Drake G.W.F. (editor) Handbook of Atomic, Molecular, and Optical Physics

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Part F 81

1233

Application

Part G

Part G Applications

82 Applications of Atomic and Molecular

Physics to Astrophysics

Alexander Dalgarno, Cambridge, USA

Stephen Lepp, Las Vegas, USA

83 Comets

Paul D. Feldman, Baltimore, USA

84 Aeronomy

Jane L. Fox, Dayton, USA

85 Applications of Atomic and Molecular

Physics to Global Change

Kate P. Kirby, Cambridge, USA

Kelly Chance, Cambridge, USA

86 Atoms in Dense Plasmas

Jon C. Weisheit, Pullman, USA

Michael S. Murillo, Los Alamos, USA

87 Conduction of Electricity in Gases

Alan Garscadden, Wright Patterson Air Force

Base, USA

88 Applications to Combustion

David R. Crosley, Menlo Park, USA

89 Surface Physics

Erik T. Jensen, Prince George, Canada

90 Interface with Nuclear Physics

John D. Morgan III, Newark, USA

James S. Cohen, Los Alamos, USA

91 Charged-Particle–Matter Interactions

Hans Bichsel, Seattle, USA

92 Radiation Physics

Mitio Inokuti, Argonne, USA

1235

Applications o

82. Applications of Atomic and Molecular Physics

to Astrophysics

The range of physical conditions of density,

temperature, and radiation ﬁelds encountered in

astrophysical environments is extreme and can

rarely be reproduced in a laboratory setting. It is

not only reliable data on known processes that

are needed but also a deep understanding so that

the relevant processes can be identiﬁed and the

inﬂuence of the conditions in which they occur

fully taken into account.

We present here a summary of the processes

that take place in photoionized gas, collisionally

ionized gas, the diffuse interstellar medium,

molecular clouds, circumstellar shells, supernova

ejecta, shocked regions and the early Universe.

82.1 Photoionized Gas ................................1235

82.2 Collisionally Ionized Gas .......................1237

82.3 Diffuse Molecular Clouds ......................1238

82.4 Dark Molecular Clouds ..........................1239

82.5 Circumstellar Shells

and Stellar Atmospheres ......................1241

82.6 Supernova Ejecta .................................1242

82.7 Shocked Gas........................................1243

82.8 The Early Universe ...............................1244

82.9 Recent Developments...........................1244

82.10 Other Reading .....................................1245

References ..................................................1245

Almost all our information about the Universe reaches

us in the form of photons. Observational astronomy is

based on measurements of the distribution in frequency

and intensity of the photons that are emitted by as-

tronomical objects and detected by instrumentation on

ground-based and space-borne telescopes. Information

about the earliest stages in the evolution of the Uni-

verse before galaxies and stars had formed is carried

to us by blackbody background photons that attended

the beginning of the Universe. The photons that are the

signatures of astronomical phenomena are the result of

many processes of nuclear physics, plasma physics and

atomic, molecular and optical physics. The processes

that modify the photons on their journey from distant

origins through intergalactic and interstellar space to

the Earth belong mostly to the domain of atomic, mo-

lecular, and optical physics, as do the instruments that

detect and measure the arriving photons and their spec-

tral distribution. The spectra are used to classify galaxies

and stars and to identify the astronomical entities and

phenomena such as quasars, active galactic nuclei, grav-

itational lensing, jets and outﬂows, pulsars, supernovae,

novae, supernova remnants, nebulae, masers, protostars,

shocks, molecular clouds, circumstellar shells, accretion

disks and black holes.

Quantitative analyses of the spectra of astronom-

ical sources of photons and of the atomic, molecular,

and optical processes that populate the atomic and mo-

lecular energy levels and give rise to the observed

absorption and emission require accurate data on transi-

tion frequencies and wavelengths, oscillator strengths,

cross sections for electron impact, rate coefﬁcients

for radiative, dielectronic and dissociative recombina-

tion, and cross sections for heavy particle collisions

involving charge transfer, excitation, ionization, dissoci-

ation, ﬁne structure, and hyperﬁne structure transitions,

collision-induced absorption and line broadening. Data

on radiative association and ion–molecule and neutral

particle reaction rate coefﬁcients are central to the in-

terpretation of measurements of chemical composition

in molecular clouds, circumstellar shells and supernova

ejecta.

82.1 Photoionized Gas

The Universe contains copious sources of energetic pho-

tons most often in the form of hot stars, and much of

the material of the Universe exists as photoionized gas.

Photoionized gas produces the visible emission from

emission nebulae, planetary nebulae, nova shells, star-

burst galaxies and probably active galactic nuclei [82.1].

Part G 82

1236 Part G Applications

Emission nebulae are extended regions of luminosity

in the sky. They arise from the absorption of stellar radi-

ation by the gas surrounding one or more hot stars.The

gas is ionized by the photons and excited and heated by

the electrons released in the photoionizing events. A suc-

cession of ionization zones is created in which highly

ionized regions give way to less ionized gas with in-

creasing distance from the central star as the photon ﬂux

is diminished by geometrical dilution and by absorp-

tion. The outer edge of a nebula is a front of ionization

pushing out into the neutral interstellar gas. The dens-

ities are typically between 100 and 10 000 cm

−3

and the

temperatures between 5000 and 15 000 K. Nebulae are

also called Hii regions. At low densities, the luminosity

is low, but the ionized regions can still be detected by

radio observations.

Planetary nebulae are smaller in extent and more

dense. They have a passing similarity in appearance to

planets. Planetary nebulae are produced by the photoion-

ization of shells of gas that have been ejected from the

parent star as it evolved to its ﬁnal white dwarf stage.

Because the core of the parent star is very hot the ir-

radiated gas is more highly ionized than are emisssion

nebulae and has a distinctive spectrum.

Photoionized gas is also found around novae. No-

vae are stars that have undergone spasmodic outbursts

and they are surrounded by faint shells of ejected

gas, photoionized and excited by the stellar radiation.

Some supernova remnants, which are what remains after

17 16 15 14 13 12

Energy (eV)

700 800 900 1000 1100

Counts / 320 s

Fe XVII

–2p

Fe XVII

–2p

VIIIO

Ly β Ly γ Ly δ

Ne IX

–1s 2p, s

Fe XVII

–2p

Ne X

Ly α

Ne IX

–1s 3p

Wavelength (Å)

Fig. 82.1 X-ray spectrum of the supernova remnant Puppis A as observed by the Einstein satellite. Note the high level of

ionization with hydrogen-like ions of oxygen and neon, suggesting a high temperature. After [82.2]

a massive star has exploded, have spectra that also ap-

pear to be emanating from photoionized gas. The source

of ionization may be synchrotron radiation. Figure 82.1

shows the X-ray emission spectrum of a supernova

remnant.

The nuclei of starburst galaxies have spectra like

those of emission nebulae. They result from gas pho-

toionized by radiation from hot stars created in a period

of rapid star formation. Active galactic nuclei, such

as quasars, have a different spectrum characterized

by broad lines indicating a large range of velocities.

Photoionized gas is the most likely interpretation.The

ionizing source may be an accretion disk around a com-

pact object such as a black hole.

The ionization structure in a photoionized gas is

determined by a balance of photoionization

(m−1)+

+ hν → X

+ e

−

(82.1)

and radiative

+ e

−

→ X

(m−1)+

+ hν (82.2)

and dielectronic

+ e

−

→



(m−1)+



∗

→ X

(m−1)+

+ hν (82.3)

recombination, and in plasmas with a signiﬁcant popula-

tion of neutral hydrogen and helium, by charge transfer

recombination

+ H → X

(m−1)+

+ H

(82.4)

+ He → X

(m−1)+

+ He

. (82.5)

Part G 82.1

Applications of Atomic and Molecular Physics to Astrophysics 82.2 Collisionally Ionized Gas 1237

Many detailed calculations of the ionization structure of

photoionized regions have been carried out [82.1].

The ionizing source spectra of hot stars can

be obtained from calculations of stellar atmospheres

Sect. 82.5. Approximate values of cross sections for

photoionization for a wide range of atomic and ionic sys-

tems in many stages of ionization are available [82.3–6].

Calculations of higher precision and reliability that

incorporate the contributions from autoionizing reso-

nance structures exist for speciﬁc systems [82.7]. They

are undergoing continual improvements as increasingly

powerful computational techniques are brought to bear

on the calculations.

The cross sections for radiative recombination are

obtained by summing the cross sections for capture into

the ground and excited states of the recombining sys-

tem. Because of the contribution from highly excited

states which are nearly hydrogenic, the rate coefﬁcients

are similar for different ions of the same excess nu-

clear charge. They vary slowly with temperature. In

contrast, dielectronic recombination is a speciﬁc process

whose efﬁciency depends on the energy level positions

of the resonant states. For nebular temperatures, the rate

coefﬁcients vary exponentially with temperature. Ex-

plicit calculations have been carried out for many ionic

systems [82.8–10]. Because the photoionization cross

sections of the major cosmic gases hydrogen and helium

diminish rapidly at high frequencies, multiply charged

ions and neutral gas coexist in cosmic plasmas produced

by energetic photons and charge transfer recombina-

tion may control the ionization structure. For multiply

charged ions with excess charge greater than two, charge

transfer is rapid. For doubly charged and singly charged

ions, the cross sections are sensitive to the details of the

potential energy curves of the quasimolecule formed in

the approach of the ion and the neutral particle. Few re-

liable data exist. Some recent calculations may be found

in the papers [82.11–13].

Photoionized gas is heated by collisions of the en-

ergetic photoelectrons and cooled by electron impact

excitation of metastable levels, principally of O

and

and N

,andS

and S

, followed by

emission of photons which escape from the nebula. Con-

siderable attention has been given to the determination

of the rate coefﬁcients [82.14]. The resulting cooling

rates increase exponentially with temperature and keep

the temperature of the gas between narrow limits. Some

contribution to cooling occurs from recombination and

from free-free emission by electrons moving in the ﬁeld

of the positive ions.

The luminosity of the photoionized gas comes from

the photons emitted in the cooling processes and from

radiative and dielectronic recombination. The radiative

recombination spectrum of hydrogen extends from the

Ly α line at 121.6 nm to radio lines at meter wavelengths.

The recombination spectrum can be predicted to

high accuracy, and calculations for a wide range of

temperature, density, and radiation environments have

been carried out for diagnostic purposes [82.15–18].

Electron impact and proton impact induced transitions

are important in determining the energy level popula-

tions and the resulting spectrum. Stimulated emission

often affects the intensities of the radio lines, espe-

cially those from extragalactic sources. Comparisons of

the predicted intensities in the visible and infrared with

theoretical predictions yield information on interstellar

extinction in the nebula and along the line of sight.

The relative intensities of the lines emitted by dif-

ferent metastable levels depend exponentially on the

temperature. The relative intensities of the lines at

500.7 nm and 436.3 nm originating in the

and

levels of O

vary as exp(33170/T ), and are commonly

used to derive the temperature T.

The electron density can be inferred from the lines

emitted from neighboring levels with different radiative

lifetimes for which there occurs a competition between

spontaneous emission and quenching by electron im-

pact. There are many possible combinations of lines.

The lines at 372.89 nm and 372.62 nm emitted by the

3/2

and

5/2

levels of N

are readily observable

and their relative intensity yields the electron density.

Radiative and dielectronic recombination lines are

often seen in the spectra, as are a few lines due to charge

transfer recombination. Fluorescence of starlight and

resonance ﬂuorescence of lines emitted in the nebula

(called Bowen ﬂuorescence by astronomers) also con-

tribute to the spectra of photoionized gases. Many data

are needed to adequately interpret the observations.

82.2 Collisionally Ionized Gas

Hot gas is found in the coronae of stars and particu-

larly the Sun, and in young supernova remnants, in

the hot phase of the interstellar medium, and in inter-

galactic space. In a hot gas the ionization is produced

by the impact ionization of the fast thermal electrons

and recombination is radiative and dielectronic [82.19].

Part G 82.2

1238 Part G Applications

The rate coefﬁcients for electron impact ionization and

for recombination for any given ionization stage are

functions only of temperature, and hence so is the

resulting ionization distribution. When ionization and

recombination balance, coronal equilibrium is attained

in which the ionization structure is speciﬁed by the

temperature.

Recombination at high temperatures is dominated

by dielectronic recombination. At high temperatures,

dielectronic recombination is stabilized by transitions

in which the core electrons are the active electrons.

The associated emission lines lie close in frequency to

that of the resonant transition of the parent ion. They

are called satellite lines. Together with lines generated

by electron impact excitation, they provide a powerful

diagnostic probe of density and temperature. In many

circumstances such as in supernova remnants, coronal

equilibrium does not hold, and the ionization and re-

combination must be followed as functions of time. The

temperature also evolves as the hot plasma is cooled by

electron impact excitation and ionization.

The recombining gas produces X-rays and extreme

UV radiation which modify the ionization structure.

There is a particular need for more reliable data on

high energy photoionization cross sections, on collision

cross sections for electrons and positive ions, and on

4.0 5.0 6.0 7.0

–log P/n

(erg × cm

/s)

log T(K)

(OII)

Fig. 82.2 Total emissivity and emissivity by element as

a function of temperature in coronal equilibrium. The heavy

solid curve is total emissivity and the lighter lines are

contributions from individual elements. After [82.20]

the energy levels and transition probabilities of highly

stripped complex ions. Figure 82.2 shows the emissivity

of coronal gas.

82.3 Diffuse Molecular Clouds

Diffuse molecular clouds are intermediate between the

hot phase of the galaxy and the giant molecular clouds

where much of the gas resides. They are called diffuse

because they have optical depths of order unity, so pho-

tons can penetrate from outside the cloud and affect

the chemical composition. The atoms and molecules

are observed in absorption against background stars.

Translucent clouds with optical depths between about

2 and 5 are intermediate between diffuse and dark clouds

where photons from the outside still affect the chem-

istry. They can be observed both in absorption against

a background source or in emission in the radio.

The temperature is 100–200 K at the edges of

a diffuse cloud with a density of about 100 cm

−3

.In

a typical diffuse cloud the temperature decreases to

about 30 K at the center while the density increases to

about 300–800 cm

−3

. The chemistry is driven by ion-

ization from interstellar UV photons and from cosmic

rays.

Interstellar UV photons ionize species which have

ionization potentials less then that of atomic hydrogen.

Atomic hydrogen is so pervasive in the galaxy that UV

photons with energies higher than 13.6 eV are absorbed

near the source. The UV ﬂux is a very important param-

eter in determining the composition of a diffuse cloud.

Photodissociation provides destruction which limits the

buildup of more complex species and so diffuse clouds

are dominated by simpler diatomic species.

Species with ionization potentials greater then hy-

drogen are mainly ionized by cosmic rays. Cosmic rays

are high energy nuclei which stream through the galaxy.

The cosmic ray ionization rate, the number of cosmic

ray ionizations per second per particle, is an important

parameter in interstellar chemistry. A lower limit to the

cosmic ray ionization rate may be set by measured high

energy cosmic rays reaching earth, giving an ionization

rate of ≈ 10

−17

−1

. More realistic estimates of the cos-

mic ray ionization rate from looking at recombination

lines suggest values of a few ×10

−17

−1

The hydroxyl radical OH is produced in a man-

ner similar to that discussed below in Sect. 82.5,and

removed by photodissociation. Thus in diffuse clouds,

Part G 82.3

Applications of Atomic and Molecular Physics to Astrophysics 82.4 Dark Molecular Clouds 1239

OH may be used to measure the cosmic ray ionization

rate, subject to uncertainties in the OH photodissociation

rate and the H

recombination rate. The OH abun-

dances give rates of several ×10

−17

−1

for many diffuse

clouds.

The carbon chemistry begins with the ionization of C

by UV photons:

C+ hν → C

+ e

−

. (82.6)

The carbon ion cannot react directly with H

+ H

→ CH

+ H , (82.7)

as this reaction is exothermic by 0.4 eV. Instead, the

chemistry proceeds by the slow radiative association

process

+ H

→ CH

+ hν. (82.8)

The CH

ion may either dissociatively recombine

+ e

−

→ CH+ H , (82.9)

or react with molecular hydrogen

+ H

→ CH

+ H . (82.10)

The CH

then undergoes dissociative recombination

+ e → C+ H

+ H (82.11)

→ CH+ H + H (82.12)

→ CH

+ H (82.13)

→ CH+ H

, (82.14)

where the products are listed in order of decreasing

likelihood. The CH is removed by photodissociation

CH + hν → C+ H

(82.15)

and by photoionization

CH + hν → CH

+ e

−

. (82.16)

CH may also be removed by reactions with oxygen or

nitrogen atoms to form CO and CN respectively.

One of the outstanding problems in diffuse clouds

is to understand the large abundance of CH

relative to

CH. The problem is producing the CH

without pro-

ducing additional CH. Since most reaction paths go

through reaction (82.7), this is the most likely candi-

date. What is needed is some extra energy to overcome

the endothermicity. This energy must come from ei-

ther hot C

or from hot or vibrationally excited H

The most popular model is gas heated by a shock,

possibly a magnetic shock in which ions stream rela-

tive to the neutrals, giving a high effective energy.

Unfortunately, though these shock models can repro-

duce the CH

abundances, they also predict relative

velocities between the CH

and CH which are not

often observed. Recently there have been suggestions

that turbulence in the cloud could account for the CH

abundance.

The most comprehensive models of diffuse clouds

are by van Dishoeck and Black [82.21]. A collection

of photodissociation rates and photoionization rates is

given in Roberge et al. [82.22].

The UV ﬂux is predominantly from stars and may be

as much as 10

times larger near an Hii region Sect. 82.1

than it is in the general interstellar medium. Regions

in which the chemistry is dominated by photons are

referred to as photon dominated regions or photodis-

sociation regions (PDR’s). In the presence of high UV

ﬂux the cloud is much warmer than in a typical diffuse

cloud. Temperatures may reach 1000 K near the edge

of the cloud and 100 K far into the cloud. The chem-

istry differs from traditional diffuse cloud chemistry

in that the high temperatures allow endothermic reac-

tions to proceed. Sternberg and Dalgarno [82.23]have

published a comprehensive model of photodominated

regions.

82.4 Dark Molecular Clouds

Much of the mass of the galaxy is in the form of dark mo-

lecular clouds. The molecular clouds are sites of forming

new stars. They are composed primarily of hydrogen,

with about 10% helium and trace amounts of heavier

elements. They have densities of approximately 10

−3

and temperatures between 10 and 20 K, and

often contain denser clumps. The clouds are optically

thick and so photons from the outside are absorbed on the

surface of the clouds. The interiors are heated and ion-

ized by cosmic rays which penetrate deep into the cloud.

The temperatures are too low to sustain much neutral

chemical activity in the clouds, and cosmic ray ion-

ization is important in driving the chemistry. In dense

clouds, the cosmic rays both initiate the chemistry and

limit it through the production of He

and through

cosmic ray induced photons.

Table 82.1 is a list of molecules that have been ob-

served in the interstellar medium, many of which have

also been found in other galaxies. It is likely that all

but H

are formed in the gas phase by ion–molecule re-

Part G 82.4

1240 Part G Applications

Table 82.1 Molecules observed in interstellar clouds

Hydrogen CH Methylidyne

Methylidyne ion OH Hydroxyl

Carbon CN Cyanogen

CO Carbon monoxide CO

Carbon monoxide ion

NH Amidogen NO Nitric oxide

CS Carbon monosulphide SiO Silicon monoxide

SO Sulphur monoxide SO

Sulphur monoxide ion

NS Nitrogen sulphide SiS Silicon sulphide

PN Phosphorus nitride HCl Hydrogen chloride

SiN Silicon nitride NH

Amino radical

O Water C

H Ethynyl

HCN Hydrogen cyanide HNC Hydrogen isocyanide

HCO Formyl HCO

Formyl ion

Protonated nitrogen H

S Hydrogen sulphide

HNO Nitroxyl OCS Carbonyl sulphide

Sulphur dioxide HCS

Thioformyl ion

O Carbon suboxide C

S Dicarbon sulphide

O Nitrous oxide H

CN Methylene amidogen

CO Formaldehyde H

CS Thioformaldehyde

Ammonia HCNS Isothiocyanic acid

HNCO Isocyanic acid HOCO

Protonated carbon dioxide

H Propynylidyne C

N Cyanoethynyl

S Tricarbon sulphide C

O Tricarbon monoxide

Acetylene H

Hydronium ion

HCNH

Protonated hydrogen cyanide C

Cyclopropenylidene

Methane H

CCC Propadienylidene

HCOOH Formic acid CH

CO Ketene

N Cyanoacetylene HNCCC Cyanoacetylene isomer

HCCNC Ethynyl isocyanide C

H Butadinyl

CH Cyanamide CH

CN Cyanomethyl radical

NH Methanimine CH

CH Methyl cyanide

CCCC Butatrienylidene CH

SH Methyl mercaptan

H Pentynylidyne HCC

HO Propynal

OH Methyl alcohol HC

Protonated cyanoacetylene

CHO Formamide CH

H Methyl acetylene

CHCN Vinyl cyanide HC

N Cyanodiacetylene

H Hexatrinyl CH

Methylamine

CHO Acetaldehyde HCOOCH

Methyl formate

N Methyl cyanoacetylene CH

H Methyl diacetylene

O Dimethyl ether CH

CN Ethyl cyanide

N Cyanohexatriyne CH

OH Ethyl alcohol

N Cyano-octatetra-yne HC

1N Cyano-decapenta-yne

action sequences initiated by cosmic ray ionization. The

fact that isomers such as HCN and HNC are seen in

approximately equal abundances suggests a low dens-

ity gas phase environment. Reactions on surfaces and

the formation of grains are not well understood, but are

surely important.

The chemistry of molecular clouds is dominated by

ion–molecule reactions driven by cosmic ray ionization.

Part G 82.4

Applications of Atomic and Molecular Physics to Astrophysics 82.5 Circumstellar Shells and Stellar Atmospheres 1241

The cosmic rays primarily ionize H

→ H

+ e

−

, (82.17)

producing both H

and fast electrons. The fast electrons

produce additional ionizations. The H

quickly reacts

with H

to form H

+ H

→ H

+ H . (82.18)

The H

reacts with other species by proton transfer,

which then drives much of the interstellar chemistry.

As an example of the production of more complex

molecules in interstellar chemistry, we examine the re-

action networks leading to the production of water H

and the hydroxl radical OH. The H

ions formed by

cosmic ray ionization react with atomic oxygen to form

O+ H

→ OH

+ H

, (82.19)

which quickly reacts with H

to form H

in an

abstraction sequence

+ H

→ H

+ H , (82.20)

+ H

→ H

+ H . (82.21)

The H

then undergoes dissociative recombination

to form water and OH

O++e

−

→ H

O+ H , (82.22)

→ OH+ H

. (82.23)

The water is removed by reactions with neutral or

ionized carbon, which eventually lead to the produc-

tion of CO. OH is primarily removed by reactions with

atomic oxygen leading to O

.TheCOandO

are re-

moved by reactions with He

.TheHe

, generated by

cosmic ray ionization of helium, does not react with H

and so is available to remove species by reactions such

CO + He

→ C

+ O+ He . (82.24)

Water and OH are also removed by UV photons gen-

erated within the cloud. The clouds are too thick for

external UV photons to penetrate, but cosmic rays ex-

cite H

into electronically excited states which decay

through emission of UV photons. These internally gen-

erated photons play an important role in determining the

composition of the cloud. Gredelet al. [82.24]havecom-

piled a list of the photodissociation and photoionization

rates for cosmic ray induced photons.

Modern chemical networks for molecular clouds

include several hundred species and several thousand

reactions. A standard set of reaction rates is provided by

the UMIST (University of Manchester Institute of Sci-

ence and Technology) dataset [82.25, 26]. The dataset

may be obtained from the UMIST Astrophysics Group

homepage (http://saturn.ma.umist.ac.uk:8000/).

The clouds also contain dust particles as evidenced

by the extinction curves for clouds and the observed de-

pletions of heavier elements. The importance of surface

chemistry on these dust particles to interstellar clouds

is still uncertain. Dust particles are the best candidate

to be the site of formation of molecular hydrogen, be-

cause known gas phase reactions fail to produce H

the quantities observed.

82.5 Circumstellar Shells and Stellar Atmospheres

The continuum emission from a star is very nearly that

of a blackbody. This emission is then absorbed and re-

distributed by the atmosphere of the star. The spectrum

of the star is thus determined by its atmosphere. In the

hottest stars, most of the material is ionized and the ab-

sorption lines are predominantly those of ions, while in

the coldest stars, molecular lines are prominent. Kurucz

has calculated models with continuum spectra and the

inclusion of a large number of absorption lines [82.27].

There are two major projects for calculating the required

atomic data. In 1984 an international collaboration

named the Opacity Project was set up to calcu-

late accurate atomic data needed for opacity calcula-

tions [82.28,29]. The other earlier project is called

OPAL. The two sources of data are compared

in [82.30,31].

Low and intermediate mass stars eject circumstellar

envelopes in their red giant phase near the end of their

evolution. Circumstellar envelopes are an important part

of astronomy and they are a likely location for dust

formation. They provide an interesting environment for

studying molecules because they represent a transition

between very high density stellar atmosphere environ-

ments to low density interstellar environments. These

objects evolve to become planetary nebulae Sect. 82.1.

We are fortunate in having one example, IRC 10216,

which is very close to the Sun. The brightest 10 µm

source beyond the solar system, IRC 10216 is a carbon-

Part G 82.5

1242 Part G Applications

rich star surrounded by dust and gas it ejected in a strong

stellar wind. The central star is so shielded that it is

almost undetectable at optical wavelengths, and was not

discovered until the 2 µm survey. IRC 10216 is where

most of the circumstellar molecules are detected, and

has greatly increased our understanding of circumstellar

envelopes.

The envelopes are ejected by the red giant in its

ﬁnal phase of evolution. The mass loss rates increase

to ∼ 10

−4

solar masses per year and temperatures in

the envelope are of order 1000 K. Close to the star, the

density is high and the chemistry is characteristic of

thermal equilibrium. The situation is quite unlike any

interstellar environments. For example in IRC 10216,

the HNC is over one hundred times less abundant than

HCN, whereas in molecular clouds, they have about the

same abundance.

If the star is oxygen-rich, large amounts of H

Oare

formed and if it is carbon-rich, C

.Thishightem-

perature environment forms both molecules and grains.

The high obscuration of the central source indicates that

grains are formed in these envelopes. Polycyclic aro-

matic hydrocarbons (PAH’s) or some similar species are

observed in carbon rich planetary nebulae. These large

molecules must have been produced when the object

was a carbon rich circumstellar envelope.

As the material ﬂows out from the star the den-

sity and temperature decrease. As the density becomes

lower, three body reactions become less important and at

some point the products of these reactions are frozen out

in a similar manner to the evolution of molecules in the

early universe (Sect. 82.10). In the outermost portions of

the circumstellar envelope, molecules are dissociated by

interstellar UV photons. The penetrating UV radiation

is shielded by dust, H

, and CO. The relative abun-

dances can vary rapidly with radius, and observations

provide abundance and radial distribution, a wealth of

data for modelers. Circumstellar chemistry is reviewed

by Omont [82.32] and recent chemical models are given

in [82.33–35].

82.6 Supernova Ejecta

A supernova, the explosion of a massive star following

core collapse, is one of the most spectacular displays in

the Universe. The explosion occurs when the iron core

of a massive star collapses to form a neutron star and the

rebound shock and neutrino ﬂux eject the outer portion

of the star. The ejected portions of the star are rich in

heavy elements produced in the interior of the progenitor

star.

We are fortunate to have had in our lifetime a super-

nova which was close and in an unobscured line of

sight. Supernova 1987A, the ﬁrst supernova observed in

1987, went off in the Large Magellanic Cloud, a small

satellite galaxy to our own. It was the ﬁrst supernova

visible to the naked eye in nearly 400 years (since

the Kepler supernova in 1604). Using the full range of

modern astronomical instruments has allowed us to get

detailed spectra of the evolving ejecta which has greatly

increased our knowledge of supernovae. We will use

SN1987A as an example of supernova ejecta.

Initially the temperature of the ejecta of SN1987A

was high, ≈ 10

K, but it quickly cooled through adi-

abatic expansion and radiation from the photosphere.

The temperature leveled off at several thousand degrees

because of heating by radioactive nuclei, ﬁrst

Ni and

then

Co, formed in the explosion. The dynamics is

homologous free expansion: the velocity scales linearly

with the radius r(t) = vt where v the velocity and t the

time since the explosion.

The ejecta at ﬁrst were optically thick and the

spectrum resembled that of a hot star continuum with

absorption lines from the surface. After a few days, the

temperature dropped, but the ejecta remained optically

thick and continued to show strong continuum emission.

As the ejecta expand, the temperature drops, the ejecta

become optically thin, and the spectrum is dominated by

strong emission lines, superﬁcially resembling an emis-

sion nebula Sect. 82.2. The emission is dominated by

neutral atoms and singly ionized species.

The gas is heated and ionized by the gamma rays

from radioactive decay. The gamma rays Compton scat-

ter, producing X-rays and fast electrons. The X-rays

further ionize the gas and produce multiply charged ions

through the Auger process. These multiply charged ions

recombine through charge transfer with neutral atoms.

Further charge transfer determines the relativeionization

of different species, with the lowest ionization poten-

tial species more ionized than the higher ionization

potential species. The development of the infrared and

optical spectrum of Supernova 1987A has been recently

reviewed by McCray [82.36].

One of the great surprises in the spectrum of Super-

nova 1987A was the discovery of molecules in the

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