H 2 Formation in the Perseus Molecular Cloud: Observations Meet Theory

Motivation (2) Theory Krumholz et al. (2009) Analytic solution for H 2 content in an atomic-molecular complex No direct comparison to individual molecular clouds in the MW! (1) Observations Strong correlation between star formation rate and H 2 surface density Constant SF efficiency in molecular clouds Ability to form H 2 controls the evolution of individual galaxies! log Σ SFR (M  yr -1 kpc -2 ) log Σ H2 (M  pc -2 ) 30 nearby spiral galaxies Bigiel et al. (2011) A high resolution study of the HI–H 2 transition across a molecular cloud Estimate R H2 = Σ H2 / Σ HI Investigate how R H2 spatially changes Perseus molecular cloud D ~ 300 pc and solar Z Low mass (~10 4 M  ) with intermediate SF

Background: Analytic Modeling of H 2 Formation in a PDR Krumholz et al. (2009; KMT) model H2H2 CNM Pressure equilibrium with WNM Sharp HI-H 2 transition Uniform isotropic ISRF Equilibrium H 2 formation: Formation on dust grains = Photodissociation by LW photons

Background: Analytic Modeling of H 2 Formation in a PDR KMT's predictions: R H2 is determined by CNM property, metallicity, gas surface density, and is independent of ISRF. log Σ HI + Σ H2 (M  pc -2 ) log Σ HI (M  pc -2 ) M H2 / M (1) Minimum Σ HI to shield H 2 against ISRF Σ HI ~ 10 M  pc -2 for solar Z (2) H 2 -to-HI ratio (R H2 ) 10 M  pc -2

IRAS 100 μm image (~4.3': ~0.4 pc at D = 300 pc) GALFA-HI N(HI) image (~4') R H2 = Σ H2 / Σ HI for Perseus Σ HI : GALFA-HI DR1 data Σ H2 : IRAS 60, 100 μm, Schelegel et al. T dust, 2MASS A V images

R H2 image 12 CO contours Dark regions Star-forming regions B5 B1E B1 IC348 NGC1333 Lee et al. (2011, submitted)

Σ HI vs Σ HI + H2 1) Uniform Σ HI ~ 6–8 M  pc -2 General results Consistent with KMT's prediction of Σ HI ~ 10 M  pc -2 for solar Z! 2) No detection of turnover HI envelopes are highly extended (> 30 pc)! Σ HI (M  pc -2 ) Σ HI + Σ H2 (M  pc -2 ) 3σ IC348 (Star-forming region) HI-dominatedH 2 -dominated Σ HI (M  pc -2 ) Σ HI + Σ H2 (M  pc -2 ) 3σ B1E (Dark region) HI-dominated H 2 -dominated

R H2 vs Σ HI + H2 5) HI–H 2 transition (R H2 ~ 0.25) at N(HI + H 2 ) = (8–10) × cm -2 Consistent with previous estimates in the Galaxy (e.g., Savage et al. 1977)! B1E (Dark region) R H2 = Σ H2 / Σ HI Σ HI + Σ H2 (M  pc -2 ) IC348 (Star-forming region) R H2 = Σ H2 / Σ HI Σ HI + Σ H2 (M  pc -2 ) 3σ General results 4) Best-fit parameter Φ CNM = 6– 10 T CNM ~ 70 K, consistent with observed CNM properties (Heiles & Troland 2003)! 3) Agreement with KMT on sub-pc scales

Discussion: Equilibrium vs Non-equilibrium H 2 Formation Equilibrium H 2 formation τ H2 = 10–30 Myr (e.g., Goldsmith et al. 2007) ≥ Lifetime of GMCs Role of turbulence: non-equilibrium H 2 formation? Time (Myr) R H2 = Σ H2 / Σ HI Mac Low & Glover (2011) Equilibrium: R H2 ~ constant Non-equilibrium: R H2 keeps increasing Turbulence may play a secondary role!

Discussion: Importance of WNM / Internal Radiation Field Importance of WNM for shielding H 2  Importance of internal RF T dust image Lee et al. (2011, submitted) T dust ~ 17 K KMT: all CNM Perseus: WNM about 50% Perseus – Uniform external RF, negligible internal RF

Summary 1) The dark and star-forming regions have uniform Σ HI ~ 6–8 M  pc -2. 2) The purely HI envelopes are highly extended (> 30 pc). 3) HI–H 2 transition occurs at N(HI) + 2N(H 2 ) = (8–10) × cm -2. 4) KMT's equilibrium model captures the fundamental principles of H 2 formation on sub-pc scales! 5) The importance of WNM for H 2 shielding, internal RF, and the timescale for H 2 formation still remain as open questions.