AstrochemistryLecture 8, Circulation of matter

Jorma Harju

Department of Physics

Friday, March 15, 2013, 12:15-13:45, Lecture room D117Course web page

http://www.courses.physics.helsinki.fi/astro/astrokemia

Interstellar dust (1)Dust grains are present everywhere: Solar System, around stars,interstellar clouds, galaxies, and intergalactic medium(picture: ESO/Yuri Beletsky)

This lecture largely follows the article of Karine Demyk “Interstellar dust within the life cycle of the interstellar

medium”, EPJ Web of Conferences 18, 03001 (2011)

Interstellar dust (2)

The interstellar extinction and the emission of diffuse clouds isreproduced by three populations:-“Big Grains” (BGs, ∼ 10− 500 nm) with silicate cores and icymantles-“Very Small Grains” (VSGs, ∼ 1− 10 nm) carbonaceous nanograins- macromolecules and PAHs (Polycyclic Aromatic Hydrocarbons)The relative populations change with the environmentDust is well mixed with gas and constitutes about 1% of the mass ofthe ISM in our Galaxy

Dust is mainly formed in the at-mospheres and circumstellar shellssurrounding evolved stars - AGBstars and Red Giants - and expelledinto the ISM

Interstellar dust (3)

I Dust is destroyed or changed by various physical and chemicalprocesses (size distribution, structure, composition)

I In collapsing clouds, dust is incorporated in protostars andcircumstellar disks which evolve into proto-planetary disks

I At the end of their evolution these stars regenerate dust fromheavy elements

I The properties of interstellar dust have been studied extensivelyin the space-missions IRAS (1983), ISO (1995-1998), Spitzer(2003-2005), and Herschel (2009-2013)

I Laboratory work to produce analog of interstellar dust grains andices and to study their spectroscopic properties

Observations of interstellar dust (1)

I The nature of dust is derived(i) from spectroscopic studies of the light absorbed andscattered by dust grains and(ii) from depletion of heavy elements in the gas phase

I The missing atoms are supposed to be bound to the dust grainsI Major dust constituents:

C, O (> 300 ppm)Mg, Si, Fe (∼ 30 ppm),

I Minor dust components: Na,Al,Ca, Ni (∼ 3 ppm),I Trace components: K, Ti, Cr, Mn, Co (∼ 0.1− 0.3 ppm)

Observations of interstellar dust (2)

I

The ∼ λ−1 dependence of the ex-tinction curve in the visual and UVimplies a grain population with a ∼λ ∼ 0.1− 0.5µm (“Big Grains”)

I BGs are in thermal equilibrium with the interstellar radiation field(ISRF)

I BGs are responsible for various spectral bands in themid-infrared (8− 70µm), both in absorption and emission

Observations of interstellar dust (3)

In the ISM, Big Grains dominate the emission of cold clouds (T ≤ 30K) at submillimeter wavelengths(pictures: Cocoon Nebula, far-IR ESA/Herschel Gould Belt Survey, photo Adam Block/NOAO)

Silicates and oxides (1)

I

Refractory grains are mainlycomposed of slicates of olivine(Mg2xFe2−2xSiO4) and pyroxenetype (MgyFe1−ySiO3)

I

Amorphous silicates are respon-sible for spectral features at ∼9.7µm (Si-O stretching mode) and∼ 18µm (O-Si-O bending mode) inabsorption or emission (picture: Molster &

Kemper 2005, Space Science Review 119, 3)

I

Most interstellar silicates are amor-phous (manifest in broad, blurredshape of the features)Exact composion difficult to deter-mine

Silicates and oxides (2)

I

Sharp crystalline silicate featuresobserved in the range 20 − 70µm,composition well-constrained (picture:

Molster & Kemper 2005, Space Science Review 119, 3)

IAlso oxides and sulphites arepresent (more difficult to detect thansilicates)

I Highly refractory oxides like TiO2, Al2O3, Al2TiO5 are believed tobe nucleation seeds on which the silicate grains growProposed carriers of MIR bands at 13µm, 19.5µm, and 28µm

I Mg and Fe oxides observed in the range 13− 22µm , may forma distinct grain population

I Sulphides common in cometary grains, not yet firmly detected ininterstellar space

Interstellar ices (1)

I In dense clouds the atoms and molecules freeze onto the grainsand form an icy mantle

I Ices are observed in circumstellar envelopes around OH/IRstars, towards field stars, and protostars

I

The richest ice mantles containneutral molecules, H2O, CO2, CO,CH3OH, CH4, H2CO, but also ionssuch as OCN− and NH+

4(picture: Gibb 2000, ApJ 536, 347)

Interstellar ices (2)

I Also more complex species are likely to be present at lowabundances, but their identification is difficult

I CO2 and CO profiles are affected by the ice composition, andthey allow trace segregation into “polar” (molecules with strongH-bonds, H2O, CH3OH, CH4) and “apolar” ices (CO, O2, N2)

I The shape of the OH stretching mode of H2O at ∼ 3µm can beused to examine(i) if the ice is crystalline or amorphous(ii) the grain size (line becomes asymmetric when the grainsgrow).

I There are also other lines that can be used to trace interactionbetween molecules and thus the chemical composition of theices

Very Small Grains and PAHs (1)

I Interstellar carbon can be in the form of graphitic or diamondgrains or amorphous carbon grains, carbon clusters, aromaticmacromolecules, etc.

I Very Small Grains, VSGs - nanometer-sized grains responsiblefor emssion at ∼ 25− 60µm.Stochastic UV heating to high temperatures (several hundreds to∼ 1000 K).Attributed to carbonaceous grains, exact composition not known

I

PAHs - possible carriers of the Aro-matic Infrared Bands (AIB) emis-sion in the MIR (∼ 3− 20µm).Caused by IR-fruorecence of FUV-pumbed PAH molecules containingseveral tens of C-atoms(picture: Dartois et al. 2005, A&A 432, 895)

Very Small Grains and PAHs (2)

I

PAH features: 3.4µm (C-H stretch),6.2µm (C-C stretch), 7.7µm (C-C stretch), 8.6µm (C-H bend inplane), etc.Strong absorption owing to π → π∗

electronic excitation near 2200 Å- alikely contributors to the UV-bump inthe extinction curve at 2175 Å(picture: Boersma et al. 2009, ApJ 690, 1208)

I Contain ∼ 10− 20% of interstellar carbon in the Galaxy

Very Small Grains and PAHs (2)

I Intensity of PAH emission proportional to the UV field - used as ameasure of star formation in external galaxiesImportant for the neutral gas heating and ionization balance inmolecular clouds

I Identification of individual PAH molecules is difficult:-C60 and C70, Cami et al. 2010-Anthracene C14H+

10?, Iglesias-Groth et al. 2010-C60+, Berné et al. 2013

Hydrogenated Amorphous Carbons, HACs

I The interstellar band at 3.4µm observed towards diffuse Galaticlines of sight is attributed to methyl (−CH3) and methylene(−CH2) groupd in the Hyrdogenated Amorphous Carbon, HAC(also designated by a:C-H)3.4µm feature and the AIBs observed in emission in someproto-planetary nebulae, suggested to caused by PAHs withmethyl groupsOther diffuse absorption features 6.85 and 7.25 µm

I Widespread grain population containing ∼ 5− 30% of thecosmic carbon abundance

I HACs may be in the form of a refractory mantle surroundingsilicate grains or it may constitute a distinct grain population

The origin of interstellar dust (1)

I Cosmic dust form in the cool gas outflowing from AsymptoticGiant Branch (AGB) stars, red giants, cool supergiants, andsupernova (SN) explosions

I Most of interstellar gas and dust originates from M-type giants,OH/IR stars, and carbon stars. These are much more commonthan massive stars (M > 10M�)

I Supernovae are important for the production of heavy species(and influence strongly the dynamics of the ISM). Their share ofthe mass of the ISM is, however, only a few percent.

I Dust is formed in dense cools shells created by stellar pulsationsor large-scale convective motions which transfer gas above thestellar surface

The origin of interstellar dust (2)

I In the dusty envelopes of stars the densities andtemperaturesare high enough for condensation of macroscopicdust particles out of free atoms and light molecules

I Dust condensation occurs in the inner parts of the envelope,T ∼ 500− 1500 K

M-tpe giants: SiO, MgS, Mg,Fe→ silicates, i.e., mineralesbuilding around SiO4 chains(picture: Reid & Menten 1997, ApJ 476, 327)

The origin of interstellar dust (2)

I When the envelope grows thick the star becomes a strongsource of thermal dust emission and radio spectral lines

I

A Mira-type variable becomesa so called OH/IR star withmasing OH, H2O, and SiOlines, along with thermal dustand line emission

I The radiation pressure and shock waves related to thepulsation push the envelope outwards (v ≤ 30 km/s)

The origin of interstellar dust (3)

I

Carbon stars: dust productiondriven by carbides (for exampleSiC) and acetylene(C2H2) whichforms chains. Carbon chainscan settle to ring-like structures →PAHsCarbon stars show thermal lines ofCO, CS, SiS, SiC2, etc.(picture: Olofsson et al. 2000, A&A 353, 583)

Plateau de Bure interferometer

I

The expanding dust envelope turnsa planetary nebula which eventuallybecomes merged with the diffuseinterstellar matter.

Ring Nebula of Lyra

Chemical processes in an AGB star

picture: Josef Hron, University of Vienna

Dust formation

I Dust grains are accelerated away from the the star by radiationpressure

I In circumstellar shells, dust forms from the original matter starformed of, and from elements produced by nuclear reactions inthe star

I The composition of dust evolves. In particular the C/O ratio isimportantC/O ≤ 1 (M-stars): all carbon is bound to CO, and excessoxygen goes to oxygen-rich dust productionC/O ≥ 1 (C-stars): carbon is in excess, O bound to CO,hydrocarbons and carbon grains are formedC/O ≈ 1 (S-stars): only small amount of dust is formed

Nucleation and growth of O-rich dust (1)

I The dust formation starts with nucleation, the condensation of astable cluster of molecules from the gas phase

I The grain grows through via the accretion of the ambient gasonto the nucleus

I The condititions for stellar dust formation: high gas temperature,high condensation temperature, low pressure

Nucleation and growth of O-rich dust (2)

I Simple molecules (H2O, CO, SiO, SO2, HCN, SiS) formed close tothe stellar photosphere (1R∗ < R < 5R∗, 2000 K > T > 1000 K)High temperature and temperature, “thermochemistry”, e.g.CO + 3H2 ↔ CH4 + H2ON2 + 3H2 ↔ 2NH3

CO + H2O↔ CO2 + H2

I Dust condensation zone (5R∗ < R < 100R∗,1000 K > T > 100 K): silicates, oxides, molecuiles maycondense on dust

I Circumstellar envelope (100R∗ < R < 20000R∗,100 K > T > 10 K): production of moleucles in ion-moleculereactions, desctruction by photo-dissociation

Nucleation and growth of O-rich dust (3)

I Crystals are condensed directly from the gas phase.One possible nucleation seed is corundum, Al2O3,observed in evolved stars

I According to experiments using solar-like abundances,silicates like forsterite (Mg2SiO4), enstatite (MgSiO3),corundum (Al2O3), spinel (MgAl2O4, ruby), and melilite((Al,Si)SiO7 , with Ca, Na, Mg,Fe) can accrete on the cores

Carbon-rich dust formation

I Chemistry around C-stars is complex because of many differentforms of carbon

I

Dust formation possibly begins withthe nucleation of small PAHs, fol-lowed by reaction with acetylene(C2H2) or the dimerization of PAHs,like pyrene to form larger PAHs andnanograins

I The scenario adopted from soot formation and combustionSo far models of AGB winds combined with the results ofterrestrial soot studies do not reproduce the observedabundances of PAHs and carbonaceous grains

Dust evolution

I Dust evolves as it is exposed to radiation, shocks, and otherprocesses in different phases of the ISM (CSEs, diffuse clouds,dense clouds, protostellar cores, etc)

I The destruction time scale in the ISM is estimated to be ∼ 6 108

and ∼ 4 108 yr for graphite and silicates, respectively.

I Some astronomers believe that most of interstellar dust isformed in situ in the ISM and not in stars. However, nucleation,i.e. the formation of grain cores in the ISM is unlikely.

I Furthermore, it is not sure if the silicate or carbonaceous dustgrains can grow in the ISM via accretion of atoms and moleculesonto pre-existing surfaces (although they can get ice coatings)

Processes modifying grains (1)

I Shock waves or supernovae replenish the ISM in smaller grainsby shattering big grains

I Grain coagulation in shielded regions increases the grain size

I In cold dense clouds atoms and molecules from the gas phasescondense onto grains to form an icy mantle

Processes modifying grains (2)

I Grain and their mantles provide surfaces on which solid-stateand surface chemistry can proceed

I Grain surfaces are exposed to UV photolysis and irradiation bycosmic rays that can induce chemistry leading to formation ofcomplex molecules and refractory carbonaceous mantles

I In hot cores thermal or explosive desorption of icy mantlesenrich the gas phase composition

I In PDRs small hydrocarbons are released into the gas phasethrough photo-evaporation of carbonaceous nanograins andPAHs

Evolution of icy mantles (1)

I Icy mantles evolve via surface reactions, UV photolysis, cosmicray irradiation, or thermal cycling

I These processes change also the gas phase compositionbecause surface molecules desorb from the surfaces (cosmicrays, thermal heating, X-rays)

I UV photolysis and cosmic raysUV photons originate in stellar radiation and in cosmic rayexcitation of H2 (in the interiors of dense clouds)

Evolution of icy mantles (2)

I Experiments have been used to simulate conditions favourablefor the formation of complex species in interstellar ices.Electronic irradiation (mimicking MeV cosmic rays) of ices haveproduced the isomers acetic acid (CH3COOH), methyl formate(HCOOCH3), and glycolaldehyde (HCOCH2OH)These are also produced by UV photolysis of ices together withacetaldehyde (CH3CHO), ethanol (CH3CH2OH), and dimethylether (CH3OCH3)

Surface reactions (1)

I In surface reactions, atomic H and O play a major role

I H abudance is determined by the destruction of H2 by cosmicrays, and is roughly independent of the cloud density n(H) ∼ 1cm−3

I O abundance increases with the cloud density O/H2 ∼ 10−4.

I In dark clouds, the first monolayer of H2O ice is suppoded toforms on bare silicate and carbonaceous grains.

I Water ice can from on grain surfaces through reactions betweenH and O atoms.

Surface reactions (2)

I New molecules from on H2O ice surface: CO2, H2CO, CH3OH,etc.

I Methanol, CH3OH, is one of the molecules the high abundanceof which cannot be explained by gas-phase ion-moleculechemistry

I Thermal annealing (heating resulting in internal stress relief) ofthe grains in protostellar objects can induce further reactions

Surface reactions (3)

picture: Burke & Brown 2010, PCCP 12, 5947

Surface reactions (4)

I The increased mobility of the ice species favours the formationof complex molecules.For example, thermal processing ofH2O : NH3 : H2CO = 10 : 5 : 0.3 ices indicates that ammonia andformaldehyde stars to react at ∼ 80 K to from an organic residuecontaining amonimethanol (NH2CH2OH).Annealing of UV photolyzed ices with oxygen and nitrogen canlead to the formation of prebiotic species.

Desorption

I Thermal desorption works in warm clouds (T ∼ 20− 100 K,depending on species)

I Also in cold dense clouds there must be mechanism returningmolecules to the gas phase, otherwise all molecues shouldfreeze out in ∼ 105 yr (contrary to observations)

I Possible processes:Shattering of grains in shocks occurring in turbulent cloudsDesorption caused by heating caused by cosmic raysUV photodesorption (mainly PDRs and cloud edges, but also ininterior parts owing to secondary photons)

I Energy from UV photons and cosmic rays can also be stored inradicals incorporated in ices. This can lead to explosivedesorption during transient heating

I Small molecules desorb more easily than heavy molecules

Desorption

picture: Pilling et al. 2010, A&A 509, A87

Evolution of carbonaceous dust (1)

I Observations towards photodissociation regions (PDRs) suggestthat a strong UV radiation field can break PAHs and small carbongrains to produce small hydrocarbons, like CCH, c-C3H2, C4H(mid-IR bands attributed to PAHs and these molecules detectedin the same region)

I According to laboratory experiments, small PAHs (< 14C-atoms) fragment mainly into acetylene, C2H2, whereas largerPAHs the carbon ring is destructued only after complete loss ofH atoms

I PAHs themselves may be produced by destruction of nanometriccarbon grains by shock waves and/or UV radiation

Evolution of carbonaceous dust (2)

I Multiwavelength photometric maps of reflection nebulae can beexplained by four template spectra: large multiply ionized PAHs,singly ionized PAHs, neutral PAHs, and carbonaceous grainswith sizes ∼ 1 nm (the smallest VSGs)

I Collisions between PAHs in shielded regions are likely to lead toPAH clustersPAH clusters dissociate in PDRs

Evolution of silicate dust (1)

I The silicate dust injected into the space from evolved stars ispartly crystalline and partly amorphous (the ratio is not known)

I

In the ISM, a large majority (∼98%) of dust is amorphous,so there is likely to be processeswhich destroy crystalline sili-cates or convert them to amor-phous(picture: Spitzer MIR spectra of evolved stars, Jiang et

al. 2013, ApJ 765, 72)

Evolution of silicate dust (2)

I Possible processes destroying crystalline lattices are the cosmicray bombardment or energetic shocks (v > 100 km/s)Also the composition may change in these processes

I Shattering, sputtering, coagulation:at relative velocities exceeding a few km/s grain-grain collisionslead to shattering of large grainssmall grains colliding at low velocities can stick to each other

I Dust coagulation in dense dark clouds is predicted by theoreticalstudies.Far-infrared observations (PRONAOS balloon, Herschel) haveprovided evidence for changes in the dust emission propertiestowards the densest parts of dar clouds. These can beexplained by grain growth.

Evolution of silicate dust (3)

I In protoplanetary disks the large width of the 9.7µm silicatefeature indicates the growth of silicate dust from thecharacteristic size of interstellar dust ∼ 0.1µm up to severalmicrons.

I In protoplanetary disks, large part of the cold dust is againcrystalline and composed mainly of forsterite (Mg2SiO4).The dust must have been processed at a high temperature(probably near the central star).The same observation has been made in cometsThis suggests that radial mixing is operating, transportingcrystallised dust from the vicinity of the star to outer cold regions

Lifecycle of the interstellar matter

picture: SESE Terahertz Laboratory

Cycle of the interstellar matter (1)

I The ionized gas in the expanding atmosphere of an evolved starrecombines and can form simple molecules, e.g. CO, SiO, C2H2

I At a certain distance from the stellar surface, the conditions arefavourable for the condensation of molecular clusters, nucleationPossible nuclei in O-rich stars: TiO2, Al2O3, Al2TiO5

In carbon stars acetylene (C2H2) chains may act as seeds

I Depending on the relative abundances of C and O, the seedsgrow to silicate or carbonaceous grains (a ∼ 0.1µm, actuallysmoke)

Cycle of the interstellar matter (2)

I When the circumstellar shell dilutes to a planetary nebula, and isexposed to the interestellar radiation field, most molecules aredissociated and the atoms are ionized.Dust grains survive but are processed by the ISRFGas and dust are well mixed

I Large scale processes drive the matter into clouds with clumpystructuresSelf-shielding against the UV radiation makes H neutral in clouds

I Cooling by atomic fine structure lines (mainly C+, O, Si+, andlater C) decreases the temperatures to a level where H2 can beproduced efficienty on grain surfaces (T < 20 K).

Cycle of the interstellar matter (3)

I Ion-molecule chemistry starts to form simple moleculesThe kinetic temperature decreases with the cloud densitybecause of a more efficient cooling (molecules) and theattenuation of the ISRF (dust converts energetic radiation tofar-infrared which escapes freely)

Cycle of the interstellar matter (4)

I In their densest parts, clouds start to feel their gravity and theyfragment into dense cores

I

Starless dense cores are typicallycold, T ∼ 5− 10 KThe grains coagulate and accreteicy mantlesMolecules are depleted in the gasphase(picture: Willacy et al. 1998, ApJ 507, L171)

Cycle of the interstellar matter (5)

I The collapse of a dense core is nearly isothermal until theformation of the so called first hydrostatic core, which quicklycollapses to a protostar

I The protostar stars to heat the surrounding core by radiation andshocks,and the core becomes warm T ∼ 100K , “hot core”

I The evaporation of icy mantles and subsequent gas-phasereactions produce complex molecules into the gas phase

I Part of the core material is survived in the proto-planetary diskand will be incorporated in planets an smaller bodies of theplanetary system

I The outer part of the protoplanetary disks can remain as areservoir of interstellar ices, which can be transported to innerparts by comets