Master course: Astrochemistry Basic principles and recent results

Transcription

1 Master course: Astrochemistry Basic principles and recent results Ewine F. van Dishoeck Leiden Observatory Spring 2008

2 Outline of course Introduction; Basic molecular processes I Basic molecular processes II Chemistry in the early Universe Chemistry in diffuse clouds, PDRs, XDRs Chemistry in shocks Chemistry in dark clouds and pre-stellar cores Chemistry in star-forming regions; disks; relation with solar system material

3 Background Reading See literature list Relevant reviews summarized for each lecture

4 Organization Lectures Friday 9:00-10:45 Oort lectures through March 14 2 afternoons of exercises in March (voluntary) Tentamen (3 ECTS) Oral exam on appointment Short presentation to group as part of exam possible Required background Radiation processes (Planck function, Einstein coeff.) Statistical physics (Boltzmann, Maxwell distribution) Quantum physics (H atom, molecular spectroscopy) Assistant: Nadine Wehres, room 507, x8411

5 1.1 Introduction Regions of space between stars not empty but filled with a very dilute gas Molecules are found in the cold, neutral component of the interstellar medium (ISM) Diffuse clouds: T kin ~100 K, n~100 cm -3 Dense clouds: T kin ~ K, n~ cm -3 Compare atmosphere at sea level: T kin ~300 K, n~ cm -3 => Conditions very different from those normally encountered in lab on Earth: molecular physics

6 Time scales Collision time: ~1 month at 10 4 cm -3 Chemical time: ~10 5 yr Star formation: ~10 6 yr Lifetime cloud: ~10 7 yr Do not expect to find many molecules since chemical reactions slow Surprise: interstellar clouds contain a very rich chemistry!

7 Horsehead nebula AAT

8 Dark cloud B68 ESO-VLT Alves et al. 2001

9 Intro (cont d) Interstellar clouds are birthplaces of new stars Evolution abundances molecules: astrochemistry Molecules as physical diagnostics: astrophysics Progress strongly driven by observations: technology => Very interdisplinary topic!

10 Lifecycle of gas and dust Based on Ehrenfreund & Charnley 2000 What are building blocks for life elsewhere in the Universe?

11 1.2. Some history Diffuse interstellar bands discovered in 1922 by Heger, 1934 by Merrill Still not identified anno 2008! Sharp bands due to simple gas phase molecules identified in CH: Swings & Rosenfeld 1937 CN: McKellar 1940 CH + : Douglas & Herzberg 1941 First astrochemical models Kramers & ter Haar 1946 Bates & Spitzer 1951 CN CH Adams 1941

12 Diffuse Interstellar Bands About 250 DIBs are known The carriers of most (all?) DIBs are unidentified DIBs are likely due to large carbon-bearing molecules C 60+ is one of the candidates for two DIBs

13 History (cont d) Development of radio astronomy H I 21 cm: Ewen & Purcell 1951; Oort & Muller 1951 OH 18 cm: Weinreb et al NH 3 1 cm: Cheung, Townes et al First polyatomic molecule! H 2 O 1 cm (22 GHz): Cheung et al Development of UV astronomy 1970: H 2 Development of millimeter astronomy 1970: CO >1970: flood of new molecules

14 Orion 230 GHz survey Sutton et al Blake et al Blake et al. 1987

15 History (cont d) Development of IR astronomy 1983: IRAS First full-sky survey at 12, 25, 60 and 100 μm Cirrus clouds and dust properties Presence of very small dust particles ( Å), large molecules (PAHs?) : Infrared Space Observatory (ISO) First complete μm spectra Nature and composition of grains (silicates, ices) and PAHs H 2 O, OH, [O I] far-ir lines Symmetric molecules: C 6 H 6, CH 3, C 2 H 4, CO 2, H 2 lines as probe of shocks and PDRs 2003-now: Spitzer Space Telescope High sensitivity imaging and mapping; limited spectroscopy Ices and silicates toward low mass protostars and disks 1980 s now: Ground and airborne IR instruments

16 Unidentified Infrared Bands due to PAHs? - Seen throughout Universe - Tracers of UV => star formation Tielens, Allamandola et al.

17 Polycyclic Aromatic Hydrocarbons PAHN: N-containing PAHs

18 Carbonaceous material

19 Inventory of ices Gibb et al. 2000

20 1.3. Composition clouds Cosmic abundances elements Element Abundance Element Abundance H 1.00 Mg He Al C Si N S O Ca Na Fe Interstellar space is continuously enriched with heavy elements from dying stars through red-giant winds, novae, supernovae

21 The Astronomers Periodic Table by number Dust grains: by number B. McCall 2001

22 Interstellar grains J.M. Greenberg - Small solid particles ~ μm in size consisting of silicates and carbonaceous material - Most of Si, Mg, Fe incorporated in silicate cores; ~30% of O; ~60% of C - Cold dense clouds (T dust ~10 K): gas-phase species condense on grains forming an icy mantle

23 1.4 Observational techniques: Diffuse clouds Diffuse clouds clouds with total visual extinction A V 1 mag Visible + UV light from background stars not completely obscured If A V 0.3 mag virtually all hydrogen in atomic form diffuse atomic clouds If 0.3 mag A V 1 mag significant fraction of hydrogen is molecular Note: A V NH cm 2 with N H =N(H)+2N(H 2 )

24 Observations of diffuse clouds Observed primarily by absorption lines at visible (since 1900 s) and UV wavelengths (since 1970 s) Classical example: line-of-sight toward ζ Oph Spectra show sharp interstellar lines super-imposed on broad stellar lines

25 Molecular spectroscopy E=E el + E vib + E rot

26 Dense clouds Opaque at visible and UV wavelengths => molecules shielded from dissociating UV radiation Millimeter emission: rotational transitions Limitation: molecule must have permanent dipole moment => cannot observe H 2, C 2, N 2, CH 4, C 2 H 2, Advantage: many molecules down to low abundances; lines in emission => map Infrared absorption: vibrational transitions Limitation: need background IR source => only info along line of sight Advantage: symmetric molecules + solid state Earth s atmosphere prevents observations of key molecules: H 2 O, O 2, CO 2

27 Infrared: absorption gas and solids 10 7 cm -3 Flux Continuum due to hot dust 10 4 cm -3 Absorption by cold dust Wavelength Vibrational transitions of gases and solids

28 1.5 Identified interstellar molecules - list dd. 2002

29 Diversity of molecules More than 130 different molecules found Ordinary molecules NH 3, H 2 O, H 2 CO, CH 3 CH 2 OH,. Exotic molecules HCO +, N 2 H +, HCCCCCCCN,. Unusual molecules (rare on Earth but not in space)

30 Millimeter spectra: myriads of lines! G327.3 Massive YSO - Not all lines identified Gibb et al. 2000

31 Some recent detections H 3+, H 2 D +, D 2 H + : cornerstones ion-molecule chemistry C 3, C 4, C 6 H 2, CH 3 CHCH 2 : new carbon chains Cyclic C 2 H 4 O: fifth ring C 6 H 6 : benzene: simplest PAH C 6 H -, C 8 H -, C 4 H - : first negative ions! D 2 CO, ND 3, CD 3 OH: doubly + triply deuterated molecules NaCN, AlCN, SiN: metal-containing species Not convincingly detected: O 2, glycine,.

32 Interstellar H 3 + Oka & Geballe 1996 McCall et al. 1999, 2003

33 H 2 D + and D 2 H + Strong H 2 D + and D 2 H + in cores 372 GHz H 2 D + D 2 H GHz Stark et al Caselli et al Vastel et al H 3+, D 3+ : IR

34 Detection of circumstellar benzene CRL 618 Evolved star Cernicharo et al. 2001

35 Interstellar antifreeze Ethylene glycol Turner, Hollis et al. 2003

36 Latest news! Negative ions Propene in TMC-1 Marcelino et al McCarthy et al. 2006

37 Molecules at high redshift: z=6.4! CO and [C II] in quasar SDSS J at z=6.4 Walter et al. 2003, Maiolino et al. 2005

38 1.6 Importance of molecules Exotic chemistry: unique laboratory Astrochemical evolution Molecules as diagnostics of temperature T kin, density n H, velocity, Molecules as coolants Radiation escapes from cloud => net kinetic energy lost => cloud cools down Collisions Radiation CO (J=0) CO(J=1) CO(J=0)

39 Molecular excitation Pure rotational transitions Collisions radiation J=1 J=0 - Population distribution not in thermodynamic equilibrium but determined by competition radiative and collisional processes - Critical density μ μ 2 ν 3 => higher frequencies probe higher densities and temperatures

40 Questions addressed What are chemical processes leading to formation and destruction of molecules? How well are basic molecular processes known from experiments or theory What is evolution of molecules in the universe, from their creation at high redshifts to interstellar clouds to incorporation in new solar systems How can molecules be used at physical and chemical diagnostics of physical structure, evolution, cosmic-ray ionization, primordial D,

41 1.7 Examples of observations that need explanation Chemical gradients in dark clouds See TMC-1 maps NH 3 vs HC 7 N Why are complex long carbon chains only found in southern part of cloud? Chemical differentiation in star-forming regions See Orion-KL vs. Orion 1.5 S Why does Orion-KL contain so many complex saturated molecules, but Orion 1.5 S only simple species, even though the sources have similar mass and luminosity? Is this an evolutionary effect?

42 Taurus molecular cloud Perseus CO 1-0 Auriga B5 L1449 IC348 NGC1333 TMC-1 T Tau Taurus Ungerechts & Thaddeus 1987

43 Chemical gradients in dark clouds 0.1 pc Long carbon chain peak Olano et al. 1988

44 Orion molecular cloud HST image HST image

45 Chemical differentiation in Orion Groesbeck Two high-mass YSOs, comparable L,M -Only 1.5 apart -Very different chemistry!

46 1.8 Basic molecular processes: gas phase Van Dishoeck 1988 in MAII Tielens, Ch. 4 Because of low temperatures and densities in clouds, chemistry is not in thermodynamic equilibrium but controlled by two body reactions => abundances depend on physical conditions (T,n,radiation field), history, Three body reactions do not become important until n>10 12 cm -3 Although models contain thousands of reactions, only few different types of processes Rate of reaction: kn(x) n(y) cm -3 s -1 Rate coefficient in cm 3 s -1

47 Types of chemical reactions Formation of bonds Radiative association: X + + Y XY + + hν Associative detachment X - + Y Æ XY + e Grain surface: X + Y:g XY + g Destruction of bonds Photo-dissociation: XY + hν X + Y Dissociative recombination: XY + + e X + Y Collisional dissociation: XY + M X + Y + M Rearrangement of bonds Ion-molecule reactions: X + + YZ XY + + Z Charge-transfer reactions: X + + YZ X + YZ + Neutral-neutral reactions: X + YZ XY + Z

48 1.9 Radiative association τ c X + Y Æ XY* ÆXY + hν τ d τ r Energy conservation => photon must be emitted, which is a very slow process τ r = s vibrational transition τ c,d =10-13 s collision time Molecule formation occurs only 1:10 10 collisions

49 Radiative association (cont d) Process becomes more efficient if electronic states available τ r =10-8 s electronic transition τ c,d =10-13 s collision time => Efficiency increased to 1:10 5

50 Radiative association (cont d) Efficiency can be enhanced if: Electronic state available: τ r shorter Entrance channel has barrier: τ d longer Molecule larger: τ d longer Radiative association is extremely difficult to measure in laboratory because 3-body processes dominate under most lab conditions. Many rate coefficients are based on theory; overall uncertainties 1-2 orders of magnitude Exception: C + + H 2 Æ CH 2+ + hν k~10-15 cm 3 s -1 within factor of 2-3 Initiates carbon chemistry

51 Experiment vs. theory k c X + Y Æ XY* k r k d XY*ÆXY + hν k m XY* + M ÆXY + M Measured in experiments k k eff ra k + k [ M] = k r m c k + k + k [ M] d r m kc = ( ) k k r d keff = kra + k 3 b[ M] Theory + required by astrochemists

52 Uncertainties in rate coefficients Herbst 1988 Comparison experiment with theory needs to take k 3b into account => Measure rate coefficients as functions of pressure

53 Experiment vs. theory: examples C 2 H 2+ + n-h 2 Experiment: Gerlich & Horning 1992 Theory: Herbst & Yamashita 1993

54 1.10 Associative detachment Usually not important in cold clouds, but can play a role in partly ionized regions and early universe Form negative ions by radiative attachment X + e Æ X - + hν slow process Form molecule by associative detachment X - + Y Æ XY + e

55 Associative detachment fast Examples: H + e Æ H - + hν H - + H Æ H 2 + e S + e Æ S - + hν S - + CO Æ OCS + e fast fast

56 1.11 Photodissociation XY + hν Æ X + Y Direct photodissociation OH, H 2 O, CH, CH 2, Predissociation CO, NO,.. Experiments available for stable molecules, but not for radicals or ions Small molecules: quantum chemical calculations of potential surfaces of excited states + transition dipole moments, followed by nuclear dynamics to obtain cross sections

57 Examples experiments H 2 O absorption followed by direct dissociation: accurate cross sections within 20% NO absorption (full) and fluorescence (dashed); mostly predissociation through discrete transitions; large uncertainties (order of magnitude)

58 Photodissociation (cont d) Spontaneous radiation dissociation Ex: H 2 Photodissociation of both H 2 and CO occurs by line absorption => self-shielding important inside clouds Photoionization: XY + hν Æ XY + Æ X + + Y

59 H 2 spontaneous radiative dissociation 90% of absorptions into B and C states are followed by emission back into bound vibrational levels of the X state 10% of the absorptions are followed by emission into the unbound vibrational continuum, leading to dissociation

60 Summary processes: small molecules Direct p.d. Ex: H 2+, OH, H 2 O Predissociation Ex: CO Coupled states p.d. Ex: OH Spontaneous Radiative dissociation Ex: H 2

61 Photodissociation rate Continuum photodissociation k pd z σ ( λ) I( λ)dλ pd 912A = where σ pd is the cross section in cm 2, I= radiation field Discrete photodissociation k pd e = π 2 2 λ f η I λ 2 line line line ( line ) mc lines e where f is oscillator strength and η is the dissociation probability

62 Interstellar radiation field CO, H 2 CH, OH 912 Å = 13.6 ev cutoff Average radiation provided by young O + B stars in solar neighborhood

63 Cosmic-ray induced radiation H 2 + CR => H 2+ + e* H 2 + e* => H 2 * + e p H 2 * => H 2 + hν λ(å) Prasad & Tarafdar 1983 Gredel et al Detailed line + continuum spectrum peaking around 1600 Å and continuing below 912 Å

64 Other radiation fields Ly-α dominated Shocks,.. Stellar blackbodies T eff = K Disks, cool PDRs, Solar radiation T eff =5500 K + Ly α Comets See van Dishoeck et al. 2006

65 1.12 Dissociative recombination Atomic ions: X + + e Æ X + hν Molecular ions: XY + + e Æ XY + hν Æ X + Y Radiative: slow Radiative: slow Dissociative: rapid at low T Need curve crossing between XY + and repulsive XY potential for reaction to proceed fast

66 Storage ring experiments CRYRING, Stockholm - Also Aarhus, Heidelberg

67 H 3+ + e DR rate McCall et al Literature values range from <10-12 to 10-7 cm 3 s -1 at 300 K over last 25 years - High rate coefficients now also reproduced by theory, even without curve crossing - Rates for other molecules known accurately

68 DR products XH n+ + e Æ XH n-1 + H Æ XH n-2 + H 2 Æ Branching ratios: Major uncertainty! Example: H 3 O + Æ H 2 O + H Æ OH + H 2 Æ OH + H + H Æ O + H 2 + H

69 -Major surprise experiments: 3-body products (e.g. OH + H + H) abundant! H 3 O + DR products Product Vejby-C. et al. 97 Williams et al. 96 H 2 O + H 33% 5% OH + H 2 18% 36% OH + H + H 48% 29% O + H 2 + H 1% 30%

70 1.13 Collisional dissociation If T high ( 5000 K), H 2 destroyed by collisions H + H 2 Æ H + H + H He + H 2 Æ He + H + H H 2 + H 2 Æ H2 + H + H H 2 no dipole moment => significant population in high v at high T => large dissociation rate CO small dipole moment => radiative stabilization rapid => not much population in high v => small dissociation rate

71 Astronomical Units pc = parsec = 206,265 AU = cm M = solar mass = g L = solar luminosity = erg s -1 ev = erg Å= 10-8 cm Jansky = erg s -1 cm -2 Hz -1 Raleigh = 10 6 photons s -1 cm -2 (4π sr) -1 Debye = esu cm