Astrochemistry 2016

Lecture 5: Shock chemistry

http://www.sron.nl/~vdtak/astrochem.html

Recap of Lecture 4 l  Physics of diffuse ISM clouds combine VIS-UV absorption with mm-wave emission T from H2, C2 absorption; n from CO rotation & CNO fine structure result: T = 50 – 100 K; n = 100 – 300 cm-3

l  Chemistry: mostly gas phase (except H2, NH) carbon: start with RA of C+ with H2 oxygen: charge transfer to O+, H abstraction to HnO

+ & OH, H2O nitrogen, deuterium: need H+ / H3

+ to start l  Comprehensive models depth dependent: UV attenuation reduces photorates sharp transitions H-H2 and C+/C/CO CH+ problem: role of turbulence? l  Photon dominated regions dense cloud nearby hot star expect layered structure as in Orion Bar extended C+, C emission: clumpy structure

Date Topic Literature

17-05-2016 I. Basic chemical processes Tielens 2013, Rev. Mod. Phys.

19-05-2016 II. Gas-phase and grain surface reactions

Smith 2011, ARAA

24-05-2016 III. Early Universe chemistry Galli & Palla 2013, ARAA

26-05-2016 (09:00)

IV. Diffuse interstellar clouds Snow & McCall 2006, ARAA

31-05-2016 V. Shock chemistry Larsson et al 2012, Rep. Prog. Phys.

02-06-2013 VI. Dense interstellar clouds Bergin & Tafalla 2007, ARAA

07-06-2016 VII. Star- and planet-forming regions Herbst & van Dishoeck 2009, ARAA

09-06-2016 VIII. Steps toward astrobiology Lineweaver & Chopra 2012, Ann. Rev. Earth Planet Sci.

17-06-2016 Presentations KB 257, 13:30-16:30

Course Schedule

Today's lecture

Shock basics J- and C-type shocks Chemistry in dissociative shocks Non-dissociative shocks Comparison with observations

Shock waves Shock: pressure-driven compressive disturbance traveling faster than the local signal speed Produces irreversible change in the state of the fluid Ubiquitous in ISM: expanding H II regions supernova explosions stellar winds protostellar outflows accretion onto star or compact object cloud-cloud collisions ...

Example protostellar outflow: HH46-47

Stellar wind 100 – 200 km/s sweeps up surrounding molecular gas: outflow momentum conservation: 20 – 30 km/s mass loss rates up to 10-4 M0/yr: short phase

Sound speed and Mach number

Sound speed: cS2 = dP / dρ

Take EoS P = Kργ γ = 5/3 for adiabatic fluid γ = 1 for isothermal fluid Adiabatic flow: cS ~ ρ1/3

sound speed is larger in denser gas For isothermal gas in ISM cS = (kT/m)0.5 ≈ 1 km/s (very small!) Mach number: M = V / cS

Jump conditions for shocks

Adopt reference frame where shock is stationary Consider plane-parallel shock conditions depend only on distance x from front Neglect viscosity, except in transition zone: large velocity gradient viscous dissipation transform bulk kinetic energy into heat irreversible change: entropy increases

Velocity & temperature profiles (in shock frame)

More about jump conditions

Thickness of shock front ≤ mean free path of particles always << thickness of radiative zone need collisions for radiative cooling Regard thickness as infinitely small discontinuity Need to find physical conditions at (2) immediately behind shock or (3) in post-radiative zone given those at (1) pre-shock and the shock velocity

Today's lecture

Shock basics J- and C-type shocks Chemistry of dissociative shocks Non-dissociative shocks Comparison with observations

Signal speeds in the ISM In the absence of magnetic fields, information travels at the sound speed M < 1: subsonic M > 1: supersonic → shocks If magnetic field present, disturbances travel along the field lines at the Alfvén speed Interstellar field strengths: empirical law B = 1 μG √nH for 10 < nH < 106 cm-3

Magnetosonic speed:

Magnetic precursors Since vA ~ 1/√ρ the Alfvén speed for decoupled ion-electron plasma can be much larger than if coupled In many cases: cS < vA,n < vS < vA,ie The ion-electron plasma sends information ahead of the disturbance and “informs” the pre-shock plasma that the compression is coming: magnetic precursor The compression is now subsonic and the transition smooth and continuous: C-type shock The ions then couple to the neutrals by collisions

Comparing J- and C-shocks

J-type (“jump”) shocks: vS > 50 km/s shock abrupt neutrals and ions tied into one fluid warm: T = 40 vS

2 [K; v in km s-1] most radiation in ultraviolet C-type (“continuous”) shocks: vS < 50 km/s gas variables (T, ρ, v) change gradually ions ahead of neutrals; drag modifies neutral flow Ti ≈ Tn; both much lower than in J-shocks most radiation in infrared

Shock structure

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Chemical processes in shocks

Endothermic reactions e.g., O + H2 → OH + O Exothermic reactions with barriers Collisional dissociation only at very high T, n Release of icy grain mantles Sputtering of grain cores (Si) enhance SiO in gas-phase Strong UV radiation field i.p. Ly α

Today's lecture

Shock basics J- and C-type shocks Chemistry in dissociative shocks Non-dissociative shocks Comparison with observations

Basic J-shock chemistry

Molecules dissociate at shock front re-form in postshock gas VS > 50 km/s: dissociation into atoms VS > 70 km/s: ionize the atoms VS > 80 km/s: photodissociation ahead of shock

Shock structure

Solve hydrodynamic equations conserving mass, momentum, energy Cooling: atomic / molecular excitation gas-grain collisions H2 dissociation H ionization Heating: kinetic thermal energy radiation from warm gas upstream exothermic chemical reactions Chemical composition determines cooling rate solve hydrodynamics and chemistry simultaneously

Chemical processes

l  T > 105 K: collisional dissociation H2 + H2 → H2 + H + H H2 + H → H + H + H CO + H → CH + O (slower than H2) CH + H → C + H2 or C + H + H l  T < 104 K: slow H2 reformation Grains too hot – gas-phase process (like early universe) H + e → H­ + hν (slow) H­ + H → H2 + e (fast) dominates if n(e)/n(H) > 0.02, results in H2/H = 10-3

which is sufficient to cool the gas l  T < 3000 K: rapid H2 reformation O + H2 → OH + H OH + H2 → H2O + H C+ + OH → CO + H+

Chemical diagnostics

l  OH plays a key role X + OH → XO + H X+ + OH → XO+ + H where X = Si, S, N ... l  SiO = good shock diagnostic especially if Si enhanced by grain core destruction l  Molecule destruction: reverse reactions with H photodissociation l  Much of shock energy emerges in Ly α 1216 Å line can dissociate OH, H2O, ... but not H2, CO, CN, ...

Grain destruction High velocity J-type shocks thermal sputtering of grain cores local enhancements of gas-phase Si, Fe, ... Lower velocity shocks: non-thermal sputtering of grain cores and ice mantles locally enhanced H2O, CH3OH, ... Si from grain cores reacts with OH to form SiO seen in highly collimated bipolar outflows abundance increases with mass loss rate Schilke et al 1997; Cabrit et al 2012 (HH 212)

Today's lecture

Shock basics J- and C-type shocks Chemistry in dissociative shocks Non-dissociative shocks Comparison with observations

Basic C-shock chemistry

Typical temperatures < 3000 K molecules survive shock Key reactions: O + H2 → OH + H OH + H2 → H2O + H C+ + H2 → CH+ + H Flower & Pineau des Forêts 2010

Predictions of shock models

Oxygen O, OH, H2O enhanced Carbon CO, C, C2H, C2H2, CH3, ... H2CO (from CH3 + O) HCO+ uncertain Sulphur H2S SO, SO2, CS Silicon SiO Nitrogen NH3, HCN Results depend on pre-shock H/H2 and C/CO ratios

Herschel images and spectra of H2O emission Protostellar environments: traces hot spots where energy is dissipated profiles: outflow dominates

Nisini et al 2010; Kristensen et al 2010

Today's lecture

Shock basics J- and C-type shocks Chemistry in dissociative shocks Non-dissociative shocks Comparison with observations

Comparison with observations: IC 443 Interaction supernova remnant / molecular cloud best studied and nearest example (1500 pc) Optical, infrared, radio, X-ray observations: several shell-like regions ring of shocked molecular gas e.g. H2 2.12 µm, IRAS 100 µm, ... Molecular emission lines: broad asymmetric velocity profiles single J- or C-type shock does not match require multiple T,n components Q: is chemistry affected by shocks?

Optical-infrared view of IC 443

Diameter 50' (>Moon) age 3000-30,000 yr Type II supernova central neutron star

2MASS: red = H2, blue = [Fe II] CFHT zoom of NE shell

Molecular emission lines

SWAS / FCRAO: Snell et al 2005

Conclusions for IC 443

No single shock explains all observations fast J-type shock (~100 km/s) + slower shock (12-25 km/s) of either type combination expected if SNR overtakes clumpy medium Fast J-shock: strong UV radiation photodissociates H2O behind shock enhances ionization: smoothes shock front Weak H2O emission seen with SWAS depleted on grain mantles in pre-shock gas

Example 2: Orion

Region of high-mass star formation distance 420 pc

H2 emission in 2 peaks explosion-type event

Broad SO2 and SiO lines

Schilke et al 1997

Abundances in Orion KL

Pre-shock density ~105 cm-3 >> 103 cm-3 in IC 443 lower H/H2 ratio UV less effective

Molecule Outflow Quiescent gas

SO 5 x 10-7 <1 x 10-9

SO2 5 x 10-7 <3 x 10-9

SiO 1 x 10-7 <3 x 10-10

Example 3: Protostellar outflow L1157

Forming solar-type star active accretion collimated outflow distance 440 pc luminosity 11 L0 Line profiles: shocked + quiescent gas shock: SiO, CH3OH, CS, H2CO, ... quiescent: C18O, N2H

+, H13CO+, DCO+, ... Strong enhancements of SiO and CH3OH but not much in other species mainly grain core/mantle destruction Strong lines: excitation or chemistry? need careful analysis Bachiller et al 2001

Molecular line emission from L1157

Some peak at star others in lobe some toward both

Summary

l  Physical models of shock in molecular clouds J-type = fast & sharp C-type = slower, smoothed by B-field l  Observations: pure J- or C-shocks are rare but combination is common

l  Chemical models: dissociative vs non-dissociative shocks UV photodissociation important for both l  Clear observational diagnostic: SiO from grain core sputtering l  Often see CH3OH enhancements and other grain mantle components l  Strong H2O emission from warm gas-phase chemistry