H2 Formation in the Perseus Molecular
Cloud:
Observations Meet Theory
Motivation
(1) Observations
•
•
Strong correlation between star
formation rate and H2 surface density
Constant SF efficiency in molecular
clouds
Ability to form H2 controls the evolution
of individual galaxies!
•
•
•
Krumholz et al. (2009)
Analytic solution for H2 content in an
atomic-molecular complex
No direct comparison to individual
molecular clouds in the MW!
A high resolution study of the HI–H2 transition
across a molecular cloud
log ΣSFR (M yr-1 kpc-2)
•
(2) Theory
•
•
Perseus molecular cloud
D ~ 300 pc and solar Z
Low mass (~104 M) with intermediate
SF
Estimate RH2 = ΣH2 / ΣHI
Investigate how RH2 spatially
30 nearby spiral galaxies
changes
Bigiel et al. (2011)
log ΣH2 (M pc-2)
Background:
Analytic Modeling of H2 Formation in a
PDR
•
Krumholz et al. (2009; KMT) model
CNM
Uniform isotropic ISRF
Pressure equilibrium
with WNM
H2
Equilibrium H2 formation:
Formation on dust grains = Photodissociation by LW photons
Sharp HI-H2 transition
Background:
Analytic Modeling of H2 Formation in a
PDR
• KMT's predictions:
log ΣHI (M pc-2)
10 M pc-2
(1) Minimum ΣHI to shield H2 against
ISRF ΣHI ~ 10 M pc-2 for solar Z
(2) H2-to-HI ratio (RH2)
RH2  fH2 / fHI
1/ 3
  s  3  125  s  3 
 1     
 
  11   96  s  
 1 where s ~
 total Z

MH2 / M
 f (nCNM , Z,  total )
RH2 is determined by CNM property, metallicity,
gas surface density, and is independent of ISRF.
log ΣHI + ΣH2 (M pc-2)
RH2 = ΣH2 / ΣHI for Perseus
• ΣHI : GALFA-HI DR1 data
• ΣH2 : IRAS 60, 100 μm, Schelegel et al. Tdust, 2MASS AV images
IRAS 100 μm image (~4.3': ~0.4 pc at D = 300 pc)
GALFA-HI N(HI) image (~4')
RH2 image
12CO
contours
Dark regions
Star-forming regions
B5
B1
NGC1333
IC348
B1E
Lee et al. (2011, submitted)
ΣHI vs ΣHI + H2
IC348
(Star-forming region)
HI-dominated
General results
1) Uniform ΣHI ~ 6–8 M pc-2
H2-dominated
ΣHI (M pc-2)
Consistent with KMT's prediction of
ΣHI ~ 10 M pc-2 for solar Z!
B1E
(Dark region)
HI-dominated
H2-dominated
3σ
ΣHI + ΣH2 (M pc-2)
ΣHI (M pc-2)
3σ
2) No detection of turnover
HI envelopes are highly extended (> 30 pc)!
3σ
3σ
ΣHI + ΣH2 (M pc-2)
RH2 vs ΣHI + H2
IC348
(Star-forming region)
General results
3) Agreement with KMT on sub-pc scales
3σ
RH2 = ΣH2 / ΣHI
4) Best-fit parameter ΦCNM = 6–
10
TCNM ~ 70 K , consistent
with observed
CNM properties (Heiles & Troland 2003)!
B1E
(Dark region)
3σ
ΣHI + ΣH2 (M pc-2)
5) HI–H2 transition (RH2 ~ 0.25)
at N(HI + H2) = (8–10) × 1020 cm-2
RH2 = ΣH2 / ΣHI
3σ
3σ
Consistent with previous estimates
in the Galaxy (e.g., Savage et al. 1977)!
ΣHI + ΣH2 (M pc-2)
Discussion:
Equilibrium vs Non-equilibrium H2 Formation
• Equilibrium H2 formation
τH2 = 10–30 Myr (e.g., Goldsmith et al. 2007) ≥ Lifetime of GMCs
• Role of turbulence: non-equilibrium H2 formation?
RH2 = ΣH2 / ΣHI
Equilibrium: RH2 ~ constant
Non-equilibrium: RH2 keeps
increasing
Turbulence may play a secondary role!
Mac Low & Glover (2011)
Time (Myr)
Discussion:
Importance of WNM / Internal Radiation
Field
• Importance of WNM for shielding H2
KMT: all CNM
Perseus: WNM about 50%
 Importance of internal RF
Perseus – Uniform external RF, negligible
internal RF
Tdust ~ 17 K
Tdust image
Lee et al. (2011, submitted)
Summary
1) The dark and star-forming regions have uniform ΣHI ~ 6–8
M pc-2.
2) The purely HI envelopes are highly extended (> 30 pc).
3) HI–H2 transition occurs at N(HI) + 2N(H2) = (8–10) × 1020
cm-2.
4) KMT's equilibrium model captures the fundamental
principles of H2 formation on sub-pc scales!
5) The importance of WNM for H2 shielding, internal RF, and the
timescale for H2 formation still remain as open questions.