Download Protoplanetary Disk Models

Inga Kamp
The role of the central star for
the structure of protoplanetary
disks
Peter Woitke, Wing-Fai Thi, Ian Tilling (all Edinburgh)
Protoplanetary Disks and Planet Formation
Dust dynamics are controlled by gas, even at late times !
• dust substructure
• dust coagulation
• dust settling
[HR4796: NICMOS 1.6 mm
Schneider et al. 1999]
[Reverse ∂p/∂r at 70 AU
Klahr & Lin 2001]
8 Myr
mJy/pixel
[Barge & Sommeria 1995,
Johansen et al. 2006, 2007]
Outline
• Various planet forming environments (BD - T Tauri - Herbig)
• Second generation disk structure modeling
• How do observations inform us about the star-disk
interaction ?
Outline
• Various planet forming environments (BD - T Tauri - Herbig)
• Second generation disk structure modeling
• How do observations inform us about the star-disk
interaction ?
Outline
• Various planet forming environments (BD - T Tauri - Herbig)
• Second generation disk structure modeling
• How do observations inform us about the star-disk
interaction ?
Protoplanetary Disk Models
Key astrophysical questions:
• What is the environment in which planets form ?
disk structure - boundary conditions for planet formation
• Can planets form around stars different from our Sun ?
star-disk interaction - impact on disk structure and dispersal
Outline
• Various planet forming environments (BD - T Tauri - Herbig)
• Second generation disk structure modeling
• How do observations inform us about the star-disk
interaction ?
Disk properties
Dust in BD disks evolves on same time scale as T Tauri disks
[2 BDs @ 2 and 10 Myr: Sterzik et al. 2004]
Inner disk clearing occurs much faster in T Tauri disks
[8mm excess @ 5 Myr: Carpenter et al. 2006]
Dust models fit BD, T Tauri and Herbig SEDs in the same way
indicating the first steps of planetesimal formation in all cases
[Pascucci et al. 2003, Allers et al. 2006, Bouy et al. 2008]
some BD SEDs require flaring disks, others flat… low # statistics
Disk properties
[Kessler-Silacci et al. 2007]
[Pascucci et al. 2009]
lower HCN/H2C2 ratios in the disks around brown dwarfs
weaker and more processed silicate features
Outline
• Various planet forming environments (BD - T Tauri - Herbig)
• Second generation disk structure modeling
• How do observations inform us about the star-disk
interaction ?
Protoplanetary Disk Models
physical
structure
radiative
transfer
chemical
composition
comparison
with observation
Protoplanetary Disk Models
physical
structure
surface density
dust temperature T = T0 (r/r0)-q
disk scale height
radiative
transfer
chemical
composition
 = 0 (r/r0)-p
-e
H = Hcomparison
(r/r
)
0
0
with observation
[e.g. Menard et al.; Pinte et al. 2006]
Protoplanetary Disk Models
physical
structure
surface density
 = 0 (r/r0)-p
dust temperature 2D continuum RT
2 r3/GM )0.5
disk
scale
height
H
=
(c
s
*
radiative
comparison
transfer
with observation
chemical
composition
[e.g. D’Alessio et al. 1998, Dullemond et al. 2002]
Gas chemical modeling on top of fixed density structure
[e.g. Aikawa et al. 2002, Semenov et al. 2008]
Protoplanetary Disk Models
physical
structure
hot flaring surface
line radiative
transfer
chemical
composition
comparison
with observation
molecule freeze-out
rich molecular chemistry
atomic/ionized
[e.g. Aikawa et al. 2002, Kamp & Dullemond 2004]
[PPV chapters: Dullemond et al. 2007, Bergin et al. 2007]
Protoplanetary Disk Models
ProDiMo
physical
structure
radiative
transfer
comparison
with observation
chemical
composition
[Woitke, Kamp, Thi 2009]
Similar approaches have been followed by other groups
[Nomura & Millar 2005, Gorti & Hollenbach 2004, 2008]
Chemistry
dn  
k n n  k n n   n   n
dt
i
ijk
jk
j
k
jik
jk
i
ij
k
j
ji
i
j
stationary solution with modified
Newton-Raphson algorithm
(switch to time-dependant)
65 species: H, H+, H-, H2, H2+, H3+
He, He+
C, C+, CO, CO+, CO2, CO2+, HCO, HCO+, H2CO
CH, CH+, CH2, CH2+, CH3, CH3+, CH4, CH4+, CH5+
O, O+, O2, O2+, OH, OH+
H2O, H2O+, H3O+
N, N+, NH, NH+, NH2, NH2+, NH3, NH3+, N2, HN2+
CN, CN+, HCN, HCN+, NO, NO+
Si, Si+, SiH, SiH+, SiO, SiO+, SiH2+, SiOH+
S, S+, Mg, Mg+, Fe, Fe+
--> plus CO, H2O, CO2, CH4, NH3 ice
Heating and cooling
 (T
k
k
Heating
Cooling
,ni )    k (Tgas,ni )
gas
k
photo-electric heating, PAH heating, viscous ( )
heating, cosmic ray, C photo-ionisation, coll. de excitation of H⋆2 , H2 on grains, H2 photodissociation,
IR background line heating
thermal accomodation on grains, OI, CII, CI finestructure cooling,
CO ro-vibrational (110 levels, 243 lines),
o/p H2O rotational (45/45 levels, 258/257 lines),
o/p H2 quadrupole (80/80 levels, 803/736 lines),
SiII, SII, FeII semi-forbidden (80 levels, 477 lines),
Ly  , MgII h&k, OI 6300Å
Heating and cooling
 (T
k
k
Heating
Cooling
,ni )    k (Tgas,ni )
gas
k
photo-electric heating, PAH heating, viscous ( )
heating, cosmic ray, C photo-ionisation, coll. de excitation of H⋆2 , H2 on grains, H2 photodissociation,
IR background line heating, (X-ray heating)
thermal accomodation on grains, OI, CII, CI finestructure cooling,
CO ro-vibrational (110 levels, 243 lines),
o/p H2O rotational (45/45 levels, 258/257 lines),
o/p H2 quadrupole (80/80 levels, 803/736 lines),
SiII, SII, FeII semi-forbidden (80 levels, 477 lines),
Ly  , MgII h&k, OI 6300Å
Protoplanetary Disk Models
interstellar UV
stellar UV
stellar X rays
soft inner and outer edges
[Hartmann et al. 1989, Hughes et al. 2008,
Woitke, Kamp & Thi 2009]
The Thermal Balance
Photoelectric heating of the gas
stellar photons
e-
small grains: a << l(e.g. PAHs)
yield ~ ??
e-
e-
e-
large grains: a >> l (micron sized)
yield ??
The Thermal Balance
Photoelectric heating of the gas
stellar photons
e-
small grains: a << l(e.g. PAHs)
yield ~ ??
e-
e-
e-
large grains: a >> l (micron sized)
yield ??
[Abbas et al. 2006]
Impact of UV Irradiation
HD104237
H2
OVI
H2, CO photodissociation
FUSE:
Apr 2000
May 2001
C ionization
STIS:
GHRS:
Oct 2001
Nov 1994
UV excess
Impact of UV Irradiation
time
10
50
80
500
5000 K
UV excess
Impact of UV Irradiation
time
-14
-10
-5
0 log n(OH)/n(tot)
Protoplanetary Disk Models
interstellar UV
stellar UV
stellar X rays
soft inner and outer edges
[Hartmann et al. 1989, Hughes et al. 2008,
Woitke, Kamp & Thi 2009]
The Thermal Balance
TW Hya
LX ~ 1030 erg/s
X-ray heating of the gas
[Kastner et al. 1997]
[Nomura et al. 2007]
Primary ionization
stellar X-rays
e124.0 Å
*
e-
1.24 Å
Secondary ionization
*
eAuger electrons
[Glassgold et al. 2004, Meijerink et al. 2008]
Protoplanetary Disk Models
Flaring disk structure “early Sun”
M* = 1 M
L* = 1 L
Teff = 5800 K
Mdisk = 10-2 M
0.5-500 AU
dust:
90% Mg2SiO4
10% Fe
0.1-10 mm (2.5)
[Woitke, Kamp, Thi 2009]
X-ray source @ 20 R to
irradiate disk midplane
Protoplanetary Disk Models
Flaring disk structure “Herbig star”
M* = 2.2 M
L* = 32 L
Teff = 8500 K
Mdisk = 10-2 M
0.5-500 AU
dust:
90% Mg2SiO4
10% Fe
0.1-10 mm (2.5)
T Tauri
Herbig Ae
Protoplanetary Disk Models
Herbig Ae
T Tauri
decreasing disk mass
Flaring disk structure “Herbig star”
M* = 2.2 M
L* = 32 L
Teff = 8500 K
Mdisk = 10-2, 10-3, 10-4 M
0.5-700 AU
dust:
90% Mg2SiO4
10% Fe
0.05-200 mm (3.5)
self-similarity
Outline
• Various planet forming environments (BD - T Tauri - Herbig)
• Second generation disk structure modeling
• How do observations inform us about the star-disk
interaction ?
Probes of UV Irradiation: [OI], OH
OH photodissociation layer:
OH + n  O* + H
O* denotes the 1D2 excited level (decay emits a 6300 Å photon)
460 AU
WFPC2
[Bally et al. 2000, Störzer & Hollenbach 1998,
Acke et al. 2005]
[Mandell et al. 2008, Salyk et al. 2008]
Probes of UV Irradiation: [OI], OH
OH photodissociation layer:
OH + n  O* + H
O* denotes the 1D2 excited level (decay emits a 6300 Å photon)
normalized flux
VLT/UVES: HD97048
velocity [km/s]
[Bally et al. 2000, Störzer & Hollenbach 1998,
Acke et al. 2005]
[Mandell et al. 2008, Salyk et al. 2008]
Probes of UV, X-ray Irradiation: CO, H2
Fluorescence excitation vs. thermal emission:
CO UV pumping of electronic states => cascade into high v levels
H2 UV (Lyman & Werner bands) or X-rays (collisions with fast e-)
[Brittain et al. 2007]
see talk by Gerrit van der Plas
[v=1-0 S(1), S(0), v=2-1 S(1) NIR:
Carmona et al. 2008]
Probes of UV, X-ray Irradiation: CO, H2
Fluorescence excitation vs. thermal emission:
CO UV pumping of electronic states => cascade into high v levels
H2 UV (Lyman & Werner bands) or X-rays (collisions with fast e-)
[Brittain et al. 2007]
see talk by Gerrit van der Plas
[S(1), S(2), S(4) pure rotational MIR
lines: Bitner et al. 2008]
Probes of X-ray Irradiation: NeII
see talk by Richard Alexander
[Glassgold et a. 2007, Pascucci et al. 2007, Lahuis et al. 2007, Meijerink et al. 2008]
Probes of gas-dust decoupling
H
D
E
0
50
100
R (AU)
D: 13CO 3-2
E: 13CO 1-0
[van Zadelhoff et al. 2001]
150
[Pietu et al. 2007]
C
200
C: 12CO 2-1
H: HCO+ 4-5
temperatures derived from gas
lines at the surface are higher
than dust temperatures
=> support for decoupling of gas
and dust in the surface layers
Probes of gas-dust decoupling
[OI] emission
MIR dust
emission
[Fedele et al. 2008]
dust is flat (self-shadowed) and
gas is flaring
=> support for decoupling of gas
and dust in the surface layers
GASPS: A Herschel Key program
GAS evolution in Protoplanetary Systems (PI: Dent)
~200 disks distributed
over spectral type, age,
and disk mass
[CII], [OI], CO and H2O
lines
Aim: probe the gas mass
evolution throughout the
planet forming phase
GASPS: A Herschel Key program
physical
structure
line radiative
transfer
chemical
composition
comparison
with observation
Monte Carlo line radiative transfer
[Hogerheijde, van der Tak 2000]
understand where the line originates
and then put resolution there !!
Kamp, Tilling, Woitke, Thi, Hogerheijde
Key points to take home
• Models predict that stellar UV and X-ray radiation shape the
disk structure and chemistry and we actually observe that !
• Individual gas lines DO NOT probe total disk mass !
but we keep trying with multiple tracers …
and we try hard with GASPS to find the desperately wanted 10-4 M disks…
• Gas lines DO probe the physical and chemical conditions in the
region where they arise (interplay between radiative transfer
and chemical structure)
Thank You !