Transcript M p - University of Rochester

Sculpting
Circumstellar
Disks
Alice Quillen
University of Rochester
Netherlands
April 2007
Motivations
• Planet detection via disk/planet interactions –
Complimentary to radial velocity and transit detection
methods
• Rosy future – ground and space platforms
• Testable models – via predictions for forthcoming
observations.
• New dynamical regimes and scenarios compared to
old Solar system
• Evolution of planets, planetesimals and disks
Collaborators: Peter Faber, Richard Edgar, Peggy
Varniere, Jaehong Park, Allesandro Morbidelli , also
Eric Blackman, Adam Frank, Pasha Hosseinbor, Amanda LaPage
Observational Background
•
Submillimeter imaging
Young Clusters: 5--20% of stars
surveyed in young clusters are TTauri stars hosting disks with large
clearings – Dozen or so now with
IRS/Spitzer spectra, more identifiedFomalhau
t
with Spitzer/IRAC photometry
Older Disks and Debris Disks:
Optical scattered
light
•
Fraction detected with disks with IR
excess depends on age, wavelength
surveyed and detection limit. 50—100
now known from Spitzer/MIPS
surveys
•
HD 100546
HR4796A
Unexplained structure: edges,
clearings, spiral arms, warps,
clumps
Beta
Pictorus
HD 141569A
Credits, ESO,Schneider Wilner, Grady, Clampin, HST, Kalas
Dynamical Regimes for Circumstellar
Disks with central clearings
1. Young gas rich
accretion disks –
“transitional disks” e.g.,
CoKuTau/4.
Planet is massive
enough to open a gap
(spiral density waves).
Hydrodynamics is
appropriate for
modeling.
Dynamical Regimes– continued
2. Old dusty diffuse debris disks – dust
collision timescale is very long; e.g.,
Zodiacal cloud.
Collisionless dynamics with radiation
pressure, drag forces, resonant trapping,
removal of orbit crossing particles
3. Intermediate opacity dusty disks – dust
collision timescale is in regime 103-104
orbital periods; e.g., Fomalhaut, AU Mic
debris disks
This Talk
i.
What mass objects are required to
account for the observed clearings,
What masses are ruled out?
Planets in accretion disks
with clearings
-- CoKuTau/4
ii.
Planets in Debris disks
with clearings
-- Fomalhaut
iii.
Embryos in Debris disks
without clearings
-- AU Mic
iv.
Total mass in planets
in older systems
-- Clearing by planetary systems
Transition Disks
Estimate of minimum
planet mass to open a
gap requires an estimate
of disk viscosity.
CoKuTau/4
D’Alessio
et al. 05
4 AU
10 AU
Wavelength μm
Disk viscosity estimate
either based on clearing
timescale or using study
of accretion disks.
Mp > 0.1MJ
Models for Disks with Clearings
1. Photo-ionization models (Clarke, Alexander)
Problems:
-- clearings around brown dwarfs, e.g., L316, Muzerolle et al.
-- accreting systems such as DM Tau, D’Alessio et al.
-- wide gaps such as GM Aur; Calvet et al.
-- single temperature edges
Predictions: Hole size with time and stellar UV luminosity
2. Planet formation, gap opening
followed by clearing (Quillen, Varniere)
-- more versatile than photo-ionization
models but also more complex
Problems: Failure to predict dust
density contrast, 3D structure
Predictions: Planet masses required to
hold up disk edges, and clearing
timescales, detectable edge structure
Minimum Gap Opening Planet
In an Accretion Disk
heat from
accretion
heat from
stellar
radiation
 =0.01, M *  M
Gapless disks lack
planets
Park et al. 07 in preparation
Minimum Gap Opening Planet Mass
in an Accretion Disk
Different mass stars
M  M *2
 =0.01
Planet
trap?
Smaller
planets can
open gaps in
selfshadowed
disks
Hole radii scale
with stellar mass
(Kim et al. in prep)
Retired A stars lack
Hot Jupiters?
(Johnson et al. 07)
Fomalhaut
’s
• eccentric
steep edge
profile
h /r ring
~ 0.013
z
• eccentric
e=0.11
• semi-major axis
a=133AU
• collision
timescale =1000
orbits based
on measured
opacity at 24
microns
• age 200 Myr
• orbital period
Free and
forced
eccentricity
e sin v
radii give you
eccentricity
efree
v
free
eforced
v
forced
e cosv
If free eccentricity is
zero then the object has
the same eccentricity as
the forced one
v longitude of pericenter
Pericenter glow model
• Collisions cause orbits to be near closed ones.
Small free eccentricities.
• The eccentricity of the ring is the same as
the forced eccentricity
e forced
b3/2 2 ( )
 1
e planet
b3/ 2 ( )
a

ap
• We require the edge of the disk to be truncated by
the planet 
 ~ 1  ering  e forced  e planet
• We consider models where eccentricity of ring and
ring edge are both caused by the planet. Contrast
with precessing ring models.
Disk dynamical boundaries
• For spiral density waves to be driven into a disk
(work by Espresate and Lissauer)
Collision time must be shorter than libration time
 Spiral density waves are not efficiently driven by a
planet into Fomalhaut’s disk
A different dynamical boundary is required
• We consider accounting for the disk edge with the
chaotic zone near corotation where there is a large
change in dynamics
• We require the removal timescale in the zone to
exceed the collisional timescale.
Chaotic zone boundary
  N 
N
D



and removal within
a  a  t
collisionless lifetime
removal
Neptun
e size
Saturn
size
What mass
planet will
clear out
objects
inside the
chaos zone
fast
enough
that
collisions
will not
fill it in?
Chaotic zone boundaries for particles
with zero free eccentricity
Hamiltonian at a first order mean motion resonance
H (  ; ,  )  a 2  b  c  d 1/ 21/p 2 cos( -  p )
 g 0 1/ 2 cos(   )  g11/p 2 cos(   p )
Poincare variables
 ~ e2 ,
 only depends on a
c

regular
resonance
 5/ 2b3/1 2
4
g 0    2 5/ 4 f 31
d

corotation
 5/ 2b3/2 2
2
g1    2 5/ 4 f 27
With secular terms only there is a fixed point at
   p , 1/ 2
secular terms
b3/2 2 1/ 2
 1  p that is the e free  0 orbit
b3/ 2
Dynamics at low free eccentricity
Expand about the fixed point (the zero free
eccentricity orbit)
H (  ; I ,  )  a 2  b  cI
same as for zero  g I 1/ 2 cos(
0
eccentricity planet
goes to zero near
the planet
  )  ( g 01/f 2  g11/p 2 ) cos(   p )
For particle eccentricity equal to the forced
eccentricity and low free eccentricity, the
corotation resonance cancels
 recover the 2/7 law, chaotic zone same
width
Dynamics at low free eccentricity is
similar to that at low eccentricity near
a planet in a circular orbit
width of
chaotic zone
different eccentricity
points
1.5 2/ 7
planet mass
No difference in
chaotic zone width,
particle lifetimes, disk
edge velocity
dispersion low e
compared to low efree
Velocity dispersion in the disk edge
and an upper limit on Planet mass
• Distance to disk edge
set by width of chaos
zone
2/ 7
da ~ 1.5
• Last resonance that
doesn’t overlap the
corotation zone
affects velocity
dispersion in the disk
edge
• Mp < Saturn
ue ~  3/ 7
cleared out by
perturbations
from the
planet
Mp > Neptune
Assume that the edge of the
ring is the boundary of the
chaotic zone. Planet can’t
be too massive otherwise
the edge of the ring would
thicken  Mp < Saturn
nearly closed
orbits due to
collisions
eccentricity
of ring equal
to that of the
planet
First Predictions for a planet just
interior to Fomalhaut’s eccentric ring
•
•
•
•
Neptune < Mp < Saturn
Semi-major axis 120 AU (16’’ from star)
Eccentricity ep=0.1, same as ring
Longitude of periastron same as the ring
The Role of Collisions
• Dominik & Decin 03 and Wyatt 05
emphasized that for most debris disks the
collision timescale is shorter than the PR drag
timescale
• Collision timescale related to observables
tcol ~  n  
1
where  n is normal optical depth
The number of collisions per orbit N c ~ 18 n
2r
n ~
f IR where f IR is fraction stellar light
dr
re-emitted in infrared
The numerical problem
• Between collisions particle is only under
the force of gravity (and possibly
radiation pressure, PR force, etc)
• Collision timescale is many orbits for the
regime of debris disks: 100-10000
orbits.
Numerical approaches
• Particles receive velocity perturbations at random
times and with random sizes independent of particle
distribution (Espresante & Lissauer)
• Particles receive velocity perturbations but dependent
on particle distribution (Melita & Woolfson 98)
• Collisions are computed when two particles approach
each other (Charnoz et al. 01)
• Collisions are computed when two particles are in the
same grid cell – only elastic collisions considered
(Lithwick & Chiang 06)
Our Numerical Approach
Perturbations independent of particle distribution:
• Espresate set the vr to zero during collisions. Energy
damped to circular orbits, angular momentum
conservation. However diffusion is not possible.
• We adopt
vr  0
v  v  v
• Diffusion allowed but angular momentum is not
conserved!
• Particles approaching the planet and are too far away
are removed and regenerated
• Most computation time spent resolving disk edge
Parameters of 2D simulations
N c collision rate, collisions per particle per orbit
- related to optical depth
 v tangential velocity perturbation size
- related to disk thickness
 planet mass ratio
- unknown that we would like to
constrain from observations
radius
Morphology of
collisional disks
near planets
  105 , N c  102 , e  0.02
  104 , N c  103 , e  0.01
radius
• Featureless for low mass
planets, high collision
rates and velocity
dispersions
• Particles removed at
resonances in cold,
diffuse disks near massive
planets
angle
Profile shapes
chaotic zone
boundary 1.5 μ2/7
  105
  104
  106
Rescaled by distance to
chaotic zone boundary
Chaotic zone
probably has a role
in setting a length
scale but does not
completely
determine the profile
shape
Density
decrement
• Log of ratio of density near planet
to that outside chaotic zone edge
• Scales with powers of simulation
parameters as expected from
exponential model
Reasonable well fit with the function
  
log10   0.12  0.23log10  6 
 10 
 Nc 
  dv 
 0.1log10  2   0.45log10 

 10 
 0.01 
Unfortunately this does not predict a
nice form for tremove
To truncate a disk a
planet must have
mass above
 n 
log10   6  0.43log10 
3 
 5  10 
 u / vK 
 1.95 

 0.07 
(here related to
observables)
Log Planet mass
Using the numerical
measured fit
Log Velocity dispersion
Observables can lead to planet
mass estimates, motivation for
better imaging leading to better
estimates for the disk opacity
and thickness
• Upper mass limit
confirmed by lack of
resonant structure
• Lower mass limit ~
lower than previous
estimate unless the
velocity dispersion
at the disk edge set
by planet
Log Planet mass
Application to Fomalhaut
Quillen 2006, MNRAS, 372, L14
Quillen & Faber 2006, MNRAS, 373, 1245
Quillen 2007, astro-ph/0701304
Log Velocity dispersion
Constraints on Planetary
Embryos in
Debris Disks
AU Mic JHKL
Fitzgerald, Kalas, & Graham
h/r<0.02
• Thickness tells us the velocity dispersion in dust
• This effects efficiency of collisional cascade resulting in dust
production
• Thickness from gravitational stirring by massive bodies in the disk
The size distribution and
collision cascade
observed
Figure from Wyatt
& Dent 2002
set by age of
system scaling
from dust opacity
constrained by
gravitational
stirring
Scaling from the dust:
1 q
 a 
d ln N
 N (a )  N d  
d ln a
 ad 
3 q
 a 
 (a)   d  
 ad 
(multiply by a 2 )
As tcol ~  
1
13
1 q
3
 a   u 
tcol  tcol ,d    * 
 ad   2QD 
Set tcol  tage and solve for a
2
The top of the
cascade
Gravitational stirring
In sheer dominated regime
2
 
1 d i
~ s2 s
 dt
i
where  s 
s 
Solve: i (t )  t 4
s
mass density ratio
M *r 2
ms
M*
mass ratio
Comparing size distribution at top
of collision cascade to that
required by gravitational stirring
size distribution
might be flatter
than 3.5 – more
mass in high end
 runaway
growth?
top of
cascade
Earth
Comparison
between 3 disks
with resolved
vertical structure
108yr
107yr
107yr
Clearing by Planetary Systems
• Assume planet formation leaves behind a
population of planetesimals which produce
dust via collisions
• Central clearings lacking dust imply that all
planetesimals have been removed
• Planets are close enough that interplanetary
space is unstable across the lifetime of the
system
• ~ 50 known debris disks well fit by single
temperature SEDs implying truncated edges
(Chen et al. 06)
Clearing by Planetary Systems
Log10 time(yr)
Instability happens at
 t
 
 1/ 4 ~ 9.5  log10  7
5 
 10 yr 10 
Mp
where planets of mass ratio  =
M*
are spaced ri 1  1    ri
 ~0.5
r
N
2 1    ~ max
rmin
Separation
Chambers et al. 96
Faber et al. in
preparation
Clearing by planetary systems
2 1    ~
N
rmax
rmin
rmin set by ice line
rmax set by observed disk temperature
Result is we solve for N and find 3-8 planets required
of Neptune size for most debris disks.
This implies a total minimum mass in planets of about
a Saturn mass
Summary
• Quantitative ties between
disk structure and planets
residing in disks
• Better understanding of
collisional regime and its relation to observables
• In gapless disks, planets can be ruled out – but we
find preliminary evidence for embryos and runaway
growth
• The total mass in planets in most systems is likely to
be high, at least a Saturn mass
• More numerical and theoretical work inspired by
these preliminary crude numerical studies
• Exciting future in theory, numerics and observations
5km/s for a
planet at
10AU
PV plot
Prospects
with
ALMA
Edgar’s simulations
Diffusive approximations
  N  Nf (r )
D

r  r  tremoval
planet
  dv 
u2
where D ~
~
N

c
tcol n 2
v
 K 
2
Consider various models for removal of particles by the planet
f (r )  1  N (r )  elr
f (r )  1  r  N (r ) is an Airy function
f (r )  e  r  N (r ) is a modified Bessel function
All have exponential solutions near the planet
with inverse scale length
1/ 2
l ~ tremove
 2 / 7 N c1/ 2 dv1 and unknown function tremove