms due to gravitational instability. These spirals
transport angular momentum outward, spreading the disc and aiding protostar accretion. The disc size oscillates as spirals emerge
and disperse.

0.125

14

0.100

Rd [au]M [M ]

16

protostar
disc
total

0.150

12

1020
1019
1018

1016
1026

1028

1029

1041

1039

pseudo-discenvelope
outflowmagnetic
gravitational
inwardoutward

1038
1037

0.050

8

1035
1026

0.025

6

質量と角運動量
1. どちらも pseudo-disk から
2. Outflow の寄与は小さい
2. はMachida&Hosokawa13や

1040

1036

0.000

1027

t [kyr]100.075

半径 20au の球殻を通過するの inflow が支配的

1017

Lz flux [g cm2 s 2 ]“Protostar” = mass at r < 1au

pseudo-discenvelope
outflow
inwardoutward

1027

1028

1029

Tomida+19の結論と異なる。
より現実的な磁場散逸を取り入れた結果 toroidal field が弱くなったためとしている。

t [kyr]

4
1027

1028

t [kyr]

1029

1027

1028

1029

t [kyr]Figure 4. Time evolution of protostar mass, disc mass and disc sizein the fiducial simulation.

Figure 6. Contribution to mass and angular-momentum flux
through a 20-au sphere from di↵erent mechanisms. Both mass flux
and angular-momentum flux are dominated by inflow (accretion)from the pseudo-disc.

14

protostar-disc system, which is shown in Fig. 8. We see a rel-

Toomre Q

Lz flux [g cm2 s 2 ]r [au]Lz [g cm2 s 1 ]radial force density [dyn cm 3 ]n at M10 = 0.01 M is close to the M –Lz relation of
Figure 9. Angular-momentum flux due to di↵erent mechanisms
protostar-disc system for M . 0.05 M , but the slope of
10 W. Xu and M. W. Kunz
through spherical shells
di↵erent radii,
averaged for t >andatevolution
9
protostar-disc
M –Lz relation becomes less steep at laterProtostellar disc formation
質量・角運動量はどこから来るか
角運動量の再分配
1028 kyr. The axisymmetric hydrodynamic flux (grey) corresponds, suggestingmagnetic
that magnetic
continues
to decrease
braking hasbraking
reduced angularmomentumby about1041
to advection
with the mass accretion and outflow and does not con1.5
an
order
of
magnitude.
14
lar momentum
after
protostarformation.
gravity (negative)
10
tributetotheradial transport of angularmomentum with respecthydro-axisymmetricmagnetic
We can also look at when and how
magneticbraking hapthermalpressure
1.0
is also worth
pointingout
whymagneticbrakingcanbehydro-nonaxisymmetricgravitational
to the gas. Inside the disc, the transport of angular momentum
pens. In Fig. 8 we also plot the Mmagnetic(< r)–Lz (< r) relationpressure
10 16
ortant in theis always
unimportant
inthewith0.5respect to the gas is dominated by gravitational instability,
atpseudo-disc
M10 = 0.01 M but, whichis around
the
epoch
of
protostarmagnetic tension
Protostellar disc formation
formation.
This
curve
has
a
slope
of
⇡2,
consistent
with
that
18
which produces the non-axisymmetric hydrodynamic flux (blue)
ostellar disc.
Thisismainlybecausethedensityandcolcentrifugal
100.0
a flattened core with uniform column
density and rotaall support
and gravitational stress (green).
density in oftheprotostellar disc are muchhigherthan in
Vorobyov & Basu 2007; also seetion. This suggests that the bulk of the pre-stellar core has
10 20
0.5
1994; Tomida et al. 2017).
pseudo-discundergone(for example,
Figsof 1magnetic
and 2),
which
makes
a similarsee
amountbraking.
This
is reaWe also comment that, historica
sonableof magnetic
brakingtakes place atimthe
22
magnetic 10field
lesssince
wellmost
coupledand less
dynamically
1.0
ity is often considered to regulate
beginningof the collapse phase; once the collapse becomes
Basu97のsemiangular-momentum transport within the disc. ature
The bysumof radiative coolant. The distinction
between a dense, non-magnetizedbalancing
dynamical (with near free-fall velocity), braking is generally
1.5 0
10 24
10
101
102
these two terms first increases then decreases backzero,
spiral to
shocksand turbulence. Th
and a thin,slower
magnetizedpseudo-dischasandalso
beenobthan the modelと類似
timescaleof collapse
angularmomenanalyticr [au]
first
proposedby Paczynski (1978)
tumisapproximately
conserved.
The
M
(<
r)–L
(<
r)
re26
z
suggestingthat
gravitational
instability
takes
angularmoed, for instance,
in Masson
et al.2 (2016). 3
10
1
100
104
model due to the adoption of a b
lationat M1010 = 0.01 M 10is close to 10
the M –Lz relationofmentumfrom the inner
the disc
and deposits
it in theFigure9. Angular-momentum
flux partdue to ofdi↵erentmechanisms[au] M . 0.05 M , but the slope ofgravitational
instability controls dthe protostar-disc systemr forthroughspherical
shellsatdi↵erent
radii,
averagedfort>outer part of the disc. This transport of angularspreading
momentum(angular-momentum tra
the protostar-disc M –Lz relation becomes less steep at later
kyr. Theaxisymmetrichydrodynamic
flux (grey)
corresponds
Figure 7. Radial force density from gravity and various supporting
Figure 8.1028
Relation
betweenthe total mass
M and angularmomenthey should
equally important.
time,suggesting
thatfrommagneticbraking
continueswhen
to decrease
leads
to
disc
spreading,
and
is
the
main
thatbealto
advection
with
the
mass
accretionandoutflowanddoes
conforcesatthemidplane,taken
ourfiducial
simulation
tum
L
of
the
protostar-discsystem(blueline).
HereM
canbe not mechanismz
VOLUTIONangular
OF DISC
SIZE
AND
MASS
state,
therateof heating and ang
momentumafter
protostarformation.
tributetoprotostar
theforradialof angularwith and
respect
M10 = 0.01 M . Althoughsomesupporting
forcesare of the sameconsidered
as a proxytime.transport
Above
⇠0.1
M when, wemomentum
see aoutflow
power-lawlowsaccretion
magneticbrakingsince both are pr
arecomparable
It
is
also
worth
pointingout
whymagneticbraking
canbeto thegas.of Inside
the reference,
disc, the transportof theangularmomentumorder of magnitude as gravity, the sum of them (grey line) remains
scaling withslope
⇡1.5. For
we also plot
relation
he previous
sectiondiscussed
the
mass
and
angularthe
disc
are
weak.
which gravitational instability ext
importantinwe
the formation
pseudo-discbut
is always
to
gasat
istwo
dominated
gravitational
instability,
insufficient
to balance
gravityin the pseudo-disc,and unimportantthepseudobetween withMin(<respectr) andLzthe(<r)
di↵erentby
epochs:the initial
Protostellar
discand
evolution
9 in thetial rotationdisc(Gammie 2001).
producestheandnon-axisymmetrichydrodynamic
flux
disc.
This is mainly because
the density
and coldiscbudget
is inprotostellar
near free-fall.
conditionwhich(solidgrey importantline)
when MM , whichis (blue)mentumof theprotostar-discsystemas a whole.
10 = 0.01
One
consequenceof
thisgravitational
Figure 10. Toomre Q (computed using azimuthally averaged proandgravitational
stress(green).
aroundprotostarformation(dashedgreyline).
umn density in the protostellar disc are much higher than infiles)
in the fiducial
3D simulation.
Each
curve
corresponds
to a
we will focus
on the redistribution of angular momenspreadingis thatthe disccannowhave significantly
higher1052
the pseudo-disc (for example, see Figs 1 and 2), which makes
panel in Fig. 3. The gravitationally unstable disc generally has Q
5.3
Predicting
disc evolution
within the theprotostar-discsystem,which
determinesthespecific
angularmomentum than the pseudo-disc inflow thatmagnetic field less well coupled and less dynamically imbetween
1 and
4 2.
tension
(see
Fig.
7).
Althoughtheinitialmass-to-flux
ratio
Magnetic
breaking
重力不安定性によるself-regulation
angular-momentum
transport within
the disc.
sum of Fig.
The mass.
distinction
between disca dense,non-magnetizedUsing
the idea aof gravitational seution fieldof discsizeresolution)and(atportant.
35-au
in a protostellar
survey.
There
feeds the protostar-discsystem.
ForTheexample,5 showsof the pre-stellar
is justsupercriticaland Bz /⌃ to zero,these two coreterms
first slightly
increases
then decreases
backa thin, magnetizedalso been obtal mass M and angular momentu
are alsodisc
someanddi↵erencesin the initialpseudo-discconditions,has
includingsignificant
in ⌦ instability
betweenthe angulardisc
pseudo-disc;does notsuggesting
decreasesignificantly
duringmostof thetakes
infall,the andthat jump
gravitational
mo- the system,
served,instance,near-criticalin Masson etal. (2016).
one can obtain a robust es
our adoptionof fora stronger,
initialmagnetic field,
3.0
pressurementum
gradientandmagnetic
forcesremainatleastafactorsimilar
jumps
in
v
are
also
visibleinthemiddle
right
panel
from
the
inner
part
of
the
disc
and
deposits
it
in
the
sity
profile,mass, and size of thewhich could
leadto
substantiallyweakeroutflowassuggested
15of a few outersmallerthanthetransport
acceleration
is &50%momentumof1051
Angularmomentumtransport
by gravitational
partof2.gravity,
theThisdisc.andThisof angularof the possible
disc is known. Here we per
2.5
by resultsfrom Machida& Hosokawa(2013). Overall,instability
it is
of
Fig.
makes
havingrelatively
large
discs
free fall leadsin thetobulk
of spreading,pseudo-disc.
Still,
thismainis notmechanism
to say that that aldisc
andis theisothermal disc (with sound speed
importantto verify in futurestudiesour result is
5 EVOLUTIONOF DISC
SIZEwhether
AND
MASS
even
whenarethepseudo-disc
low thespecific
angularmomen10 forcesthese retarding
unimportant.
On thehas
contrary,
time evolution
afterapplicablethe formation
of
the
first
hydrostatic
core/torus,
lows
protostaraccretionwhen
outflowandmagneticbraking
Since we know most of the disc sh
2.0
over a broader range of initial parameters. Finally,
formationoftheself-similarcolumndensity
profilerequiresainitialcondition
thealthoughpreviousobservations
section we discussed
thehighmass
andandangulardue
magnetic braking.
intum
the disc
areto
weak.
we may assume that the whole discnote Inthat,of relatively
specific
densewetorus
becomesgravitationallyunstable,
reprotostarsystemformation
small but One
nontrivial
ratio between these forces and self gravmomentum
budgetof theprotostar-discasa whole.
disc
5 important consequence of this gravitational 1.5
simplicity, we also assume
angularmomentuminoutflows(e.g.,
Bjerkeli
et
al.
2016)
are
p that the(M10 = 0.01M ) Gravitans so for therest
offocus
our fiducial
simulation.
ity (similar
to theisrolepressure
in thecollapse higher
Nowusedwe as
will
the outflows
redistributionof angularmomenspreading
thatofthe
disc cannow sphericalhave significantly
M
GM/R3 . T
? ⇡ M and (R) ⇡sometimes
evidenceonthatcanremovea
sig1.5Lz M
model ofspecificLarsonangular1969), momentum
and this setsa characteristic
accre50
10tum
withintheprotostar-discsystem,which
determinestheal instability
plays
animportant
role
for
disc
growth,Spiralshockによるディ
thanthepseudo-disc
inflow
that
1.0
of
the
disc
is
then
nificant amount of angular momentum, this argument is nottion ratefeeds
thatthedepends
only 1027.5
on system.
the isothermal
sound
speed
1026.51027.01028.01028.51029.02revolutionof discsize
and mass.
10
10 1outflowand
protostar-disc
For example,
Fig.
5 shows a
5.2
Gravitational
self-regulation
groundedfirmlyunlessthe
mass-loss
rate
by
and thetransports
angularmomentumoutwardspreads
c
GM
t
[kyr]
cs and Gsignificant(for a given
mass-to-flux
ratio).
The
mass
accretions0 効果は、本
スク加熱の
M
[M
]
jump in ⌦ between the disc and the pseudo-disc;⌃(R) ⇡.
of thedisc arealso known.
disc. accretion
To seerate
this
e↵ect
clearly,
we plot the contriburate calculated
fromtheself-similar
model
of
Basu
(1997)
is⇡GQ
R30
similarjumps
in v 1are also
visible in oftheamiddle
right panel thin isothermalThe
gravitational
stability
geometrically
For the remainder of this section, we focus only on the
35
シミュレーションにはAngularmomentumtransport
by gravitational
instability
13cs /G of⇠ 2Fig.
⇥ 102. This
M makesyr , similartoour result.
ng
Figure8.5.1
Relation
betweenthemass M andangular momenfrompseudo-discdi↵erentstresses
to totalthehavingrelatively
largeevolution
discs possible
Figure 11.
Similarto Fig.
10 butshows timeof theinflow(accretion),
as itangular-momentum
largely determines the flux
Thisdefined⌃ / R 3/2
disc
can
be
described
bytheToomreQparameter,
asprofile is also obse
entum Lz of the protostar-disc system (blue line). Here M can be
even
when
the
pseudo-dischas
lowspecific
angularmomenToomre
Qprofile.
The
disc
size
is
also
plotted(whitecurve)
for
含まれていない。
Soon
after
the
formation
of
the
first
hydrostatic
core/torus,
mass
andangular-momentum
budgetoftheprotostar-disc
regulateddisc evolution of Voroby
ugh spheres
ata proxydi↵erent
duringdiscinmeconsidered as
for time.radii
Above ⇠0.1
M , wesee a expansion
power-lawtum
due to magnetic braking.
reference.
the
dense
becomes
unstable,re- Angularsystem.
Theslope
exclusion
thereference,othermechanisms,
which
are and 4.3
size Rd and disc inner boundary R
momentum budget: magnetic brakingnsscalingwithof torus⇡1.5.ofFor
also plot
9. Theangular-momentum
fluxgravitationally
iswethesumtheofrelationhydrodymains
forLzgreatly
therestof ourfiducial
simulation.
oftenhighlythe problem.
angular momentum of the disc are
Q
=
1
–
2
の領域は非軸対称
obetween
M (<variable,r)soand(< r)
atsimplifies
twodi↵erent
epochs:the initialGravitac
s
ic advection(Reynoldsv =),0.01
magneticbraking
For the angular-momentum
budget, we focus on the rela15
r10
Q⌘.
instability
plays
an⇢v
importantrole
discconditiontional(solidgrey line) stress
and
when
M
M ,forwhich
is growth,
Z Rd (3)Thethesteep
dependenceof the angular-momentum
trans⇡G⌃
4cs0 Mtion
between
total
mass
and
angularmomentumofthe5.2
Gravitational
self-regulation
重力不安定性が発現aroundprotostarformation(dashedgreyline).
asit
transports
angularmomentumoutwardandspreadsxwell stress Br B /4⇡), and gravitational torque (gravMd ⇡2⇡RdR ⌃(R) ⇡port rate on disc column density (through Q) implies a self1/
Gd pros to ahas Qof thee) forransa selfleadsng (astable
rease
o-discining
ll butld belies a
d 11,
tween
milar
(e.g.,

5.3 Predicting disc evolution
Using the idea of gravitational self-regulation, for given total mass M and angular momentum Lz of the protostar-discsystem, one can obtain a robust estimate for the surface density profile, mass, and size of the disc if the thermal profileof the disc is known. Here we perform this estimate for an
isothermal disc (with sound speed cs0 ) as an example.
Since we know most of the disc should be marginally stable,
we may assume that the whole disc has constant Q = Q0 . For
simplicity, we also assume
p that the disc is not very massive, so
M? ⇡ M and (R) ⇡ GM/R3 . The surface density profileof the disc is thenrcs0GM⌃(R) ⇡
.
(4)
⇡GQ0
R3
This ⌃ / R 3/2 profile is also observed in the gravitationallyregulated disc evolution of Vorobyov & Basu (2007). For discsize Rd and disc inner boundary Rd,in ⌧ Rd , the mass andangular momentum of the disc are then
Z Rd
1/2
4cs0 M 1/2 Rd
Md ⇡2⇡RdR ⌃(R) ⇡,(5)
G1/2 Q00
Z Rdp2cs0 M Rd
Ld ⇡2⇡RdR ⌃(R) GM R ⇡
.
(6)
Q0
0
Since most angular momentum of the protostar-disc system
is within the disc, we have Ld ⇡ Lz , which gives
Lz Q 0,2cs0 M✓◆1/2
8cs0 Lz
Md ⇡
.
GQ0
Rd ⇡

(7)
(8)

Using the M –Lz relation of our fiducial simulation (extrapolated assuming a power-law slope of 1.5), the above estimate
gives disc mass ⇡0.15 M and size ⇡80 au when the protostar
reaches ⇡1 M .
In reality, a marginally gravitationally unstable disc is of-16results, and show that properly treating the inner boundary condition is crucial for achieving convergence at an acceptable
The Perseus
computationalcost. ALMA Chemistry Survey (PEACHES). I. The Complex Organic Molecules in Perseus
338#37 Embedded Protostars

The Perseus
ALMA Chemistry Survey (PEACHES). I. The Complex Organic Molecules
2021 ApJ…910…20Y
in Perseus Embedded Protostars

Yao-Lun Yang, Nami Sakai, Yichen Zhang, Nadia M. Murillo, Ziwei E. Zhang, Aya E. Higuchi, Shaoshan
Zeng, Ana López-Sepulcre, Satoshi Yamamoto, Bertrand Lefloch, Mathilde Bouvier, Cecilia Ceccarelli, Tomoya
Hirota, Muneaki Imai, Yoko Oya, Takeshi Sakai, Yoshimasa Watanabe F To date, about two dozen low-massembedded protostars exhibit rich spectra with complex organic molecule (COM) lines. These protostars seem to possess
diﬀerent enrichment in COMs. However, the statistics of COM abundance in low-mass protostars are limited by the scarcity ofobservations. This study introduces the Perseus ALMA Chemistry Survey (PEACHES), which aims at unbiasedly characterizing
the chemistry of COMs toward the embedded (Class 0/I) protostars in the Perseus molecular cloud. Among the 50 embedded
protostars surveyed, 58protostars have CH3 OH, CH3 OCHO, and N-bearing COMs, respectively. The detectability of COMs
depends on neither the averaged continuum brightness temperature, a proxy of the H2 column density, nor the bolometricluminosity and the bolometric temperature. For the protostars with detected COMs, CH3 OH has a tight correlation with
CH3 CN, spanning more than two orders of magnitude in column densities normalized by the continuum brightness temperature,suggesting a chemical relation between CH3 OH and CH3 CN and large chemical diversity among the PEACHES samples at thesame time. A similar trend with more scatter is also found between all identified COMs, hinting at common chemistry forthe sources with COMs. The correlation between COMs is insensitive to the protostellar properties, such as the bolometricluminosity and the bolometric temperature. The abundance of larger COMs (CH3 OCHO and CH3 OCH3 ) relative to thatof smaller COMs (CH3 OH and CH3 CN) increases with the inferred gas column density, hinting at an eﬃcient production ofcomplex species in denser envelopes.

Low-metallicity Young Clusters in the Outer Galaxy. III. Sh 2-127

Chikako Yasui, Naoto Kobayashi, Masao Saito, Natsuko Izumi, Warren Skidmore F In deep near-infrared imaging of
the low-metallicity ([O/H] = 0.7 dex) H II region Sh 2-127 (S127) with Subaru/MOIRCS, we detected two young clusters with
413 members (S127A) in a slightly extended H II region and another with 338 members (S127B) in a compact H II region. The
limiting magnitude was K = 21.3 mag (10 ), corresponding to a mass detection limit of ⇠0.2 M . These clusters are an orderof magnitude larger than previously studied young low-metallicity clusters and larger than the majority of solar neighborhood
young clusters. Fits to the K-band luminosity functions indicate very young cluster ages of 0.5 Myr for S127A and 0.1-0.5 Myr
for S127B, consistent with the large extinction (up to AV ' 20 mag) from thick molecular clouds and the presence of a compact
H II region and class I source candidates, and suggest that the initial mass function (IMF) of the low-metallicity clusters is
indistinguishable from typical solar neighborhood IMFs. Disk fractions of 28% ± 3% for S127A and 40% ± 4% for S127B are
significantly lower than those of similarly aged solar neighborhood clusters (⇠50%-60%). The disk fraction for S127B is 17
higherthan those of previously studied low-metallicity clusters (<30 %), probably due to S127B’s age. This suggests that a large

substantial. The synthetic spectra were calculated for the size ofcontinuum emission (Table 2) to include the effect of beam
dilution. We allowed the line width to vary between 1.2 and3.5 km s−1 for a better ﬁtting quality and chose a range of thecolumn density for each molecule according to the strength ofthe emission.
The ﬁtting function in XCLASS includes several optimization
algorithms that can be used in sequence to reduce biases. We
conﬁgured the algorithm chain that starts with the genetic
algorithm followed by the Levenberg–Marquardt χ2 minimization. The genetic algorithm iteratively searches for the best-ﬁtting
parameters with generations that evolve like a natural selection,

where the better ﬁtting models are modiﬁed less over generations.
We set up the genetic algorithm to search for the top two bestﬁtting models with 30 generations. Then, the Levenberg–
Marquardt χ2 minimization was applied to the two best-ﬁtting
models for 20 iterations to determine the best-ﬁtting models. The
genetic algorithm aims to ﬁnd global minima, and the Levenberg–
Marquardt χ2 minimization further ﬁnds the best-ﬁtting models in
the global minimums. The two best-ﬁtting models found by thegenetic algorithm are often very similar, suggesting that there is
only one minimum. In the rare cases of two distinct global
minima, we selected the model with the lower χ2 values from thetwo best-ﬁtting models constrained by the Levenberg–Marquardt

19

Figure 6. The column densities of Per-emb 44 (SVS 13 A) ﬁt using the hybrid algorithm, genetic and Levenberg–Marquardt χ2 minimization, and the MCMC
algorithm. The molecular species are annotated, and the gray line indicates the equality.

we used 500 walkers to better sample the parameter space. The
ﬁtting column densities from the hybrid (genetic and Levenberg–
Marquardt χ2 minimizations) method are consistent with that ﬁt

with the MCMC optimization (Figure 6). Moreover, our hybrid
method yields comparable uncertainties to those from the MCMC
method, validating our modeling approach.

20

4. Results
4.1. Detection Statistics
We summarize the detection statistics in Figure 7, which
includes COMs, carbon-chain molecules, and simple molecules, such as CS, H13CN, SO, 34SO, and SO2. To derive thedetection fraction, we included the three sources that were
excluded from modeling because their source velocity was notreliable (Section 3.1), making a total of 50 sources. The COMs

discussed here are derived from the spectra taken from the areasurrounding the continuum peak, as explained in Section 3.1.
We focus on the chemistry of COMs toward the disk-forming
regions in this study. Therefore this study excludes any
molecules that are only detected outside the region ofcontinuum emission. A comparison of the chemical composition in protostellar envelopes (100–1000 au) requires observations with a larger maximum recoverable scale (θMRS).
The PEACHES sample shows a diverse chemistry of COMs,
from no molecule detected (e.g., B1-b N and L1455 IRS 2) torich spectra of COMs (e.g., Per-emb 13). Most protostars of the
PEACHES sample have simple molecules, such as SO, CS,
H13CN, and HDCO, and ∼60% of the sources have SO2 and34
SO. Emission of C2H can be easily identiﬁed from thespectra. However, C2H toward the continuum sources often
shows irregular line proﬁles together with velocity offsets andabsorption (Appendix E, Figure 21). Warm environments, suchas the outﬂow cavity wall, easily enhance the abundance ofC2H because of the elevated abundance of C+ (e.g., Zhang
et al. 2018; Imai et al. 2019). Thus, C2H emission tends toextend along with the outﬂow cavities, which means that 1D
spectra toward the continuum source are no longer representative of the abundance of C2H. Future studies will investigate
the comparison between unsaturated species and COMs using
follow-up observations with a higher dynamic range of thespatial distribution (from ∼100 au to several 1000 au scale).
Previous observations with single-dish telescopes have a beam
size of ∼3000 au (e.g., Higuchi et al. 2018). Thus, observations
that ﬁll in the spatial scale from several 100 au to 1000 au arenecessary.
Several sources have broad SiO emission, signiﬁcantly
contaminating the emission of CH3CH2CN and CH3CHO. In
comparison, the emission of COMs has relatively narrow line
width of a few km s−1. In the quantitative discussion below,
we exclude the spectral windows that are contaminated by the
SiO emission. For a few sources, such as Per-emb 26, we canstill distinguish the emission of CH3CH2CN and CH3CHO
from the broad SiO emission and include these identiﬁcation in
the detection statistics.
Figure 7 shows the number of COMs detected toward thePEACHES sample. Twenty-nine (58%) sources have detected
COMs. CH3OH is detected in 28 sources (56%), CH3OCHO is
detected in 14 sources (28%), and N-bearing COMs aredetected in 20 sources (40%). Compared to the COMs in the
CALYPSO survey (Belloche et al. 2020), the fraction ofsources that have methanol, ∼50%, is similar to that for the
PEACHES sample. Furthermore, 30% of the CALYPSO
sources have at least three species of COMs, while 28% of
the PEACHES sample have at least three species of COMs. In asmaller sample selected from several molecular clouds, vanGelder et al. (2020) found that 3 of 7 protostars show emissionof COMs.
To investigate the impact of protostellar properties on theappearance of COMs, we divided the PEACHES sample into
four groups according to the number of detected COMs from no
COM, at least 1, 5, and 10 species of COMs (Figure 8). The
compared protostellar properties include the bolometric luminosity (Lbol), the bolometric temperature (Tbol), and the meancontinuum brightness temperature (Tb,cont), where the Tbol and
Lbol are collected from Tobin et al. (2016) and Murillo et al.
(2016). Tbol traces the evolutionary stage of protostars withhigher Tbol for more evolved protostars (Myers & Ladd 1993;

21

Figure 16. The column densities of CH3OH toward the PEACHES sample (excluding Per-emb 17) normalized by the averaged continuum brightness temperature as afunction of the minimum distance to the 10σ contour in the Planck 217 GHz observations. Due to the lower brightness at 217 GHz toward Per-emb 5 and Per-emb 25,
their minimum distances are calculated with 5σ and 3σ contours, respectively. The solid circles show the sources with detection of CH3OH, while the open circles
show the sources without detection of CH3OH.

Figure 17. The ratios of the COMs commonly found in the PEACHES and the CALYPSO survey (Belloche et al. 2020), as well as a few selected hot corinos. The
CALYPSO sources are categorized into three groups according to their COM abundance to CH3OH. Group 3 has no detection of CH2OHCHO. The column densities
of COMs toward IRAS 16293−2422 B are taken from Jørgensen et al. (2016) for CH3OH, Jørgensen et al. (2018) for C2H5OH, CH3OCH3, CH3OCHO, CH3CHO,
NH2CHO, and CH2OHCHO, and Calcutt et al. (2018) for CH3CN and CH3CH2CN. The column densities of COMs toward IRAS 16293−2422 A are taken from
Manigand et al. (2020) for CH3OH, C2H5OH, CH3OCH3, CH3OCHO, CH3CHO, NH2CHO, and CH2OHCHO, and Calcutt et al. (2018) for CH3CN and CH3CH2CN.
The column densities of COMs toward L483 are taken from Jacobsen et al. (2019). The column densities of COMs toward BHR 71 are taken from Yang et al. (2020),
assuming a Tex of 100 K for CH3OH. The purple hatched regions indicate the range of the peak abundance ratios to CH3OH from the warm-up models in Garrod(2013) that have three different speeds of the warm-up process.

abundance ratios in the warm-up models by Garrod (2013).
The models underestimate the abundance ratios for most of
the O-bearing COMs, while the ratios of CH2OHCHO andN-bearing COMs agree with the ratios derived from thePEACHES sample (Figure 17). Because we may underestimate
the column density of CH3OH due to the optical depth, theabundance ratio to CH3OH could be overestimated. If the ratios
derived from the PEACHES sample were lower by a factor offor the ﬁve COMs in the left, but the model would thenoverestimate the abundance ratios of CH2OHCHO, NH2CHO,
and CH3CH2CN. This disagreement may point to a dominant
gas-phase formation of (some) COMs other than CH3OH, assuggested by other studies (e.g., Balucani et al. 2015; Skouteris
et al. 2018; Vazart et al. 2020).
The well-correlated abundance of N-bearing COMs andO-bearing COMs (Figures 12 and 13) indicates no “N-/O-bearing
COMs differentiation,” which has long been discussed for23

luminosity and the bolometric temperature. The abundance of larger COMs (CH3 OCHO and CH3 OCH3 ) relative to thatof smaller COMs (CH3 OH and CH3 CN) increases with the inferred gas column density, hinting at an eﬃcient production of
Low-metallicity Young Clusters in the Outer Galaxy. III. Sh 2-127
338
complex species in denser envelopes.

#38

2021 AJ…161..139Y

Low-metallicity Young Clusters in the Outer Galaxy. III. Sh 2-127

Chikako Yasui, Naoto Kobayashi, Masao Saito, Natsuko Izumi, Warren Skidmore F In deep near-infrared imaging of
the low-metallicity ([O/H] = 0.7 dex) H II region Sh 2-127 (S127) with Subaru/MOIRCS, we detected two young clusters with
413 members (S127A) in a slightly extended H II region and another with 338 members (S127B) in a compact H II region. The
limiting magnitude was K = 21.3 mag (10 ), corresponding to a mass detection limit of ⇠0.2 M . These clusters are an orderof magnitude larger than previously studied young low-metallicity clusters and larger than the majority of solar neighborhood
young clusters. Fits to the K-band luminosity functions indicate very young cluster ages of 0.5 Myr for S127A and 0.1-0.5 Myr
for S127B, consistent with the large extinction (up to AV ' 20 mag) from thick molecular clouds and the presence of a compact
H II region and class I source candidates, and suggest that the initial mass function (IMF) of the low-metallicity clusters is
indistinguishable from typical solar neighborhood IMFs. Disk fractions of 28% ± 3% for S127A and 40% ± 4% for S127B are
significantly lower than those of similarly aged solar neighborhood clusters (⇠50%-60%). The disk fraction for S127B is higher
than those of previously studied low-metallicity clusters (<30 %), probably due to S127B’s age. This suggests that a large
fraction of very young stars in low-metallicity environments have disks, but the disks are lost on a very short timescale. These
results are consistent with our previous studies of low-metallicity star-forming regions, suggesting that a solar neighborhood
IMF and low disk fraction are typical characteristics for low-metallicity regions, regardless of cluster scales.

1324338 The Hawaii Infrared Parallax Program. V. New T-Dwarf Members and Candidate#39 Members of Nearby Young Moving Groups
The Hawaii
Infrared Parallax Program. V. New T-Dwarf Members and Candidate
2021 ApJ…911….7Z
Members of Nearby Young Moving Groups

Zhoujian Zhang, Michael C. Liu, William M. J. Best, Trent J. Dupuy, Robert J. Siverd F We present a search
for new planetary-mass members of nearby young moving groups (YMGs) using astrometry for 694 T and Y dwarfs, including
447 objects with parallaxes, mostly produced by recent large parallax programs from UKIRT and Spitzer. Using the BANYAN⌃ and LACEwING algorithms, we identify 30 new candidate YMG members, with spectral types of T0 T9 and distances of
10 43 pc. Some candidates have unusually red colors and/or faint absolute magnitudes compared to field dwarfs with similar
spectral types, providing supporting evidence for their youth, including 4 early-T dwarfs. We establish one of these, the variableT1.5 dwarf 2MASS J21392676+0220226, as a new planetary-mass member (14.6+3.2
1.6 MJup ) of the Carina-Near group (200 ± 50Myr) based on its full six-dimensional kinematics, including a new parallax measurement from CFHT. The high-amplitude
variability of this object is suggestive of a young age, given the coexistence of variability and youth seen in previously known
YMG T dwarfs. Our four latest-type (T8 T9) YMG candidates, WISE J031624.35+430709.1, ULAS J130217.21+130851.2,
WISEPC J225540.74 311841.8, and WISE J233226.49 432510.6, if confirmed, will be the first free-floating planets (⇡ 2 6
MJup ) whose ages and luminosities are compatible with both hot-start and cold-start evolutionary models, and thus overlap theproperties of the directly-imaged planet 51 Eri b. Several of our early/mid-T candidates have peculiar near-infrared spectra,
indicative of heterogenous photospheres or unresolved binarity. Radial velocity measurements needed for final membershipassessment for most of our candidates await upcoming 20 30 meter class telescopes. In addition, we compile all 15 known
T7 Y1 benchmarks and derive a homogeneous set of their eﬀective temperatures, surface gravities, radii, and masses.

26