Download - THE CHEMICAL COMPOSITION OF THE INTERSTELLAR MEDIUM

THE CHEMICAL COMPOSITION OF THE INTERSTELLAR MEDIUM*

Lawrence H. Aller University of California Los Angeles, Calif.

The physics and chemistry of the interstellar medium present one of the most challenging problems in contemporary astrophysics. The problems are intimately intertwined: we must know the physical situation, temperature density, and SO on in the emitting material, but these in turn depend upon its chemical composition.

For a long time progress was painfully slow. It is nearly seventy years since Hartmann found the interstellar H and K lines of ionized calcium, and more than half a century has elapsed since Miss Heger discovered the sodium D lines. In the twenties, it was customary to speak of the interstellar “calcium clouds” - a terminology quantitatively as accurate as referring to rain clouds as nitrogen oxide clouds. In those days it was also customary to speak of the “calcium chromo- sphere.” The lines of ionized calcium were well placed for observation in the optical spectrum and were prominent in many celestial sources. Astrophysics moved slowly out of the “calcium age.”

By the end of the thirties, what appears to be the correct picture, began to emerge. The studies of Struve and his collaborators of the faint H I1 regions, by means of a specially designed nebular spectrograph, revealed the overwhelming importance of interstellar hydrogen; Dunham’slO careful investigation of the curve of growth for interstellar lines observed in the spectrum of 5 Ophiuchi yielded additional clues to the physical nature of the medium, dilution of the radiation field, and so forth, while W y ~ e ’ s ~ ~ classical investigation of the spectrum of the Orion Nebula indicated that it had essentially the same chemical composition as stars such as the sun. Since then, enormous progress has been made in interpreting interstellar lines of the optical spectral region,lg in studies of bright H I1 regions such as the Orion Nebula, and particularly by the opening of new windows on the electromagnetic spectrum in the infrared, radio-frequency region, and ultraviolet.

Observations of interstellar lines yield rather specialized information. In the optical region they tell something about the kinematics and motion of the material in the direction of specific bright stars. In the ultraviolet region of the spectrum, where fall the resonance absorptions of several abundant elements and also important molecular lines, the information yield is considerably greater.

Radio-frequency observations have provided a wealth of intriguing data. Types of interstellar lines include high-level, n - 100, atomic recombination transitions of hydrogen, helium, and ionized carbon, and transitions of diatomic molecules and of polyatomic molecules. The latter appear to originate mostly in small condensations with dimensions comparable with that of the solar system. They show remarkable time variations and maser effects and, for the time being at least, they are more likely to cast light on processes pertinent to the origin of solar systems and stars than directly to abundance problems, that is, except for data pertaining to isotope ratios that can be found from these interstellar molecular microwave lines.

For example, Snyder and B ~ h 1 ~ ~ have detected both H12C14N and H13C14N at wavelengths of roughly 3 mm in Orion A and Sgr A. They found N ( W) /N(I3C)

* This research was supported in part by National Science Foundation grant GP-23460.

45

46 Annals New York Academy of Sciences

= 8.9 for Orion and 4.7 for Sgr A. Penzias has detected CO in three isotopic forms, 12C160, 13C160, 12C1*0, for example in Sgr B2, but the isotopic ratio is a function of radial velocity indicating possible saturation effects.

Formaldehyde has been detected in two forms, H212C1E0 and H Z ~ ~ C ~ ~ O , and has given isotope ratios N( 12C) /N(13C), equal to 8.6 for Sgr B2, and 50 for W51 as compared with the terrestrial value of 89. Optical depth effects may complicate things. To date, there have been few interferometer observations which could verify the optical depth explanation of the discrepancy. Possibly, the 13C/12C ratio in the interstellar medium may actually differ from the terrestrial ratio, although a number of workers have suggested it is the same.

Infrared spectroscopic observations are extremely important in that they pro- vide data on certain forbidden lines of [NeII], [SII], and other ions that cannot be observed in conventional spectral regions. The ultraviolet interstellar extinction law also gives some information on the probable qualitative chemical composition of the solid grains-graphites, magnesium, silicates-but the data are mostly in- direct and can cast little light on the quantitative chemical composition of the interstellar medium.

Bright H I1 emission regions, that is, diffuse nebulae, offer some very special ad- vantages. A representative portion of the interstellar medium is heated to incan- descence and caused to fluoresce by the radiation of hot stars embedded within it. The emergent radiation is principally that of abundant permanent gases- hydrogen, helium, nitrogen, oxygen, neon, and argon-and also of sulfur, carbon, and occasionally weak lines of silicon, magnesium, and iron. An understanding of the relevant physical processes was developed in the twenties and thirties, and the necessary quantitative ideas and equations were expounded by Menzel and his associate^.^'

Later, I shall discuss the uses that may be made of gaseous nebular data in more detail, but I must mention some of the restrictions and limitations that are en- countered. First, the chemical elements whose abundances can be found from their emission line spectra in gaseous nebulae are limited to the more abundant ones. Second, most elements are observed in only one or two stages of ionization although they may actually exist in several (invisible) stages, and third, the analysis is enormously complicated by the clumping tendency of the material. There exist huge fluctuations in density, and possibly also in temperature, from point to point in most gaseous nebulae.

On the other hand, it must be remarked that the He /H ratio can probably be established much more accurately for the interstellar medium than for any star. The He/H ratio is an extremely important datum for many popular cosmological theories. It is fortunate that it can be measured in so many places throughout our universe.

Another trick is to use the Fabian approach, that is, to study the chemical compositions of the atmospheres of stars that have but recently been formed from the interstellar medium. Almost any early type B star will do, since its age can scarcely exceed a few tens of millions of years. Stars of spectral class BO and 0 stars are likely to be less useful for our purposes-since we cannot even pretend to know enough about their atmospheric structures-and spectral lines must certainly be formed under conditions departing severely from local thermodynamic equilibrium. We must refer here to investigations by Underhill, Strom, Avrett, and their associates, and by Mihalas and Auer.l*

An account of the principles of abundance determination from stellar atmospheres would be quite extraneous to this presentation. A few general statements

Aller : Composition of the Interstellar Medium 47

may suffice. We interpret the spectrum of a star with the aid of a model atmosphere, often by an iterative process involving the chemical composition; however, in any event, we must know the effective temperature and surface gravity. The model atmosphere predicts the emergent radiation from a star, in the continuum as well as the lines. As one example, I might mention the work of Mrs. Petersz8 on the B3V star iota Herculis, for which available energy distributions, the measured ultraviolet flux, and determinations of angular diameters of similar stars indicate an effective temperature of 17,000° K. The surface gravity is near 1 04cm-2. The derived relative abundances are remarkably similar to those found for the sun. Somewhat similar results have been found for other stars by a variety of investigators, supporting a claim made some years ago that there is no established difference between the present chemical composition of the interstellar medium and that of the sun which was formed 4.5 X lo9 years ago. I feel that this statement can be taken only as a general guide. We must look for small but perhaps significant differences.

The questions to which we would like to have answers are the following: 1. Does the interstellar medium have the same mixture of elemental abun-

dances now as it had several aeons ago? Has its chemical composition changed appreciably with time?

2. Does the chemical composition of the interstellar medium change from point to point within the galaxy? How does the interstellar medium of our galaxy compare with that of other stellar systems?

3. What can be said about elemental abundances including isotope ratios in the interstellar medium in the neighborhood of the sun?

Of course, we can give no direct answer to the first question since we cannot find an uncontaminated sample from that period. All we can do is to examine the spectra of stars that are believed to be of very great age. I shall not discuss this topic here other than to recall that some very ancient clusters such as the globular cluster 47 Tucanae and the galactic cluster M67 (also of very great age) appear to have essentially the same chemical composition as the sun, while other globular clusters such as M92 appear to be metal-deficient. Apparently, the composition of the interstellar medium changed very rapidly sometime during the early history of the stellar system, and relatively slightly since the time of the formation of the sun. But I think it is very important to find out just how much the interstellar medium changed from the epoch of formation of the solar system to the present time. Of course, if the composition of the medium changes from point to point in the galactic system, the pertinent datum is a comparison of the composition of the present-day interstellar medium at the point where the sun was formed. Since the sun is moving in a nearly circular orbit and since the composition of the interstellar medium changes slowly, if at all, with distance from the galactic center, it seemed legitimate to regard the local medium as corresponding to the medium of the region of the formation of the solar system, wherever that was.

The He/H ratio can be measured from line ratios in the optical region of the spectrum, e.g., 1(5876)/ (HP) , or 1(4471)/I(HP) or from radio-frequency lines. In each instance we are concerned with lines produced by recombination and subsequent cascade. TABLE 1 gives the He /H ratio for several diffuse galactic nebulae regions as measured by various observers using both optical and radio methods.

What about the stellar results? We note that here the values depend on what you assume about theories of line broadening. Shipman and S t r ~ m ~ ~ obtained


N(He/N(H) = 0.10 from the spectra of B stars, while Leckrone,15 from a series of B-star spectra carefully obtained for this specific purpose at Lick Observatory with the 120-inch CoudC, found 0.12. From main sequence eclipsing binaries in the spectral range A-G, Popper and coworkers30 obtained 0.12 and concluded that the ratio is unlikely to lie outside the range 0.09-0.14. They identified seven of thirteen of these main sequence binaries as unevolved and used the theoretical mass luminosity relationship to get the helium/ hydrogen ratio.

For the sun, one cannot determine the He/H ratio by spectroscopic methods. From a solar model with improved opacity and nuclear-reaction data, Bahcall and Ulrich' found N(He)/N(H) = 0.089. On the other hand, solar cosmic ray data predict a rather low value, -0.06 * 0.01 for this ratio. At the present time, we should accept probably a value of 0.1 1 as the most likely value for the helium/ hydrogen ratio in our part of the galaxy, remarking that the solar data are not yet properly explained .

Turning, to external galaxies, numerous measurements have been made of the He/H ratio in the local group of galaxies (TABLE 2) .

Peimbert and SpinradZ8 find a mean ratio N(He)/N(H) = 0.10 in several more remote external galaxies, NGC 4449,5461, 5471, and 7679. Thus, we reach the interesting conclusion that helium/hydrogen ratio is remarkably similar from one galaxy to another in our local region of the universe.

What about other elements? From a study of the nuclei of M51 and M81, Peim- bert argues that both galaxies have an overabundance of nitrogen in their central regions. He attributes most of the variation of the [NII] / H a intensity ratio from the spiral arms to the nuclei to an actual variation in the nitrogen abundance rather than to a change in physical conditions. Peimbert and Spinrad conclude

TABLE 1 THE H E L I U M ~ ~ Y D R O G E N RATIO FOR SEVERAL GALACTIC NEBULAE.

Object N(He)/N(H) Method Reference

Orion 0.12 optical 11, 4a, 16 0.12 2 0.01 optical 23 0.083 I?r 0.004 radio 23 0.11 optical 27

M 8 0.113 optical 16 0.11 optical 27

M17 0.107 optical 21 0.09 I?r 0.01 radio 23

M20 0.097 optical 16 Carinae 0.122 f 0.02 optical 11

W5l 0.081 I?r 0.008 radio 23

NGC 6357 0.088 k 0.018 radio 23

IC 1795 0.081 k 0.016 radio 23

Notice that ratios measured by radio-frequency methods tend to be smaller than those found by optical methods. One must, of course, allow for the fact that in some H I1 regions much of the helium exists as neutral gas and, therefore, its recombination lines are not observable. Particularly in the extended regions measured by the radio techniques, there are likely to be large zones in which the helium is neutral so that the measured ratio will tend to be too low. We prefer, therefore, the results of the optical studies and adapt them for our local region of the galactic system: N(He)/N(H) = 011.

Aller : Composition of the Interstellar Medium 49 TABLE 2

THE HELIUM/HYDROGEN RATIO IN EXTERNAL GALAXIES

Nebulosity Galaxy Ratio Reference

NGC 604 M33 0.106 16 0.12 1. 0.13 28

NGC 346 SMC 0.076 f 0.02 11.

30 Doradus LMC 0.082 f 0.02 11. *The value 0.12 has been augmented to 0.17 by Peimbert and Spinrad to allow for the

presence of neutral helium, but I believe this correction to be too large. Possibly, the results for NGC 346 and 30 Doradus might have to be similarly increased. However, it will be difficult to make such corrections until detailed structural models of these nebulae have been developed. Johnson1* obtained 0.1 3 for 30 Doradus by photographic photometry.

that the irregular galaxy NGC 6822 is deficient in nitrogen and oxygen by factors of 6 and 1.7 respectivel,, whereas the helium/hydrogen ratio is similar to that found in other systems. Further investigation of the physical conditions in such systems is necessary because the [NII] lines are enhanced in regions of low excitation and high density. They are very sensitive to the clumpiness of the material-as exhibited for example in a number of planetary nebulae (Wyse 1942, Aller and Walker 1970), and also apparently in objects such as 30 Doradus. The [SII] lines are even more sensitive indicators of density and excitation conditions. It is probably inappropriate to lay much stress on abundance ratio differences amounting to less than a factor of two until the scale and character of the clumpiness is carefully evaluated.

The Chemical Coinpositions of H I1 Regions

The most thorough and satisfactory discussion of galactic H I1 regions is that by Peirnbert and Costero27 for Orion, M8 and M17, OsterbrockZ2 has given a summarizing report on the abundance of elements in several prominent types of gaseous nebulae. Some of the relevant formulas and methods are described in Aller and Liller.48 It is possible to extend these studies to additional elements by including weak lines measured by photographic photometry and calibrated by photoelectric methods. For a number of these ions, the necessary atomic parameters have been computed recently.

First, let us consider ionized carbon. The only line of ionized carbon that is usually observed is the 3d-4f transition X4267. If b4 ( tc) measures the deviation of the population of the fourth level from that appropriate to local thermodynamic equilibrium at a temperature T, = 104tc, then:

where &,2(tH) is a function given by Clarke (1965), tH = T,(H+)/~o,ooo. Here T , (H+) denotes the electron temperature in the regions where hydrogen is ionized, and t = T,(C+ + ) / 10,000 where T,(C+ +) is the electron temperature in the region where carbon is doubly ionized.

For the range of temperatures with which we are normally concerned, Clarke's El2 may be represented by

E;,z(tH) = 1.21t~-'.''. (2)


Using Bednarek and Clarke'sa values for b4, we have the following approximate eauation:

Likewise, we can write down expressions for chlorine, sulfur and argon, com- paring the ionic concentrations with those of ionized hydrogen or doubly ionized oxygen, viz. :

For 0.75 < t < 1.25

Outside of this range of temperature and density, it is necessary to use the expres- sion :

where b2(t) is obtained from the work of Aller and coworkers2 (See TABLE 3). Here t = T, (C l+ +) / 104, where T,(Ct++) is the electron temperature in the region of doubly ionized chlorine and tH = 10-4T,(H+) pertains to the electron temperature in the region where the hydrogen lines are radiated predominantly. It should be emphasized that this choice of differing electron temperatures for the ionized hydrogen and ionic zones represents an attempt to allow for the fact that there may exist large temperature differences, as has been discussed by Peimbert.

In the Orion Nebula we also observe the auroral transition A5191 lS - ID of [ArIIIl and the nebular lines of [ArIV]. We may compare the concentrations of Ar++ and Ar+ + + ions with that of O+ + ions, with the aid of the following equations:

(6) I(5191)

I0 exp( -2.89/t0) ~

N(Ar+ + ) N(O+ + )

where tA2 = 10-4T,(Ar++), and to = lO-*T,(O++) and

wheretAB = 10-4T,(Ar+++). In most instances, these temperatures are probably very nearly the same as

those pertaining to doubly ionized oxygen.

TABLE 3 LOG C(X,T) FOR [CIIII]

t -2 -1 -0.5 0 +o.s +1.0 log x

0.75 -4.35 -4.38 -4.59 -4.45 -4.45 -4.33

1 .o -4.80 -4.82 -4.89 -4.88 -4.87 -4.75

1.25 -5.07 -5.09 -5.12 -5.16 -5.14 -5.00

Aller: Composition of the Interstellar Medium 5 1

TABLE 4

LOG P(X,T) * FOR [SII] A4068

log x T- -2 -1 4 . 5 -0 +0.5 1 .o

0.5 0.28 0.18 0.05 -0.16 -0.35 -0.53 0.75 0.27 0.18 0.05 -0.14 -0.34 -0.50 1 .o 0.27 0.17 0.04 -0.14 -0.33 -0.48 1.25 0.26 0.17 0.04 -0.13 -0.32 -0.47 1.50 0.26 0.17 0.04 -0.13 -0.31 -0.46 1.75 0.26 0.17 0.04 -0.13 -0.30 -0.45 2.0 0.26 0.17 0.05 -0.13 -0.29 -0.44

-

* t = I O - * T , ( ~ + ) ; x = 10-*N&Tc

For ionized sulfur, one may, for example, use the transauroral line A4068. Then we have

(8) N(S+) - 0 .00373~ I(4068) -

N(O+ + ) ( ~ + 6 7 ) b ~ ( t ) ~ ~ ~ ( 3'53'ts) exp( -A) I, where b5(t) = b(2P3/2)/b(4S3/2) is tabulated by Czyzak and associatese (TABLE

1) . Alternately, we may write

(9) N(s+t+' = p,(X,t) exp ( y ) e x p - (T)- 2.89 I(4068)

N(O 1 I, where p,(X,t) is given in TABLE 4. Over most of the relevant temperature range, the factor varies slowly with density.

TABLE 5 compares the results for the Orion Nebula with the sun and B stars. The Orion results for helium, nitrogen, oxygen, and neon are quoted directly from Peimbert and C o ~ t e r o . ~ ~ Their result for sulfur is slightly modified because I have used the newerg theory for the excitation of the [SII] lines, which takes into account the collisional exchanges between the levels of the 2D and 2P terms. For carbon, I have used the Bednarek and Clarke theoretical results. The final value depends on the correction one employs for the numbers of atoms in other stages of ionization; with Peimbert and Costero's intensity measurement for I(4267) and a 10% correction, one obtains log N = 8.24. I have obtained a higher value by assuming a larger correction factor and a larger value for the intensity of I(4267).

For [CCIII] lines, I've adopted t = 0.75, x = 0.50 and have assumed that the same correction factor holds for [CCIII] as for sulfur as in Peimbert and Costero's

TABLE 5 A COMPARISON OF THE CHEMICAL COMPOSITIONS OF THE ORION NEBULA,

THE SUN. AND Two TYPICAL B STARS.

H He C N 0 Ne CI S Ar

Onon 12.00 11.04 8.37 7.63 0.79 7.06 4.94 7.47 5.95

Sun 12.00 - 8.55 7.93 8.86 7.97 5.63 7.21 7.3

Bstars s12.00 11.00 8.45 8.03 8.92 8.05 6.0 7.02 6.08 2.25 2 .22 f . 2 4 2 .14 k.5 e . 3 1 2 . 3 2

The quantity tabulated is log N on the scale log N(H) = 12.00.


region Orion 11, namely, 2.40. I've adopted the argon abundance as the sum of the abundances of Ar+ + and Ar+3. For the chlorine and argon lines, it is necessary to use photographic measurements of intensity.13*6 These intensities have been calibrated as carefully as possible by comparison with the available photoelectric measurements.

The solar values are based on measurements by OMara,*O Lambert and Warner (1968) and others, as compiled recently (Aller 1968). There are no photo- spheric values for helium or argon. The chlorine value is taken from Lambert and Mallia,14 the argon value is from the solar corona.31 It seems to be definitely high.

As a set of representative numbers for young, early type stars recently formed from the interstellar medium, I chose an average of the two B stars 22 Orionis and Pegasi, which Mrs. Peters has analyzed using Mt. Wilson coudd data. She has estimated the errors in log N for each element. I have not attempted to estimate the corresponding errors in log N for the Orion Nebula. These errors depend on the clumpiness of the material, temperature fluctuations, and improper allowance for distribution of atoms among different stages of ionization.

The role of density and temperature fluctuations in gaseous nebulae is very important. Some nebulae show domains of very high density and low excitation. The Mg I A4571 and [Mg I] A4562 enable us to identify regions of very low ionization and, presumably of low temperature. Recently, Van Blerkom calcu- lated the target areas for the collisional excitation of the upper 3P term of these transitions. Assuming the deactivation coefficients to be the same for both the 3 P ~ and 3P1 levels and that there is no substantial collisional interchange between the two levels of the same term, I find

N(Mgo) - 4.3 t-0.79e~p(3.14/tM) I(4571) ~- N ( H + 1 tH0.87 I(H/3)

and

In the high excitation planetary NGC 7027, Mg I A4571 and [Mg I] A4562 are observed with intensities 1.9 and 0.07 respectively. The relative intensities of the two lines which must be produced in the same regions are thus given by

I(4562) - 1.52 -~ - - I(4571) N,t"."X 10-4+0.92

which would imply an electron density of the order of 4 X lo5 in the regions where the Mg I lines are produced. That such high-density regions exist in NGC 7027 has been indicated in recent work by Goldberg ( 1971). In 1954, I suggested that the spectroscopic results for NGC 7027 ". . . can be understood in terms of a nebula consisting of numerous filaments, knots and tenuous regions such that the density may range from less than lo4 ions/cm3 to perhaps something like 200,000 ions/cm3."

The 4562 and 4571 lines are also observed in other planetaries, notably NGC 40, but we have not found these lines in the Orion Nebula or other diffuse nebulae. The conclusion is that in the luminous, ionized hydrogen zones of Orion and similar nebulae there are few, if any, dense cool regions such as those which characterize the planetaries. Hence, we are probably better justified in deriving

Aller: Composition of the Interstellar Medium 53 TABLE 6

ADOPTED ~ N T E N S I T I E S IN NGC 604' ~~~~ __

A Ion Log I Reference A Ion Log I Reference

3727 roiii 2.42 1.28 4363 rom 0.03 28 3722 [SIllj 0.62 1' 4471 'He.1. 0.82: 1 3865 [Nel I I ] 1.16 1 3968 [ N e l l l l 0.64 1 5007 l O I l I l 2.29 1,28 .~ 4068 [SII] 0.35 I 5876 He1 1.12 28 4076 [SII] 0.13 1 6563 Ha 2.46 28 4026 He1 0.45 1 6584 "111 1.60 28 4267 C I I -0.5: 1

* Corrected for space absorption. .-

chemical abundances for these objects than we are in similar treatments of planetary nebulae-or at least those planetaries that show high-density condensations.

One of the best studied nebulosities observable in the local group of galaxies is NGC 604 in the Triangulum Spiral, Messier 33. Measurements of the principal nebular lines have been made by Peimbert and SpinradZ8 and by Peimbert,28 who used photoelectric spectrophotometry, and by Aller and collaborators,' who employed both photoelectric and photographic techniques. Peimbert and Spinrad suggested a space absorption correction C amounting to

TABLE 6 gives the intensities, thus corrected for space absorption for lines used in the abundance analysis described below.

For the analysis, I have chosen a two-component model wherein the lines of [OIII], CII, [NeIII], and [SIII] are all produced in a zone whose electron temperature is taken as 10,000" K. The lines of [OH], [NII], [SII] and most of the H emission are presumed to be produced in a lower temperature region where T, = 8,000" K. Except for CII, where I've used the results of Bednarek and Clarke, the equations and atomic parameters employed are essentially those of Peimbert and Costero. Correction of ionic abundances to allow for concentrations of atoms in unobserved stages of ionization is carried out also by essentially the same pro- cedure as that used.by Peimbert and Costero, except that I've increased the helium abundance by only 10%. I've also carried out calculations on the assumption that the temperature of the HP region is 7,000' K instead of 8,OOOO K. TABLE 7 gives

TABLE 7

LOG N ABUNDANCES IN NGC 604 IN TRIANGULUM SPIRAL

He C N 0 Ne S

T, = 7000 11.08 8.48 8.14 8.81 8.19 7.64 8000 11.08 8.45 7.71 8.51 7.88 7.57

TABLE 8 ABUNDANCES IN 30 DORADUS

He C 0 Ne S

Log N 11.00 7.91 8.49 7.90 7.46


the results. A determination of the electron temperature fluctuations is urgently needed for this object.

I've employed a similar model for 30 Doradus. Here, the electron temperature defined by the [OIII] lines, viz., 11,000" K is assumed to apply to regions wherein are produced the lines of [Ne 1111, [OIII], CII, and [SIII], whereas the lines of [OII] and [SII] are presumed to originate in a zone with T, = 8,000" K, which is assumed to produce most of the hydrogen emission. The results are shown in TABLE 8.

The results for neon and sulfur are in good agreement with those obtained for Orion. I attach no significance to the oxygen abundance being a factor of two lower. This nebula is of lower excitation than Orion and apparently much more inhomogeneous, so that the correction for atoms in unobserved ionization stages (mostly the neutral stage) may be inadequate. Also, the carbon abundance turns out to be too low. Since what we measure from the A4267 line interpreted as a recombination feature is the amount of doubly ionized carbon, it is probable that we have underestimated the correction for the amount of neutral and singly ionized carbon.

The conclusion to be drawn is that within the limitations of the analysis, the chemical composition of nebulosities NGC 604 and 30 Doradus in our neighbor- ing galaxies is probably very nearly the same as in our own galaxy. Due to the severe limitations imposed by the inhomogeneity of the radiating material, we cannot now obtain abundance estimates accurate to better than a factor of two, if indeed we can do that well.

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1969. Ap. J. Letters 157; Letters 167.

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