Light-emitting gaseous nebulae in the Galaxy∗

H.C. van de Hulst, M. Minnaert, J.H. Oort and L. Rosenfeld

Originally published in Dutch in Nederlands Tijdschrift voor Natuurkunde, 12, 153 (1946)

I. Introduction (J.H. Oort) Among the luminous nebulae, two groups must be distinguished which differ radically in their origin, viz. the so-called diffuse nebulae, such as the Orion nebula and the nebula around η Carinae, and the planetary nebulae, so christened, I believe, because of their often planet-like appearance (cf. the accompanying figure). In the first ones we are dealing with condensations in the interstellar gas, probably caused by its turbulent motions and brought to light by the, as a rule accidental, proximity of a star of high surface temperature and great intrinsic luminosity. The second, on the other hand, must be understood as shells of light ejected from a star of particularly high temperature, which are also excited by this star. The character of this emission and the connection with the eruptions observed in nova eruptions are discussed in a later chapter. The spectra of nebulae consist of bright lines and, at least in the case of plan- etary nebulae, a usually faint continuum. As in Sun and stars, hydrogen is the predominant element. The intensity ratios vary greatly from nebula to nebula and also from one part of a given nebula to another. The same is true of the ratio of emission lines to continuum. In the case of diffuse nebulae, the continuum is probably formed by radiation scattered by small solid particles from the exciting star. In the case of planetary nebulae, it can perhaps be attributed to nearby pas- sages of electrons past protons. The following, somewhat schematic data, show something about character of the excited objects:

∗Astronomical Colloquium of the Netherlands Astronomers Club, No. 6.

1 Diffuse emission nebulae Planetary nebulae Radius in light years 1–10 0.05-1.0 Density in 10−24 g/cm3 10–1000 10,000 Spectrum central star normal B0, O, W dwarf O and W Density central star (unit solar density) 0.1 500 Luminosity central star, in Sun as unit 5000-10,000 50 Temperature central star 25,000◦-70,000◦ 30,000◦-100,00O◦

The problems that demanded a solution include the following: 1. How does the light emission of the atoms in the nebulae originate? 2. Where do the strong, until 1928 unidentified so-called nebulium lines come from? 3. How do the gas shells of planetary nebulae and their peculiar structure origi- nate? 4. What is the relationship between the planetary nebulae and other stellar types? 5. What is the composition of the nebulae and specially that of the interstellar medium? 1, 2 and 5 will be treated in chapters I, II and III, 3 and 4 in chapters IV and V. The process by which the emission of light is predominantly determined is the following: The uv radiation of the star ionizes H, He, O and other atoms. Passages of free electrons and ions give rise to continuous emission, capture of electrons to line emission and emission bands at the limits of the series. If we assume that the nebula is so thick that it can absorb all the uv energy from the star, then it is clear that from the Balmer lines integrated over the entire nebula the total uv energy emitted by the star can be calculated. Indeed, if we restrict ourselves to hydrogen alone, then for every absorbed uv quantum that detaches an electron, a recom- bination of an electron with a proton must take place as compensation, and one can estimate at which fraction of these recombinations a particular Balmer line, e.g. Hγ, will be emitted. Zanstra, who first performed calculations of this nature for the planetary nebulae, estimated this way the uv radiation of central stars; by comparing this radiation with the directly observed radiation in the photographic range, he was also able to estimate the temperatures (based on the hypothesis that the stars are ‘black bodies’). A refinement of these calculations naturally led to theoretical results for the intensity ratios of the different Balmer lines and from these to the Balmer continuum. Van de Hulst will present you details about the last and most comprehensive calculations of Menzel and his collaborators. Pre-

2 liminary, somewhat more simplified calculations by Cillie´ and by Carroll, did not lead to agreement with the observations; but this appears to be mainly due to neglect of the of the influence of interstellar absorption. The explanation of the unknown nebular lines, the so-called nebulium lines, was found in 1928 by Bowen. The strongest of these lines are so-called forbidden transitions in the ions of O, N, Ne and some other elements. That these forbidden lines can prevail under the conditions which occur in gas nebulae is understand- able. Namely, we have here to do with gases which are so rarefied and occur in such a rarefied radiation field that the electrons are hardly removed from the metastable levels by collisions or by absorption of radiation and thus finally es- cape only along the narrow, so-called forbidden path. This, however does not yet make comprehensible why these lines have such a large intensity, much greater than that of the ordinary lines of the atoms concerned and even mostly greater than the lines of the certainly 100 a` 1000 times more frequent H-atom. This is because the forbidden lines have a particularly low impact-energy (for the green nebulium lines e.g. only 2.5 eV, whereas the ionization voltage of the ion in- volved is 55 eV). The lines can thus be excited by rather slowly moving electrons, which are present in abundance. Compared to these the ionizations by uv quanta or the associated capture of free electrons are rare events. These processes, which play the leading role in ordinary lines, contribute only an imperceptible fraction to the forbidden lines, which are generated almost entirely induced by collissional excitation. The accurate calculation of the process of the excitation of these forbidden levels is rather complicated; it is important, however, especially since without such calculations it is not really possible to make estimates of the relationship between O, N and H in planetary nebulae and in interstellar space. Knowledge of this ratio is of the utmost importance for the investigation of the state of interstellar gas and smoke. The following two chapters show how this problem has been tackled by Menzel and his collaborators. An interesting by-product of these calculations is the determination of the electron density and hence density in the planetary nebulae. In addition to the two above-mentioned excitation mechanisms also fluores- cence processes play an important role in the generation of certain emission lines in nebulae (in particular the fluorescence caused by the very close coincidence of the resonance line A303,780˚ of HeII with the ground-level of the absorption line A303,799˚ of OIII). In some luminous nebulae the emission may come about through an entirely different process, viz. by heating as a result of collissions of two interstellar gas masses.

3 The starting point of the colloquium was the series of articles by Menzel and his collaborators Baker, Aller, Shortley, Hebb and Goldberg, which were pre- sented under the title ‘Physical Processes in Gaseous Nebulae’ and published in the Astrophysical Journal. From 1937 to 1941, 17 articles were published in this series, which is still in progress; it is still being continued1. The subject mat- ter dealt with in the following article, however, is extending this further. On the one side a physical introduction has been added, on the other side a discussion of forms of and movements in nebulae is added also.

Caption to the figure: 1. Central part of the Orion nebula (60-inch Mt Wilson). 2. Ring Nebula in the Lyra (100-inch Mt Wilson. 15 min. and 120 min. exposures, respectively). 3. Drawings of the so-called N1-line of OIII in N.G.C. 2392, resp. with slit in O- W- and N-S-direction (Lick Observatory). The portion in which the line is double corresponds to the inner ring of the nebula (cf. panel 4).The spikes in which the lines extend upward and downward relate to the weaker outer shell. 4. Two images of the planetary nebula N.G.C. 2392 (60-inch Mt Wilson). 5. Planetary nebula N.G.C. 7009 (60-inch Mt Wilson). 6. Crab nebula in light of wavelength 3600-5000 A˚ (100-inch Mt Wilson). 7. Two slitless spectrograph recordings of the planetary nebula N.G.C. 7009 (Lick Observatory; cf. also panel 5). The brightest image in the center is that of the line λ 3869 of NIII, to its left λ 3727 of OII. The two brightest interlocking images at the far right are the so-called N1- and N2-lines of OIII (λ 4959 and 5007). Just to the left of these are Hβ and a little further left λ 4686 from HeII. Striking is the difference between the light distribution in this last line and the with lower ionization degree corresponding OII-line λ 3727.

1The concluding article, on the chemical composition of planetary nebulae, has in the mean- time also appeared (Aller & Menzel, Ap. J. 102, 289, 1945).

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