Photoevaporation of Disks around Young Stars D. Hollenbach NASA Ames Research Center From Stars to Planets University of Florida April, 2007 Collaborators: F. Adams, K. Dullemond, U. Gorti, D. Johnstone, S. Lizano, F. Shu Recent reviews: Hollenbach et al (2000), Clarke (2002), Hollenbach & Adams (2004), Johnstone et al (2004), Yorke (2004), Hester & Desch (2005), Bally et al (2005), Richling et al (2006), Dullemond et al (2006)

I. Intermediate Disk Evolution t ~ yr

I. Effect of Gas Dispersal on Planet Formation 1. Formation of gas giant planets via core accretion. 2. Extent of planet migration 3. Lifetime of disks around stars of various mass 4. Ability of planetesimals to form by gravitational instability. 5. Existence of water and volatile-rich objects in outer regions of planetary systems 6. Eccentricity and size of terrestrial planets 7. Transitions disks, CTTS ->WTTS, dust rings

Outline I. Introduction II. Viscous Evolution III. Photoevaporation by Central Low Mass Star IV.Disk Dispersal: Applications of Models V.Summary

II. Mechanisms for Dispersal: Viscous Spreading and Accretion Hartmann et al (1998), Calvet et al (2000), Muzerolle et al (2005), (Clark et al (2001), Matsuyama et al (2003), Ruden (2004), Alexander et al (2006a,b), Hueso & Guillot (2006) Most spirals in but small amount spreads out

II. Viscous Spreading and Accretion Alpha Disk Evolution

III. Photoevaporation by Central Star Hollenbach et al (1994), Shu, Johnstone, & Hollenbach (1994), Yorke & Welz (1994, 1996), Richling & Yorke (1997), Kessel et al (1998), Richling & Yorke(2005), Alexander et al (2006), Gorti & Hollenbach (2007) Flow to ISM | cold disk (green), hot surface (red) | Flow to ISM

III. Photoevaporation by Central Star A. Photoevaporation: Basic Physics Need to self-consistently calculate the chemistry, heating and cooling, radiative transfer, vertical and radial structure, and dynamics of flow. Approximations are made! or nearby star UV

A. Physics of Photoevaporation Heating, Cooling A. EUV (h > 13.6 eV) Heating by photoionization of H, Cooling like HII regions, T=10 4 K B. FUV (h < 13.6 eV) Heating by grain photoelectric mechanism or FUV pumping of H 2 Cooling by O, C +, H 2, grains. T~ K. C. Xrays (h > 100 eV) Heating by photoionization of H, He, C, O, etc and secondaries Cooling by OI (6300), [OI] 63um, [NeII] 12.8um, etc. T~ K

A. Physics of Photoevaporation Adams, Hollenbach, Laughlin, & Gorti (2004).

III. Photoevaporation by Central Star B. Sources of FUV, EUV, Xrays 1. Magnetic Activity (coronal, chromospheric): Alexander et al (2005) 2. Accretion (Hartmann, Calvet) Xrays EUV FUV Ribas et al (2005) 100 Myr G star

B. Photoevaporation by Central Star (EUV, FUV, Xrays)

C. Photoevaporation by Central Star: Model Results (FUV, EUV+Xrays) M * = 1 M , L * = 2.3 L , L X =6x10 -4 L , L EUV = L X, L FUV = 0.28 L  EUV photons photoevaporate from 1-50 AU, especially strong from 1-10 AU

C. Photoevaporation by Central Star: Model Results (FUV, EUV+Xrays) M * = 1 M , L * = 2.3 L , L X =6x10 -4 L , L EUV = L X, L FUV = 0.28 L  Beyond 100 AU optical photons photoevaporate

IV. Applications: Model Spectra and Observations Verification of hot gas capable of evaporating Model Spectrum: 20 AU optically thin disk M * = 1 M  M gas = 0.1M J M dust = M J NeII and NeIII verify hot, T~ K gas FeI SiII ss H 2 S(0)

IV. Applications: Model Spectra and Observations Observational Diagnostics of ionized (HII) surface and Xray heated region just below it (partially ionized) Infrared Fine Structure Lines ([NeII] 12.8 um, [NeIII], etc) Hollenbach & Gorti (in prep), Glassgold, Najita, and Igea 2007, Pascucci et al (2007) for observation of [NeII] 12.8um, Geers et al (2007)

IV. Applications: Model Spectra and Observations Herczeg et al 2007: [NeII] has FWHM = 23 km/s (TW Hya using Michelle)

IV. Applications: Model Spectra and Observations Model Spectra: M * = 1 M , varying L EUV and L X Observed NeII

IV. Applications: Model Spectra and Observations See Poster P3-11, Espaillat et al

IV. Applications: Lifetimes of Disks M * = 1 M , M d = 0.1 M  L * = 2.3 L , L X =5x10 -4 L , L FUV = 5x10 -4 L  Grains have 0.1 of UV/optical opacity Of ISM grains.

IV. Applications: CTTs, Transition Disks, and WTTS : Clarke et al (2001), Armitage, Clarke, & Palla (2003), Alexander et al (2006 a,b), Hollenbach, Dullemond, Gorti (in prep)

V. Summary and Conclusions 1. Photoevaporation is likely the dominant dispersal mechanism for the outer disk (beyond several AU). 2. Viscous spreading/accretion disperses the inner disk. 3. EUV photoevaporation creates gap at a few AU, then an inner hole, then rapidly evaporates outer torus from inside out. 4. Photoevaporation by FUV photoevaporates from outside in. 5. [NeII] and [NeIII] will provide excellent constraints on the photoevaporation rates. 6. Photoevaporation and viscous accretion may explain the rapid transition of CTTS to WTTS, the small population of transition disks, and the large inner holes of some transition disks.

End

III. Photoevaporation by Central Star B. Sources of FUV, EUV, Xrays 1. EUV from accretion shock does not make it through accretion column (Alexander et al 2004) 2. In addition, any EUV (accretion or magnetic activity) will not make it through a strong (> M  yr -1 protostellar disk wind (Gorti & Hollenbach 2007) 3. Xrays not very important: Alexander et al (2004) 4. We take X ray and FUV during accretion phase from observations of accreting stars of ages ~1 Myr, accretion rates of order 3x10 -8 M  yr -1. For magnetic activity, we use observations of either the FUV or X rays, and scale the EUV by the Ribas et al spectrum

A. Physics of Photoevaporation Gorti & Hollenbach (2007)

C. Photoevaporation by Central Star: Model Results (FUV, EUV+Xrays) M * = 1 M , L * = 2.3 L , L X =6x10 -4 L , L EUV = L X, L FUV = 0.28 L 

III. Photoevaporation by Central Star (EUV) Hollenbach et al (1994)

C. Photoevaporation by Central Star (FUV, EUV+Xrays) M * = 1 M , L * = 2.3 L , L X =6x10 -4 L , L FUV = 0.28 L  Log yr

IV. Applications: Lifetimes of Disks Disk Dispersal versus Central Star Mass, M disk = 0.03 M *

IV. Disk Dispersal and Planets: Gas Giants Planet Migration 1.Prevention of Type II Migration onto the star, and possible outward migration by photoevaporating the outer disk. Allows gas giants at 2-50 AU. Lecar & Sasselov ( 2003), Veras & Armitage (2004) 2. Prevention of Type II Migration onto the star by photoevaporation by EUV from central star, which creates a central hole out to several AU. Allows gas giants at several AU. Matsuyama et al (2003)

IV. Disk Dispersal and Planets: Gas Giants Armitage (2000), Hollenbach & Adams (2004), Adams et al (2004), Adams et al (2006) Time to form critical mass core vs evaporation time M * at 0.2 pc Note: M * = 20 M  is N * ~ 2000 (Orion) M * = 5 M  is N * ~

Very new results of Dullemond, Hollenbach, & Gorti  g cm -2 r AU M * =1 M , M d = 0.1 M ,  = no FUV photoevaporation 1 Myr contours

Very new results of Dullemond, Hollenbach, & Gorti  g cm -2 r AU M * =1 M , M d = 0.1 M ,  = FUV photoevaporation 1 Myr contours

IV. Disk Dispersal and Planets: Kuiper Belt Cutoff Chiang et al (2003)

IV. Disk Dispersal and Planets: Kuiper Belt Cutoff ( Hollenbach & Adams 2004 a,b) Time to form large grains vs evaporation time

IV. Disk Dispersal and Planets: Disk Survival versus M * Gorti & Hollenbach (2006) L FUV excess scaled from Xrays

IV. Disk Dispersal and Planets: Disk Survival versus M * Gorti & Hollenbach (2006) L FUV excess = L *

IV. Disk Dispersal and Planets: Planetesimal Formation Throop & Bally (2005)

IV. Disk Dispersal and Planets: Late Evolution of Gas and Dust Takeuchi, Clarke, & Lin (2005)

IV. Disk Dispersal and Planets: Late Evolution of Gas and Dust Takeuchi, Clarke, & Lin (2005)

Photoevaporation by External Star Observed!