Sp.-V/AQuan/2000/02/14:16:18 Page 523 Chapter...

Chapter 21

Circumstellar and Interstellar Material

John S. Mathis

21.1 Overview of the Interstellar Medium . . . . . . . . . . 523

21.2 Galactic Interstellar Extinction . . . . . . . . . . . . . 527

21.3 Abundances in Interstellar Gas . . . . . . . . . . . . . 529

21.4 Line Emissions from the ISM . . . . . . . . . . . . . . 530

21.5 H2 and Molecular Clouds . . . . . . . . . . . . . . . . 532

21.6 Neutral Gas; Clouds; Depletions . . . . . . . . . . . . 534

21.7 H II Regions, Ionized Gas, and the Galactic Halo . . . 536

21.8 Planetary Nebulae (PNe) . . . . . . . . . . . . . . . . . 538

21.9 Supernova Remnants . . . . . . . . . . . . . . . . . . . 540

21.10 Cosmic Rays (Excluding Photons and Neutrinos) . . 541

21.1 OVERVIEW OF THE INTERSTELLAR MEDIUM

The bulk of the material presented here was submitted in October 1992, but some updating has beendone.

21.1.1 Characteristic Pressures and Energy Densities

P(gas) = nkT ; P(gas)/k ≈ 3 × 103 cm−3 K for phases in approximate mechanical equilibrium;P(magnetic)/k = B2/8πk = 7200(B/5µG)2 cm−3 K; B ≈ 5µG (diffuse ISM) [1, 2]; energydensities ≈ 0.5 eV cm−3 (starlight); 1.5 eV cm−3 (cosmic rays); 0.3 eV cm−3 (thermal pressureof gas in clouds); ≈ 1 eV cm−3 (gas kinetic motions); 0.6 eV cm−3 (magnetic field). All are highlyvariable spatially.

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21.1.2 Structures Within the Interstellar Medium (ISM) [3–6]

1. Molecular clouds: Defined by hydrogen being in H2, but observed by many molecules,especially CO; n(H2) > 100 cm−3; T ≈ 10–20 K. Geometrical structure is very complex and fractalin nature. Contains half of mass of ISM in the Galaxy (the rest in H I, diffuse and in clouds). Scaleheight in Galaxy: varies with DG = distance from center of Galaxy: ≈ 40 pc (DG ≈ 2 kpc) to 60 pc(DG ≥ DG� = 8.5 kpc).

2. Diffuse interstellar clouds: 10 cm−3 ≤ n(H) ≤ 1000 cm−3; T ≈ 80 K. Defined by H beingatomic, and observed by 21 cm line in emission and absorption. Also observed by narrow (≤ 10 kms−1) absorption lines of C II, Si II, Fe II, Mn II, Zn II, Mg II, Ni II, Na I, Ca II, etc., in the spectra of moredistant stars. Also seen as dark clouds against background stars because of the dust absorption, or asreflection nebulae when a nearby star illuminates the cloud, or from far-infrared emission (especiallyin 60 and 100 µm emission as seen by Infrared Astronomical Satellite, IRAS).

3. Warm intercloud medium: 0.1 cm−3 < n(H) < 10 cm−3; T ≈ 8000 K. Seen by 21 cmemission (but not absorption) and the narrow lines in stellar spectra (see above). Hydrogen is atomic.

4. Warm ionized medium: 0.3 cm−3 < n(H+) < 10 cm−3; T ≈ 8000 K. Seen by nebularemission lines, such as Hα, [N II], [S II], etc. The total power radiated is very large, ≈ 1041 erg s−1 ≈total power of all supernovae. The electrons from this medium are responsible for most of pulsardispersion measures. Extremely faint, covering the sky.

5. Hot ionized medium: n(H+) < 0.01 cm−3; T > 105.5 K. Seen by interstellar absorption linesof high stages of ionization (O VI, N V, perhaps Si IV) that could be formed at the interfaces betweenthis gas and the warm ISM. Almost surely seen directly by the soft X-ray background, produced in acavity (“the local bubble”) surrounding the Sun. This medium was thought to be pervasive in the diskof the Galaxy, but is probably confined to local disk regions that have been heated by hot-star stellarwinds and supernova shocks. This material has a large scale height (≈ 3 kpc) and is an important partof the galactic “halo.”

6. Supernova remnants: n(H+) > 1 cm−3; T ≈ 104–107 K. Seen by radio synchrotron emission,X-rays, optical spectrum from discrete filaments. Sometimes very peculiar composition, but older onesshow ISM swept up by the expanding remnant. Spectra are H, He recombination lines, [O II], [O III],[S II], etc., with relative intensities produced by shocks rather than photoionization. Some show strongfar-infrared radiation from dust. Widely varying conditions reflecting strong expansion and coolingthroughout evolution.

7. Large-scale nonthermal structures: Arcs seen in synchrotron radio emission and X-raysarising from large-scale ordered magnetic fields interacting with cosmic rays. An example is “Loop I,”the inner edge of which is the “North Polar Spur,” with an angular radius on the sky of about 30◦ anda physical radius of about 125 pc.

8. H II regions: 10 cm−3 < n(H+) < 104 cm−3; T ≈ 8000 K. Much denser than warm ionizedISM, so easily detectable. Strongest lines are [O II], [O III], H recombination lines; 52 and 88 µm[O III] fine-structure lines. Large ones powered by clusters of stars.

9. Ultracompact H II regions: 104 cm−3 < n(H+) < 3 × 105 cm−3; T ≈ 8000 K; excitationdominated by a single type O or early B star within the parent molecular cloud. Radio surfacebrightnesses much larger than classical H II regions mentioned above. Not seen optically because ofdust extinction.

10. Ring nebulae around Wolf–Rayet (WR) stars: Ejecta from the rapid mass loss from WRs,but containing some swept-up ISM as well. Sizes ≈ few pc; compositions He/H ≈ 0.1–0.2; N/O ≈ 1.

11. Planetary nebulae: 100 cm−3 < n(H+) < 104 cm−3; T ≈ 104 K. Excited by hotcentral star that has left the Asymptotic Giant Branch on its way to becoming a white dwarf, with3×104 ≤ T∗ ≤ 2×105 K. Spectrum varies with T∗, ranging from very strong [O II], no He II to strongHe II and Ne V. Nebula expands at 20–50 km s−1. M(neb) ≈ 0.2M�.

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21.1.3 Interstellar Gas and Radiation [7–11]

Table 21.1 gives details of the interstellar gas near the Sun.

Table 21.1. Interstellar gas densities in solar vicinity [1].

〈n(0)〉a Frac. T v (rms) h0Component (cm−3) of vol. (K) (km s−1) n(z)b (pc) σgas

c N (H)d

H2 0.6 0.001 ≈ 10 5 Gau 70 3.3 1.6Cold H I 0.3 0.02 80 6 Gau 135 3.2 1.5Warm H I

(clouds) 0.07 0.05 8000 9 Gau 135 0.7 0.3Warm H I(intercloud) 0.10 0.3:e 8000 ≥ 9 Expt 400 0.9 0.4Warm H+(diffuse) 0.025 0.2: 8000 ≥ 9 Expt 900 1.5 0.7Warm H II

regions f 0.015 0.04 8000 9 Expt 70 0.04 0.02Hot H+ 0.002 0.5: 106 Expt 3000 0.2 0.1Total 1.1 1 · · · · · · · · · · · · 10 5

Notesa〈n(0)〉 = midplane density, averaged over the area. Local density at z = 0 is 〈n(0)〉/(Frac. of vol.).bn(z) = type of distribution: Gau: n(z) = n0 exp(−z2/2h2

0); Expt: n(z) = n0 × exp(−z/h0).cσgas = surface density of gas, including a factor of 1.29 for He, on plane of Galaxy, in M� pc−2.d N (H) = Ave. column density of H (all forms), z = 0 to z = ∞, in 1020 cm−2.eColons in this and other tables denote uncertain quantities.eH+ detected by radio free–free emission, with density ≥ 1 cm−3.

Reference1. Boulares, A., & Cox, D.P. 1990, ApJ, 365, 544

Table 21.2 gives this hydrogen gas distribution in the Galaxy. All masses include He. z = heightabove the midplane of the Galaxy. Masses are in M�; surface densities in M� pc−2.

Table 21.2. Distribution of hydrogen in the Galaxy.

DGa (kpc) = < 0.4 0.4–3.5 3.5–7 7–8.5 8.5–14 Total

Mass of gas in H2 2 × 108 2.9 × 108 1.4 × 109 3.7 × 108 2 × 108 2.5 × 109

% of total H2 9% 11% 57% 15% 8% 100%FWHW in z of H2 (pc) 65 75 120 120 300 · · ·n(H2, z = 0) (cm−3) 100 2 1.6 0.8 0.03 0.5Surface dens. in H2

b 300 7 9 4 0.4 4Surface dens. in H Ib 7 1–5 6 6 6 5Mass of H I (M�) 5 × 106 5 × 107 6 × 108 4 × 108 2 × 109 3 × 109

Surface dens. in H+b ,radio H II regions 1.2 1.5 3.2 0.9 0.2 0.9

Surface dens., diffuse H+b · · · 2.7 2.7 1.7 · · · · · ·

Notesa DG = distance from center of Galaxy. DG (Sun) is assumed to be 8.5 kpc.bUnits of surface densities are M�/pc2.

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21.1.4 Diffuse Ionized Gas [12]

The average distribution of electrons with height in local Galaxy: ne = 0.025 cm−3 exp(−z/900 pc);N (H+) = column density of H+ or electrons from midplane to infinity = 1×1020 cm−2; filling factorof diffuse H+ ≈ 0.2.

The intensity of Hα in the galactic plane ≈ 10 Rayleighs, where 1 Rayleigh = 106/(4π)

photons cm−2 s−1 sr−1, or 2.4 × 10−7 erg cm−2 s−1 sr−1 at Hα. At the galactic poles, I(Hα) ≈1 Rayleigh = 2 × 10−7 erg cm−2 s−1 sr−1.

The total recombination rate producing this Hα ≈ (2–7) × 106 recombinations cm−2 s−1 (best fit,4 × 106), requiring a power of at least 5 × 10−5 erg cm−2 s−1 and probably twice that. Total powerproduced by supernovae is ≈ 10−4 erg cm−2 s−1, which is inadequate because of the large radiativelosses of SNRs. O, B stars produce about 3×107 ionizing photons cm−2 s−1 (Abbott, D.C. 1982, ApJ,263, 723; Mezger, P.G. 1978, A&A, 70, 565), which is adequate if ≈ 15% of the photons can reach thehigh-latitude gas.

Recombination time of H+ at z = 1 kpc ≈ 1.0 × 106 yr.Spectrum of galactic diffuse ionized gas: [N II]6583/Hα ≈ 0.3–0.4; [S II]6717/Hα ≈ 0.3–0.5;

[O III]5007/Hα ≤ 0.06; [O I]6300/Hα ≤ 0.02.The spectrum of the local photon radiation near the Sun is given in Table 21.3. The X-ray spectrum

is given in Table 21.4.

Table 21.3. Local interstellar radiation field [1–3].

λ (µm) 4π Jλa λ (µm) 4π Jλ

a λ (µm) 4π Jλa λ (µm) 4π Jλ

a

0.091 1.07(−2) 0.216 9.17(−3) 1.20 9.26(−3) 25 6.0(−5)0.10 1.47(−2) 0.23 8.25(−3) 1.8 4.06(−3) 60 4.6(−5)0.11 2.04(−2) 0.25 7.27(−3) 2.2 2.41(−3) 100 7.3(−5)0.13 2.05(−2) 0.346 1.30(−2) 2.4 1.89(−3) 200 2.6(−5)0.143 1.82(−2) 0.435 1.50(−2) 3.4 6.49(−4) 300 5.4(−6)0.18 1.24(−2) 0.55 1.57(−2) 4 3.79(−4) 400 1.72(−6)0.20 1.04(−2) 0.70 1.53(−2) 5 1.76(−4) 600 3.22(−6)0.21 9.61(−3) 0.90 1.32(−2) 12 1.7(−4) 1000 7.89(−6)

NoteaUnits of 4π Jλ: erg cm−2 s−1 µm−1.

References1. Mathis, J.S., Mezger, P.G., & Panagia, N. 1983, A&A, 128, 2122. Cox, P., & Mezger, P.G. 1989, A&AR, 1, 493. Wright, E.L. et al. 1991, ApJ, 381, 200

Table 21.4. Mean intensity of galactic X-rays [1].

Energy Name(keV) of band λ (µm) I (E)a 4π Jλ

a

0.100 Be 1.24(−2) 118 1.9(−5)0.163 B 7.61(−3) 73 3.1(−5)0.232 C 5.34(−3) 52.5 4.5(−5)0.625 M1 1.98(−3) 28.1 1.8(−4)0.799 M2 1.55(−3) 20.7 2.1(−4)1.04 I 1.19(−3) 14.1 2.5(−4)1.56 J 7.95(−4) 9.78 3.8(−4)3.3b · · · 3.76(−4) 6.82 1.2(−3)

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Notesa I (E) in keV cm−2 s−1 sr−1 keV−1; 4π Jλ in erg cm−2 s−1 µm−1.bFor E > 3.3 keV, I (E) = 11.0E (keV)−0.4 keV cm−2 s−1 sr−1 keV−1.

Reference1. McCammon, D., & Sanders, W.T. 1990, ARA&A, 28, 657

Sources of the observed gas and dust in the Galaxy are given in Table 21.5.

Table 21.5. Sources of gas and dust in the Galaxy [1, 2].

d M(gas)/dt d M(dust)/dtStellar type No. in Galaxy (M� yr−1) (M� yr−1)

M stars (Miras) 1.3 × 105 0.01–0.03 (1–3) × 10−4

OH/IR stars 104 0.1–0.5 (1–5) × 10−3

C stars (3–6) × 104 0.1–0.5 (1–5) × 10−3

Supernovae 0.02–0.03 yr−1 0.1–0.3 0.001–0.006M supergiants 5211 0.05–0.5 (2–50) × 10−4

Wolf–Rayets: WN, WC7 2744 0.05 0WC8, WC9 484 0.01 10−4

Planetary Nebulae 1.5 × 104 0.02–0.2 (0.7–7) × 10−5

Novae 30–50 yr−1 (0.5–1) × 10−4 10−5–10−4

RV Tauri stars 600–1200 0.006–0.01 (3–5) × 10−6

O, B stars (2.5–5) × 104 0.03–0.3 0Total in Galaxy · · · 0.3–1.5 0.003–0.015Star formation rate · · · −(3–10) −(0.03–0.1)

References1. Gehrz, R.D. 1989, in Interstellar Dust, edited by L.J. Allamandola and A.G.G.M.

Tielens (Kluwer Academic, Dordrecht), IAU Symp. 135, p. 4452. Jura, M., & Kleinmann, S.G. 1992, ApJS, 79, 123

21.2 GALACTIC INTERSTELLAR EXTINCTION

21.2.1 Extinction

If E(B − V ) = A(B) − A(V ), then N (H)/E(B − V ) = 5.8 × 1021 atoms cm−2 mag−1 [13]. HereA(λ) is the extinction, in magnitudes, or 1.086τ(λ), where τ is the optical depth in dust. The meanextinction law for interstellar dust can be described [14] as depending upon the optical parameterRV = A(V )/[A(B) − A(V )]. The diffuse ISM has a typical value RV = 3.1; in dense clouds, atypically RV = 4–5. Table 21.6 gives mean values for 3.1 and 5. There is considerable uncertainty inthe infrared extinction for (λ > 5 µm), perhaps a factor of 2 or more for λ ≥ 20 µm.

A(V )/N (H) ≈ 5.3×10−22 cm2 mag, where N (H) = column density of (H+H+ +2H2). Johnsonfilters are indicated in parentheses in Table 21.6.

Table 21.6. A(λ)/A(V ) at various wavelengths for RV = 3.1 and 5.a

λ (µm) RV = 3.1 5 λ (µm) RV = 3.1 5 λ (µm) RV = 3.1 5

250b 4.2(−4) 4.9(−4) 5 0.027 0.031 0.24 2.54 1.68100 1.2(−3) 1.3(−3) 3.4 (L) 0.051 0.059 0.218 3.18 1.9760 2.0(−3) 2.3(−3) 2.2 (K ) 0.108 0.125 0.20 2.84 1.74

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Table 21.6. (Continued.)

λ (µm) RV = 3.1 5 λ (µm) RV = 3.1 5 λ (µm) RV = 3.1 5

35 3.7(−3) 4.2(−3) 1.65 (H) 0.176 0.204 0.18 2.52 1.5225 0.014 0.016 1.25 (J ) 0.282 0.327 0.15 2.66 1.4920 0.021 0.025 0.9 (I ) 0.479 0.556 0.13 3.12 1.6018 0.023 0.027 0.7 (R) 0.749 0.794 0.12 3.58 1.7415 0.015 0.017 0.55 (V ) 1.00 1.00 0.091c 4.85 · · ·12 0.028 0.032 0.44 (B) 1.31 1.20 0.073 5.38 · · ·10 0.054 0.063 0.365 (U ) 1.56 1.33 0.041 2.58 · · ·9.7 0.059 0.068 0.33 1.65 1.35 0.023 2.06 · · ·9.0 0.042 0.051 0.28 1.94 1.42 0.004 0.96 · · ·7.0 0.020 0.023 0.26 2.15 1.50 0.002 0.38 · · ·

NotesaExcept as noted below, entries are from [1]. Values of A(λ)/A(V ) for other values of RV can be

determined from that paper.bFor λ > 250 µm, multiply entry for 250 µm by (250 µm/λ)2.cFor λ < 0.1 µm, entries are from [2], increased by 1.15 for continuity at 0.12 µm.

References1. Cardelli, J.A., Clayton, G.C., & Mathis, J.S. 1989, ApJ, 345, 2452. Martin, P.G., & Rouleau, F. 1990, in Extreme Ultraviolet Astronomy, edited by R.F. Malina and S. Bowyer

(Pergamon, Oxford), p. 341

21.2.2 Opacity in X-Ray Region of Spectrum [15]

The cross section per H atom is given by σ = (c1 + c2 E + c3 E2)E−3 × 10−24 cm2 (H atom)−1. Thereare breaks in σ at various energies as detailed in Table 21.7.

Table 21.7. X-ray opacity of interstellar gas and dust.

E range (keV) Edgea c1 c2 c3

0.030–0.100 · · · 17.3 608.1 −21500.100–0.284 C 34.6 267.9 −476.10.284–0.400 N 78.1 18.8 4.30.400–0.532 O 71.4 66.8 −51.40.532–0.707 Fe–L 95.5 145.8 −61.10.707–0.867 Ne 308.9 −380.6 294.00.867–1.303 Mg 120.6 169.3 −47.71.303–1.840 Si 141.3 146.8 −31.51.840–2.471 S 202.7 104.7 −17.02.471–3.210 Ar 342.7 18.7 0.03.210–4.038 Ca 352.2 18.7 0.04.038–7.111 Fe 433.9 2.4 0.757.111–8.331 Ni 629.0 30.9 0.08.331–10.00 · · · 701.2 25.2 0.0

NoteaThe element whose absorption produces a discontinuity at

the upper energy of the range.

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21.2.3 Diffuse Interstellar Bands (DIBs)

Over 100 interstellar features [16, 17], some broad, are still unidentified. Some profiles are asymmetricto the red; some are symmetric. Observed properties of strongest diffuse interstellar bands. InTable 21.8, W is the equivalent width in nm; � = central depth; the FWHM is in nm.

Table 21.8. Selected diffuse interstellar bands.

λ (nm) W a �a FWHM λ (nm) W a �a FWHM

443.0 0.34 0.16 2.0 617.7 0.18 0.07 3.0488.2 0.13 0.06 1.7 628.3 0.20 0.38 0.38544.9 0.56 0.045 1.4 661.38 0.036 0.24 0.013577.83 0.095 0.06 1.7 722.40 0.037 0.21 0.016578.0 0.088 0.37 0.26 792.7 0.036 0.026 1.4579.7 0.039 0.22 0.13 862.08 0.042 0.084 0.43

Notea W , depth are for HD 183143, a well-observed star [1].

Reference1. Herbig, G.H., & Leka, K.D. 1991, ApJ, 382, 193

21.3 ABUNDANCES IN INTERSTELLAR GAS

Table 21.9 gives the elemental abundances for the gas in many objects where observations can be made.

Table 21.9. Chemical composition of interstellar gas.

106 100 100 100 1000Object He/H O/H C/O N/O Ne/O S/O Fe/O Reference

Orion Neb a

(t2 = 0) 0.096 384 0.73 13 17 3.0 10b [1–4](t2 = 0.035) 0.100 550 0.62 13 16 2.8 7.8b [5]

M 17 (t2 = 0) 0.105 324 1.69 16 19 3.0 1.3b [5]30 Doradus (LMC) 0.89 200 0.15 2.8 22 2.3 · · · [6–8]Other LMC H IIs 0.091 230 0.35 5.0 20 3.2 · · · [6, 7, 9]SMC H IIs 0.083 135 0.14 2.6 36 5.0 · · · [6, 7, 9]Cas A SNR, FMKc c c < .003 < .007 < 2 6.5 < 1 [10]Crab Neb 0.7 2000 0.75 7 32 2 · · · [11]LMC SNRs · · · 177 0.26 16 18 2.3 93 [9]SMC SNRs · · · 85 · · · 9 9 3 110 [9]Local B stars 0.100 390 0.42 16 19 3.4 110 [12, 13]Sun 0.098 850 0.43 12 14 1.91 55 [14]

Notesat2 = temperature fluctuation parameter = [〈T 2〉/〈T 〉2 − 1] [15].bGas-phase Fe only. Most of the Fe is in solid grains.cFMK = Fast-moving knots. H, He not detected; H/O ≤ 0.3, He/O ≤ 1.7.

References1. Baldwin, J.A. et al. 1991, ApJ, 374, 5802. Rubin, R.H., Simpson, J.P., Haas, M.R., & Erickson, E.F. 1991, ApJ, 374, 5643. Osterbrock, D.E., Tran, H.D., & Veilleux, S. 1992, ApJ, 389, 3054. Peimbert, M. 1987, in Star Forming Regions, edited by M. Peimbert and J. Jugaku (Reidel,

Dordrecht), p. 111

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5. Peimbert, M., Torres-Peimbert, S., & Ruiz, M.T. 1992, Rev. Mex. Astron. Astrof., 24, 1556. Dufour, R.J., Shields, G.A., & Talbot, R.J., Jr. 1982, ApJ, 252, 4617. Dufour, R.J. 1984, in Structure and Evolution of the Magellanic Clouds, edited by S. van den Bergh

and K.S. de Boer (Kluwer Academic, Dordrecht), p. 3538. Mathis, J.S., Chu, Y.-H., & Peterson, D. 1985, ApJ, 292, 1559. Russell, S.C., & Dopita, M.A. 1990, ApJS, 74, 93

10. Chevalier, R.A., & Kirshner, R.P. 1979, ApJ, 233, 15411. Pequinot, D., & Dennefeld, M. 1983, A&A, 120, 24912. Gies, D.R., & Lambert, D.L. 1992, ApJ, 387, 67313. Kilian, J. 1992, A&A, 262, 17114. Anders, E., & Grevesse, N. 1989, Geochim. Cosmochim. Acta, 53, 19715. Peimbert, M. 1967, ApJ, 150, 825

21.3.1 Isotopic Abundances [18]

The hydrogen, carbon, nitrogen, and oxygen isotopic abundances are given for some places in theGalaxy in Table 21.10.

Table 21.10. Isotopic abundances in the Galaxy.

Isotope Galactic Carbon star: 5 Dusty Solar Local Gradienta

ratio center IRC10+216 carbon stars system ISM (dex kpc−1)

105 2H/1Hb · · · · · · · · · (1–3) 1.65 · · ·12C/13C 20 40 ± 8 ≥ 30 89 70 6.2 ± 1.614N/15N 900 > 515 · · · 270 380 28 ± 1416O/18O 250 > 2700 320–1300 490 490 66 ± 2418O/17O 3.2 < 1 0.6–0.9 5.5 3.2 −0.03 ± 0.05

Notesad(log10 X)/d DG in range 4–14 kpc.b2H/1H from Encrenaz, T., & Combes, M. 1983, Icarus, 52, 54; for local ISM from Linsky, J.L.

et al. 1992, ApJ, 402, 694.

21.4 LINE EMISSIONS FROM THE ISM

Relative line intensities, I (λ)/I (Hβ), for various objects corrected for reddening are given inTable 21.11. These objects are: Orion Nebula = NGC 1976, position near θ1Ori C; 30 Doradus =NGC 2070 = N157, giant H II region in LMC; Crab Nebula = NGC 1952, remnant of SN1054, aHe-poor and He-rich filament listed; Cygnus Loop = NGC 6960, NGC 6979, NGC 6992/5, a medium-aged SNR, with high-ionization and low-ionization filaments listed; NGC 7662 = high-excitation PN.

Table 21.11. Line emission fluxes relative to Hβ.

λ (nm) f (λ)a

OrionNeb[1–4]

30Dor[1, 5]

Crab Nebula

He-poor He-rich[6, 7] [7, 8]

Cygnus Loop

High ion Low ion[9, 10] [9]

NGC7662[11, 12]

155.5 C IV −1.17 · · · 0.012: 6.00 · · · 5.1 0.76 13.9164.0 He II −1.11 · · · · · · 4.80 · · · 3.1 0.50 6.28166.3 [O III] 1.10 · · · 0.15 · · · · · · 4.6 1.5 0.53

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λ (nm) f (λ)a

OrionNeb[1–4]

30Dor[1, 5]

Crab Nebula

He-poor He-rich[6, 7] [7, 8]

Cygnus Loop

High ion Low ion[9, 10] [9]

NGC7662[11, 12]

175.9 [N III] −1.09 · · · · · · < 1.7 · · · 1.4 0.8 0.33190.8 [C III] −1.20 0.19 0.15 5.6 · · · 13.4 6.5 4.23232.6 [C II] −1.32 0.09 0.07: < 6.5 · · · 2.91 3.86 · · ·247.0 [O II] −0.99 0.056 · · · · · · · · · 0.69 0.68 · · ·279.8 Mg II −0.64 < 0.012 < 0.02 < 1.10 · · · · · · · · · · · ·372.7 [O II] −0.31 0.94 1.38 10.0 12.0 16 12.8 0.10386.9 [Ne III] −0.28 0.20 0.44 1.72 0.187 2.0 0.72 0.75434.0 Hγ −0.15 0.49 0.46 0.48 0.48 0.48 0.46 0.48436.3 [O III] −0.14 0.015 0.041 0.15: 0.26 0.92 0.22 0.161447.1 He I −0.11 0.046 0.044 0.23 0.24 · · · · · · 0.028468.6 He II −0.04 · · · · · · 0.68 0.88 · · · · · · 0.79486.1 Hβ 0.00 1.00 1.00 1.00 1.00 1.00 1.00 1.00495.9 [O III] 0.025 1.70 3.50 4.63 4.21 1.25 3.66500.7 [O III] 0.037 3.43 4.87 11.0 15.1 13.9 3.38 10.9587.6 He I 0.20 0.137 0.124 0.74 0.77 · · · 0.07 0.076630.0 [O I] 0.26 0.007 0.007 0.98 0.45 0.35 0.31 · · ·656.3 Hα 0.29 2.71 2.92 2.98 3.22 3.00 3.00 2.87658.4 [N II] 0.30 0.44 0.058 3.0 3.79 3.11 2.98 0.051667.8 He I 0.31 0.036 0.035 0.19 0.12 · · · · · · 0.022671.7 [S II] 0.31 0.019 0.050 1.61 1.42 1.26 1.15 0.0039673.1 [S II] 0.31 0.034 0.042 1.90 1.72 0.85 0.78 0.0060713.6 [Ar III] 0.37 0.154 0.12 0.26 0.44 0.44 0.17 0.0945732.5 [O II] 0.39 0.119 0.023 0.77 0.22 0.62 0.44 0.014906.9 [S III] 0.48 · · · · · · 1.09 0.94 · · · · · · · · ·953.2 [S III] 0.61 1.45 · · · 1.90 1.73 · · · · · · · · ·C(Hβ)a · · · 0.84 0.65 0.74 0.7 0.12 0.12 0.26

NoteaObserved flux ratio F(λ)/F(Hβ) = dex[C(Hβ) f (λ)][I(λ)/I(Hβ)], where [I(λ)/I(Hβ)], the de-reddened

flux ratio, is tabulated.

References1. Dufour, R.J., Shields, G.A., & Talbot, R.J., Jr. 1982, ApJ, 252, 4612. Baldwin, J.A. et al. 1991, ApJ, 374, 5803. Torres-Peimbert, S., Peimbert, M., & Daltabuit, E. 1980, ApJ, 238, 1334. Bohlin, R.C., Harrington, J.P., & Stecher, T.P. 1978, ApJ, 219, 5755. Mathis, J.S., Chu, Y-H, & Peterson, D.E. 1985, ApJ, 292, 1556. Davidson, A.F. et al. 1982, ApJ, 253, 6967. Henry, R.B.C., MacAlpine, G.M., & Kirschner, R.P. 1984, ApJ, 278, 6198. Fesen, R.A., & Kirshner, R.P. 1982, ApJ, 258, 19. Raymond, J.C. et al. 1988, ApJ, 324, 869

10. Blair, W.P. et al. 1991, ApJ, 379, L3311. Barker, T. 1986, ApJ, 308, 31412. Aller, L.J., & Czyzak, S.J. 1983, ApJS, 51, 211

21.4.1 Far-Infrared Line Intensities [19]

Luminosities given in Table 21.12 are estimates for the power emitted inside the solar circle only.

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Table 21.12. Far-infrared line luminosities.

λ (µm) Source 1084π J a log LG λ (µm) Source 1084π J a log LG

1302 CO (J = 2–1) 11 ± 2 4.9 370.4 [C I] 44 ± 6 5.5867.2 CO (J = 3–2) 15 ± 4 5.1 205.3 [N II] 630 ± 34 6.7650.4 CO (J = 4–3) 16 ± 4 4.1 157.7 [C II] 6550 ± 270 7.7609.1 [C I] 24 ± 5 5.3 121.9 [N II] 990 ± 140 6.9519.8 CO (J = 5–4) 12 ± 3 5.0 Cont. Dust 2.2 × 106 10.26

NoteaUnits of J (= mean intensity) are erg cm−2 s−1 sr−1; LG is in L�.

21.4.2 “Unidentified Infrared Bands”

The total power radiated from bright, well-observed objects [20] is > 0.05 of far-infrared. The totalpower from Galaxy is unknown. These unidentified infrared bands are given in Table 21.13.

Table 21.13. The unidentified infrared bands.

Average strengthrelative to 7.7 µm [1]

λ ν FWHM PNe Refl. H II

(µm) (cm−1) (cm−1) (C-rich) Neb. regions Assignment and comments [2]

3.29 3040 30 0.063 0.057 0.13 Aromatic C–H stretch5.2 1960–1890 30 0.035 0.054 0.075 C–H out-of-plane and in-plane bend (?)5.6 1785–1755 40 0.073 0.090 0.058 Overtone of 11.3 µm band; C = O

stretch (?)6.2 1615 30 0.41 0.68 0.58 Aromatic C–C stretch6.9 1450–1470 30 · · · 0.16 · · · Aromatic C–C stretch; aliphatic C–H

deformation7.7 1250–1350 70–200 1 1 1 Blending of several aromatic C–C

stretches11.2 890 30 0.29 0.22 0.30 Aromatic out-of-plane bend for

nonadjacent, peripheral H atoms

References1. Cohen, M. et al. 1989, ApJ, 341, 2462. Allamandola, L.J., Tielens, A.G.G.M., & Barker, J.R. 1989, ApJS, 712, 733

21.5 H2 AND MOLECULAR CLOUDS

The conversion from CO strength to mass: X = N (H2)/∫

T (12CO) dv ≈ 1.8 × 1020 cm−2 K−1

(km/s)−1 in solar neighborhood and inner Galaxy [21, 22]; 1/(4 ± 2) as much in outer Galaxy [23];twice as much for DG < 1 kpc. The value and constancy of X is controversial; some authors use3.6 × 1020 throughout Galaxy.

Molecular clouds are hierarchical in nature and have very poorly defined edges, structures, andmasses. For a large range of sizes they seem to be fractal in nature. The following are statisticalrelationships that are commonly used, but there is a considerable spread about the values derived fromthese relations:

Let R = half the average diameter of a cloud; n = number density of H atoms = 2n(H2); σ =velocity dispersion for a Gaussian distribution (FWHM2 = 8 ln 2σ 2). Then [24]: n(H atoms cm−3) ≈

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21.5 H2 AND MOLECULAR CLOUDS / 533

3.9 × 103 R (pc)−1; N (H) = 2n R ≈ 3 × 1022 cm−2; σ(CO; km s−1) ≈ 0.15(M/M�)1/4;M(M�) ≈ 540R (pc)2; σ(CO; km s−1) ≈ 0.71R (pc)1/2; M(M�) ≈ 40L(CO)0.8.

For the forms given above, clouds approximately obey the virial theorem, but they may besupported magnetically. Other workers have found somewhat different relations; for instance,Larson [25] suggested σ (km s−1) = 1.4R (pc)0.38. Mass spectrum of massive molecular clouds [26]:d N/dM ∝ M−1.5(7 × 104 ≤ M/M� ≤ 2 × 106). This relation implies that most of the mass is inthe massive complexes, but it is controversial [27].

Typical parameters for the observed molecular clouds are given in Table 21.14.

Table 21.14. Typical parameters of molecular clouds [1].

Size Density FWHMa TType (pc) (cm−3) (km s−1) (K)

Giant MC complex 50 100 10 10Giant MC 4 1000 4 25Core of MC 1 4000 2 40Clump within MC 0.5 > 105 4 100Dark cloud complex 10 400 3 10Dark cloud 0.5 104 1 12Dark cloud core 0.2 4 × 104 0.3 10Circumstellar envelope 0.2 102–107 20–40 10–100

NoteaFull width half maximum in 12CO.

Reference1. Goldsmith, P.F. 1987, in Interstellar Processes, edited by D.J. Hol-

lenbach and H.A. Thronson, Jr. (Reidel, Dordrecht), p. 51

21.5.1 Global Properties, Solar NeighborhoodGiant Molecular Cloud Complexes [28, 29]

Mass, (1–2) × 105M�; volume averaged n(H2) within cloud ≈ 50 cm−3; mean N (H2) ≈ (3–6) ×1021 cm−2; mean optical absorption A(V) ≈ 1–3 mag.; local surface density of complexes ≈ 4 kpc−2;mean separation ≈ 500 pc.

21.5.2 Individual Molecular Clouds(Some of Which Coincide with H II Regions)

Some individual molecular cloud positions (equinox 2000) are given in Table 21.15.

Table 21.15. Individual molecular clouds.

Name α (2000) δ (2000) Name α (2000) δ (2000)

Barnard 5 3 48 32 54 NGC 2024 05 42 −1 56L1551 4 31 18 03 ρ Oph 16 26 −24 24TMC-1 4 42 25 41 M17 SW 18 17 −16 14OMC-1 5 33 −5 26 Cep A 22 56 62 02OMC-2 5 36 −5 10 NGC 7538 23 14 61 28

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21.5.3 Molecules Found in Clouds andCircumstellar Envelopes [30, 31]

Line frequencies are given in [32].Inorganic stable molecules: H2, CO, CS, NO, NS, SiO, SiS, HCN, PN, NaCl∗, AlCl∗, KCl∗,

AlF(?)∗, H2O, H2S, SO2, OCS, NH3, SiH4∗.

Organic stable molecules: CH3OH, C2H5OH, H2CO, CH3CHO, H2CCO, (CH3)2CO (?), HCN,HCOOH, HNCO, C2H2

∗, C2H4∗, CH4, NH2CHO, NH2CH, NH2CN, NH2CH3, HCOOCH3, (CH3)2O,

H2CS, HNCS, CH3SH, CH3CN, C2H5CN, HC3N, CH3C2H, CH3C4H, CH2NH, CH2CHCN.Free radicals: CH, CN, OH, SO, HCO, C2

∗, C2H, C3∗, C5

∗, C3H, C3N, C3O, C4H, C5H, C6H,C2S, CH2CN, SiC, HNO, HCC2HO.

Ions: CH+, HCO+, N2H+, HOCO+, HCS+, H3O+(?), HCNH+, H2D+, SO+.Rings: SiC2

∗, C3H2, C3H.Carbon chains and isomers: C3S, HC5N, HC7N, HC9N, HC11N, CH3C3N, CH3C4N, CH3C5N(?),

HNC, CH3NC.[∗Found only in circumstellar material, such as in IRC+10 216 (a carbon star).]

21.5.4 Interstellar and Circumstellar Masers

Molecular masers are observed in both interstellar and circumstellar locations as listed in Table 21.16.

Table 21.16. Interstellar and circumstellar masers [1, 2].

Molecule Freq. (GHz) Transition Where found

OH 1.612231 3�3/2, J = 3/2, F = 1–2 M stars1.665402 3�3/2, J = 3/2, F = 1–1 MCs, M stars1.667359 3�3/2, J = 3/2, F = 2–2 MCs, M stars1.720530 3�3/2, J = 3/2, F = 2–1 MCs4.66042 3�1/2, J = 1/2, F = 0–1 MCs4.765562 3�1/2, J = 1/2, F = 1–0 MCs

H2O 22.235080 616–523 MCs, M starsH2CO 4.829 110–111 MCsCH3OH 12.178 20–3−1 E MCsSiS 18.155 1–0 C starsNH3 23.870 J, K = 3, 3 MCsSiO 43.122 v = 1, J = 1–0 MCs, M stars, S stars

86.243 v = 1, J = 2–1 M stars, S starsHCN 89.087 v = 2, J = 1–0 C stars

References1. Reid, M.J., & Moran, J.M. 1981, ARA&A, 19, 2312. Genzel, R. et al. 1981, ApJ, 247, 1039

The size of the OH masering region in the W51MAIN molecular cloud [33] ≈ (1–3) × 1013 cm.They are ≈ 1014 cm in compact H II regions, and 1015 cm in circumstellar masers. A catalogue ofHerbig–Haro Objects is given in [34].

21.6 NEUTRAL GAS; CLOUDS; DEPLETIONS

Standard cloud parameters for neutral clouds are given in [35]. It is assumed that there are twotypes of clouds, “standard” and “dense,” embedded within a low-density warm intercloud component.

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21.6 NEUTRAL GAS; CLOUDS; DEPLETIONS / 535

Varying mean densities 〈n(H)〉 arise from varying proportions of the three constituents. The suggestedproperties of the components are: (a) Intercloud component, n(H) ≈ 0.14 cm−3, T ≈ 8000 K.(b) “Diffuse” cold cloud, N (H) ≈ 3.5 × 1020 cm−2; number clouds/kpc ≈ 6.2; reddening in eachcloud, E(B − V ) ≈ 0.060 mag.; mean contribution of clouds ≈ 0.70 cm−3. (c) “Large” cold clouds,N (H) ≈ 1.7 × 1021 cm−2, E(B − V ) ≈ 0.30, number/kpc ≈ 0.8; mean contribution ≈ 0.45 cm−3.

21.6.1 Interstellar Depletions from Gas Phase

The depletion, δ, of an element A is defined by

log δ(A) ≡ log(N (A)/N (H))gas − log(N (A)/N (H))�,

where N (A) and N (H) are column densities of the element and H in all forms (H2, H0, and H+),respectively. In practice, usually various stages of ionization are observed, and theory and velocityinformation are used to decide whether the stage of ionization is mainly found in the H0 or H+ phases.There is no depletion information from deep inside molecular clouds because suitable background starscannot be observed through the heavy extinction. Depletions are measured by ionic absorption lines.See [36] for an extensive list.

As shown in Table 21.17, depletions vary with 〈n(H)〉, the mean gas density along the line of sightto the star: n(H) = N (H)/d , where N (H) is the observed column density and d is the distance to thestar whose spectrum shows the interstellar absorption lines of the ion. The 〈n(H)〉 is, presumably, onlya crude measure of local density. A formula for δ is

log δ(A) = d0 + m[log〈n(H)〉 + 0.5].

Table 21.17. Parameters of depletion of element A in H I regions [1].

Elem. log(A/H)� d0 m

C −3.44 −0.34 ± 0.15 aN −4.04 −0.12 ± 0.05 −0.01 ± 0.09O −8.08 −0.36 ± 0.07 −0.01 ± 0.11Fe −4.50 −1.65 ± 0.02 −0.38 ± 0.05Ti −7.26 −1.86 ± 0.04 −0.84 ± 0.06Cr −6.30 −1.58 ± 0.09 −0.50 ± 0.11Zn −7.40 −0.38 ± 0.07 −0.11 ± 0.09P −6.50 −0.43 ± 0.03 −0.32 ± 0.06Si −4.45 −1.09 ± 0.11 −0.49 ± 0.15Mg −4.40 −0.44 ± 0.02 −0.28 ± 0.05Mn −4.50 −1.65 ± 0.02 −1.06 ± 0.02S −4.80 ≈ 0 ≈ 0Ca −5.66 −2.4 ± 0.3 −1.2 ± 0.3Al −4.57 −1.85 ± 0.2 −0.85

NoteaAt time of writing (Sept. 1992) C II, the dominant ion for carbon in

the neutral ISM, has only been measured in three stars [2], all of whichsample relatively dense gas.

References1. Jenkins, E.B. 1987, in Interstellar Processes, edited by D.J. Hollen-

bach and H.A. Thronson, Jr. (Reidel, Dordrecht), p. 5332. Cardelli, J.A., Mathis, J.S., Ebbets, D.C., & Savage, B.D. 1992, ApJ,

402, L17

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21.6.2 Reflection Nebulae

Lists, with exciting stars and references, are given in [37–40].

21.6.3 Halo Equivalent Widths

Table 21.18 gives major ions, oscillator strengths ( f ), and the equivalent widths toward two stars at verylarge distances above and below the plane: HD18100 (z = −3.4 kpc) and HD100340 (z = 5.2 kpc),averaged by [41].

Table 21.18. Equivalent widths of interstellar absorption lines.

Ion λ (nm) f W (nm) Ion λ (nm) f W (nm)

H I 121.567 0.41 1.8 Si III 120.651 1.65 0.055C I 165.693 0.135 0.007 Si IV 140.277 0.26 0.012

156.031 0.081 0.006 139.376 0.53 0.023127.721 0.008 P II 115.28 0.236

C II 133.4532 0.128 0.070 130.19 0.017232.540 6.0(−8) · · · 153.25 7.61(−3)

C IIa 133.571 0.115 0.012 S II 125.95 0.016 0.035C IV 154.819 0.194 0.034 125.38 0.011 0.025

155.076 0.097 0.018 125.06 5.45(−3) 0.014N I 119.99 0.266 0.030 Ca II 396.847 0.3145

115.98 8.51(−6) · · · 393.366 0.635O I 130.217 0.049 0.060 Cr II 205.62 0.140 0.007

135.560 1.25(−6) · · · 206.22 0.105O Ia 130.486 0.049 · · · Mn II 257.688 0.351 0.016Mg I 285.213 1.97 0.030 259.450 0.271 0.014

202.648 0.11 0.006 260.646 0.193 0.012Mg II 123.993 2.67(−4) · · · Fe II 160.8451 0.062 0.030

124.039 1.34(−4) · · · 234.421 0.110 0.080279.635 0.612 0.135 237.446 0.028 0.040280.351 0.305 0.125 238.277 0.301 0.080

Al II 167.079 1.89 0.050 258.665 0.0653 0.075Al III 186.279 0.27 0.013 260.017 0.224 0.080

185.472 0.54 0.025 226.078 0.0037 · · ·Si II 180.801 0.0055 0.013 224.988 0.0025 · · ·

152.671 0.230 0.045 Ni II 131.722 0.146130.437 1.473 0.047 137.013 0.131126.042 1.007 0.060 174.155 0.103 0.008119.329 0.499 0.059 Zn II 202.614 0.515 0.012119.042 0.250 206.266 0.253 0.006

Si IIa 119.450 0.623 · · ·126.474 0.903 · · ·153.343 0.229 · · ·

NoteaArising from an excited state within a few kelvins excitation from ground; population

sensitive to density.

21.7 H II REGIONS, IONIZED GAS, AND THE GALACTIC HALO

21.7.1 Some Galactic H II Regions

W designations are found in [42], and the DR designations are from [43]. Catalogs of H II regions arefrom [44, 45]. Table 21.19 entries are from [46, 47].

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21.7 H II REGIONS, IONIZED GAS, AND THE GALACTIC HALO / 537

Table 21.19. Galactic H II regions.

Size Sν

G desig. (′) 2R 6 cm log(nion)

Name(s) (l, b) α δ (pc) log(EM)a (Jy) (s−1)

W 3 133.7+1.2 1.7 1.5 1.3 6.56 80 49.89M 42, Ori ANGC 1976 209.0−19.4 3.8 4.3 0.6 6.86 430 48.84M 43, NGC 1982 30.8−0.04 2.5 2.5 0.7 5.42 16.7 47.63NGC 2024, Ori B 206.5−16.4 3.6 3.0 0.6 6.08 65.6 48.10NGC 2237-8;2244, 2246, Rosette 206.4−1.9 80 80 34 3.88 350 49.88Sgr B2 0.7−0.0 1.8 1.8 9 6.28 21.4 50.40M 20, NGC 6514Trifid 7.0−0.2 5.4 5.8 6.2 4.76 13 50.40M 8, NGC 6523Lagoon 6.1−1.2 8.5 7.6 2.2 5.15 85 48.36M 17, NGC 6618Omega 15.0−0.7 4.1 6.7 5.2 6.65 900 50.31W49 43.2−0.0 1.5 2.0 7.0 6.23 50 50.96W 51 49.1−0.4 complex · · · · · · 400 50.37DR 21 81.7+0.5 0.3 0.4 0.2 7.70 19 48.62NGC 7538 111.5+0.8 2.3 1.9 3.0 5.78 26 49.73

NoteaEM = Emission measure = ∫

n2e ds (cm−6 pc); nion = number of ionizations s−1 within

nebula.

21.7.2 Galactic Halo and High Stages of Ionization [48–51]

High ionization stages for interstellar elements may be in coronal (collisional) ionization equilibrium.Table 21.20 gives typical column densities, with H I and electrons for comparison.

Table 21.20. High stages of ionization in the galactic halo.a

IP(i − 1) IP(i) n (z = 0) log N hIon (eV) (eV) Tmax fmax (cm−3) (cm−2) (kpc)

Si IV 33.5 45.1 0.6 × 105 0.26 4.1 × 10−9 13.54 2.8C IV 47.9 64.5 1.0 × 105 0.35 1.1 × 10−8 14.20 4.7N V 77.5 97.9 1.8 × 105 0.12 5.1 × 10−9 13.40 1.6O VI 113.9 138.1 3.0 × 105 0.22 2 × 10−8 > 13.27 > 0.3H I · · · 13.6 b b 0.10 20.39 1.5e− · · · · · · b b 0.04 20.06 0.4

Notesa fosc = oscillator strength of transition; IP(i − 1) = ionization potential of lower

species of ionization, from which this species is produced in coronal (i.e., purely collisional)ionization; IP(i) = ionization potential of this species, above which the species can bedestroyed by collisions; Tmax = temperature at which the maximum amount of speciesis produced; fmax = fraction of element in this stage of ionization when T = Tmax;n(z = 0) = density in plane; N = total column density, from mid-plane of Galaxy;h = exponential scale height.

bNot produced by collisional ionization.

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21.7.3 Ultracompact H II Regions

Ultracompact H II regions are ionized nebulae surrounding newly formed stars still embedded withintheir natal molecular clouds. Optically invisible because of the dust in the cloud, they are observed inthe radio and far-infrared spectral regions. About 1700–3000 are seen in the IRAS far-infrared surveyof the Galaxy, representing ≈ 20% of the total O-star population. Their shapes (from radio maps) are:cometary (an arclike structure), 20%; core-halo (bright central source, the core, surrounded by a morediffuse halo), 16%; shell (brighter at rim than in center, circular shape), 4%; irregular and clumpy, 17%,and spherical or unresolved by VLA, 43%. They probably usually represent bow shocks of hot starsmoving ≈ 10 km s−1 through the molecular gas, as viewed from various directions [52, 53]. Some ofthese ultracompact H II regions are listed in Table 21.21.

Table 21.21. Bright ultracompact H II regions.

D τν 10−7 EMa 10−4ne log nion �sb

Name (l, b) (kpc) (15 GHz) (pc cm−6) (cm−3) (s−1) (pc)

G5.89−0.39 2.6 2.67 244.5 23.7 48.65 0.044G10.62−0.38 6.5 1.27 116.0 15.6 48.96 0.048G29.96−0.02 9.0 0.23 21.1 4.4 49.34 0.109G34.26+0.15A 3.7 3.19 291.7 29 48.74 0.035W51d 7.0 3.71 340 20 49.42 0.084

NotesaEM = emission measure = ∫

n2e ds; nion = ionizing photons s−1.

b�s = path length through the emitting region.

21.8 PLANETARY NEBULAE (PNe)

A major uncertainty with PNe is their distances. Many authors refer to the scale of Seaton [54]; othershave found distances similar to those of Cudworth [55], but see also [56], which are statistically 1.47times those of Seaton. Let K = the statistical average of the true distances relative to the Seaton scale.Then 〈M〉 ∝ K 5/2, n = number density ∝ K −3, birth rate per volume ∝ K −4.

21.8.1 General Statistics [57]

The local number density of PNe is 90 kpc−3; scale height above plane of Galaxy = 150 pc; lifetime(as visible PN) ≈ (2–4) × 104 yr; rate of production, locally, = 2.4 × 10−12 pc−3 yr−1; total numberin Galaxy (radii ≤ 0.6 pc) = 3×104; rate of production = 1.3 yr−1; average mass of ionized gas [58]:0.217M� [58] for optically thin nebulae.

21.8.2 Main Catalogs of PNe

Perek, L., & Kohoutek, L. 1967, The Catalogue of Galactic Planetary Nebulae (Praha, AcademiaPress, CSSR) (1036 objects); Acker, A. et al. 1982, Catalogue of the Central Stars of True andPossible Planetary Nebulae, Publication Speciale du C.D.S. No. 3, Observatoire de Strasbourg, andComplements I (1982), II (1983), III (1984) (480 stars) [12]. Line intensities and wavelengths of brightPNe: Kaler, J.B. 1976, ApJS, 31, 517 [59]. Radio observations of PNe: Higgs, L.A. 1971, Publicationsof the Astrophysics Branch, National Research Council of Canada, Vol. 1, No. 1, NRC 12129.

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21.8 PLANETARY NEBULAE (PNE) / 539

21.8.3 General Information Regarding Some Bright PNeand Their Central Stars [60]

Table 21.22 gives positions (eqinox 2000) and some data for selected bright planetary nebulae. Othercommonly known ones are: NGC 3587 = “Owl”; NGC 6853 = “Dumbbell Nebula”; NGC 7009 =“Saturn Nebula.”

Table 21.22. Bright planetary nebulae.

α (2000) δ (2000) Radius Excit. Fν (6 cm)Namea Desig.b (h m) (◦ ′ ′′) (′′) classc log F(Hβ)d (mJy)

NGC 40 120 +9◦ 00 13 72 31 16 18.2 2 −10.64 460NGC 246 118 −74◦ 00 47 −11 52 40 118 8p −10.20 247IC 418 215 −25◦ 05 27 −12 41 43 6.2 3 −9.59 1550NGC 2392 197 +17◦ 07 29 20 54 48 23.0 8p −10.39 251NGC 2440 234 +2◦ 07 42 −18 12 28 16.2 9 −10.45 410Abell 30 208 +33◦ 08 47 17 52 37 64 · · · −12.19 · · ·NGC 6210 43 +37◦ 16 44 23 47 56 8.5 5 −10.08 260NGC 6543 96 +29◦ 17 58 66 37 59 10.0 5 −9.58 850NGC 6572 34 +11◦ 18 13 06 51 15 7.3 6 −9.81 1307NGC 6720 63 +13◦ 18 53 33 01 36 35 6p −10.06 360BD+30◦3639 64 +5◦ 19 32 30 24 18 2.5 2 −10.01 600NGC 7027 84 −3◦ 21 07 42 14 06 4.0 10p −10.17 6370NGC 7293 36 −57◦ 22 29 −20 50 30 360 6 −9.35 1292NGC 7662 106 −17◦ 23 26 42 32 05 7.6 · · · −9.98 600

NotesaCommon designations: NGC 6720 = “Ring Nebula”; NGC 7293 = “Helix Neb.”bThe designation in the Perek–Kohoutek (1967) catalog, or galactic longitude and latitude in degrees,

numbered from north to south within each 1◦ × 1◦ bin.cExcitation classes range from 1–10, increasing with the excitation as shown in line ratios [O II]/[O III];

[O III]/Hβ; He II/Hβ; Ne V/Ne III. See Aller, L.H., 1956, Gaseous Nebulae (Chapman & Hall, London),p. 66. p = peculiar.

d The reddened Hβ flux, in erg cm−2 s−1.

Other information regarding these nebulae [60] are given in Table 21.23.

Table 21.23. Further information on bright galactic PNe.

10−3ne; Star vexpan vrad Dist. RNamea E(B − V ) 10−4T mV , typeb (km s−1) (km s−1) (kpc) (pc)

NGC 40 0.51, 1.3 0.85 11.65, WC8 29 −13 2.3 0.20NGC 246 0.0 0.09, · · · 11.95, OVI 38 −46 0.50 0.29IC 418 0.20 14.0, 0.85 9.93, Of 12 +43 1.25 0.038NGC 2392 0.12 3.4, 1.3 10.43, O7f 54 +64 2.0 0.22NGC 2440 0.3 2.5, 1.4 19 23 +45 2.2 0.17Abell 30 0.0 14.3, OVI 40 2.2 0.68NGC 6210 0.04 7.5, 1.0 12.90, O3 21 −18 3.3 0.14NGC 6543 0.05 4.0, 0.83 11.44, O7+WR 20 −51 1.9 0.092NGC 6572 0.27 20, 1.05 13.0, Of+WR 16 +10 1.6 0.057NGC 6720 0.07 0.6, 1.0 15.00, Cont. 30 0. 0.65 0.11BD+30◦3639 0.24 10.0, 0.8 9.95, WC9 26 −13 1.1 0.013

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10−3ne; Star vexpan vrad Dist. RNamea E(B − V ) 10−4T mV , typeb (km s−1) (km s−1) (kpc) (pc)

NGC 7027 0.93 80, 1.4 19.4 18 +24 0.78 0.015NGC 7293 0.0 13.43 14 −26 0.30 0.5NGC 7662 0.11 4.5, 1.25 13.3, Cont. 26 −5 1.95 0.072

NoteaWR and WC = Wolf–Rayet spectrum (wide emission lines of H, He, C, N, and O; OV I = a WR spectrum

plus strong O+5 lines at 381.1, 383.4 nm; Of = spectrum like young Of stars with emission of H, He II, N III,and usually C III; Cont. = continuous, but often taken with a low dispersion spectrogram.

21.9 SUPERNOVA REMNANTS

Large catalogs of SNRs: Green, D.A., Ap. Space Sci., 148, 3; Saken, J.M., Fesen, R.A., & Shull, J.M.1992, ApJS, 81, 715.

In general, SNRs have synchrotron radio emission showing linear polarization, with Sν ∝ να . Thespectral index varies with type, as given below. Types of SNRs:

1. Balmer-dominated: Steep radio spectrum (a ≤ −0.3); radial magnetic field when young;thermal optical spectrum of strong Balmer lines, weak [O III], [S II]; full or partial shell arrangementof filaments; thermal X-ray emission with both shell and filled forms possible. Ex.: Kepler’s SN(SN 1604); Tycho’s SN (1572); SN 1006.

2. Oxygen-rich: Steep radio spectrum (a ≤ −0.3); radial magnetic field when young; opticalspectrum of strong Balmer lines, strong [O III], high-velocity dispersion; full or partial shellarrangement of filaments; strong thermal X-ray emission; presumably caused by shock within peculiarinternally processed stellar material. Ex: Cas A, Puppis A.

3. Plerionic-composite: Extended, filled (centrally concentrated) centers, possibly with surround-ing shells (“composite”). Flat radio spectrum (a ≥ −0.3), possibly steeper (a ≤ −0.3) in shell, regularin inner part. Thermal optical spectrum plus optical synchrotron emission. Nonthermal X-rays, γ -rays.Inner part probably driven by pulsar. Ex: Crab (SN 1054), 3C 58.

4. Evolved: Partial shell form; steep radio spectrum (Fν ∝ να; α ≤ −0.3); mixed magnetic fielddirection; thermal optical line radiation from shocked ISM, with strong [O I], [S II]; thermal X-rays;slow shocks (50–200 km s−1) from shocked cloudlets of ISM. Most galactic SNRs in this class becauseof long lifetime (final stage of the other types). Ex: Cygnus Loop.

21.9.1 Continuum Spectrum of Crab Nebula(with Pulsar Spectrum Excluded) [61, 62]

The spectrum generally follows a power law between the tabulated points. There are breaks in theexponent of the power law at log(ν) ≈ 12.7, 16, and 19 for the nebula itself, ≈ 8, 14, and 19 for thepulsar. Below log(n) ≈ 8, the pulsar flux increases as ν3.8, peaking at log(ν) ≈ 8. Details are given inTable 21.24.

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Table 21.24. Spectrum of the Crab Nebula; ν in Hz, Sν in Jy.

log ν 8 9 10 12 14 14.7 16 19 21 23log Sν

a 3.3 3.0 2.7 2.13 1.10 0.88 −0.60 −4.05 −7.0 −10.0log Sν

b 1.1 −1.95 < −4.1 · · · −2.58 −2.50 −3.15 −4.92 −7.24 −9.56

NotesaFlux from the nebula, excluding the pulsed radiation.bPulsed flux from the system.

21.9.2 Galactic SNRs [63, 64]

For line spectra of the supernova remnants Crab and Cas A, see the section on spectra. Other details ofsome supernova remnants are given in Table 21.25.

Table 21.25. Galactic supernova remnants.

Date or Typea vexp 2R E(B − V )

Name (age in yr) d (kpc) (km s−1) (pc) (mag)

Cas A 1658 II, O, 2.8 6000 4.1 1.4Kepler’s SNR 1604 I, B 4.4 2000 3.8 1.1Tycho’s SNR 1572 I, B 2.3 1800 5.4 0.7Crab Nebula 1054 II?, P 2.0 1500 2.7 × 4.2 0.5SN 1006 1006 I, B 1.0 > 3000 8.8 0.2Pup A (W44, 3C392) (3700) II?, O 2 400–1600 35 0.12Cygnus Loop (15,000) ?, E 0.77 130 50 0.08

NoteaFirst number, type of SN from light curve and H content of ejecta (I = no H; II = H

present); second, type of remnant (see above: B = Balmer, O = oxygen-rich, P = plerion,PS = plerion/shell, E = evolved). R (pc) from radio image.

21.10 COSMIC RAYS (Excluding Photons and Neutrinos)by D.P. Cox and J.F. Ormes

This information was compiled in October 1992.Cosmic rays observed at Earth are of five types: solar CR, anomalous CR probably arising from

the interaction of the solar wind with the interstellar medium, galactic CR with local origins, galacticCR with distant origins, and CR with extragalactic origins.

Galactic CR have several components [65, 66]:1. Primary nuclei, with source abundances paralleling solar (or local galactic) values, but with a

tendency for elements with higher first ionization potential to be under-represented [67].2. Primary electrons, with a flux about 1% that of protons at the same rigidity (same momentum

per unit charge, and gyroradius).3. Secondary nuclei (easily identifiable by their excess abundances and steeper spectra) generated

by spallation of primaries in collisions with interstellar matter.4. Positrons and an equal number of additional electrons, probably daughters of charged pion

secondaries, with fluxes about 10% of the primary electrons at 10 GeV [68, 69].5. Antiprotons, < 5 × 10−4 as abundant as protons, probably arising as products of high-energy

CR collisions on interstellar matter [68–70].6. Heavier antinuclei, none of which have been found, consistent with the strong limit from γ -

rays [71] on antimatter content of the Galaxy.

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From measurements of radioactive species produced on Earth and in meteorites by CR impact, theaverage CR flux has been fairly constant, at a value like that observed at present, over time scales from103 to 107 yr, with less certain evidence to about 5×108 yr, beyond which it may have been somewhatlower [70, 72]. Except at the shorter time scales, the techniques are insensitive to potential fluctuationsfrom nearby (transient) galactic sources, and none have been unambiguously identified (but see [73]).From the brightness of the γ -ray background, it appears that the CR intensity elsewhere in the Galaxyis not very different from that here, with about a factor of 2 decrease between the inner and outerGalaxy [74].

Given this temporal and spatial uniformity, it is commonly assumed that the CR sample just outsidethe Solar System is representative of that in the larger neighborhood, i.e., that atypical CRs with originswithin a few hundred parsecs are not a major component. See, however, [75–77].

Assuming our sample to be typical of the galactic population, one infers a number of properties,some of which depend on the choice of the model used for propagation and escape (values given hereare ranges for leaky box and simple disk diffusion models at 1 GeV):

1. From the ratio of secondary to primary abundances, mean grammage X along path [78] ≈(9 ± 1)(v/c) g cm−2. Grammage decreases with increasing rigidity above ≈ 4.5 GV (GeV/ec).

2. From the ratio of radioactive to stable secondary isotopes (principally 10Be/ 9Be), the meantrapping (= disk storage) time t ≈ (0.8–2) × 107 yr [79, 80]; mean path length within the disk = ct ≈2–6 Mpc.

3. From (1) + (2), the mean density within the trapping volume n ≤ 0.3 cm−3. This value mayrise due to the cross sections quoted in [78, 80], but is reduced in diffusion models or when there is anested disk-halo box.

4. From (3)+ the surface density of ISM (σ ≈ 2×10−3 g cm−2 ≈ 10M� pc−2 [81]), the verticalextent of trapping region, ±h; h ≥ 0.5 kpc. (This is for n < 0.3 cm−3.)

5. From the synchrotron surface brightness perpendicular to the galactic plane, the effective half-thickness of emission region h ≈ 1–2 kpc [81–83].

6. From the weakness of the radial gradient of γ -ray distribution versus probable cosmic raysources, the vertical extent of the diffusion region h > 15 kpc [74], or sources have weak gradientand h < 4 kpc [84]. Models with a radially decreasing diffusion coefficient or an increasing escapetime encounter no problem with a shallow radial gradient.

7. From comparison of the trapping volume scale with total path before escape, the scatteringmean free path λ ≈ 1 pc ≈ 106 gyro radii at 1 GeV. (With enough convection, the mean free path fordiffusion could be smaller. In a true leaky box, diffusion is not the confining mechanism and the meanfree path could be much longer.)

8. From the anisotropy (0.1%) of the CR flux, increasing above 1015.5 eV, the bulk flow speedof CR past the solar system < 65 km/s [85], also providing limits on nearest sources, the diffusioncoefficient, and the distance of the Sun from the galactic midplane in various models.

9. From the measured flux and spectrum (corrected for solar modulation), energy density u ≈2.4 × 10−12 erg cm−3 [86]; pressure ≈ 1.0 × 10−12 dyn cm−2 [81]. Significant contributions to eachare made by subrelativistic particles which are subject to large solar modulation.

10. From the energy density, grammage, and ISM surface density, the CR source power per unitarea of the galactic disk in the solar neighborhood must equal escaping energy flux: P/A ≈ σuc/X ≈(1–2) × 10−5 erg cm−2 s−1.

11. Summed over the Galaxy, depending on method, the total CR power P ≈ (0.5–2)×1041 erg−1,compared to a total supernova power ≈ (1051 erg/30 yr) ≈ 1042 erg s−1.

12. From the abundances of products of ion-driven chemistry in diffuse clouds, and the γ -rayintensities from diffuse and molecular clouds, CR are present at fairly normal intensity throughoutthe bulk of interstellar matter, where they appear to induce ionization at a rate about 5 × 10−17 s−1

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21.10 COSMIC RAYS (EXCLUDING PHOTONS AND NEUTRINOS) / 543

(H atom)−1 [74, 87]. It is less clear whether they penetrate easily into the densest cores of molecularclouds (a small fraction of the total mass), although ions present there suggest that exclusion, if any, isnot extreme.

Models more physical than the simple leaky box, including diffusion, convective transport, and theinteraction between CR pressure and the confining magnetic field tend to imply a larger galactic halofor particle storage. This leads to rather large differences in interpretation of the CR lifetime, meandensity in the trapping volume, and scale height [88].

“All particle” spectra exist up to 1020 eV, and for individual nuclei and groups to lower energies(≈ 1012 eV/nucleon). For energies between 109.5 and 1012 eV/nucleon, primary nuclei have powerlaw flux spectra: d J/d E ∝ E−γ , with γ ≈ 2.6. The secondary nuclei have a steeper slope(γ + δ) with δ ≈ 0.6, interpreted as a decreasing trapping time scale with increasing energy, andimplying a primary source spectrum with index (γ − δ) ≈ 2 (up to ≈ 2.4) [65]. This spectrumcould derive from acceleration by shocks with a compression factor of about 4 (e.g., strong adiabaticshocks in a nonrelativistic monatomic gas), but most models have difficulty reaching sufficiently highenergy [89, 90], but see [91, 92].

A downward turn or “knee” appears at 1015.5 eV in the all-particle spectrum, similar to the energyat which the anisotropy begins to rise. This has been attributed to a breakdown of magnetic trappingor an upper limit on the energy achieved by a major acceleration mechanism [90].

The correlation of the elemental abundances (for primaries, relative to solar) with the first ionizationpotential, an absence of depletion of those elements common in interstellar dust, an underabundanceof nitrogen, and an overabundance of the neutron-rich isotope 22Ne (and perhaps 25Mg and 26Mg) arestrong contraints on the injectors of primary CR particles. Stellar flares are a potential source, witha small contribution required from He-burning zones, possibly Wolf–Rayet stars [67], but see [78].Injection at MeV must then be followed by powerful acceleration, perhaps by a supernova remnantshock, which generates the full source spectrum. Further accelerations in the ISM must be ratheruncommon [93], but see [91, 94].

Cosmic rays with energies in excess of 1019 eV have high anisotropy, poor trapping in the galacticmagnetic field and probably are extragalactic in origin, perhaps from the Virgo supercluster [90].

Scattering of cosmic rays in the ISM is thought to occur via pitch-angle diffusion, possibly onresonant waves created by the cosmic rays themselves [85]. Conditions in the diffuse ISM are toopoorly known for the wave damping rate (and thus the scattering rate) to be known with certainty.

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