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P. C. MYERS

E., andhild,

R.

ion.THE ATOMIC COMPONENT

ted paper Shriniva$ R. Kulkarnit & Carl Heiles*

ional Park. tDepartment of Astronomy, 105-24r. Caltech, Pa$adena, CA 91125, U. S. A.

A.U.ht. *Department of Astronomy.c .Univ. of California, Berkeley, CA 94720, U. S. A.

Black andABSTRACT. We review the physical conditions and the distribution of the three phases)

Ap. J. which consitute the atomic component: the cold clouds, the warm, neutral medium andthe warm, ionized medium. Together, these three pha$es occupy about ~ 40% of theinterstellar volume and contain half the interstellar mass. The size of the H I disk is

~tt.) 295 comparable or exceeds that of the stellar disk. In the outer Galaxy, a spiral pattern is, , clearly discernable. We discuss extensively the distribution of the three phases near the sun

d L as well a$ the physics and the limitations of various probes that have been used to study ther, .atomic component. The ionization of the pha$es is studied in depth and we conclude that

most of the interstellar electrons reside in the warm, ionized medium. We review variousdeterminations of interstellar pressure and its scale height. Finally, we derive the fillingfactors of the atomic phases a$ a function of z. [Unless otherwise stated, .Ro = 10 kpc.]

Introd uction,

The interstellar medium (ISM) is composed of three main components: the ~olecular,the atomic and the hot component. 'These are acted upon by cosmic rays, electromagtleticradiation, gravity and magnetic field. The atomic component is composed of neutral atomichydrogen (H I) and the diff-lIse, Ionized hydrogen (H II). The hot component, while certainly'atomic', follows such a different cycle a$ compared to H I and the diffuse H It that it is

considered to be an entirely different component. 1The Galaxy contains about 4.8 X 109 M0 of H I (Henderson, Jackson and Kerr 1982). '

Estimates for the total amount of the H2 range from one nearly equal to the H I estimate(Scoville and Sanders, this volume), to one that is only 25% of the H I estimate {Bloemen,this volume). The surface density distribution of H I is roughly constant from about 4 kpcto 20 kpcj however H I dominates H2 in ma$s beyond Galactocentric radius 8 kpc (Blitz,Fich and Kulkarni 1983). Not much is known about the radial distribution and the massof the diffuse, ionized hydrogen. Ba$ed on extrapolation of local data, we estimate at leasttv 109 M0 for the total ma$s of the diffuse ionized ga$. Estimates of the fraction of Galactic

,. interstellar space occupied by the atomic component, range from 40% to 80%. All these; factors establish the preeminence of the atomic component in the dynamics and evolution

.of the ISM....87

D. J. Hollenbach and H.'A. Thronson, Jr. (eds.J, Interstellar Processes, 87-/22.@ 1987 by D. Reidel Publishing Company.

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88 S. R. KULKARNI AND C. HEILES

The mass estimates of the interstellar components depend upon the assumed rotationcurve and ~, the distance of the Sun to the Galactic center. The above estimates assume:~ = 10 kpc, the standard Schmidt rotation curve for the inner Galaxy and a flat rotationcurve for the outer Galaxy. Recent analyses suggest ~ '" 8.5 kpc. To first order, the massestimates will have to be decreased by a factor'" (8.5/10)2 = 0.72. After correcting fordifferences in ~, the atomic component constitutes 5% of the mass of the visible matter(Bahcall, Schmidt and Soneira 1983).

This review is not an exhaustive review. Much of the material presented here isadapted from a review on a similar topic that we wrote (Kulkarni and Heiles 1986). Thereader is advised to consult the aforementioned review for a more detailed treatment andSpitzer (1978) for standard derivations not presented here.

The atomic component is thought to consist of three phases: (1) cold (T '" 80 K),dense, neutral H I (cold medium or CM), (2) warm (T '" 8000 K), neutral H I (the warm,neutral medium or WNM) and (3) a warm (T '" 8000 K), ionized hydrogen (the warm,ionized medium or WIM). The CM is supposed to be distributed in the form of discretestructures, the so-called 'diffuse clouds'. In contrast, the WNM is widely distributed sinceit can be seen along most lines of sight. The cold clouds, the warm neutral and the ionizedphases and the hot component (see Savage, this issue) are supposed to be in rough pressureequilibrium. Whether the interstellar space is largely occupied by the hot component or bythe warm, atomic phases is uncertain and controversial. So is the relation of the WNM tothe CM. Finally, the topology of all the phases is unknown.

1. Probes of the Atomic Component.

We now discuss the probes employed to study the atomic phases. Our knowledge ofthe properties of the CM and the WNM have mainly resulted from the study of the 21-cmline of HI (§1.1). The CM and the WNM are distinguished by H I absorption observations.The limitations of such observations are discussed in §1.2. The WIM can be studied usingmany probes: Ha emission, pulsar dispersion measures and the optical metastable lines ofsulfur, oxygen, nitrogen etc. and these are briefly discussed in §1.3.

1.1 The 21-cm Line of H [.

In the H I atom, the magnetic moment of t~e protcn interacts with the combinedmagnetic field generated by the crbiti~g electrcn and the magnetic moment of the electron;this interaction leads to th~ 'hyperfine' splitting of all the energy levels. The 21-cm line isthe result of the hyperfine splitting of the ground state of H I. This hyperfine transition,be.ing a magnetic dipole tranEition, is a forbidden transition. This, together with the lo\vfrequency of the 21-cm line, makes the Einstein A-coeffecient exceedingly small: 2.85 X 10-15

-1S .Consider an H I cloud in the ISM. Let nl and n2 be the density of the H I atoms in

the lower and upper hyperfine states, respectively. We define, Tz, the so-called 'excitationtemperature' by

~ = e-E/(kTs). (1.1-1)nl/91

For the 21-cm line, the statistical weights 91 and 92 are 1 and 3, respectively. The energyseparation E/k is tiny: a mere 70 mK. In keeping with general convention for the 21-cmline, we will henceforth refer to the excitation temperature as the spin temperature (T.).

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THE ATOMIC COMPONENT 89

The population ratio or Ts is determined by the excitation of H I atoms by collisionswith particles, the 2.75 K cosmic background radiation and indirectly by Ly a pumping (seeField 1959 and references therein). If the cosmic background radiation is the only sourceof excitation, then, clearly, Ts = 2.75 K. The other two mechanisms drive Ts towards Tk.the kinetic temperature of the gas. Neglecting atomic collisions with electrons, which isjustified when Xe f:- 0.05, collisions dominate over radiative excitation if

nHI > ncrit = 4.7 X 10-3(Tk/1000K)oo77 cm-3. (1.1-2)

This density is sufficiently low that for all Galactic situations, one can safely assume Ts = Tk.

1.1.1 Radiative Transfer of the 21-cm line.

At cm wavelengths, the Rayleigh-Jeans approximation is valid and it is conventionalto use TB, the brightness temperature instead of I, the specific intensity. TB is linearlyproportional to I via the Rayleigh-Jeans relation: 1= 2v2 kTBC-2. It is convenient to useTB since TB, unlike I, is independent of v. The equation of radiative transfer in the H Iline can then be written as

dTB(v)/dr(v) = Ts -TB(V), (1.1-3)

where r is the optical depth along the line of sight. This equation shows yet anotheradvantage of using TB: in the limit of large r, TB = T s i.e. the brightness temperature ofan optically thick H I cloud is equal to its kinetic temperature, as it should.

Though equation (1.1-3) appears simple, the solution to the general case of both Tsand r varying from position to position is non-trivial. For the simple case of an isolated,single, isothermal H I cloud, equation (1.1-3) can be solved to yield

TB(V) = Thge-'T(tI) + Ts[l- e-'T(tI)], (1.1-4)

where Thg is the brightness temperature of the radiation incident on the far side of thecloud. In an actual measurement, vie measure the spectrum with respect to the continuumi.e. we measure ~TB(V) = TB(V) -Thg:

~.TB(V) = (Ts -Tbg)[l- e'T(tI)]. (1.1-5)

r( v) is related to N( v), the number of atoms with a velocity v in a cylinder of base 1 cm-2(i.e. the column density) by

N(v)r(v) = .(1.1-6)C X Ts

Here v is assumed to be in km S-l and the constant Cis 1.83 X lot8 cm-2 K-1 (km S-l )-1.In order to gain a physical feeling for these equations we consider two extreme limits

of the single cloud with no background radiation (i.e. Thg «Ts):Optically thin case (r(v) « 1). In this case, TB(V) = Tsr(v) = N(v)/C and themeasured brightness temperature is proportional to the column density per unit velocity.Physically what this means is that almost all the spontaneously emitted 21-cm photonsescape the c~oud without being absorbed. The emission rate is practically independent of, .

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90 S. R. KULKARNI AND C. HElLES

T. since E / kT. is exceedingly small for any re3$onable T.. Thus the number of photonsleaving the cloud is a direct me3$ure of the H I column density.Optically thick case (T(V) » 1). In this C3$e, TB = T., i.e. the brightness temperatureis simply the spin temperature. Since T » 1, the 21-cm photons emitted within theinterior of the cloud get absorbed within the cloud. Only the photons emitted by g3$ withinT ~ 1 of the front surface manage to escape from the cloud. Thus the observed brightnesstemperature is independent of the column density and depends only on the temperature ofthe cloud.

1.2 H I Emission and Absorption Observations.

From equation (1.1-6) it should be apparent that H I is seen in 'emission' or in'absorption' depending whether T. is bigger or smaller than Tbg, respectively. In the formerc3$e, the emission from the foreground H I more than makes up for the attenuation of thebackground radiation and vice-versa for the latter case. For the special situation of a cloudhaving T. = Tbg, the emission from the foreground is exactly compensated by the absorptionof the background and there is no net emission. For example, an H I cloud at the sametemperature 3$ the 2.75 K cosmic background radiation would not be detectable. Howeverthis is not a problem for Galactic H I since T. is never so small at any Galactic site.

Emission data have been extensively used to map the distribution of Galactic H I.Absorption data is needed to separate cold H I from warm H I and to me3$ure the tem-perature of H I. For this re3$on, absorption data are crucial to our understanding of thethermodynamics and the structure of diffuse clouds 3$ well 3$ the warm medium.

1.2.1 H I Absorption Techniques.

Both single dishes and interferometers can be employed to obtain H I absorptionspectra. We discuss both these techniques and show that in general, interferometric mea-surements are more reliable than single dish techniques.

Consider an isothermal H I cloud and a radio source of flux density S behind it. Inthe single dish technique, we assume that the beamwidth of the telescope is much smallerthan the angular size of the cloud but larger than that of the radio source. Then (equation1.1-5), the antenna temperature me3$ured by the telescope in the direction towards thesource (the 'on-source' spectrum) is

ATA.on(V) = (T. -T.rc)[1- e-'I"(1J)]. (1.2-1)

Here we implicitly 3$sume that the beam efficiency of the telescope is unity so that, for anextended source such 3$ the H I cloud, the antenna temperature is equal to the brightnesstemperature. T.rc is the antenna temperature of the compact source and is given by S X Gwhere G is the gain of the telescope. G for a single 25-m antenna such 3$ the ones used at theVery Large Array (VLA) is "" 0.12 K Jy-l and "" 8 K Jy-l for the giant Arecibo reflector.We have neglected background contributions such 3$ the 2.75 K cosmic background radiationsince such contributions are present in both the 'on' and the 'off' spectra and do not affectthe determination of the optical depth spectrum.

The single dish technique consists of obtaining an 'off-source' emission spectrum,ATA,oJ J( v), in a direction displaced by "" beamwidth so 3$ to obtain an independent me3$ure

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of Te. Thus, an important assumption is being made: the cloud has little structure overone beamwidth. Under this assumption, the absorption spectrum is readily obtained as:

e-T(tI) = 1 + ~TA.on(v) -~TA.Off(V). (1.2-2)Terc

In practice, the off-source spectrum is formed from several off-spectra taken aroundthe source in a pattern such as a cross or hexagon. Linear gmdients in TB( v) are can-celled out by the use of such symmetrical patterns. However, structure on the scale of thebeamwidth is not, which is a fundamental limitation on single-dish work. For example, evenfor the giant Arecibo dish, a fluctuation of 10% in adjacent emission spectra is typical forIbl > 10°. This variation increases at low latitudes, towards intermediate and high-velocityclouds, and of course is worse for larger beamwidths and weaker background sources. Forthis reason, essentially all single dish H I absorption data towards weak or low-latitudesources are unreliable.

Intereferometers surmount these difficulties because of their ability to act as spatialfilters. An interferometer with a baseline B responds only to structures in the sky withspatial frequency'" B f.x, where .x is the wavelength. Most of the H I features in the sky donot have much structure on angular scales below about an arc minute (see Crovisier, Dickeyand Kazes 1985), whereas most background radio sources are generally smaller than 1 arcminute. Hence an interferometer with a fringe spacing smaller than 1 arc minute but biggerthan the size of the background source resolves the foreground H I emission and respondsonly to the background sources. Thus interferometers measure e-T(tI) directly. The abilityof interferometers to form images (aperture synthesis) is not critical for H I absorptionmeasurements. Thus the sensitive Arecibo interferometer is as useful as the multi-elementVery Large Array (VLA). In fact, the most extensive low-latitude survey was done at theVLA in the so-called 'phased array' mode in which the VLA is essentially reduced to a3-element interferometer (Dickey et al. 1983).

Interferometric surveys are always more reliable than single dish surveys. However,the Arecibo single dish surveys provide reliable absorption data because of the very narrowArecibo beam (4 arc minute) and the associated high antenna gain. Kulkarni et al. (1985)compare Arecibo single dish data with Arecibo interferometric data and conclude, that, atleast for Ibl > 10° and strong sources (5 > 3 Jy), the Arecibo single dish data are reliable.

1.2.2 Limitations of the Derived Spin Temperatures.

Our discussion of the excitation temperature of the 21-cm line was pursued in thecontext of a single, isothermal H I cloud. In practice, this simple situation is rarely encoun-tered. When more than one H I cloud. is present along a given line of sight, the measuredspin temperatures is related in a complicated way to the true spin temperatures.

Consider an isothermal H I cloud with a spin temperature Te and a column densityN H( v). The velocity dependence of N H is a consequence of both thermal and macroscopicmotions. Observationally we can measure T(V) and TB(V). In th~ spirit of equation (1.1-5),and assuming Tbg « Te, we define Tn(v), the naively-derived spin tempemture, in terms ofmeasured quantities as:

Tn(v) = TB(V)f[l- e-T(tI)j. (1.2-3)

For the simple case of a single, isothermal cloud that we considered in §1.1.1, Tn(v) is indeedequal to T e;.

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Complications arise when there is more than one parcel of H I with overlapping ve-locities along the same line of sight. For example, consider the case of two isothermal cloudswith the same velocity distribution, but with two different spin temperatures. Applicationof equation (1.2-3) yields

Tn(v) = T.,l[1-e-1"1]+T.,2[1-e-T2]e-1"1( 1.2-4

).[1 -e-1"1-1"2]

where the subscript '1' refers to the cloud nearer to the observer (the 'foreground' cloud) andthe velocity dependence of T has been dropped out for clarity. The derived spin temperature,Tn is now a function of v and from T..l through T.,2. Thus, literally interpreting Tn(v) asspin temperature would lead us to the false conclusion that H I at temperatures intermediateto the T1 and T2 exists.

Consider three illustrative cases to appreciate the pitfalls and the limitations ofTn(v):(1) Optically thick foreground. Then Tn(v) ~ T.,l and we have no information about thebackground cloud.(2) Optically thick background and optically thin foreground. Then Tn( v) ~ T..l Tl (v) +T..2and the naively derived spin temperature of the background cloud is increased by the fore-ground contribution. This is the usual case with the WNM being the foreground and a coldbackground cloud.(3) Optically thin background and foreground. Then, the naively deduced temperature,Tn(v) = [N1(v) + N2(v)]/[(N1(v)/T.,l) + (N2(v)/T.,2)] i.e. Tn(v) is simply the columndensity weighted harmonic mean temperature. This leads to a severe bias when we try tomeasure the temperature of the WNM e.g. a tiny cold (T "" 80K) cloud with only one-tenth the column density of the WNM (T "" 8000K) will decrease the naively derived spintemperature of the WNM from 8000 K to "" 800 K!

1.3 Probes of the Ionized Gas.

Ionized gas in the ISM is revealed by a variety of observations, each sensitive to acombination of the electron density and temperature:(a) Dispersion of Pulsar SignaLs. The group velocity of an electromagnetic signal of fre-quency II travelling through ionized gas is Vg = c(l -1I~/1I2)1/2, where Vp = 8.97n~/2 kHzis the plasma frequency. Owing to this dispersion in the ISM, a pulse at higher radio fre-quency arrives earlier than that at lower frequencies. The observed rate of change of delaywith frequency provides a measurement of DM == J nedl, the integrated column density ofelectrons to the pulsar. The conventional units of DM are cm-3 pc.(b) Optical Recombination Emission. In the ISM, a free electron eventually recombineswith a proton and in the process emits a host of recombination line photons. Nearly alluseful work has employed a wide beam (5' to 50') Fabry-Perot spectrometer to observethe Ha >.6563 A recombination line (see Reynolds 1984). The observed Ha intensity, 1=0.36 J n;T4-0.9 dl Rayleighs (R), where T4 is the temperature in units of 10. K and dl is thepath length in parsecs. One Rayleigh is the isotropic intensity emitted by 106 recombinationsper second. In steady state, I is a measure of the ionization rate. For an assumed constanttemperature, the emission measure, EM == J n;dl, can be obtained. The usual units ofEM are cm-6 pc.


(C) Low-frequency Radio Absorption. In ionized gas, the encounters of free electrons withions leads to the emission of radiation by the 'Bremsstrahlung' mechanism. Free-free ab-sorption is the inverse of this process. The optical depth Tff OC 9fffn;T-3/2v-2dl. Atradio frequencies, after taking account of the variation of the Gaunt factor, 9f f, we findTff oc v-2.1T-l.35 X EM. In order to see significant absorption by the Galactic ISM onehas to go down to frequencies as low as a few MHz! Tasmania is one of the few locationswhere, occasionally, the ionosphere becomes transparent to let radio astronomers peek atthe low-frequency heavens. Almost all the ground-based observations in this field has beendone in Tasmania by the pioneering radio astronomer Grote Reber and his colleagues. Ob-servations below 2 MHz have been obtajned from spacecraft; the space observations havealmost no angular information due to the small size of the space antennas.

2. Distribution of the Atomic Component.

Locally, the H I layer and the diffuse H II layer can be approximately describedby a stratified horizontal atmosphere model, i.e. NH(b) = NH,.Ljsin(lbl), where NH(b) isthe integrated column density towards galactic latitude band NH..L = (NH(b)sin(lbl)} isthe mean column density projected to the pole; the total vertical column is thus 2NH..L.Hereafter, we will use the subscript" .L" to indicate any mean column density projected tothe Galactic pole.

The distribution of the neutral atomic component has been traced solely from H Iemission data. H I emission data are a measure of the column density and hence do notdiscriminate between the cold medium (CM) and the warm, neutral medium (WNM). Thusthe discussion of the large-scale (§2.1), the H I halo (§2.2) and intermediate-scale (§2.3) referto both the neutral phases. The warm, ionized medium (WIM) has been mostly studiedusing optical recombination and metastable lines. Obscuration in the galactic plane hasthus far hindered the determination of the large-scale distribution of this medium, althoughsome meager information has been obtajned using pulsars (§2.4).

Before we discuss the Galactic distribution of the various atomic phases, a briefsummary of the local gas (distance < 1 kpc) is in order. Extensive H I absorption surveyshave been used to derive the distribution of cold clouds. Detajled emission line studies of HIkinematics have been successfully used to infer the physical conditions and the distributionof the WNM. The CM appears to be organized into discrete structures in contrast to theWNM which appears to be widely distributed. The scale height of the CM layer is about 100pc, half that of the WNM; layer. The cloud-cloud velocity dispersion of cold clouds is about6.9 km s-1 and that of the WNM is slightly larger: 9 km s-l. Typical cloud temperaturesare around 80 K whereas various observations suggest a temperature of '"'" 8000 K for theWNM. A detajled discussion of the physical conditions and the distribution of the local H I(the CM and the WNM) can be found in §5 and §6, respectively.

From emission data, Heiles (1976) finds NH..L '"'" 3.7 X 1020 cm-2 which translates toa mass surface density of 8 Me pc-2 more or less evenly divided between the CM and theWNM. A small caveat: it appears to be the practice of H I observers not to include He whenestimating the mass surface density. In this article all the quoted mass surface densitieshave been increased by 1.36 to include contribution from He. From pulsar dispersion data,Ne..L '"'" 0.9 X 1020 cm-2 (see §2.4.1). The surface density of the WIM is a non-neglible 2Me pc-2. Thus, altogether, the surface mass density of the atomic phases is 10 Me pc-2,which may be, compared to the total disk surface-density of 75 Me pc-2 (Bahcall, Schmidt~ .

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and Soneira 1983) and the molecular mass surface density of '" 3 Me pC-2 (Scoville andSanders, this volume). [The molecular material is exponentia.lly decreasing around R = ~.Consequently it is difficult to accurately estimate the local mass surface density.] Note thatNH..L and Ne,.L, are based on local measurements and hence are unaffected by changes in~.2.1 Large-scale distribution of H [.

The large-scale distribution of H I is derived in a straightforward manner: one firstobtains H I emission spectra which are a measure of the column density of H I per unitvelocity interval; with the help of a rotation curve, the volume density as a function ofvelocity and therefore of distance can be derived. Various quantities of interest such as thesurface density and the scale height are easily derived from the volume density distribution.The derived distribution thus depends upon the choice of the rotation curve including thevalue of~, the distance to the Galactic center. For example, changing Ro by a factor achanges the derived densities by a-I, the scale height by a, the surface density remainsunchanged and the mass by a2 (the additional dependence on 1, the Galactic longitude, hasbeen ignored here).

Distribution of H I in the Inner Galaxy (R < Ro). In the inner Galaxy, the so-ca.lled'distance ambiguity' problem prevents a unique relation between velocity and distance.However, the radial (i.e. galactocentric) distribution .of H I can be determined since thedifferential velocity depends only on the Galactocentric radius, R. Studies of scale heightare best done at the 'tangential points'. These are locations for which the line of sightis tangential to a circle centered at the Galactic center. Tangential points are desirablein two ways: (a) these are easily identified as extreme velocity gas and (b) the distanceto tangential points is independent of the assumed rotation curve and determined strictlyfrom geometrYj it is simply Rocos(l) where 1 is the Galactic longitude.

The surface density of the H I layer in the inner Galaxy is constant for R > OA~and decreases rapidly for R < OARo. The precise value of the surface density dependsupon the optical depth correction. Assuming the 21-cm line is optically thin, the meansurface density of H I is about 4.5 X 1020 cm-2 (Lockman 1984). No rigorous estimate ofthe correction due to the optical depth in the 21-cm line has been made. Using Kulkarni's(1983) estimate of a correction factor of 1.25, the mass surface density is '" 6 Me pc-2 -apparently somewhat smaller than the local value. Given the uncertainties, this discrepancymay not be significant.

The vertical structure of the H I layer appears to be independent of R for OARo <R. < Ro (Lockman 1984). (z), the mean of the H I layer is close to 0 pc throughout theinner Galaxy. The FWHM (full width at half maximum) of the H I layer is about 365 pc,independent of Rj this value is comparable to the local value. The constancy of the shapeand the width of the H I layer is a great puzzle. Ignoring the effect of cosmic rays andmagnetic fields,

u z <X ul1P.(0)-1/2 (2-1)

(Spitzer 1978); here Uz is the rms scale height of the gas layer, Ul1 is the velocity dispersionof the gas and P.(O) is the stellar mass density at z = OJ this expression is valid whenu. » UZ where u. is the scale height of the stellar disk. From studies of H I in ourown Galaxy and in external galaxies, Ul1 appears to be independent of R. Assuming anexponential stellar disk with a scale length of 4.3 kpc (Bahca.ll, Schmidt and Soneira 1983),

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we find p.(4kpc)/p.(10kpc) = 4. So naively we expect the H I scale height at R = 4 kpcto be twice as small as locally. Perhaps, pressure from magnetic fields and cosmic rays ishigher in the inner Galaxy and conspire to keep the scale height constant. It remains tobe demonstrated whether the mild gradients in magnetic field (Heiles, this volume) and

Icosmic rays (Bloemen, this volume) can actually account for the constancy of the H I z- ,distribution. In contrast, the width of the molecular layer in the inner Galaxy is smaller ,than the local value (Scoville and Sanders, this volume).

/';Distribution of H I in the Outer Galaxy (R > ~). The structure of H 1 in the outer ,IGalaxy has been successfully determined since there is no distance ambiguity problem. ~Recently, there have been two comprehensive analyses: one assuming a flat rotation curve ~(Henderson, Jackson and Kerr 1982) and the other using a rising rotation curve (Kulkarni, ~Blitz and Heiles 1982). Ii

In both these analyses, three major coherent features are seen very clearly. These tfeatures spi:al out as one follows the~ along their length and for this reason we refer to I~them as SpIral arms. The three major arms have been named as Perseus, Cygnus and :Carina after the constellations where the major portion of their length lie. There is a hint "!of a fourth 'arm across the Galaxy towards I = O. Locally the sun appears to be situated Iin a minor feature popularly referred to as the Orion arm. These three (and perhaps four)

Imajor arms appear to have constant surface density and are also confined to the main H Idisk. High-density molecular clouds and their associated H II regions dot these features ;(Blitz et al. 1983), just like the spiral arms in external galaxies.

It is not clear whether the large-scale features we see in the H I map arise due todensity variations or systematic non-circular velocity fields or both. In fact, if spiral armsare due to non-axisymmetric potential perturbation then we expect systematic non-circularvelocity fields. Incidentally, the non-circular velocities distort the H I distribution whenviewed in the R,9 plane; this is one reason (and perhaps the only one!) why the H Istructure is best viewed in 'I-v' diagrams. Thus the derived parameters such as pitch angleand the arm-interarm contrast are uncertain. The other parameters such as the surfacedensity and the size of the H I disk depend upon the choice of the rotation curve.

For the sake of 'hard' numbers we briefly summarize the results of one analysis(Kulkarni et al. 1982). In this analysis, the spiral arms have a pitch angle of , 25°,a length of about 20 kpc, an arm-interarm surface-density contrast of about 4 and areconfined to R < 20 kpc. The H I surface density is roughly constant to R = 20 kpc andfalls exponentially beyond that.

Verschuur (1975) has suggested that there may be more distant and fainter arms.This intriguing suggestion has not been investigated further.

The H I disk in the outer Galaxy is warped. Large sections of the H I disk in thefirst two quadrants (0 < I < 180°) are systematically above and that in the other twoquadrants systematically below the plane defined by the the H I layer in the inner Galaxy.The large-scale warp is fairly small to a distance of , 18 kpc from the center and thenrises very sharply. At R = 20 kpc the peak-to-peak displacement of the warp is about 3kpc, which amounts to more than 6° in galactic latitude! The outer edge of the H I diskoscillates i.e. the outer edge is scalloped. The wave number of this oscillation, m , 10 andthe amplitude is about 1 kpc. If one views the large-scale warp as an m = 2 oscillation thenperhaps it is not unreasonable to expect additional oscillations with values of m between 2and 10. Such oscillations have not been carefully searched for.

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In the outer Galaxy, unlike in the inner Galaxy, the thickness of the H I layer increaseslinearly with R. In view of equation (2-1) this is understandable since the stellar densityis rapidly decreasing and u" is unchanged. Since an exponential disk is inconsistent witheither a flat or a rising rotation curve, a spherical massive 'dark halo' has been postulated.The postulated dark halo does not affect the vertical structure of the H I disk because ofits small z-gradient.

2.2 The H I Halo.

More than a decade ago, W. W. Shane pointed out the existence of H I at high-z.Recently, the distribution of this H I (hereafter referred to as the 'H I halo') has beenstudied by Lockman (1984) by observing the vertical structure of H I at the tangentialpoints. Lockman found (a) about 13% of H I exists at Izl > 500 kpc and (b) the high-z H Iappears to be corotating with the disk. The H I halo disappears for R ~ 0.3~.

Lockman found a three component fit with the following parametE:rs provides anadequate description of the H I layer (0.3~ < R ~ 0.9~): (1) a Gaussian with an rmsscale height, Uz '" 100 pc, (2) a Gaussian with Uz '" 250 pc, and (3) an exponential withan exponential scale height, Hz '" 500 pc; the total vertical column densities without anycorrections are 1.2 X 1020 cm-2, 1.7 X 1020 cm-2 and 1.6 X 1020 cm-2 respectively. Weidentify component 1 with the CM from the similarity. of the scale heights. Thus, it is notunreasonable to suggest that components 2 and 3 together make up the WNM (see also§4.1 for further justification). The existence of more than one component implies that H Ias a whole cannot be characterized by one single velocity dispersion (c/. equation 2-1).This brings us to an important point viz. are the three components merely mathematicaldecompositions or are they physical components?

There are two biases which affect the parameters for the components quoted above.One is that there was no attempt to apply optical depth correction. Component 1 is mostaffected by this bias. This bias results in an underestimation of the measured width andthe column density of component 1; crudely, we estimate that the measured column densityof component 1 has to be increased by a factor of 1.3. The second bias is that in orderto convert the observed brightness temperature to volume density of H I in the tangentialregion, an estimate of Llr, the 'length' of the tangential region must be made. Clearly Llrincreases with the assumed velocity dispersion of the gas. However, as discussed above thereis more than one velocity dispersion. In fact, it is the higher dispersion gas that reaches thehigher Izi. In short, Llr is itself a function of z. This bias results in an overestimation ofthe measured width and column density of component 3. Clearly, careful modelling of thetangential point data is needed to reduce these biases.

What is the evidence for a local H I halo? Lockman, Shull and Hobbs (1985) havecompared the amount of H I inferred from Lya absorption towards high-latitude OB starswith that inferred from 21-cm observations. Such comparisions, though meager when com-pared to the 21-cm data, confirm the exist~nce of a local halo similar to that in the innerGalaxy.

There is no study of the halo in the outer Galaxy. Studying the halo in the outerGalaxy will require careful modelling since there are no convenient regions like the tangentialpoints. H I from many regions will contribute to the brightness temperature at a givenvelocity. Additionally, the modelling is complicated by the existence of more than onevelocity dispersion and the flaring of the H I disk. Ignoring these complications, Kulkarni

~~~~


et al. (1982) found that the shape of H I layer in the outer Galaxy as characterized by theratio of the 50-percentile width to the rms width appears to be invariant with R. Clearlymore work is needed here.

2.3 Distribution of H I on Intermediate Scales.

In this section we discuss the distribution of H I on scales below about 1 kpc. Thisdiscussion is brief and cursory. The reader is referred to the review by Kulkarni and Heiles(1986) for a complete discussion. The structure of H I has been traced solely from H Iemission data. The emission data come mainly from the following surveys: in the north atthe Hat Creek Observatory (Heiles and Habing 1974) and Weaver and Williams (1973), inthe south at the Instituto Argentino de Radioastronomia, Argentina (Colomb, Poppel andHeiles 1980) and at Parkes, Australia (Cleary, Heiles and Haslam 1979).

In the solar neighbourhood, there are three major regions containing young stellarobjects (OB associations and H II regions): Ophiuchus, Perseus/Taurus and Orion. Everyone of these ~egions is enveloped by a large H I concentration ('H I complexes') with highcolumn density. Approximate typical physical properties of these complexes are: NH ""'1.0 X 1021 cm-2, (nH) ""' 2.5 cm-3 (volume average; actual value probably higher due toclumping), linear diameter""' 120 pc, and mass""' 1.0 X 105 M 0.

These parameters vary widely from one regions to another. Orion, for example, isenveloped by about 7 X 104 M 0 of H I with a linear diameter of about 125 pc (Gordon 1971),while the associations I Mon and II Mon are enveloped by H I masses of about 1.5 X 105 M 0and 2 X 104 M 0, respectively (Raimond 1966). At least in some cases, an equivalent massis contained in molecular clouds; in Orion, for example, Thaddeus (1982) reports a totalof about 3.3 X 105 M 0 in H2 (a somewhat uncertain figure, because it is derived from COobservations using a controversial CO/H2 conversion figure). In the early days of 21-cmastronomy, Deiter(1960), using a small telescope (and hence a large beam), reported thepresence of such complexes in 31 out of 40 OB associations investigated by her. We no\vknow that star forming regions are intimately associated with the giant, molecular cloudcomplexes (see Scoville and Sanders, this volume). The time is now ripe for a systematic,high angular-resolution study of H I around all molecular complexes.

Apart from these complexes, the H I sky abounds in filaments on all angular scales.The largest of these filaments appear to be parts of expanding shells. The nearest suchobjects include the North Polar Spur, the Eridanus Loop, a shell encircling the NorthCelestial Pole and another running close to Radio Loop II. Some of these shells such as theNorth Polar Spur and the Eridanus Loop appear to be filled with hot gas emitting X-raysand some are associated with intense radio continuum emission e.g. North Polar Spur andRadio Loop I. Less spectacular expanding shells have been inferred from high resolutionobservations of H I around H II regions. In the latter case, the H I is the H2-dissociationfront.

Heiles (1979 and 1984) searched for distant expanding shells such as the ones foundlocally and discovered quite a large number of them. Most suprising was the discoveryof a population of extremely large shells, 'supershells', ranging up to 2 kpc in diameter.The Sun may be located just inside the boundary of a supershell (Lindblad et al. 1973).These supershells are truly remarkable. If produced by the instantaneous release of explosiveenergy, the required energies range from 4 X 1053 erg -equivalent to 400 supernovae. Energyreleased from !i large number of stars in the form of stellar winds and supernovae explosions

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may, in fact, be the energy source for most or all supershells. However, the large energiesand the fact that usually only one hemisphere of an expanding supershell is visible mayimply another mechanism in some cases -specifically, the collision of infalling gas with thegas in the Galactic plane.

About half of the H I in the solar vicinity appears to be falling towards the Galacticplane (see Kulkarni and Fich 1985). Most of this H I lies at LSR velocities within ~40 kms-l. However, some of the gas lies at higher negative velocities: 'intermediate velocity gas(IVG)', extending to about -90 km s-1 and 'high velocity gas (HVG)', extending to moreextreme velocities (Giovanelli 1980). Both the IVG and the HVG were originally discoveredby Dutch radio astronomers. The Dutch work on IVG (Wesselius and Fejes 1973) remajnsthe best and the most comprehensive.

Space and other considerations restrict our discussion of the IVG and HVG. Theprincipal result of the Dutch study is that IVG is located primarily at high positive Galac-tic latitudes and coincides with a 'hole' in the low velocity gas (LVG) distribution. Theanticorrelation between LVG and the IVG is very striking. In some regions, the low velocitygas and the IVG fit together like pieces of a jigsaw puzzle. The IVG column density columndensity is close to what the LVG column density would be if the LVG 'hole' did not exist.Thus, Wesselius and Fejes argue that the LVG has been 'displaced' to form the 'IVG' -that some agent affected the LVG and changed its velocity without changing the total H Icolumn density.

The physical cause for the origin of the IVG and the HVG is not clearly established.Heiles (1984) has presented photographs of the IVG distribution which makes it look like aportion of an expanding circular shell. Cohen (1981) has shown that the HVG lies preciselyon top of the LVG in the direction of Radio Loop II. Nichols-Bohlin and Fesen (1986) suggestthat the star HD 50896 is in the middle of an expanding, intermediate velocity shell. Thecircumstantial evidence i.e. expanding shells and a cluster in the middle of the shell (as inthe shell surrounding HD 50896) suggest a supernova origin. Heiles (1984) showed that theIVG in the Galactic anticenter region is organized in the form of a large (30°) one-sidedshell. Kulkarni (1983) has investigated this structure and favors an expanding/moving ringmodel. Mirabel (1982) has argued that this structure has resulted from interaction of highvelocity clouds with the LVG. Kulkarni and Mathieu (1986) have placed a lower limit on thedistance (2.5 kpc) and favor Mirabel's suggestion. This structure deserves to be investigatedfurther since it is probably the best example of a structure formed by the interaction ofHVG with LVG.

2.4 Distribution of the Ionized Gas.

In §7 we show that most of the interstellar electrons come from the warm, ionizedmedium (WIM). Thus the distribution of the WIM is the same as the distribution of theelectrons. The various probes of interstellar electrons are reviewed in §1.3. Locally, theelectrons appear to be wide spread as evidenced by diffuse Ha emission along most linesof sight. The Ha line and optical metastable lines have provided us much insight aboutthe physical conditions of the interstellar electrons. Unfortunately, owing to obscuration,we cannot use these optical probes to study the distribution of the interstellar electrons onGalactic scales. At the present moment our knowledge of the electrons on non-local scalescomes from analysis of pulsar dispersion measure (DM) data.

,2.4.1 Local Distribution

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The vertical column density distribution of the electrons can be estimated from thepulsar DM data. H I absorption spectra towards 40 low-latitude pulsars have been usedto establish the distance-DM calibration for pulsars (Lyne, Manchester and Taylor 1985).(lzl) for these 40 pulsars is "" 200 pc and hence the exact form of the vertical profile of(ne) cannot be directly inferred from observations. Using a carefully selected sub-sample ofpulsars and superimposing the requirement that the spatial distribution of pulsars deducedfrom the dispersion measures should be consistent with cylindrical symmetry around thegalactic center, Vivekanand and Narayan (1982) tested a number of simple models for thedistribution of interstellar electrons. Their principal conclusions are: (i) (ne) = 0.037!g:gi~cm-3 and (ii) the exponential scale height of electrons is greater than 300 pc and probably 1kpc or larger. Electrons from the diffuse medium as well as discrete H II regions contributeto (ne). After allowing for the contribution from the H II regions we get (ne(O») "" 0.030cm-3. The lower limit on the scale height of electrons is comparable to the scale heightof pulsars (400 pc). Thus (DMsin(b») obtained from pulsar data is a lower limit to theelectron vertical column density. Manchester and Taylor (1977) have attempted to correctthis by obtaining the vertical distribution of pulsars and electrons self-consistently and findDM.L ~ 30'cm-3 pc. If the scale height of interstellar electrons is significantly larger thanthat of the pulsars then even this determination should be considered a lower limit. If weassume an exponential distribution: (ne(z») = (ne(O»)e-lzl/H. then we have He "" 1000 pcto reproduce the observed DM.L.

Almost all the observations in the Ha A6563A line have been done by wide-beam (5arcmin to 50 arcmin) Fabry-Perot spectrometers and mainly by Reynolds (a good referenceis Reynolds 1984). The Ha emission appears to be widely spread with a disk-like distribu-tion. The intensity projected to the pole i.e. I.L is found to be between 0.5 Rand 1.7 R.Adopting a mean value of 1 R, this translates to an EM.L "" 2.8 cm-6 pc for T4 = 1 (§1.3).From latitude scans in the Perseus arms, Reynolds (1986) finds that the Izl-distribution ofEM is well represented by an exponential with a scale height of 300 pc. If we assume thatthe local scale height of the local WIM is also 300 pc, then the observed EM.L requires(n~) ~ 0.00ge-lzl/300 cm-6.

The observational situation of the low-frequency absorption is unfortunately uncer-tain. From observations of extragalactic sources at 10 MHz, Bridle and Venugopal (1969)find T.L(10 MHz) = 0.1 :i: 0.02. From observations of the diffuse backgroud, Ellis (1982)derived a value five times smaller towards the south Galactic pole; this is approximatelyconsistent with space-based observations extending to frequencies below 1 MHz. We favorEllis's result because the absorption effects are more pronounced at the lower frequencies.The optical depth per kpc in the Galactic plane, "-0, can be estimated by measuring Ttowards well-known Galactic sources such as the Crab Nebula etc. Comparision of T.L with."-0 yields a scale height of about 1 kpc (Bridle and VenugopalI969).

Irregularities in the interstellar electrons give rise to the phenomenon of interstellarscintillation (ISS). ISS dilates images of compact extragalactic sources, broadens pulsarpulse profiles and cause pulsar signals to vary in time and radio frequency. Since ISSdepends upon irregularities (i.e. 8ne) and not on ne it is not clear whether the irregularitiesarise in the WIM or in the hot component. Either medium has more than enough electronsto account for the deduced 8ne. However, quite surprisingly, Readhead and Duffet-Smith(1975) obtain a scale height of about 1 kpc from measurements of ISS scatter-broadening ofcompact extragalactic sources. This scale height is similar to that inferred from pulsar DM

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100

data. Thus, there is some circumstantial evidence suggesting that the irregularities indeedarise in the WIM. We favor this interpretation but at the same time agree that more datais needed to firmly prove this point.

2.4.2 Distribution on Non-Local Scales.

Our knowledge of the distribution of the interstellar electrons on non-local scalesis meager and mainly comes from pulsar dispersion measure (DM) data. H I absorptiondata have been used to obtain kinematic distances for about 40 pulsars. These data, inconjunction with the DM data, have been used to establish the Galactic mean electrondensity distribution (Lyne, Manchester and Taylor 1985):

{ne} = [0.025 + 0.015e-IZI/70][ R2/ ]cm-3. (2.4-1)1+ Ro

The first term describes the extended component with a scale height of about 1000 pcand discussed above. The second term describes, in a statistical way, the contribution bydiscrete, bright H II regions and hence is of not great interest in this review. [In equation2.4-1, Ro has been assumed to be 10 kpc.]

Equation (2.4-1) is valid only within 2 kpc. of the sun since most of the pulsarsare nearby. Thus there appears to be an increase in the interstellar electron density inthe inner Galaxy. This is only to be expected since the density of young stars which areprobably responsible for most of the ionization in the ISM (§7.1) increases dramatically inthe inner Galaxy. Recently, Clifton and Lyne (1986) have discovered 40 new pulsars, allwithin 300 of the Galactic center. Determination of distances to these pulsars is crucial andwill enable us to quantitatively determine the increase of {ne} in the inner Galaxy. Indeed,if the WIM is powered by radiation from hot stars (§7.1), it is possible that the WIMcould become as important as H I in the inner Galaxy! It is crucial to test this reasonablesuggestion. This interplay b'etween star formation and interstellar medium should also bestudied by observations of diffuse Ha and H I in external galaxies. However, the detectionand measurement of the WIM in external galaxies are challenging observations.

3. Pressure of the ISM.

In steady state models of the ISM in which gas is heated by interstellar radiationfield or cosmic rays and cooled by collisional deexcitation, interstellar gas can exist either'as cold clouds or warm, neutral hydrogen. This is essentially the heart of all steady statemodels (Field, Goldsmith and Habing 1969). However, theoretically these two phases cancoexist over a narrow range of pressures. In the supernova-dominated model of McKeeand Ostriker (1977), interstellar pressure is predicted, given input parameters such as thesupernovae rate etc. Clearly, interstellar pressure is a critical parameter for any model ofthe ISM.

A note on nomenclature: pressure is usually defined as P = knT in standard physicstextbooks; here, it is more convenient to drop k, the Boltzman constant and redefine P =nT. Thus the units of pressure will be cm-3 K. Note that n is the total particle densityand thus for a completely ionized medium P = 2neT.

There are three methods by which pressure has been determined:a. Excitation of the C I fine-structure lines.


The ground electronic state of C I is split into three fine-structure levels; the transi-tions between these levels lie in the far-infrared. Only recently has one of these transitionsbeen detected (see Phillips, this volume). Jenkins, Jura and Lowenstein (1983) measuredthe fractional population of these levels by observing the ultraviolet transitions betweenthe ground and some excited electronic states. The level-population is determined by abalance between collisional excitation and radiative decayj the collisional rate depends bothon temperature and density. From the fractional level-populations, Jenkins et al. obtainedan estimate of interstellar pressure. Most of the gaseous carbon in the ISM is found as C IIsince the ionization potential of C I is less than 13.6 e V. Thus, this pressure 'probe' is biasedtowards high density regions where the fractional abundance of C I is higher than in lowerdensity regions. In short, the pressure determinations from the C I lines are insensitiveto the WNM and the WIM and apply mainly to the cold clumps in the CM (see §5.3 fordiscussion about the clumps). Jenkins et al. found that the pressure in most clumps isbetween 103 cm-3 K and 104 cm-3 K. About 6% of the gas is at pressures greater 104cm-3 K and one third is below 103 cm-3 K. We adopt a representative median pressure of4000 cm~3 K.

b. Excitation of the C II fine-structure line.The ground electronic state of the C II ion is split into two fine-structure levels with

a separation (in units of temperature) of 92 K or a wavelength of 157.7 p.m. Just like the C Ilines, this transition is a nice probe of pressure. Additionally C II, unlike C I, is a majorityspecies and hence is found in all the phases of the atomic ISM. Thus it is a better probe ofthe average pressure of the ISM.

The C II 157.7p.m line is the primary cooling line of diffuse clouds. Not surprisingly,the fraction of excited C II ions is usually stated in terms of the cooling rate per unithydrogen nucleus: lc = f21XcA21(hv) erg S-1 nucleon-l where 121 is the fraction of C IIin the excited state, Xc is the fractional abundance of gaseous carbon, A21 = 2.36 X 10-6s-1 is the Einstein A-coefficient for the transition and v is the frequency of the C II line.The cosmic abundance of carbon is thought to be 4 X 10-4 (Spitzer 1978). We will assumethat about half of this is depleted onto grains and thus Xc '" 2 X 10-4.

In the diffuse ISM, C II is excited by collisions with electrons and hydrogen atomsand cools by radiative decay. In the WIM, collisions with electrons dominate and f21 '"10-2(ne/0.33)(T/104)-1/2. For our nominal parameters of the WIM (ne '" 0.25 cm-3,T'" 8000 cm-3 Kj see §7), we find f21 '" 8 X 10-3 or lc '" 0.5 X 10-25 erg S-1 nucleon-I.In the CM, both electronic and atomic collisions are important. A proper calculation off21 for the C II in the CM has to take into account of the existence of CM material over asubstantial range of temperature (§5.3). Assuming a cosmic ray ionization rate, ( '" 10-16S-1 (§6.1) an estimate of lc as a function of temperature for various values of pressure isshown in Figure (3-1). This figure shows two important features of the C II line: (i) fora given cosmic ray ionization, (CR, lc cx: P, the gas pressure, and (ii) at any given P, lcvaries by less than a factor of 2 for 40 < T < 400 K, the temperature range spanned bythe CM. Thus the C II line is a robust estimator of pressure. Surprisingly, for our nominalmedian pressure (4000 cm-3 K), lc for the CM is nearly the same as that for the WIM.We do not expect much C II emission from the WNM since both nH and ne are small.

There are two measurements of 121: one by direct observations of the far-infraredline itself (Stacey et al. 1985) and the other from ultraviolet absorption observations of theexcited C!II ions (Pottasch, Wesselius and van Duinen 1979). Stacey et al. detected C II

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~~

"ie0...II

"iIn

to)"Qj

T0...~ P = 3xl01 CI -I~ / ~-=-:~: -=m K

~ P = lxl01 CI -0/ ~.:.~ -=m K

Figure 3-1. Plot of the C II cooling rate per nucleon in cold clouds as a function of temperaturefor four different values of ras pressure (P). Carbon is assumed to be depleted by 50% and a cosmic rayionization rate of 10-16 s- is assumed.

emission from the warm interfaces of giant molecular clouds and thus their results are notrelevant to our discussion here.

Pottasch et al. found the surprising result that lc "oJ 10-25 erg s-1 nucleon-1 alongeight different lines of sightj th,e variation between the lines of sight was less than a factor of2. Some of these lines of sight are nearby ("oJ 100 pc) and are dominated by the WIM whereasothers are quite distant with many cloud components. lc "oJ 10-25 erg s-1 nucleon-1corresponds to p"oJ 8000 cm-3 K for both the CM and WIM. This is in excess by a factorof 2 over the mean pressure of 4000 cm-3 K inferred from the C I measurements. Thediscrepancy gets worse when a correction is made for the low C II emission from the WNMjthe precise factor depends upon the z-distance of the star and is between 1.4 and 2.

The mean pressure deduced from the C II measurements appears to be larger thanthe median pressure obtained from C I data by a factor between 2 and 4. Agreement canbe obtained (barely) by assuming that Xc "oJ 4 X 10-4 which essentially means that thereis very little depletion of carbon in the diffuse ISM. The existence of grains can be usedto show that the CNO elements (in some combination) must be depleted by at least 25%(see Jenkins, this volume). Independent estimates of the depletion of carbon are needed toresolve this discrepancy. However, the C II UV lines are optically thick and there is somedoubt about the accuracy with which the column density of the excited C II ions can bemeasured (see pitfall # 1 in Jenkin's contribution; also Pottasch et al. do caution us aboutthis problem). All this discussion shows that the best way to probe the pressure in the ISMis by observations of the 157.7JLm line directly in the far infrared.c. Comparison of Emission Measure and Dispersion Measure.


We postpone the detailed discussion of this method to §8. The basic idea is that theemission measure per unit volume is equal to </>n~ whereas the dispersion measure per unitvolume is equal to </>ne; here </> is the filling factor of the WIM and ne is the true volumedensity of the WIM. From the observed emission and dispersion measures, one obtainsne "'" 0.25 cm-3. Using the nominal temperature of 8000 K for the WIM and assuming itis fully ionized (see §7) we derive a mean WIM pressure"", 4000 cm-3 K.

To summarize, the median pressure in the CM material appears to be about 4000cm-3 K -about the same as the mean pressure in the WIM. There is no direct measurementof the pressure in the WNM.

4. The Warm, Neutral Medium (WNM).

Broad (velocity dispersion, Uti "'" 9 km S-I) H I emission is present along all linesof sight at intermediate and high latitudes; in contrast, narrow emission (Uti ~ 5 km S-I)features are seen in only one third of the directions. Only these narrow features show de-tectable absorption. Since the optical depth is inversely related to temperature, the narrowcomponent can be identified with the cold medium (CM) and the broad component withthe warm, neutral medium (WNM). The existence of the WNM is beautifully illustratedin Figure (4-1) which shows a set of high latitude emission and absorption spectra. Thenarrow absorption features arise in cold H I 'clouds' i.e. the CM.

It ...

1

)_T~

Figure 4-1. A set of high latitude emission and absorption spectra from the Parkes survey (Rad-hakrishnan et al. 1972).

In Figure (4-1), pairs of vertical lines mark velocity ranges within which detectableabsorption (T ~ 0.1) is present. There is clearly considerable H I outside this velocityrange which does not show absorption. The broad dashed line is a fit to this optically thincomponent; this is the operational definition of the WNM. Measurement of the temperatureof this broad component is difficult because the opacity is so low. Signal-to-noise ratiosbetter th3.!1 loa are required to measure the optical depth of the WNM and has been

,

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achieved only towards a handful of the strongest background sourcesj these results arediscussed in §4.2.

Most of our information about the WNM have come from H I emission studies. TheWNM can be best probed by high-resolution UV absorption lines of warm, neutral speciessuch as H I, 0 I, N I, etc. and by far infrared cooling lines in emission (C II, Si II, Fe IIetc.) -probes which, unfortunately, are not accessible from ground-based facilities. Owingto the low density and low ionization in the WNM, even the far infrared cooling lines willbe faint and detectable only by extremely sensitive far infrared telescopes.

Locally, the WNM is widely distributed and constitute about 40% of interstellarhydrogen. The local WNM has been extensively studied from H I emission data, the detaileddistribution of which is discussed in §4.1. On Galactic scales, not much is known about thedistribution of the WNMj the meager data we have suggests that the WNM is a significantconstituent everywhere in the Galaxy outside", 8 kpc Galactocentric radius (Heiles 1980).Inside R = 8 kpc, there are no data.

4.1 Dispersion and Scale Height of the Local WNM.

Mebold (1972) decomposed about 1200 emission spectra at b '" 300 and 00 < 1 < 3600into narrow (o-v < 5 km s-l) and wide (5 < o-v < 17 km s-l) Gaussian componentsand found: (1) the wide-o-v H I has a scale height of about 200 pc and (2) an intrinsicvelocity dispersion of about 8.8 km s-l, corresponding to an upper limit of 9600 K for thetemperature of the WNM. These are the 'classical' parameters for the WNM. However,these values are at variance with those found by Lockman (1984) in the inner Galaxy(§2.2). To recapitulate, Lockman found a three component fit is necessary to adequately

.describe the vertical structure of the H I layer. Lockman's component 1 has the same scaleheight as the CM and component 2 corresponds to the 'classical' WNM layer. Analysesof the kinematics of the H I emission spectra, such as by Mebold (1972), are sensitive toperturbations of the local velocity field and hence the failure to detect the long exponentialtail (i.e. component 3) is, in hindsight, understandable. Since there is no evidence fora long exponential tail in the various H I absorption surveys we have to conclude thatthe temperature of Lockman's component 3 is significantly higher than that of the clouds.The few temperature measurements of the WNM that are available (see §4.2, below) donot distinguish between H I belonging to component 2 and component 3. Hence, it is notunreasonable to assume that the temperatures of the two components are the same i.e.T ~ 6000 K (§4.2).

It is not clear whether Lockman's three components are merely mathematical artifactsor physically meaningful components. Perhaps, the WNM is made of gas with a range ofvelocity dispersions and hence a range of scale heights. From an analysis of local H I inemission, Kulkarni and Fich (1985) found that NH(V) cx: V-2 where NH(V) is the columndensity of H I with velocity v (note that since they used H I data at high latitudes, thereis negligible contribution to v from Galactic rotation). Curiously, there appears to be anequipartition of kinetic energy per unit velocity interval, i.e. ~[dNH(v)/dv]v2 is a constant!

From equation (2-1) we expect that the higher dispersion gas should have the largerscale height. In Figure (4-2) we have plotted the rms angular scale height, (b2}1!2 of HIat a tangential point. Note the clear presence of some H I well in excess of the tangentialvelocity. The velocity in excess of the tangential velocity is simply a reflection of the

I


"

0'0 120

Figure 4-2. A plot of the angular scale height of H I, (b2)1/2 as a function of velocity along a line ofsight in the inner Galaxy. The H I emission profile in that direction is superimposed as an aid in locating

the tangential velocity.

velocity dispersion of the gas. Figure (4-2) clearly demonstrates that (b2)1/2 increases withu", implying a continuum of scale heights.

4.2 Temperature of the WNM.

At present the temperature of the WNM can be estimated using two different tech-niques: 21-cm emission and absorption studies and UV-line absorption studies of neutralspecies. So far, the 21-cm technique has contributed most of the information. The ba-sic method is to measure T(V), the optical depth of the WNM and use equation (1.2-3)to estimate the temperature. However, the measurement is difficult because T CX T-1, thetemperature. This difficulty is compounded by two biases: (i) the presence of a tiny amountof cold H I will bias the derived spin temperature, Tn to smaller values (§1.2.2) and (ii)stray radiation will increase Tn (Kalberla, Mebold and Reich 1980). The latter problemis only a technical hinderance and is surmountable. Since nothing can be done about theformer bias, the measured Tn should always be considered as a lower limit.

Many attempts have been made to measure the optical depth of the WNM. However,in most cases no absorption has been detected so that only upper limits to T, and thusmany lower limits to Tn are available. Many of these lower limits lie in the neighbourhoodof 3000 K or below (e.g. Mebold et al. 1982). The highest measured temperature of theWNM is "" 6000 K towards the strong, low-latitude source Cygnus A (Kalberla et al. 1980).The highest 3u lower limit is 104 K, towards 3C 123 (Kulkarni, Dickey and Heiles 1985);however this needs to be confirmed by independent observations before it can be accepted.

From a statistical analysis of the sensitive Arecibo survey Payne, Salpeter and Terzian(1983) find Tn "" 5000 K for gas not explicitly associated with absorbing gas. Even theArecibo survey did not have enough sensitivity to directly detect the WNM in absorptionalong anyone single line of sight. So Payne et al. averaged all the velocity channels whichcould not be explicitly associated with an absorption feature. Payne et al. found that this

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gas amounted to 50% of the local gas and had a mean spin temperature, Tn '" 5000 K. Wewould like to remind the reader, again, that Tn (§1.2.2) is the harmonic mean temperatureand hence is a lower limit to the true temperature. Thus the straightforward conclusionof the analysis of Payne et al. is that'" 50% of the local gas is the WNM and the other50% is the CM. The temperature of the WNM is probably not greater than 104 K for tworeasons: (a) a good fraction of the WNM has a dipersion of '" 9 km s-l, corresponding toa maximum temperature of ~ 104 K and (b) cooling by hydrogen increases dramaticallyfor T ~ 104 K (see discussion of the cooling curve in Black's contribution).

The study of UV absorption lines of various ions of the WNM can, in principle, yieldthe temperature. Several papers presented at the IAU Colloquium on the Local InterstellarMedium (e.g. York and Frisch 1984) find T ~ 6000 K for a handful of nearby (distance~ 200 pc) stars. This technique appears to be reliable and needs to be applied in a largenumber of directions.

To summarize, the WNM cannot be adequately described by a single component(i.e. the 'classical' model). Instead, two components with differing scale heights and dis-persion appear to provide a reasonable fit. It is not clear whether these two components aremathematical fits or physically meaningful components. Indeed, one may argue that theWNM is best described by a continuum of scale heights and dispersion. Several observationssuggest that the temperature of the WNM is ~ 6000 K. Theoretically the temperaturecannot exceed 104 K. However, temperature data are sparse and we definitely need moremeasurements.

.5. The Cold Medium (CM).

The cold medium (CM) is easily seen in absorption spectra. The narrow-O'v featuresseen in absorption in Figure (4-1) have large optical depths and hence are cold. Eachabsorption dip is assumed t~ arise in an 'H I cloud' or a 'diffuse cloud'. The nomencla-ture probably arose from the similarity of the H I absorption spectrum with the opticalabsorption spectrum such as that of N a I and Ca II.

The nomenclature is justified since the emission features corresponding to the narrowabsorption features appear to be confined to discrete structures in the sky; this shouldbe contrasted with the nearly featureless distribution of the broad-O'v, low-optical depthfeatures which constitute the warm, neutral medium. In one of the earliest models of theISM (Field, Goldsmith and Habing 1969), the cold medium was distributed into clouds'with the WNM filling the space between the clouds. For this reason the WNM is sometimesreferred to as the 'intercloud medium'.

At high latitudes, where the emission spectra are simple, absorption features canbe typically associated with the corresponding emission features. The absorption featuresare found to be well represented by Gaussian functions in T( v), the optical depth. Thedecomposition of emission features is somewhat subjective; some observers fit Gaussianprofiles whereas others find Gaussian fits inadequate. The distribution function of O'v( T),the velocity dispersion of the absorption features, peaks at about 0.75 km S-l; (O'v(T)} '" 1.7km S-l and is significantly larger than the peak value because the distribution function hasa long tail. In contrast, the emission features are twice as broad as the correspondingabsorption features; they peak at '" 2.2 km S-l (Crovisier 1981). No cloud has beenobserved with O'v( T) < 0.4 km s-l and very few with O'v( T) > 4 km s-l. The lower limit isprobably unaffected by observational biases.

~,;!'iIt" ,." _.-~ --


For every 'emission-absorption' pair we can obtain the naively-derived spin tempera-ture as a function of velocity (see equation 1.2-3). In almost all cases, Tn(v) is not a constantas would be expected for an isothermal cloud; instead, Tn changes increases on either sideof the velocity at which maximum absorption occurs. This is a simple consequence of thediffering line widths of the emission and the absorption features. The simplest and in factthe correct interpretation of the velocity dependence of Tn (v) is that the H I clouds are notisothermal blobs. The implications of this result are discussed in §5.3. Traditionally, thesmallest Tn(v), i.e. Tn at the velocity of maximum absorption, is called the spin tempera-ture of the cloud. We will denote this temperature by Tn min and the peak absorption by

, ,, Tmax.

The radial velocity of diffuse clouds can be measured to great precision because ofthe narrowness of the absorption features. The observed radial velocity of clouds is dueto Galactic differential motion and cloud-cloud velocity dispersion. The differential motiondepends upon the distance to the cloud as well as the direction. With sufficient data, themean height, (lzl) of the clouds can be extracted from the absorption data. Belfort andCrovisier (1984) find (lzl) '" 100 pc. The scale height appears to be a function of the opticaldepth: (lzl) '" 881: 13 pc for Tmax > 0.1, while that for Tmax < 0.1, (lzl) '" 2291: 48 pc.The T -T relation enables us to conclude that the clouds get warmer at higher Izi (see §5.3for an alternative explanation).5.1 The T -T ' Relation.

Lazareff (1975) first noted that Tn,min'S are inversely related to Tmax. This thefamous 'T -T' relation. A fit to the sensitive Arecibo data yields:

Tn.min = To(l -e-Tmos )-a (5-1)

with To = 551:7 K and a = 0.341:0.05 (Payne, Salpeter and Terzian 1983). Usually Tn,minand Tmax are plotted on a log-log plot and one does not realize that the scatter in thisrelation is quite large. Values of Tn.min range from 20 K to about 250 K and those of Tmaxfrom 0.01 to '" 2. The lowest observed Tmax is limited by sensitivity. There is probably areal cutoff since there is a dearth of emission features with widths between the WNM andthe cold clouds.

5.2 Statistics of Diffuse Clouds.

The procedure to derive the statistics of absorption features has been nicely presentedby Crovisier (1981) and applied to Arecibo data by Payne, Salpeter and Terzian (1983).The probability of finding an absorption peak with Tmax > T along a line of sight, reducedto Ibl = 90°, is well represented by P(T > Tmax) = 0.3T';;~:. At b = 0°, this translates to

N(T > Tmax) = 3.0T';;~:, (5-2)

where N(T > Tmax) is the number of absorption features (i.e. clouds) per kpc with themaximum optical greater than Tmax. In order to convert the optical depths to temperatureand column density, two assumptions are made viz. (i) the T -T relation is used to yieldthe temperature and (ii) the emission line width is assumed to be 1.3 times the thermalwidth of the absorption feature. With these assumptions, the corresponding probabilitiesfor Tn.min '3:fid NH are: P(T < Tn,min) cx Tn-:~~n and P(N > Ncloud) CX Ncl~...~. Both

.iI

I108

of these derived probabilites are consistent with the ones derived directly by Payne et al.(1983). At b = 0°, the latter relation translates to the number of clouds per kpc withNH > Ncloud, P(N > Ncloud) = 4.7Ncl2~~. Here Ncloud is the column density expressedin units of 1020 cm-2. These probability relations are valid over the range Tma., ~ 0.02 to1.0; this corresponds to ranges of temperature and column density of 210 K to 84 K and0.32 X 1020 cm-2 to 2.2 X 1020 cm-2, respectively. The median cloud has Tma., = 0.11,Tn,min = 115 K and NH = 0.8 X 1020 cm-2.

The distribution of column densities of diffuse clouds has been obtajned using twoother methods: interstellar reddening and optical absorption lines. These methods are ill-suited for cloud statistics. The observed reddening is simply the integrated column densityof dust. For this reason, the analysis is necessarily crude: the spectrum of diffuse cloudsis approximated by 'standard' and 'large' clouds (Spitzer 1978). However, in view of thelarge range of N H I that is actually observed, we seriously question the value of this crudeapproximation. For example, the median H I cloud has N H I = 0.8 X 1020 cm -2, nearly fourtimes smaller than a 'standard' cloud. Thus the 'standard' cloud is not even representativeof the cloud population! The reddening of a median H I cloud is so small (EB- v = 0.017mag.) that reddening data cannot be effectively used to study diffuse clouds.

Hobbs (1974) has attempted to use K I absorption lines to measure the columndensity distribution function. There are two problems in using K I lines for this purpose:(1) K I lines are weak and hence are tracers of large column density clouds; with typicaldetectors, a median cloud is not detectable in the K I line. (2) K I absorption lines do notmeasure N H I directly. Hobbs uses an empirical quadratic relation between N H I (obtajnedfrom Ly a measurements) and N K I to determine N HI. We question this scheme sincethe Ly a-derived NH I refers to the total column density of H I, which may include morethan one diffuse cloud and definitely includes a substantial contribution from the WNM.In order to obtajn NH I from K I data properly, a knowledge of depletion of metals andelectron density in clouds is J:leeded.

It has been popular among theorists and observers alike to derive the cloud 'size-spectrum' by using the observed column density distribution function together with theassumptions of (i) constant volume density and (ii) spherical clouds. Assuming pressureequilibrium, the observed spectrum of temperatures translates to a range of densities -making assumption (i) highly questionable. In §5.4 we show that assumption (ii) cannot bejustified observationally. Thus one should reject results based on these false assumptions.

.5.3 Internal Structure of Clouds.

In §5.1 we noted that emission widths are typically twice as large as the correspond-ing absorption widths. A strajghtforward consequence of this is that the najvely-derivedtemperature, Tn (v), will vary across the absorption feature. Two explanations are possible:(i) different portions of the cloud are being viewed in emission compared to absorption or(ii) H I clouds are not isothermal and that the variation in Tn( v) is real. The first hypothesishas been tested quite thoroughly and rejected (see Liszt 1983 for a summary). So we have toconclude that H I clouds are indeed not isothermal structures (Payne, Salpeter and Terzian1983; Liszt 1983). This is very puzzling since theoretically we expect H I to be either warm(T '" 8000 K) or cool (T '" 40 K; see Drajne 1978 or Field, Goldsmith and Habing 1969).The sound-crossing time across 1 pc of a diffuse cloud is only 106 y. Unless diffuse cloudsare disturbed more frequently than every million years, we expect pressure equilibrium indiffuse clouds. Hence, temperature variations should give rise to density variations.

I


A number of observers have used double-lobed radio sources as background sourcesto study the internal structure in diffuse clouds and find very little H I structure on scalelengths below 1.5 pc (see Crovisier, Dickey and Kazes 1985). Internal structure can alsobe studied by high-resolution mapping of H I in emission. This is best accomplished byaperture synthesis techniques, although the process is laborious and difficult. Kalberla,Schwarz and Goss (1985) have mapped a field containing the bright radio source 3C 147.Even at the small angular scales of their map, extended filaments and/or sheets dominatethe distribution of H I. Occasionally, there are sharp edges, probably indicative of shocksin the gas. Also seen are small clumps within the filaments and/or sheets.

The presence of 3C 147 allowed Kalberla et al. to derive temperatures independentlyfor the various H I components. They conclude:(1) The clumps, which are imbedded in the filaments and/or sheets, are primarily respon-sible for H I absorption. They have densities Tcl "" 20 to 50 cm-3 and spin temperaturesbetween 30 K and 80 K, averaging 40 K.(2) The filaments and/or sheets are associated with lukewarm (say, T ~ 500 K) H I en-velopes; the envelopes account for 80% of the H I emission.The 3C 147 field is non-ideal in some respects. It is at low latitude (b "" 10°) and in theanti-center direction. Thus the distance estimate is highly uncertain. We urgently needmore high resolution studies of H I in emission to understand the topology of the internalstructure of diffuse clouds.

The non-isothermal and inhomogeneous nature of diffuse clouds has profound effecton the interpretation of not only 21-cm data but other data as well. Consider two examples:(1) The distribution of majority species (such as C II) and minority species (such as C I)will be different. In particular, C I will be biased towards the colder regions. Thus theratio, nc l/nc 11 will vary throughout the cloud, complicating the interpretation of theobserved ratio. Optical and UV absorption line studies should pay particular attention tothis problem.(2) The good agreement between the rotation temperature ofH2, derived from UV studies (Spitzer 1978), with the apparent spin temperatures, Tn,min should be considered fortuitoussince we now know that the temperature of the cold clumps, Tcl "" 40 K. This implies thatsubstantial amounts of H2 exists outside the clumps.

Given the fact that clouds have temperature structure, how should the T -T relation beinterpreted? Liszt (1983) has simulated various cloud models and concludes that basicallythe T -T relation does not .constrain any specific cloud model. It is fair to say that the T-Trelation is basically satisfied by any class of reasonable, inhomogeneous cloud models. Onesuch model, possibly the simplest, is the 'clump-envelope' model advocated by Payne et al.(1983). In this model, the clumps are assumed to be at a fixed spin temperature (T a,cl)and variable optical depth (Tmax) and surrounded by a lukewarm envelope of constantcolumn density, contribution a fixed brightness temperature (TB,enll). This model fits theobserved data as good as the T -T relation does! The best fit parameters are Ta,CI = 55 Kand TB,enll = 4 K. Incidentally, in this simple model, the parameter TB,enll determines theexponent a of the T-T relation (equation 5-1).

In §5 we learned that the mean height of the cloud layer, (lzl) is a function of Tmox. Inview of the 'clump-envelope' model, two equally plausible interpretations are (i) the columndensity of the envelope increases with Izl or (ii) the temperature of the clumps increaseswith Izi (or Roth!).

,

...

1I


All this discussion emphasizes that we need an unambiguous verification of the'clump- envelope' model. One way is through high-resolution H I maps. The best way,which must await proper spaceborne instrumentation, would employ mapping of the far-IRcooling lines.

If we believe in thermal equilibrium models (Field, Goldsmith and Habing 1969) thenthe existence of lukewarm H I is a mystery. Thus either we have to give up steady statemodels altogether or most of H I is in a transient state. In our opinion, this challenging setof problems needs immediate theoretical and observational attention.

5.4 Shapes of H I clouds.

A survey of the literature shows that most workers assume that clouds are ~pherical.However, even a cursory look at the 21-cm sky will convince the reader that this is patentlyuntrue. Filamentary or sheet-like features abound in the 21-cm sky. Since the evaporationrate depends upon the topology and the surface area of the clouds, the topology of clouds isan important parameter in global models of the ISM. The standard assumption of sphericalclouds results in artificially decreasing the evaporation rate since the evaporation rate ofspheres is smaller than that of flat objects such as sheets.

Given the complexity of the 21-cm sky, a quantitative study of H I cloud shapes isnot easy. There are only two quantitative studies -one by Heiles (1967) and the otherby Schwarz and van Woerden (1974). Heiles stuided a 400 X 40 region of sky. The gas inthis region could be resolved into three physical components: a diffuse, smooth backgroundand two large sheets with a rift in between. The velocity of each sheet is highly orderedover many tens of degrees on the sky. Small concentrations of gas exist within these sheets;however, the density within these concentration is only twice as much as in the sheets.

In another study, Schwarz and van Woerden (1974) mapped a 200 X 100 region of skywith a beam of 36 arcmin a~d one degree spacing. The emission profiles were decomposedinto Gaussian components. A cloud was then defined to be a concentration in (I, b, v, 0')space; here V and 0' are the mean velocity and the dispersion of the Gaussian components.The above procedure furnished a total of 88 clouds, of which 10 appeared to belong tolarge-scale features including the intercloud medium. The most striking characteristics ofthe clouds are their filamentary shapes (see Figure 1 of their paper). The axis ratio variedfrom 2 to ~ 30; in fact, in a substantial number of cases, the minor axes of the filamentswere unresolved.

Clearly, both these studies are subjective to some extent: the first one to visualimpressions and the second one to the dangers associated with Gaussian analysis. Whilewe do not offer any specific panacea, we do wish to emphasize that the detailed studiesof specific regions of the sky have not yielded the postulated spherical clouds. Diffuseclouds are not self-gravitating objects. Thus the observed filamentary shapes suggest thatmagnetic fields (see Heiles, this volume) and/or evaporation play an active role in shapingthe cloud.

6. Ionization in the Atomic ISM.

The presence of free electrons in the ISM is clearly revealed by the dispersion of pulsarsignals and by wide spread diffuse Ha emission. The ionizing power needed to balance theobserved recombination rate is comparable to the total Galactic supernovae power! Clearly,the ionizing agent has a great impact on the ISM. In this section, we discuss the ionizationin the cold medium (CM) and the warm, neutral medium (WNM). We come to the very

I

THE ATOMIC COMPONENT III

important conclusion that most of the electrons reside not in the CM or WNM but in thewarm, ionized medium (WIM). The energetics and distribution of WIM are discussed inthe next section (§7).

If interstellar electrons predominantly arise in the WIM then why should we studythe ionization in the CM and the WNM? One good reason is that the cooling rate dependssensitively on xe, the fraction of electrons: in the CM, cooling by electronic excitationsdominate over atomic collisions for Xe ~ 10-3. Also quantitative measures of xe are neededto interpret the abundances of ions. In the past, it was thought that the atomic ISMwas heated and ionized by a single agent (e.g. Field, Goldsmith and Habing 1969) andhence the origin of interstellar electrons was a very important issue. Increasingly, there ismuch evidence, albeit indirect, that the neutral phases are probably heated by dust via thephotoelectric process.

First we discuss the two traditional ionizing agents: energetic particles (i.e. cosmicrays; §6.1) and energetic radiation below 912 A, the Lyman nmit (i.e. soft X-rays; §6.2).The intensities of soft X-rays and the low energy cosmic rays are either difficult or evenimpossible to measure directly. Consequently, it is important to clearly understand thebasis on which estimates of the ionization rates are made; these issues are discussed in §6.3(the CM) and §6.4 (the WNM). A brief note on notation: traditionally, the symbols (CRand (x R are used to denote the ionization rate per H I atom. The number of ionizationsper unit time per unit volume is simply (nH I.

6.1 Cosmic Rays.

The cross-section for ionization of hydrogen by cosmic rays varies inversely as thesquare of the velocity of the ionizing nucleon. Hence the ionization rate is dominated bynon-relativistic protons and heavy ions (E « 1 GeVfnucleon). Unfortunately the solarwind, with its trapped magnetic field, tends to sweep the lower-energy particles back outinto interstellar space, necessitating a large upward extrapolation to the observed intensityof cosmic rays incident at earth. Spitzer and Tomasko (1968) made the ad hoc assumptionthat the interstellar cosmic ray intensity is twice that measured at Mars (the solar windis weaker at Mars as compared to Earth and hence a more desirable location for suchobservations). Additionally, the interstellar cosmic ray spectrum was assumed to decreasesharply for energies smaller than 40 MeV fnucleon -the lower energy limit of the detectoron the space probe at Mars. With these assumptions, they derived a "minimum" ionizationrate'"" 10-17 s-1.

Historically, (CR '"" 10-15 S-1 was invoked to explain the heating of diffuse clouds(see Spitzer and Tomasko 1968). This rate can be obtained by extending the observed slope(i.e. the relativistic portion) down to a lower limit of 40 MeV fnucleon. The cosmic raypressure would then be 2x 107 cm-3 K, nearly 104 the mean ISM pressure! In order to avoidthis embarassing situation, very low energy (2 MeV fnucleon) cosmic rays were invoked (seeSpitzer and Tomasko 1968) which would contribute negligible pressure but provide sufficientionization. The very low energy cosmic rays were supposed to originate in the expandingshells of fast moving Type I supernova shells. However, there is no firm observational basisfor expecting such a high (CR and the justification for a high (CR was weak, even whenviewed in the historical context. Additionally, the short lifetime of such low energy particleslimits thei,r range (Spitzer and Jenkins 1975) -making low energy cosmic rays useless forionizing large regions of the ISM. From energetic arguments alone, (CR cannot be muchhigher than 10-15 S-1 if cosmic rays are produced by supernovae related processes.

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.I


Cosmic ray ionization is the chief source of protons in the interiors of diffuse clouds.Hence by measuring (CR the abundance of certain molecules whose formation criticallydepends upon the proton concentration can be estimated. Values of (CR ranging from 10-17S-l to 10-15 S-l, implying a typical value of 10-16 S-l are inferred from the UV -derivedabundance ratio [HD/H2] in diffuse clouds (Barsuhn and Walmsley 1977). An upper limitof 3.5 X 10-16 s-l has been derived from the radio-derived abundance ratio [DCO+ /HCO+]in three molecular clouds (Guelin, Langer and Wilson 1982). In a comprehensive modellingof four diffuse clouds, van Dishoeck and Black (1986) had to invoke (CR ,..., 7 X 10-17 S-l toexplain the observed OH abundance. From the absence of a low-frequency recombinationline of hydrogen towards 3C 123, Payne, Salpeter and Terzian (1984) obtain an upperlimit of 4 X 10-17 S-l for gas associated with the Taurus dust cloud and an upper limit of1.6 X 10-16 s-l for the diffuse clouds in that direction.

These astrophysical estimates apply to either cores of diffuse clouds or molecularclouds. Many authors have questioned whether low energy cosmic rays can penetrate clouds.Cesarsky and Yolk (1978) have shown that cosmic rays with energies as low as tens ofMeV /nucleon are not hindered by instabilites in penetrating dense molecular clouds. Indeed,cosmic ray ionization forms the cornerstone of gas-ion chemistry in dense molecular cloudsand a value approximately equal to the 'minimum' i°t:lization rate is invoked to explain theobserved abundances of some molecules (see Watson 1978).

To summarize, the minimum value of (CR based on a conservative (and somewhatad hoc) correction of the observed cosmic ray spectrum and an arbitrary low energy cutoffis ,..., 10-17 S-l. Indirect evidence suggests that an ionization rate close to this minimumis present even in the cores of dense molecular clouds. (CR in diffuse clouds appears to behigher than this minimum by about an order of magnitude. There is no firm estimate of(CR in the WNM; however (CR cannot be larger than 10-15 S-l if cosmic rays are producedby supernova related processes.

6.2 Soft X-rays.

The cross-section for photoionization of H I atoms is 6 X 10-18('>'/'>'0)3 cm-2 where>.0 = 912 A, the Lyman edge; thus low energy X-rays ('soft' X-rays) dominate ionization. Acolumn density of 1.6 X 1017 cm-2 will attenuate photons at the Lyman edge significantly.Given the mean density of ISM of,..., 1 cm-3, this column density translates to a distanceof only 0.056 pc! No wonder our knowledge of the soft X-ray radiation field is small.

In the three-phase model (McKee and Ostriker 1977), most of the interstellar volumeis occupied by the hot, ionized medium. This medium, by virtue of its high temperature(T,..., loG K), radiates X-rays (see Savage, this volume). Draine (1978) estimates, at z = 0pc, a soft X-ray ionization rate of 5 X 10-16 s-l at z = 0 pc with a mean photon energyof,..., 74 eV. This rate has been calculated assuming a filling factor of nearly unity for thehot component. Hence, it is a maximum value since in §8 we show that the filling factor ofthe atomic phases (i.e. the WNM, the WIM, and the CM) is ,...,50%. Also note that softX-ray flux will decrease with increasing Izl; the scale height of (XR will depend upon thescale height of the hot component.

6.3. Ionization within Diffuse clouds.

There are no reliable measurements of :le, the fractional density of electrons in diffuseclouds. Traditionally, analyses of optical and UV absorption line data simply assume thatthe electrons come from starlight photoionization of elements with ionization potential less

I

~~


than 13.6 eV. The main contribution is from carbon and hence Xe "" 4 X 10-46c where 6cis the depletion of carbon. Some analyses ignore the depletion of carbon which is clearlyincorrect since we do observe grains in the ISM (see Jenkins, this volume).

Including cosmic ray ionization of hydrogen and helium, the density of electrons is

ne = ni + ~[(1 + ~~ )1/2 -1] (6.3-1)2 an.

I

where a is the recombination coeffecient for hydrogen and is 4 X 10-13 (T /6000 K)-1/2 cm3s-1 and ni is the number density of metal-ions (viz. gaseous carbon; ni/nH "" 2 X 10-4assuming 50% depletion of carbon). For diffuse cloud conditions: nH "" 40 cm-3, T"" 80K we estimate ne "" 0.015 cm-3 for ( "" 10-17 s-1 with about equal contribution fromphotoionization of carbon and cosmic ray ionization of hydrogen. If ( "" 10-16 s-1 assuggested by van Dishoeck and Black (1986) then cosmic ray ionization dominates andne "" 0.0.38 cm-3. Xe is independent of nH for photoionization of carbon whereas Xe ""

n1i1/2 for cosmic ray ionization of hydrogen. In §5.3 we learned that diffuse clouds areinhomogeneous with a considerable fraction of hydrogen at temperatures greater than 40K. Assuming pressure equality we see from equation (6.3-1) that cosmic ray ionization playsan increasingly dominant role in the warmer regions of the clouds. Thus, contrary to thestandard assumptions made in UV and optical absorption studies, cosmic rays probablydominate the ionization in clouds, even if (R is as low as the minimum ionization rate.

The ionization in clouds is so small that they hardly contribute to the observedemission measure, the dispersion measure or the free-free absorption opacity (§2.4) -thisis left as an exercise to the reader.

6.4. Ionization in the Warm, Neutral Medium.

In principal, the WNM can be ionized by soft X-rays and cosmic rays. Whethersoft X-rays can actually be effective will depend upon the topology of the WNM. The softX-ray flux with a mean energy of 74 eV (§6.2) cannot penetrate a column density more than3x 1019 cm-2 and will be effective only if the WNM is organized into blobs smaller than "" 20pc. The filling factor of the WNM is estimated to be between 30% and 60% and hence it isnot unreasonable to assume that the WNM is distributed more like a pervasive 'intercloud'medium. Theoretically.such a morphology also appears to be favored (see Cowie, thisvolume). To summarize, if the WNM is distributed like an intercloud medium, then theeffective ionization rate is equal to (CR and Xe "" 0.007 for typical WNM conditions. If, theWNM is distributed into small blobs then soft X-rays dominate and the effective ionizationrate is "" (XR and Xe "" 0.05.

Regardless of this uncertainty two points should be clear. First, as argued in §7.1,on general energetic grounds, any ionizing agent which is a byproduct of supernovae cannotbe responsible for the ionization of gas seen in Ha emission. Thus neither soft X-raysnor cosmic rays can account for the observed EM.L (§2.4). Second, the WNM cannot becontributing more than 10% of the interstellar electrons. By equating the recombinationrate to the ionization rate we find Xe = 0.05[(-15/nH IP/2 where (-15 is the ionization ratein units of 10-15 s-l. For an assumed pressure of 4000 cm-3 K, nH I "" 0.5 cm-3 andXe "" 5% for (-15 = 0.5, the maximum possible ionization rate. Integrating this over z, theexpected maximum Ne..L "" 1.5 X 1019; this estimate includes a factor of 2 to account fordecreased pressure of the WNM at high z. This value of Ne..L is only 15% of the observed

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114

vertical column density (§2.4). For these two reasons, a model in which the Ha emission isproduced by regions of high ne and the pulsar dispersion measure by electrons in the lowdensity WNM is not viable. In addition, the estimated EM.L and the free-free absorptionopacity are orders of magnitude below the observed values (§2.4); this is left as an exerciseto the the reader). Thus, we conclude that neither the CM nor the WNM can accountfor the observed DM, EM or free-free absorption. On energetic grounds, we argued thatneither the CM nor the WNM can account for the observed Ha emission. So the conclusionis that we need another component -a warm, ionized medium and this medium is discussedbelow.

7. The Warm, Ionized Medium (WIM).

Two lines of reasonings lead us to postulate yet another atomic phase -the warm,ionized medium (WIM). First, a pervasive, largely ionized medium is required to explainthe wide spread, diffuse Ha emission. Several optical metastable emission lines such asN II, S II etc are also observed from the same gas as the diffuse Ha emission. From theoptical line data, we infer that this medium is warm (T ~ 8000 K) and substantially ionized(see Reynolds 1985). Second, in §6.3 and §6.4 we concluded that neither the warm, neutralmedium (WNM) nor the cold medium (CM) can satisfactorily account for the variousobservations of interstellar electrons. Thus we are .forced to invoke yet another phase toaccount for the interstellar electrons.

The simplest model is one in which most of the interstellar electrons arise in a singlephase. This phase must warm since line widths of the metastable lines suggest temperatures'" 8000 K; this is consistent with the detection of N II, the excitation of which requires Te,the temperature of the electrons to be > 3000 K. The ratios [N I/N II] and [0 I/O II]are tied to the [H I/H II] ratio. The lack of detection of N I suggests that the medium issubstantitally ionized: Xe > 0.75 (Reynolds, Roesler and Scherb 1977). We now proceed toexamine if other observational measures of interstellar electrons viz. the dispersion measuredata [DM.L], the emission measure data [EM.L], and the low-frequency free-free opacity[r.L(10 MHz)], can be consistently explained in this model. The subscript'.l' refers to aprojection of an integrated quantity such as the electron column density (for DM.L) ontothe Galactic pole; the physics of these probes is discussed in §1.3 and the observational datamay be found in §2.4.

The Ha intensity and the free-free opacity are proportional to the emission measurebut have different temperature dependences (§1.3). From the ratio EM.L/r.L(10 MHz) we

.derive T '" 4400 K -close enough, given the observational uncertainties, to the 8000 Kobtained from the optical observations of the metastable lines. If the hypothesized mediumis in pressure equilibrium with the rest of the atomic phases then P = 2neT '" 4000 cm-3 K(§3). We have assumed that the hypothesized medium is completely ionized to account forthe [N I]/[N II] ratio and hence the factor of 2 in the pressure equation. For T = 7500 K, weget ne '" 0.26 cm-3. The filling factor is simply (ne)/ne where (ne) is the volume-averagedmean density. From pulsar observations(§2.4.1), (ne) '" 0.03 cm-3 and thus the filling factoris about 11%. Thus, we find that a warm (T '" 8000 K), fully ionized medium occupying asubstantial volume of interstellar space can adequately account for all the present data ofinterstellar electrons. This important conclusion appears not to be well appreciated in thethe interstellar medium literature. Henceforth, we will refer to this new phase as the warm,ionized medium (WIM).

I


The solar system appears to be immersed in the WIM. Studies of scattering of SolarLy a (e.g. Bertaux et al. 1985) have established a temperature of 8000 f: 1000 K and apressure of about 3000 cm-3 K -similar to the parameters deduced for the WIM, above.

The local mass surface density of the WIM is estimate to be tV 2MG pC-2 (§2 or §2.4).Assuming an effective radius of 12 kpc and assuming no increase in the surface density ofthe WIM in the inner Galaxy, we estimate a total mass of 109 MG. Clearly, this is aconservative value. This mass estimate is comparable to the mass of the molecular materialobtained from observations of the diffuse ,-ray radiation (see Bloemen, this volume). Thiscomparision is meant to illustrate the importance of the WIM. [A small note of worry: the,-ray analyses do not include the contribution of the WIM. Thus, the ,-ray estimate of themolecular material is truly an upper limit and in fact including our estimate for the totalmass of the WIM may reduce the ,-ray estimate rather drastically. We suggest a careful

investigation of this point to resolve this important issue.]

7.1 Energy Source of the WIM.The power required to ionized the WIM is very high. Locally, the number of re-

combinations in a column perpendicular to the Galactic plane is tV 4 X 106 cm2 S-I. Tomaintain this over the whole Galaxy requires a minimum power of 1042 erg s-1 -compara-ble to or exceeding all the power injected by Galactic supernovae (Reynolds 1984)! Clearly,supernova-related processes are inadequate. There are two other possibilites: shocks andstellar photoionization. Shocks are unattractive because the various ratios of metastablelines are sensitive to shock speeds whereas, observationally, the ratios appear to be ratherconstant across the sky. Also, we have to address the issue of the agent which createdthe shocks in the first place. Stellar photoionization is an attractive possibility. Mathis(1986) has a model of the WIM with steady state equilibrium between recombination andphotoionization by a diffuse Lyman continuum. The model explains the optical data quite

well.However, the principal unanswered question is where does the diffuse Lyman con-

tinuum come from? The ionizing power of the 0 stars alone is more than five times thatrequired to explain the observed Ha emission (Reynolds 1984). However, 0 stars are rareand hence it is not clear how their Lyman continuum radiation can reach every nook andcorner of the Galaxy to produce the widespread WIM. Mathis (1986) estimates that about10% of the sky as seen from a typical 0 star must have a column density of H I < 1017cm-2 so that the ionizing radiation can escape into the general ISM. There are two problemswith this requirement: (1) there is not a single direction from the Sun where the columndensity of H I is so low and (2) a theoretical study by Elmegreen (1976) showed that 0stars tend to destroy clouds in their vicinity, the debris then forms an ionization-boundedH II region. Indeed there appears to be some observational data supporting this theoreticalpoint (Elmegreen 1975; McKee, van Buren and Lazareff 1984). Thus we conclude that 0

stars are ineffective within the cloud layer i.e. Izi < 100 pc.Outside the cloud layer, 0 stars can indeed ionize large regions of the interstellar

volume as evidenced by observations of diffuse H II regions around high-z 0 stars (Reynolds1982). A fraction of 0 stars, the so-called runaway 0 stars, are found above Izi > 100 pc.Accurate statistics of these stars are urgently needed to confirm whether there are sufficientnumber of runaway 0 stars to ionize the bulk of the WIM. In this scenario, the morenumerous B stars, which are confined to Izi < 100 pc, are the principle ionizing agents

within the cloud layer.I

."


Another possibility is that the WIM is ionized by a variety of hot stars such asplanetary nebula nuclei, hot white dwarfs, B stars (operative only in the cloud layer) andQSO light (operative at the edges of the WIM layer). The collective ionizing energy fromthese sources is just about enough to account for the observed recombination rate. In thisscenario, these stars simply ionize the H I nearest to them, which on the average is theWNM. Unlike the 0 stars, these stars are numerous and hence the WIM would appearwidespread and smooth -just as the observations indicate. High-sensitivity Ha maps ofregions around potential sources of ionizations are needed to confirm this scenario.

8. Filling factor of the Atomic Component.

The filling factors of the various phases of the ISM are very important numbers. Thecrucial question is whether most of the interstellar space is occupied by the hot componentor the atomic component viz. the warm, neutral medium (WNM) and the warm, ionizedmedium (WIM). If the atomic component dominates then the ISM is best described by the'two-phase' model (Field, Goldsmith and Habing 1969). If the hot component dominatesthen the ISM is best described by the 'three-phase' model (McKee and Ostriker 1977). [Thereader should be aware that there are really four phases in the McKee and Ostriker model -the CM, the WNM, the WIM and the hot component. The model of Field, Goldsmith andHabing in its purest form has only two components -" the CM and the WNMj including asufficiently intense, diffuse Lyman continuum will produce the WIM.]

First we derive the filling factor and its z-dependence c/J(z) of the WIM. The basicingredients are the z-distribution functions of (ne), from DM observations, and (n~), fromEM observations. The exponential scale heights are 1000 and 300 pc, respectively (§2.4).The true electron density, ne(z), is related to the volume-averaged electron density, (ne), by(ne(z)) = c/J(z)ne(Z)j similarly, (n~(z)) = c/J(z)ne(zf. A little algebra yields:

.ne(Z) '" 0.27e-lzl/428cm-3, (8-1)

andc/J(z) '" 0.11elzl/748. (8-2)

In §7 we concluded that the WIM is fully ionized and has a temperature of", 7500 K. Thusthe run of the pressure of the WIM with z is

P(z) '" 4000e-lzl/428cm-3K. (8-3)

It is important to note that the above equations are only statistical in nature and shouldnot be interpreted too literally. At any given Izl, the pressure in the ISM probably variesby at least a factor of 2.

Given our assumption of pressure balance among all gas components in the ISM,equation (8-2) should also describe the run of ambient pressure of the ISM. It is, then, a veryimportant equation and hence needs to be inspected critically. It is remarkably consistentwith our meagre, independent knowledge about the interstellar pressure. It predicts themean pressure within the cloud layer (izi ~ 100 pc) to be '" 3200 cm-3 K -in excellentagreement with the observational determination of the mean pressure in the CM from UVobservations of C I (§3). This agreement is remarkable because equation (8-2) was derivedwith very little or no input physics but only some basic algebra.

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We now use the assumption of gas pressure equality among the various phases, to-gether with equation (8-2), to derive the filling factors of the CM and WNM. The methodis straightforward: knowing the pressures and temperatures of these components, we derivethe true densities and then, from the observed volume-averaged densities, the filling factors.

For the CM, the determination of the filling factor is complicated by the inhomoge-neous structure of clouds (§5.3). For example, if we assume that the CM contains equalamounts of H I at 40 K and 400 K, then the filling factor of the former is a negligible'" 5 X 10-3 and of the latter a non-negligible 0.05. Since the mean ISM pressure hardlyvaries over the width of the cloud layer (izi < 100 pc), these filling factors are independentof z.

WIM + WNM~-~~ /':e: ...~~ ~-~~~ ~~~ /

~ :: /--~ ~-_::::::::::::o<::::::::= Figure 8-1 Filling factors of the WNM and the WIM as a function of Izi. The mean ISM pressure

is assumed to be 4000 cm -3 K at z = 0 pc. Pressure equality between the WIM and WNM is assumed.

For the WNM, instead of the three component fit discussed in §4.1, we use a simplifiedtwo component fit: a Gaussian of rms scale height 250 pc and an exponential of scale height480 pc. Under our assumption of pressure equality, the true density of the WNM is twicethat of the WIM since the temperatures are nearly equal. The sum of the filling factors ofthe WNM and the WIM at z = 0 pc is '" 0.4 and gradually rises to about 0.6 at z = 1 kpc

(Figure 8-1).Thus, under the assumption of gas pressure equality, we find", 50% of the interstellar

volume to be occupied by the WNM and the WIM. If these numbers were absolutely reliable,we would conclude that the hot component occupied the other 50% and thus the ISM wouldbe borderline between the Field, Goldsmith and Ostriker model and the McKee-Ostrikermodel. However, the gas pressure in the WIM could be systematically different comparedto the WNM. In §7, one of the scenarios discussed for the formation of the WIM is one ofionization, of the WNM by stray stellar Lyman continuua. If this scenario is correct then

! :1


WIM + WNM

~:==~"""'-.~ / ~

~

Figure 8-2 The mean ISM pressure is assumed to be 2000 cm-3 Kat z = 0 pc. The pressure in theWIM is a.ssumed to be twice the mean ISM pressure at a.ll z.

the WIM is a dynamically changing medium and could be overpressured by a factor of 4with respect to WNM (Elmegreen 1976). The C I measurements (§3) suggest that the meanISM pressure is unlikely to be as low as 1000 cm-3 K. However, a mean ISM pressure of2000 cm-3 K corresponding to an overpressure in the WIM by a factor of 2 is well withinthe uncertainties of observations. In this case, the filling factor of the WNM would doubleand the atomic component would occupy'" 80% of the interstellar volume (see Figure 8-2).

9. Summary and Future Work.

In this contribution, we have reviewed the phases which constitute the atomic com-ponent of the ISM: the cool diffuse clouds, the warm neutral medium and the warm ionizedmedium. Together, these three phases occupy a sizable fraction of the interstellar volumeand contain at least half the interstellar mass. The emphasis has been on synthesizingrelevant observations, from MHz radio measurements to i-ray measurements, in order toarrive at a comprehensive picture of the atomic component. Rather than simply summarizewhat we discussed in the previous sections, we will highlight some recent advances, reiteratemajor questions of current interest and discuss some future goals.

The warm, ionized medium (WIM) has been a neglected phase despite the fact that,locally, it has a filling factor'" 0.12 and a surface mass density'" 50% of the local surfacemass density of the molecular material. Conservatively we estimate a total mass of ~ 109M0. Surely with all these attributes this phase deserves more of our attention! The mostimportant recent result is the recognition that most of the interstellar electrons reside inthis medium. This result removes a great obstacle hampering all steady state models viz.the agent(s) which heat the atomic gas need not necessarily ionize the gas. The radial

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distribution of the WIM is unknown; this will be useful in constraining possibl~ ionizingagents and in obtaining a firmer total mass estimate. From pulsar D M data, we infera vertical scale height of '" 1 kpc. However, this is strictly a lower limit since the scaleheight of the pulsars is only 400 pc, significantly smaller than 1 kpc. Clearly, we needother probes of interstellar electrons. Observations of interstellar scattering of extragalacticsources appear promising; such observations may also establish in which medium do theirregularities reside: the WIM or the hot component? The power needed to sustain the WIMexceeds the power injected by supernovae. Thus the conclusion that the WIM is powered byradiation from hot stars seems unavoidable. How do the ionizing stellar photons leak intothe diffuse ISM, given the high opacity of the atomic component towards the ionizing stellarphotons? High resolution Ha observations are needed to determine which class of hot starsare the dominant sources of ionization. Finally, the existence of this widely distributedphase is indicative of the great influence of hot stars on the structure of the ISM. Thusglobal models of ISM should not only include the influence of supernovae on the ISM butalso that -of hot stars.

The warm, neutral medium (WNM) is an enigma. The presently available data,though sparse, suggests that the temperature of the WIM is indeed", 8000 K, in accordancewith our t~eoretical expectation. Space-based ultra violet and far infrared observations areurgently needed. Ground based observations of metastable lines of neutral species such asN I, 0 I etc. would be very valuable; unfortunately, the signals are expected to be low. Thefilling factor of this medium is estimated to be ~ 0.3. If the pressure in the WNM is 2000cm-3 K then the filling factor would rise to 0.6. Clearly, this phase must play an importantrole in the evolution of the ISM. Yet, we do not understand why about 50% of the local H Iis in the WNM phase. What determines this fraction? Does this fraction change with R?

The diffuse clouds which constitute the cold medium (CM) appear to be more compli-cated than predicted by theory. One might think that diffuse clouds, in some sense, are thesimplest structures in the ISM. We have good measurements of the radiation field incidenton these clouds and the atomic physics of the cooling collisional processes are well under-stood. So supposedly we know the inputs and the ouptuts. DESPITE this, observationallywe find that substantial portions of the clouds are at temperatures significantly larger thanthe theoretical expectation of 40 K. Further more there is a great deal of internal structureviz clouds are not isothermal. Do we understand diffuse clouds at all? Are clouds best de-scribed by an onion model with a cold core and successively warmer skins? What determinesthe internal structure? The inhomegeneity of diffuse clouds complicates interpretation ofall absorption line studies, including the 21-cm line and especially saturated optical andUV lines. Is the WNM wrapped around each cloud or does it exist independently of theclouds? Is the large scale filamentary structure of clouds due to cooling instabilites in thepresence of magnetic fields and/or evaporation?

The heating agents of the WNM and the CM have still not been identified. Thereis plenty of energy in the stellar radiation field to heat the diffuse medium. The principalobstacle has been identifying a mechanism to tap this vast reservoir. So far, the mostpromising mechanism appears to be photoelectric heating by dust grains. However, evenwith the most optimistic parameters, substantial amounts of lukewarm (T > 100 K) H Icannot be produced by this mechanism. At this meeting there were some suggestions thatPAHs co\lJd. be the principal heating agents via the photoelectric process. Clearly, it isworth pursuing this suggestion. From time to time, non-radiative heating methods such as

.!

'I...'"


mechanical heating, Alfven wave dissipation etc. have been suggested. Such suggestionsare difficult to model theoretically and equally difficult to verify observationally.

The pressure scale height is a crucial parameter for models of the ISM. If the ISMis dominated by the hot component, then the pressure scale height should be large ( ~ 2kpc). Our determination of the local pressure scale height of 500 pc is significantly belowthis prediction. In fact, it is comparable to the width of the WNM layer. What is thesignificance of this? Anyway, an independent determination of the pressure scale heightis urgently needed. What determines the scale height of the CM? Cloud-cloud velocitydispersion or some other parameter? Do magnetic fields keep clouds tied together and thussomehow restrict the scale height? Are there really three distinct components of H I assuggested by Lockman's work? Or should we consider H I to have a continuum of velocitydispersions with corresponding scale heights? In either case, can we theoretically explajnthe observed vertical structure? In the inner Galaxy, why is the shape and the thickness ofthe H I layer independent of R in contrast to the molecular and the stellar layer?

The topology of the various phases are not understood. In the supernova-dominatedMcKee-Ostriker model, the hot component pervades most of the interstellar space with theatomic component occupying the remajning volume. In the quiescent Field, Goldsmith andHabing model, the WNM occupies most of the interstellar space. Which picture describes

.the ISM? The applicable model almost certainly changes with R and it is not clear wherein the Galaxy the transition occurs. Clearly, the time is ripe for a systematic study of thevariation of the different phases of the ISM as a function of Rand z -including the studyof external galaxies.

Locally, both the WNM and the hot component appear to have similar filling factor.Is the hot component distributed in discrete structures? Is the WNM organized into smallblobs or more like an 'intercloud' medium? What is the relation between the WNM andthe WIM? Is the WIM simply the ionized edges of the WNM, or distributed completelyindependently of the WNM? Moving to smaller scales, we find that diffuse clouds are rarelyspherical. What determines their filamentary shapes? Magnetic fields or evaporation? In ahot component-dominated ISM, the topology of the CM governs the total evaporation rate.For this reason, it is important to determine whether diffuse clouds are filaments or sheetsand whether the WNM is distributed like an 'intercloud' medium.

This work was supported in part by an NSF grant to CH and a Millikan Fellowshipto SRK. We gratefully acknowledge help of S. Vogel (CIT) for constructive criticisms.

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