AST 1501 presentation
1 Nov 2005
Pawel Artymowicz
University of Toronto,
UTSC and St. George
1. Structures vs. reasons (planets,…)
2. Dust, avalanches, fIR, gas, and the classification of disks
3. HD 141569A as an example
4. Non-axisymmetric features without planets
New edgeon disk
NICMOS/
HST
(Schneider
et al 2005)
STIS/Hubble imaging
(Heap et al 2000)
Modeling
(Artymowicz,unpubl.):
parametric, axisymmetric disk
cometary dust phase function
Optical
thickness 
Dust
density
Possible mini-projects like this!
Radius r [AU]
Height z [AU]
The danger of overinterpretation of structure
Are the PLANETS responsible for EVERYTHING
we see? Are they in EVERY system?
Or are they like the Ptolemy’s epicycles,
added each time we need to explain a new
observation?
Also, do we really need a new type of particle
for every bandpass [optical, sub-mm]?
FEATURES in disks: (9)
ORIGIN: (10)
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm)
￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps
￭
spirals, quasi-spirals￭
tails, extensions
￭
￭ instrumental artifacts,
variable PSF, noise,
deconvolution etc.
￭ background/foreground obj.
￭ planets (gravity)
￭ stellar companions, flybys
￭ dust migration in gas
￭ dust blowout, avalanches
￭ episodic release of dust
￭ ISM (interstellar wind)
￭ stellar UV, wind, magnetism
￭ collective eff. (selfgravity
radiative instab.)
(Most features
additionally depend on
the viewing angle)
FEATURES in disks:
ORIGIN:
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm)
￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps, incl. disks ￭
spirals, quasi-spirals￭
tails, extensions
￭
￭ instrumental artifacts,
variable PSF, noise,
deconvolution etc.
FEATURES in disks:
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm)
￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps
￭
spirals, quasi-spirals￭
tails, extensions
￭
ORIGIN:
￭ background or
foreground objects
AB Aur : disk
or no disk?
Fukugawa et al. (2004)
another “Pleiades”-type star
no disk
?
Source: P. Kalas
AU Microscopii
& a less inclined cousin
This is a coincidentally(!) aligned
background galaxy
.
FEATURES in disks:
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm)
￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps
￭
spirals, quasi-spirals￭
tails, extensions
￭
ORIGIN:
￭ planets (gravity)
Some models of structure in dusty disks rely on too limited
a physics: ideally one needs to follow: full spatial distribution,
velocity distribution, and size distribution of a collisional system
subject to various external forces like radiation and gas drag -that’s very tough to do! Resultant planet depends on all this.
Beta = 0.01
(monodisp.)
Dangers of fitting
planets to individual
frames/observations:
Vega has 0, 1, or 2 blobs,
depending on bandpass.
What about its planets?
Are they
wavelength-dependent too?
HD 141569A is a Herbig emission star
>2 x solar mass, >10 x solar luminosity,
Emission lines of H are double, because
they come from a rotating inner gas disk.
CO gas has also been found at r = 90 AU.
Observations by Hubble Space
Telescope (NICMOS near-IR camera).
Age ~ 5 Myr,
a transitional disk
Gap-opening PLANET ?
So far out??
R_gap ~350AU
dR ~ 0.1 R_gap
Hubble Space Telescope/ NICMOS infrared camera
HD 14169A disk gap confirmed by new observations
(HST/ACS)
HD141569+BC in V band
HST/ACS Clampin et al.
HD141569A deprojected
Why & how do birds migrate?
Bird Migrations
To cope with changing seasons, most birds migrate, few hibernate.
In high arctic regions (northern Alaska, northern Canada, and Greenland), the entire
population of birds often consists of migratory birds (they stay for summer only).
In the forest and open country of United States, over 80% of the nesting land birds
are migratory. However, on the Pacific Coast, more species are non-migratory;
in tropical regions at least 80% of the birds are non-migratory.
In the Rockies and Sierras of the West, migration often consists of moving from the high
to low elevations. Rosy Finches, Townsend's Solitaires, and Mountain Quail perform
these movements quite regularly whereas others, such as Clark's Nutcracker, are much
more erratic. The annual fall migration of the Townsend's Solitaire may consist merely of
descending a few thousand feet from a high mountain forest to the shelter of a wooded
valley. Some migration schedules do not always closely follow seasonal changes in the
weather. For example, since the vegetative food supply of nomadic species such as the
crossbills, redpolls, and Pine Grosbeaks fluctuates in abundance from year to year,
these birds migrate in some winters and not in others. In contrast, insect-eating birds
such as warblers, vireos, and flycatchers that live in the far north have no choice but to
migrate. Their migration therefore tends to involve long distances and regular timing.
Planetary Migrations
Do planets migrate? How? How fast?
Are bird & planet migrations similar?
Do they migrate long-way or locally?
Do they migrate regularly or erraticaly?
Do planets migrate alone or in flocks?
Where and how do they stop migrating?
Migration Type I :
embedded in fluid
Migration Type II :
more in the open (gap)
Migration Type I :
embedded in fluid
Migration Type II :
in the open (gap)
Migration Type III
partially open (gap)
Type III
Outward migration of protoplanets to ~100AU
or
outward migration of dust to form rings and
spirals
required to explain the structure in
transitional (5-10 Myr old) dust disks
and perheps also the
(12-20Myr old) Beta Pictoris-type disks
DISK-PLANET interaction
and migration, including outward
migration
It used to be just type I and II...
now we study a new mode of migration:
type III
Migration:
type 0
type I
type II & IIb
type III
Interaction:
Gas drag +
Radiation press.
Timescale of
migration:
from ~1e2 yr to disk lifetime
(~1e7 yr)
Resonant
excitation of
waves (LR)
> 1e4 yr
Tidal excitation
of waves (LR)
> 1e5 yr
Corotational
flows (CR)
> 1e2 - 1e3 yr
……………………………………………………………………….
N-body
Gravity
> 1e5 yr (?)
Planets were thought to always shepherd planets…or was it
the other way around?
Pan opens Encke gap in A-ring of Saturn
Shepherding by
Prometheus and Pandora
A gap-opening body in a disk:
Saturn rings, Keeler gap region (width =35 km)
This new 7-km satellite of Saturn was announced in May 2005.
To Saturn
Prometheus (Cassini view)
(Mini-project!
Rings as a laboratory
to study possible
type III migration?)
Variable-resolution
PPM (Piecewise
Parabolic Method)
[Artymowicz 1999]
Jupiter-mass planet,
fixed orbit a=1, e=0.
White oval = Roche
lobe, radius r_L= 0.07
Corotational region out
to x_CR = 0.17 from
the planet
disk
gap
(CR region)
disk
Type I -- III migration
Figure From: “Protostars and Planets IV (2000)”;
Artymowicz (this talk).
Simulation of a Jupiter-class planet in a constant surface density disk
with soundspeed = 0.05 times Keplerian speed.
PPM = Piecewise Parab.
Method
Artymowicz (2000),
resolution 400 x 400
Although this is clearly
a type-II situation (gap
opens), the migration
rate is NOT that of the
standard type-II, which
is the viscous accretion
speed of the nebula.
Consider a one-sided disk (inner disk only). The rapid inward migration is
OPPOSITE to the expectation based on shepherding (Lindblad resonances).
Like in the well-known problem of “sinking satellites” (small satellite galaxies
merging with the target disk galaxies),
Corotational torques cause rapid inward sinking.
(Gas is trasferred from orbits inside the perturber to the outside.
To conserve angular momentum, satellite moves in.)
Now consider the opposite case of an inner hole in the disk.
Unlike in the shepherding case, the planet rapidly migrates outwards.
Here, the situation is an inward-outward reflection of the sinking satellite problem.
Disk gas traveling on hairpin (half-horeseshoe) orbits fills the inner void and moves the planet
out rapidly (type III outward migration). Lindblad resonances produce spiral waves and try to
move the planet in, but lose with CR torques.
xCR
NO MIGRATION:

In this frame, comoving with the planet,
gas has no systematic radial velocity
V = 0, r = a = semi-major axis of
orbit.
0
a
disk
r
Symmetric horseshoe orbits,
torque ~ 0

Librating Corotational (CR) region
protoplanet
Librating Hill sphere (Roche lobe) region
xCR = half-width of CR region,
separatrix distance
SLOW MIGRATION:

In this frame, comoving with the planet,
gas has a systematic radial velocity
V = - da/dt = -(planet migr.speed)
0
asymmetric horseshoe orbits,
torque ~ da/dt
a

r

FAST MIGRATION:
CR flow on one side of the planet,
disk flow on the other
0
Surface densities in the CR region and the
disk are, in general, different.
a
r
Tadpole orbits, maximum torque
Saturn-mass protoplanet in a solar nebula disk
(1.5 times the Minimum Nebula,
PPM, Artymowicz 2003)
Azimuthal
angle (0-360 deg)
Type III outward
migration
Condition for FAST migration:
disk mass (in CR region)
similar to planet mass.
Notice a carrot-shaped bubble of
“vacuum” behind the planet.
Consisting of material trapped
1
2
in librating orbits, it produces
radius
CR torques smaller than the matrial
in front of the planet. The net CR torque powers fast migration.
3
Summary of type-III migration
 Extremely rapid (timescale < 1000 years). CRs >> LRs, disks
do not shepherd planets. Requires sufficient disk density
 Direction depends on prior history, not just on disk properties.
 Supersedes a much slower, standard type-II migration (&type I ??)
 Migration stops on disk features (rings, edges and/or substantial
density gradients.) Such edges seem natural (dead zone
boundaries, magnetospheric inner disk cavities, formationcaused radial disk structure)
 Offers possibility of survival of giant exoplanets at intermediate
distances (0.1 - 1 AU),
 ...and of terrestrial planets during the passage of a giant planet on
its way to the star (last Mohican scenario)
 STRUCTURE in OUTER REGIONS of dusty transitional &
debris disks
Next Steps: Toward a better LR/CR perturbation theory
 Previous perturbation theories started from
circular unperturbed orbits [those do not exist]
 and assumed infinitesimal perturbations (Fourier
decomposition allowed) [not always!]
 Alternative way: unperturbed state adjusted for
perturbation. Trajectories of all essential types
(disk orbits, corotational hairpin/horeseshoes,
closed orbits around planet)
 On that set of unperturbed flow lines, 1st order
perturbation should give a better approximation
 Migration and additional effects can be
incorporated
Guiding center
trajectories in Hill
problem
Unit of distance =
Hill sphere
Unit of da/dt =
Hill sphere radius
per dynamical time
Animation by Eduardo Delgado
Examples of simple
orbital sets obtained
from the simplification
of Hill’s equations of
motion.
FEATURES in disks:
ORIGIN:
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm)
￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps
￭
spirals, quasi-spirals￭
tails, extensions
￭
￭ stellar companions,
flybys
Stellar flyby (of an elliptic-obit companion) explains some features
of HD 141569A
Augereau and Papaloizou (2003)
Application of the idea to Beta Pictoris less certain...
Quillen et al. (2004)
HD 141569A
Ardila et al (2005)
H/r = 0.05
Flyby+planetesimals --> dust production & outflow
H/r =0.1
LTE
= 4
Mgas = 50 ME
Ardila et al (2005)
No planet
Flyby+plane+planetesimals
5 MJ, e=0.6 planet
 =4
H/r = 0.1
Mgas = 50 ME
Best model, Ardila et al (2005)
5 MJ, e=0.6, a=100 AU
planet
Beta = 4
H/r = 0.1
Mgas = 50 M
HD 141569A
Room for improvement in theory
(more than a
mini-project?)
Wyatt (2005) - planetesimal evolution under seclural
perturbation from an eccentric planet,
initial time evolution of pericenter glow.
1. No gas drag 2. No dust 3. Planet
acts on gas disk to produce spiral waves (in gas
and dust) at Lindblad resonances.
Ardila et al (2005)
1. Sharp outer edge at 1*pericenter distance of flyby *
2. No pre-existing dust in disk, only the dust produced
after perturbation (no time for that?)
3. Single beta value,
4. No dust-dust collisions or avalanches
FEATURES in disks:
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm)
￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps
￭
spirals, quasi-spirals￭
tails, extensions
￭
ORIGIN:
￭ dust migration in gas
Planetary systems:
stages of decreasing dustiness
In the protoplanetary disks (tau)
dust follows gas.
Sharp features due to associated
companions: stars, brown dwarfs and planets.
1 Myr
These optically thin transitional disks (tau <1)
must have some gas even if it's hard to detect.
5 Myr
Warning: Dust starts to move w.r.t. gas!
Look for outer rings, inner rings, gaps
with or without planets.
Pictoris
These replenished dust disk
are optically thin (tau<<1)
and have very little gas.
Sub-planetary & planetary bodies can be detected via spectroscopy,
spatial distribution of dust, but do not normally expect sharp features.
12-20
Myr
Extensive modeling including dust-dust collisions and radiation pressure needed
Migration:
Type 0
 Dusty disks: structure
from gas-dust coupling
(Takeuchi & Artymowicz 2001)
 theory will help
determine gas distribution
Predicted dust
distribution:
axisymmetric ring
Gas disk tapers
off here
Weak/no PAH emission
Neutral (grey)
scattering from
s> grains
Size spectrum
of dust has lower cutoff
Repels ISM dust
Disks = Nature, not
nurture!
Radiative blow-out of grains
(-meteoroids, gamma meteoroids)
Instabilities
(in   1 disks)
Radiation pressure
on dust grains in disks
Dust
avalanches
Quasi-spiral
structure
Orbits of stable meteoroids elliptical
Color
effects
Enhanced erosion;
shortened dust lifetime
Dust migrates,
forms axisymmetric
rings, gaps
(in disks with gas)
Short disk lifetime
Age paradox
FEATURES in disks:
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm)
￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps
￭
spirals, quasi-spirals￭
tails, extensions
￭
ORIGIN:
￭ dust blowout
avalanches,
￭ episodic/local dust
release
Weak/no PAH emission
Neutral (grey)
scattering from
s> grains
Size spectrum
of dust has lower cutoff
Repels ISM dust
Disks = Nature, not
nurture!
Radiative blow-out of grains
(-meteoroids, gamma meteoroids)
Instabilities
(in   1 disks)
Radiation pressure
on dust grains in disks
Dust
avalanches
Limit
on fir
Quasi-spiral
structure
Orbits of stable meteoroids elliptical
in gas-free
disks
Enhanced erosion;
shortened dust lifetime
Color
effects
Dust migrates,
forms axisymmetric
rings, gaps
(in disks with gas)
Short disk lifetime
Age paradox
How radiation pressure induces large eccentricity:

= F_rad / F_grav
Dust Avalanche
(Artymowicz 1997)
Process powered by the energy of stellar radiation
N ~ exp (optical thickness of the disk * <#debris/collision>)
N
= disk particle, alpha meteoroid (  < 0.5)
= sub-blowout debris, beta meteoroid ( > 0.5)
f IR     2rdr /(4r )     dr /(2r ) *Qabs(V)*Qabs(IR)
2
   ( / s ) dr  (r / z )    dr /(2r )
so
  (r / z ) f IR / (Qabs(V)*Qabs(IR))
For instance, in HD141569A, a prototype transitional disk
  (0.1)  0.018  0.2/Qabs
1
N  ~ 10 2
(midplane optical thickness)
(number of sub-blowout debris per collision)
dN   N   N
N / N 0  exp( N  ) ~ exp( 20) ~ 106
Transitional disks MUST CONTAIN GAS or face self-destruction.
Beta Pic is almost the most dusty, gas-poor disk, possible.
Bimodal histogram
of fractional
IR luminosity fIR
predicted by disk
avalanche process
source: Inseok Song (2004)
ISO/ISOPHOT data on dustiness vs. time
-1.8
Dominik, Decin, Waters, Waelkens (2003)
uncorrected ages
ISOPHOT ages, dot size ~ quality of age
fd of beta Pic
corrected ages
ISOPHOT + IRAS
OK!
Gas-free modeling
leads to a paradox
==> gas required
or
Age paradox!
episodic dust
production
fIR =fd
disk dustiness
transitional systems
5-10 Myr age
Weak/no PAH emission
Neutral (grey)
scattering from
s> grains
Size spectrum
of dust has lower cutoff
Repels ISM dust
Disks = Nature, not
nurture!
Radiative blow-out of grains
(-meteoroids, gamma meteoroids)
Instabilities
(in   1 disks)
Radiation pressure
on dust grains in disks
Dust
avalanches
Limit
on fIR
Quasi-spiral
structure
Orbits of stable meteoroids are
elliptical
in gas-free
disks
Enhanced erosion;
shortened dust lifetime
Color
effects
Dust migrates,
forms axisymmetric
rings, gaps
(in disks with gas)
Short disk lifetime
Age paradox
Grigorieva, Artymowicz and Thebault (to be subm. to A&A 2005)
Comprehensive model of dusty debris disk (3D) with full treatment
of collisions and particle dynamics.
￭ especially suitable to denser transitional disks supporting dust avalanches
￭ detailed treatment of grain-grain colisions, depending on material
￭ detailed treatment of radiation pressure and optics, depending on material
￭ localized dust injection (e.g., planetesimal collision)
￭ dust grains of similar properties and orbits
grouped in “superparticles”
￭ physics: radiation pressure, gas drag,
collisions
Results:
￭ beta Pictoris avalanches multiply debris x(3-5)
￭ spiral shape of the avalanche - a robust outcome
￭ strong dependence on material properties
and certain other model assumptions
Model of (simplified) collisional avalanche with substantial
gas drag, corresponding to 10 Earth masses of gas in disk
Main results of modeling
of collisional avalanches:
1. Strongly nonaxisymmetric,
growing patterns
2. Substantial exponential
multiplication
3. Morphology depends on the
amount and distribution of gas,
in particular on the presence of
an outer initial disk edge
FEATURES in disks:
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm.) ￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps
￭
spirals, quasi-spirals￭
tails, extensions
￭
ORIGIN:
￭ ISM (interstellar wind:
gas + dust bombardment)
Artymowicz & Clampin (1997)
FEATURES in disks:
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm)
￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps
￭
spirals, quasi-spirals￭
tails, extensions
￭
ORIGIN:
￭ stellar influence:
photoevaporation,
wind, magnetism
Weak/no PAH emission
Neutral (grey)
scattering from
s> grains
Size spectrum
of dust has lower cutoff
Repels ISM dust
Disks = Nature, not
nurture!
Radiative blow-out of grains
(-meteoroids, gamma meteoroids)
Instabilities
(in   1 disks)
Radiation pressure
on dust grains in disks
Dust
avalanches
Quasi-spiral
structure
Orbits of stable meteoroids elliptical
Color
effects
Enhanced erosion;
shortened dust lifetime
Dust migrates,
forms axisymmetric
rings, gaps
(in disks with gas)
Short disk lifetime
Age paradox
FEATURES in disks:
ORIGIN:
blobs, clumps
￭
streaks, feathers
￭
rings (axisymm)
￭
rings (off-centered) ￭
inner/outer edges ￭
disk gaps
￭
warps
￭
spirals, quasi-spirals￭
tails, extensions
￭
￭ collective effects
(e.g., disk selfgravity,
radiative instability)
In gas+dust disks which are optically thick in the radial direction
there may be an interesting set of instabilities. Radiation pressure
on a coupled gas+dust system that has a spiral density wave with
wave numbers (k,m/r), is analogous in phase and sign to the force
or self-gravity. The instability is thus pseudo-gravitational in nature
and can be obtained from a WKB local analysis.
Forces of
selfgravity
Forces of radiation pressure in the
inertial frame
Forces of rad. pressure relative
to those on the center of the arm
In gas+dust disks which are optically thick in the radial direction
there may be an interesting set of instabilities. Radiation pressure
on a coupled gas+dust system that has a spiral density wave with
wave numbers (k,m/r), is analogous in phase and sign to the force
or self-gravity. The instability is thus pseudo-gravitational in nature
and can be obtained from a WKB local analysis.
   0 exp(    dr )  0 exp(  )
 0  effective coefficient for coupled gas+dust
 0 ~ 0.1....10
   0  1 ei ( kr m t )

   0 (r ) 
r
1
ik
ei ( kr m t )
(this profile results from
dust migration)
  Step function of r or constant
1 i ( kr m t )
   0 (r ) 
e
ik
2
f rad   K r  0 e
f self  gravity
 0
i1 i ( kr m t )
(1 
e
)
k
i 4G1 i ( kr m  t )

e
k
(WKB)
 2   4 G  Poisson eq.
   f f  f1 exp(...)   f1  4 G 1
 ikf1  4 G 1
4 G 1
f1  i
k
  Step function of r or constant
1 i ( kr m t )
   0 (r ) 
e
ik
2
f rad   K r  0 e
f self  gravity
 0
i1 i ( kr m t )
(1 
e
)
k
i 4G1 i ( kr m  t )

e
k
(WKB)
 2   4 G  Poisson eq.
   f f  f1 exp(...)   f1  4 G 1
 ikf1  4 G 1
4 G 1
f1  i
k
G
Q 
;
Q 1  1  ( grav.) instability
 orb cs
G 1
1
Q 
 2  0 e  ( r )0 r 
 orb cs
G 1
 ( r )
1
Effective Q number
Q 
 2  0e
(r 0 / r )
(radiation+selfgravity)
 orb cs
1
0
0
 1
0 
 1
1
r
Analogies with gravitational instability ==> similar structures (?)
Not only planets but also
Gas + dust + radiation =>
non-axisymmetric features
including regular m=1
spirals, conical sectors, and
multi-armed wavelets
FEATURES in disks:(9)
ORIGIN:(10)
48 POSSIBLE CONNECTIONS !
blobs, clumps
￭(5)
streaks, feathers
￭(4)
rings (axisymm)
￭(2)
rings (off-centered) ￭(7)
inner/outer edges ￭(5)
disk gaps
￭(4)
warps
￭(7)
spirals, quasi-spirals￭(8)
tails, extensions
￭(6)
￭ instrumental artifacts,
variable PSF, noise,
deconvolution etc.
￭ background/foreground obj.
￭ planets (gravity)
￭ stellar companions, flybys
￭ dust migration in gas
￭ dust blowout, avalanches
￭ episodic release of dust
￭ ISM (interstellar wind)
￭ stellar wind, magnetism
￭ collective eff. (self-gravity)
While observing: don't try to prove one theory
(like, that there MUST BE planets in your still
poorly-observed disk. They may be there,
but making such a claim requires good evidence.)
While modeling: take good care!
Don't claim success easily. Your model does NOT fit all
the data. Include all relevant physics/dynamics.
Use multi-wavelength sets of data to dramatically
improve uniqueness of the model.
THE END