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21: Photodissociation Regions(PDRs)

James R. Graham

UC, Berkeley

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Reading & Background

• Tielens Ch. 9

• “Dense Photodissociation Regions”,Hollenbach & Tielens, 1997, AARA, 35,179

• “Photodissociation Regions in theInterstellar Medium of Galaxies”,Hollenbach and Tielens, 1999, Rev ModPhys, 71, 173

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Outline

• Introduction

• Example PDRs M42 & M16

• Overview of PDR structure

• PDR emission

• Photodissociation and self-shielding

• The three regimes of PDR chemistry

• Heating & cooling

• Models

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PDRs in the Galaxy & Beyond

• The most prominent PDRs are associated withgas that lies just outside of dense, luminous HII regions, and:– The pervasive WNM– Diffuse and translucent clouds– Reflection nebulae– Neutral envelopes around planetary nebulae– Photodissociated winds from red giant & AGB stars– The ISM in starburst galaxies & AGNs (NRL)

• PDRs include all ISM regions where the gas ispredominantly neutral but where FUV photonsplay a significant role in the chemistry and/orthe heating

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M42 HIIRegion

[O III]Hα

[N II]

VV

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M42/Orion A

• [C II] 158 µm and [O I] 63 µm emission from the PDR in theOrion A molecular cloud behind the Trapezium (Herrmann etal 1997)– & indicate the location of θ1C Ori and IRc 2– Significant [O I] shock emission in a ~ 1' region around IRc 2

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M16

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Introduction

• All HI & much of the H2 in the Milky Way is inPDRs

• Most of the non-stellar IR & mm emission fromgalaxies is from PDRs– On cloud surfaces (AV < 1-3 mag) absorption of

FUV photons (6–13.6 eV) leads to intense• [C II] 158 µm, [O I] 63, 146 µm• H2 rovibrational transitions• IR dust continuum and PAH emission

– Deeper in PDRs, CO and [C I] 370, 609 µm linesoriginate

• FUV-induced feedback may regulate starformation rates

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M16

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CO 1-0 & H2 1-0S(1): Cold & HotMolecular Gas

CO J=1-0 H2 1-0S(1)

Levenson, G

raham et al. 2000, A

pJL, 553, L53

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H2 & HIILevenson

, Graham

et al. 2000, ApJL

, 553, L53

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The M16 PDR

• The temperature in the HII is 9500 K and ne =5200 cm-3

• The H2 emission is fluorescent emissionexcited by nearby O stars is the primaryexcitation source– H2 3-2 S(5) and 7-5 O(4) are detected

• The molecular gas has a mean density ~ 3 x105 cm-3

• The kinetic temperature of the H2 gas is 930 ±50 K– The observed density, temperature, and UV flux

imply a photoelectric heating efficiency of 4%

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PDR Structure

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PDR Structure

• Physical conditions depend on the intensity ofUV from the interstellar radiation field (ISRF) orfrom nearby hot stars and the density

• The incident FUV flux is denoted G0• Measured in units of the canonical ISRF (1.6 x 10-3 erg cm-

2 s-1; Habing 1968)

– Locally the average ISRF G0 ~ 1.7 (Draine 1978)– Close to (0.1 pc) an O star to G0 ~ 106

• Densities range from– n ~ 0.25 cm-3 in the WNM (n is the H nucleus

density)– n ~ 10–100 cm-3 in diffuse clouds– n ~, to 103–107 cm-3 in PDRs associated with

molecular clouds

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PDR Structure• If ionizing radiation is present PDRs are overlaid with

a layer of H+ gas– A thin H+/H0 interface absorbs the Lyman continuum radiation

• The PDR proper is characterized by layers– H0 extending to AV ~ 1–2

mag. (NHI ~ 2–4 x 1021 cm-2)from the ionization front

– C+ extending to AV ~ 2–4mag.

– O0 extending to AV ~ 10mag.

• These layers aremaintained by FUVphotodissociation ofmoleculesphotoionization of C

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Why PDRs are Important

• PDRs include cases where H is H2 and Cis mostly in CO but where the FUVdetermines the ionization fraction andstrongly affects the chemistry of O (e.g.,photodissociating OH, O2, and H2O)

• Most Galactic molecular gas is found atAV ~ 10 mag. in GMCs– All of the atomic and most of the

molecular gas in the Galaxy is in PDRs.

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PDR Emission

• Incident FUV is absorbed primarily by PAHsand dust grains at AV < 1 mag.– Most of the absorbed energy excites the PAHs and

heats the grains• FUV is converted to PAH IR features & far-IR continuum

– Typically 0.1–1% of the absorbed FUV is convertedto photoelectrons (~ 1 eV) which heats the gas

• Gas cools via lines, including [C II] 158 µm & [O I] 63 µm

– Much of the observed far-IR fine-structure and H2

rovibrational emission in galaxies originates fromPDRs

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PDR Emission

• The COBE FIRAS map of the Milky Way (Wrightet al. 1991) is dominated by PDR emission– C+ 158 µm traces the PDRs associated with the the 3

kpc molecular ring and the star forming GMCs Cygnus,Ophiuchus, Carina, Vela & Orion

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PDR Chemistry

• PDR chemistry is distinct from dark cloud ion-molecule chemistry–H2 still forms on grains

• The high FUV field emphasizes– Photoreactions: photoionization & photodissociation– Reactions with H0

– Reaction with vibrationally excited H2, i.e., H2*

• Activation barriers of H2 + atom/radicalreactions are overcome in warm gas (500 K)

• e- recombination & charge exchange areimportant for the ionization balance

• The FUV flux keeps O0 abundant & oxidationreactions are important

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Photodissociation & Self-Shielding of H2

• Photodissociation of H2 follows FUV absorptionin the Lyman & Werner transitions between912–1108 Å– Followed by fluorescence to the vibrational

continuum of the ground electronic state about10–15% of the time

– H2 photodissociation rate follows from a summationover all UV lines

• For N(H2) > 1014 cm-2, the H2 FUV linesbecome optically thick and self-shieldingbecomes important– The photodissociation rate then depends on the H2

abundance & level population distribution as afunction of depth in the cloud

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Photodissociation of H2

• 1Σg+ → 3Σu

+ is not dipolepermitted

– Photodissociation ofH2 in the ISM mustproceed via highenergy photons

– H2 cannot form viaradiative associationof two HI atoms in 1s

Σ

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Photodissociation & Self-Shielding of H2

• When G0/n ~ 0.01–0.1 cm-3 self-shielding alonedominates H2 dissociation and sets the locationof the H/H2 transition– When self-shielded, the H2 column increases rapidly

and the H/H2 transition is sharp– G0/n ~ 0.01–0.1 cm-3 includes diffuse clouds

exposed to the ISRF & clumps in bright PDRs

• PDRs associated with bright FUV sources (HIIregions) typically have G0/n ~ 1 cm-3

– The location of the H/H2 transition is determined bydust and occurs at AV ≈ 2mag.

– Dust reduces the H2 dissociation rate so that a self-shielding column of H2 can build up, and the H/H2transition is sharp again

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Photodissociation & Self-Shielding of CO

• Like H2, CO photodissociation occurs throughline absorption– CO is affected by self-shielding– Mutual shielding arises due to line coincidences with

HI and H2 (van Dishoeck & Black (1988)– For low G0/n CO self-shielding can lead to isotopic

fractionation at cloud edges, where the rareisotopes are preferentially dissociated

• The location of the C+/C0/CO transition in brightPDRs is largely governed by dust– Because of the low abundance of CO relative to H2,

the CO rarely builds up sufficient column to self-shield in the AV ≈ 1 mag. layer and the C+/C0/COtransition is much less sharp than that of H/H2

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PDR Chemistry

• PDR chemistry has three distinct regimes– AV ≈ 2 mag.: vibrationally excited H2 is

abundant and reactions with H2* drive COformation

– AV ≈ 2 mag.: “burning” of the radicals CHand CH2 in the C+/C/CO transition zone formCO

– AV ≈ 8 mag.: deep in the clouds standardion-neutral chemistry is driven by CRionization of H2

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PDR ChemistryCR + H/H2

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AV ≈ 2 mag. PDR Chemistry

• Reactions with vibrationally excited H2 (H2*)have fast reaction rates– O + H2*, OH + H2*, and C+ + H2* reactions have

large enhancements of the correspond H2 reaction

• OH & H2O are built from reactions with H2*– Most OH is photodissociated, but a fraction reacts

with C+ to form CO+, then CO+ + H → CO + H+

– H2O reacts with C+ to form HCO+ → CO + H+

• CO is destroyed by photodissociation• C/C+ balance is set by recombination/photoionization

• A small fraction of C+ reacts with H or H2* toform CH+

– Leads to the formation of small hydrocarbons

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PDR Structure

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AV ≈ 4-8 mag. PDR Chemistry

• As the depth increases, CO formation byoxidation or “burning” of the neutral radicalsCH and CH2 becomes more important than theOH-driven channel– This change indicates the start of the C+/C/CO

transition zone

• Eventually, PDR chemistry gives way tostandard dark cloud, ion-molecule chemistry– Formation of OH and H2O is initiated through the

reaction of H3+ with O

– Reactions of O with OH then forms O2

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PDR Cooling

• Cooling is by FIR fine structure lines– [C II] 158 µm, [O I] 63, 146 µm, [Si II] 35 µm,

[C I] 609, 370 µm

– H2 rovibrational transitions

– Rotational lines, especially CO

• For large n & G0, the gas temperature atthe edge of the PDR reaches ~ 5000 K– Significant forbidden line cooling

• [Fe II] 1.257, 1.644 µm, [O I] 6300 Å, & [S II]6717, 6731 Å

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PDR Heating

• Two main mechanisms couple the gas to theFUV photon energy of stars– The photoelectric effect on PAHs and small dust

grains– FUV pumping of H2

• Other heating mechanisms play only a limitedrole in the heating or become important atgreat depth in the PDR– Gas collisions with warm grains– CR ionization & and excitation– Photoionization of C– Pumping of gas particles to excited states by the

FIR radiation field of the warm dust

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PDR Heating

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Photoelectric Heating

• Photoelectric heating is dominated by thesmallest grains present– Large PAHs are an important component of the ISM

and these molecule-sized particles may play animportant role in heating interstellar gas

• FUV photons can liberate energetic (~ eV) e-’s– These e- may diffuse in the grain, reach the surface,

overcome the work function W and Coulombpotential φc (if positively charged), and inject KE =hν - W - φc into the gas

– The efficiency of PE grain heating also depends onthe probability, Y, that the electron escapes

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Photoelectric Heating• The PE heating efficiency ε is the ratio of gas heating

to the grain FUV absorption rate

ε = Y (hν - W - φc)/hν– For large grains and photon energies well above threshold, the

photons are absorbed ~ 100 Å inside the grain and thephotoelectrons rarely escape (Y ~ 0.1)

• Y is high for planar PAHs– Some (~50%) of the photon energy may remain behind as

electronic excitation

• The IP of a charged PAH can be > 13.6 eV, and FUVphotons cannot create a photoelectron– E.g., the ionization potential of C16H10

+ (pyrene) is 16.6 eV– The PE heating efficiency for small PAHs is then reduced by

the fraction of PAHs that can still be ionized by FUV photons

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Photoelectric Heating

• For an MRN distribution extending into thePAH domain, about half the gas heating is dueto grains with sizes < 15 Å– The other half originates in grains with sizes

between 15–100 Å– Grains > 100 Å contribute negligibly to the

photoelectric heating of the interstellar gas

• The PE heating efficiency depends on thecharge of a grain– A higher charge implies a higher φc to be overcome

and a smaller fraction of the electrons escape– Those that do escape carry away less KE

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Photoelectric Heating

• For PAHs, the charge determines whetherfurther ionizations can occur– The PE efficiency depends on the ratio of the

photoionization rate over the recombination rate ofelectrons with grains/PAHs (∝ G0/ne)

• When G0/ne is small, grains/PAHs are mostlyneutral and the photoelectric heating has thehighest efficiency (ε ≈ 0.05)– As G0/neincreases, the grains/PAHs become

positively charged and the photoelectric efficiencydrops

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H2 Heating

• Absorption of FUV photons pumps H2

molecules to a bound excited electronic state– Fluoresces back to the vibrational continuum of the

ground electronic state & dissociate (10–15%)

– Fluoresce back to an excited vibrational state in theelectronic ground state (85–90%)

• At low n, the bound, excited vibrational statescascade down to the ground vibrational statethrough the emission of IR photons– Gives rise to a characteristic rovibrational spectrum

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H2 Heating

• At high density, n ~ 104 cm-3 collisionaldeexcitation by atomic H is important– Leads to heating of the gas and thermalization of

the rovibrational states– The heating efficiency of this process is

approximately εH2 ≈ (Evib/hν) fH2 ≈ 0.17 fH2

– fH2 is the fraction of the FUV flux pumping H2 anddepends on the location of the H/H2 transition zone

• For G0/n < 10-2 cm3, H2 self-shielding regime– H2 competes effectively with dust for FUV photons

• For n > 104 cm-3 & G0/n < 10-2 cm3, fH2 ≈ 0.25,and εH2 ≈ 0.04

• H2 heating provides an efficient coupling to the FUV field

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Steady-State, Static PDR Models

• Most models assume thermal & chemicalbalance and ignore any flow through thePDR– Equivalent to assuming that the time scale

for H2 formation on grains, τH2 ≈ 105 /n4 yr,which dominates the chemical time scales, isshort compared to the dynamical time scalesor the time scales for significant change inthe FUV flux

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Steady-State PDR Models

• This example representsthe structure of the PDR inOrion (n=2.3 x 105 cm-3; G0= 105; G0 /n = 0.4) vs. AV(Tielens & Hollenbach1985b).

• The illumination is to the left– Top two panels: Abundances

relative to total H– Third panel: Gas & dust

temperatures– Bottom panel: Cooling in the

various gas lines

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Steady-State PDR Models

• Models consider a slabilluminated from one side– Penetrating FUV photons

create an H0 surface layer– The transition from H0 to H2

occurs at AV ≈ 2 mag.– Because of rapid photo-

destruction, H2*, does notpeak until AV ~ 2 mag.

– Balance shifts from C+ toC0/CO at AV ≈ 4 mag.because of dust attenuation

– Dark cloud conditions arereaches at AV ≈ 8 mag.

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Steady-State PDR Models– Second peak in the C0

abundance results fromcharge exchange between C+

and S (10.36 eV vs. 11.26 eV– Except for the O locked up in

CO, essentially all O is in O0

until very deep in the cloud atAV ~ 8 mag.

– Because of low ionizationpotentials, trace species, e.g.,S, can remain ionized througha substantial portion of thePDR

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Comparison with Observations• Comparison of observations of [C II]

158 µm & [O I] 63 µm and modelsshown as a function of G0 (Hollenbachet al 1991)

• G0 is inferred from the FIR continumm– : HII regions– : dark clouds, reflection nebulae &

planetary nebulae– : inner 45–60'' of galaxies

• Models are labeled by their density– Dashed line shows 5% conversion

efficiency of FUV into gas cooling—themaximum expected for the PE effect

– The arrow shows effects of beamdilution

• [O I] dominates cooling at G0 > 103

– Typically, the total cooling line intensityis in the range 10-3–10-2 of the total FIRcontinuum intensity consistent withexpected PE heating efficiency

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PDR Excited by a Super Star Cluster in the Antennae

Whitmore et al. 1999

3 kpc

• NGC 4038/4039 mergerwith ongoing starformation– Old red nuclei– Multiple blue SSCs

• Sizes ~1-5 pc• Ages 5-100 Myr

–70% < 20 Myr

– Dark interaction region• Young (embedded)SSCs

• Massive GMCs• Warm dust (ISO peak)

PANIC

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PDR Excited by a Super Star Clusterin the Antennae

• Mid-IR tracesPAHs & smallgrains heated byhot stars

• Peak opticallyinconspicuous

• K band revealsyoung SSC– Low extinction

AK = 0.1 AV

– 0.''4 seeing

Mirabel et al. 1998

Gilbert, Graham et al. 2001

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K-Band Spectrum of SSC A

Gilbert et al. 2001

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H2 RESULTS FOR SSC A

χ2 contours22 H2 lines

• Best-fit Draine & BertoldiPDR model– Fluorescent H2 excitation– nH = 105 cm-3, T = 900 K,

G0 = 5000

• Weak high-v transitions:– From less dense gas in

weaker FUV field– PDR is very clumpy

• Feedback in Action– Extended Br γ and H2 lines

imply a single SSC heats &disrupts the ISM over ~ 100pc

– Significant fraction of Lycand FUV (912-1108 A)escapes & illuminates thelocal ISM

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FUV Regulation of Star Formation• The constancy of AV in molecular clouds may follow

from the regulation of low-mass star formation in aPDR (McKee 1989; Bertoldi & McKee 1996)– The rate of low-mass star formation is governed by ambipolar

diffusion & newly formed stars inject mechanical energy intothe cloud, supporting it against gravity

– The star formation rate increases as the cloud collapses, AV

increases, the FUV-produced ionization level decreases, andthe ambipolar diffusion rate increases

– Cloud collapse is halted as the increased star formation injectsturbulent energy and equilibrium is achieved when AV ~ 7.5mag.

– The external FUV flux, in controlling the ionization fraction inmost of the cloud, regulates the low-mass star formation rateand the cloud column density