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University Microfilms 300 North Zeab Road Ann Arbor Michigan 46106 A Xerox Education Company 73 - 11,602
WARNER III, John Ward, 1944- A SPECTROSCOPIC STUDY OF GALACTIC NUCLEI.
The Ohio State University, Ph.D., 1972 Astronomy
University Microfilms, A XEROX Company , Ann Arbor, M ichigan
THIS DISSERTATION HAS BEEN MICROFILMED EXACTLY AS RECEIVED. A SPECTROSCOPIC STUDY OF GALACTIC NUCLEI
DISSERTATION
Presented in Partial Fulfillment of the Requirements for the Degree Doctor of Philosophy in the Graduate School of the Ohio State University
By John W. Warner III
The Ohio State University 1972
Approved by
' \ Advisor Department of Astronomy
Advisor / Department of Astr Some pages may have i nd i st inet print. Filmed as received. University Microfilms, A Xerox Education Company ACKNOWLEDGMENTS Above all I would like to express my deep appreciation and thanks to Professors E. R. Capriotti and John D. Kraus for their guidance, fruitful discussions and constant optimism. I would also like to thank Drs. Vera C. Rubin and W. Kent Ford of the Carnegie Institution of Washington for their time, effort and discussion and Dr. Sandra Faber for her useful suggestions re garding this work. In addition helpful discussion was obtained from Drs. Gerald Newsom, Paul Byard, William Protheroe, George Collins, Robert Wing, Terry Roark, and fellow student E. R. Craine. Computer time was made available by the 1RCC of Ohio State University. Finally, the constant and loving help of my wife Ginnie is gratefully acknowledged. ii VITA October 17, 1944 ...... Born - Dubuque, Iowa 1966 ...... B.A., Kalamazoo, College 1967-1969 ...... Research Assistant, Department of Astronomy, OSU In stellar atmos pheres 1969-1970 ...... Teaching Assistant, Department of Astronomy, OSU. 1970-1971 ...... Research Associate, OSU Radio Observatory Hay 1971 ...... Attended Scuola dl Fislca Cosmlca, Erice, Sicily PUBLICATIONS "An Extraordinary Object in Sagittarius", Ap.J., 167, L53, 1971 (with R.F. Wing) . "Radio Sources in the Ohio Survey and Zwicky Compact Galaxies", Astrophys. Letters, 11, 83, 1972. "On the Nature of IRC-20385", Ap.J. (in press), Jan. 1, 1973 (with R.F. Wing and M.G. Smith). "Optical Variations of OJ287, ON231, and OQ208", Ap.J. (Letters) (in press), Jan. 15, 1973 (with E.R. Cralne). FIELDS OF STUDY Major Field: Astrophysics Studies in Radio Astronomy. Prof. J.D. Kraus, Prof. H.C. Ko Studies in Gaseous Nebulae. Prof. E.R. Caprlotti, Prof. S.J. Czyzak Studies in Stellar Atmospheres. Prof. G.W. Collins, II, Prof. E.R. Caprlotti Studies in Spectral Classification. Prof. P.O. Keenan in TABLE OF CONTENTS Page ACKNOWLEDGMENTS...... 11 VITA ...... Ill LIST OF T A B L E S ...... vi LIST OF F IGURE S...... * ...... vlll LIST OF P L A T E S ...... x INTRODUCTION ...... 1 Chapter I. THEORY AND METHODS ...... 3 Observable and Derivable Parameters Methods of Population Synthesis Emission Line Studies in the Nuclei of Galaxies II. OBSERVATIONS AND REDUCTION OF D A T A ...... 19 Choice of Galaxies Observations Working Data Line Identification Drawing the Continuum Equivalent Widths Emission Line Intensities Velocity Dispersions III. PHYSICAL CONDITIONS AS DERIVABLE FROM THE EMISSION LINES ...... 44 Electron Temperatures and Densities Abundances M51 MSI Conclusions Iv TABLE OF CONTENTS (Continued) Page IV. MODELS OF STELLAR POPULATION ...... 84 The Method Calculations Results and Discussion NGC 5195 M51 V. PROPERTIES OF GALACTIC N U C L E I ...... 100 Sunmary of Conclusions Summary of Known Data A Composite Model Comparison With Active Objects BIBLIOGRAPHY ...... Ill v LIST OF TABLES Table Page 1. Observed and Derived Quantities . * ...... 13 2. Spectrograph Parameters...... * * 20 3. Galaxy Plates U s e d ...... 23 4. 2.5 log for M51 and NGC 5195 ...... 34 5. Observed Galaxy Equivalent Widths ...... 36 6. M81 and M51 Line S t r e n g t h s ...... 41 7. Constants for Temperature Calculations ...... 47 8. The Ratio 1(6717)/l(6731) for M81 and M51 for Various Slot S i z e s ...... 50 9. Observed Wavelengths of Oil, Ha, and SII in M 8 1 ...... 53 10. Hydrogen Recombination Coefficients ...... 63 11. Constants in Equation 2 6 ...... 64 12. H81 (2V5) Abundances ...... 67 13. M51 Nitrogen Abundances...... 68 14. M51 Oxygen Abundances ...... 69 15. M51 Sulfur Abundances...... 70 16. Adopted Stellar Equivalent Widths ...... 87 17. Adopted Standard Stellar Continua ...... 88 18. Characteristics of NGC 5195 M o d e l s ...... 94 vi LIST OF TABLES (Continued) Table Page 19. Characteristics of M51 Models ...... 100 20. Summary of Properties of the Nuclei of Several Galaxies ...... 102 21. Properties of A Composite Galactic Nucleus ...... 10S vii LIST OF FIGURES Figure Page 1. Intensity Traces of NGC 5195 and a G2V Standard . . . 28 2. Inverse Averaged Response of the System ...... 32 3. 2.5 log F, for M51, NGC 5195 and Two Standard Spectral Types ...... 35 4. Energy Level Diagrams for Ions of Interest ... 45 5. The Intensity Ratio I(3726.2)/I(3728.9) vs Observed Wavelength...... * ...... 52 6. Log Te vs log Ne for M51(2 7 5 ) ...... 55 Log Te vs log Ne for M51 (6 7 0 ) ...... 56 8. Log Te vs log Ne for M81(2 7 5 ) ...... 57 9. Log Te vs log Ne for M81(6 7 0 ) ...... 58 10. M81 (275) Abundances...... 71 11. M51 Abundances...... 72 12. M51 (275) Abundances Corrected for Reddening .... 7b 13. M51 (275) Abundances Corrected for Temperature G r a d i e n t ...... 78 14. M81 (275) Abundances Corrected for Reddening .... 81 15. Schematic H-R Diagram ...... 89 16. Comparison of Observed and Calculated Continua for NGC 5 1 9 5 ...... 92 viii LIST OF FIGURES (Continued) Figure Page 17. Stellar Models for NGC 51 9 5 ...... 93 18. Stellar Models forM5 1 ...... 97 19. Comparison of Observed and Calculated Continua for M 5 1 ...... 98 ix LIST OF PLATES Plate Page 1« The Nucleus of M 3 1 ...... 4 2. Representative Spectra of Galactic Nuclei ...... 26 3. Equivalent Width - Spectral Relation for Standard Stars ...... * ...... 38 x INTRODUCTION Recent developments In astronomy In the last ten years have indicated that a scope of activity exists In the universe which was previously not thought possible. Radio, infrared, and X-ray emission associated with the nuclei of galaxies and with star-like objects suggest that comparatively small regions contain immense sources of energy which could play a major role in the evolution of galaxies and which may be connected to the very framework of the universe (Robinson, Schild and Schucking 1965, Burbidge 1970, O'Connell 1971). With several purposes in mind this thesis takes a moderately all-encompassing point of view with respect to the study of galactic nuclei. It studies a number of their properties using spectrographic techniques and then compiles the results of this and other investi gations in an attempt at producing a composite picture of a "normal" galactic nucleus. These broad alms are modest in scope, the purpose being to gain perspective on the problem. A preliminary summary of some conclusions from this work is given in Warner (1972). The initial investigation carried out in this thesis consists of the following parts: 1) Visual and red image-tube spectrograms were taken of the inner nuclear regions of M51, M81 and NGC 5195. A skylight suppres sor was used to isolate areas surrounding the nucleus to obtain averaged spectral information as a function of distance from the 2 center. 2) Emission line intensities, relative continua, and equi valent widths for the major absorption lines are calculated (Chapter 2). 3) The physical conditions in M81 and M51 are studied in an attempt to understand variations in line intensities with distance from the center (Chapter 3). 4) Population models and synthetic continua are calculated for M51 and NGC 5195 based on measured continua and equivalent widths of standard spectral types (Chapter 4). Chapter 5 then presents a compilation of data for seven normal galaxies and one Seyfert galaxy: M31, M32, M51, M81, NGC 1052, NGC 5195, the Galactic Center, and NGC 1068. A simple composite model is presented along with a comparison of its properties to those of Seyfert nuclei and quasi-stellar sources. CHAPTER 1 THEORY AND METHODS Observable and Derivable Parameters A galactic nucleus (see Plate 1) is presumably a complex aggre gate of stars, gas, dust, radiation and, possibly, very compact forms of matter. Each of these components produces radiation and inter acts in some way with the radiation produced by each of the other components. Stars. Stars produce continuum radiation with approximately a Planckian distribution. The stellar component of radiation from the nucleus is usually strongest in the visual and near infrared and consists of a superposition of continua of various spectral types. Light in the ultraviolet region is usually produced from the early types, while middle types contribute to the blue through red regions and late types to the infrared. The absorption lines seen in galaxies are also composite. They are formed from averages fcf the lines characteristic of each of the stellar components, weighted according to the fraction of light contributed by the spectral type at the given wavelength. The observed continuum and the equivalent widths of the lines thus contain information about the number of stars of each spectral 3 4 CAPTION FOR PLATE Plate 1. 200-inch photograph by Baade illustrating the nucleus of H31 (scale 1.5 arc sec/mm). The cloudy absorption structure is real, and it appears that much obscuration is apparent in the nucleus. (Photograph courtesy of W. C. Miller, Hale Observatories). 6 type and luminosity and abundance class which make up the luminous stellar mass of the galaxy. Sufficiently accurate measurements of standard spectral groups can be combined into models of the stellar population (de Vaucouleurs and de Vaucouleurs 1959, Wood 1966, McClure and van den Bergh 1968, Spinrad and Taylor 1971, Andrlllat et al. 1971) . Difficulties exist which are Inherent to this method. Intrinsically faint stars contribute little to the luminosity at any wavelength and are not easily detected. Errors in observation, in the calibration of line indices and in the model-making process produce fairly large errors in the resulting distribution. Also, in most cases only the nuclear regions are bright enough for accur ate measurement, so that the distribution with radial distance from the center is not normally observed (see however Spinrad et al. 1971 and Welch and Forrester 1972). Another method for roughly determining the distribution of populations is that of composite photography (Zwicky 1955, 1957 and Sharpless and Franz 1963). The distribution of light in a galaxy in both blue and yellow colors is recorded on separate photographic plates. Superposition of various combinations of the negative and positive plates during printing reveals the distributions of the stars in large regions of the H-R diagram designated as "blue" and "red" populations. Absorption features also can reveal information about the dynamics and masses in nuclei. The line shapes depend upon the intrinsic shapes of the lines of the component stars and their 7 velocity dispersions. These dispersions have been derived In several instances (Minkowski 1962) and have been used to obtain the mass and study the dynamics of the nucleus. The redshlfts of the lines as a function of angular distance from the center also give a measure of the nuclear mass (Walker 1962). Gas. Both ionized and neutral gas exist in galaxies. Neutral gas (primarily hydrogen) produces no continuous radiation. Twenty- one centimeter line radiation from neutral hydrogen is observed in both our own galaxy and in external galaxies and serves to define the distribution of nonionlzed gas (Roberts 1971). Studies have shown its extent to be much larger than the distribution of bright stars and spiral arms. Elements besides hydrogen which exist within neutral regions produce absorption features such as Call X3933, X3968 and the G band of CH. They are observed in our own galaxy and have possibly been detected in M31 (Rubin 1971). Molecular features which must necessarily be associated with neutral regions have been detected using radio astronomical methods in the nucleus of our galaxy and in MSI (Rougoor and Oort 1960, Robinson and McGee 1967). The widths, strengths and redshlfts of these absorption features can provide information about the abundance of the elements and the dynamics of the interstellar medium. The ionized part of the gas component has not been observed to contribute large amounts to the continuum of a galaxy. Its strength depends on the source of ionization and the electron density and is thus coupled to the other components of the nucleus. Emission lines, most notably the forbidden lines of 01, Oil, 0111, Nil, SII and the hydrogen recombination lines, contain information about the physical state of the ionized gas component. The strengths of the lines de pend upon the type of ionization mechanism, the method of excitation, as well as the density of the gas and the abundance of the elements. Electron temperatures, electron densities and abundances can be de rived as can limits to the source of ionization of the gas. The mass of the ionized gas and the fractional volume of the nucleus filled with gas can also be derived (Spinrad and Peimbert 1970, Alloin 1971). Velocity dispersions of the emission lines give an estimate of the kinetic energy of the gas (Minkowski and Osterbrock 1959). Displaced components to the lines indicate large scale mass motions and possible loss of gas from the nucleus. The radial velocities of the lines as a function of position angle and angular distance from the nucleus give the velocity field of the gas and a mass for the galaxy (Burbidge et al. 1962). The shapes and the velocity widths of the lines reveal information about turbulence or about scatter ing mechanisms (Pacholczyk and Weymann 1968). The average redshift of the galaxy as a whole can be used to study the Interactions of the galaxy with companions in a cluster. The distribution of ionized gas in a galaxy can sometimes be studied with narrow band photography in the light of hydrogen alpha or one of the forbidden lines. Such studies can detect the pre sence of strong nuclear activity as has been done in M82 (Lynds and Sandage 1963). Dust■ Presumably contiguous with the neutral gas component is the dust component. Through the process of scattering it can pro duce a redistribution in direction of the radiation field, and through absorption, a redistribution in frequency. Reradiation of the absorbed fraction of total luminosity in the far infrared is then quite likely, although the excesses found here are not necessar ily all due to this process (Jones and Kellogg 1972). The dust component could conceivably provide information con cerning dynamics of the nucleus. Correlations between positions of highly obscured areas and of velocity residuals have been found in M31 (Rubin and Ford 1971). Circumstellar dust shells have been found to play an important role in the infrared radiation from super-giant stars, associations, nebulous objects such as Eta Carina our galactic nucleus, the nuclei of Seyfert galaxies, and possibly planetary nebulae. It is apparently an Integral part of the ener getics of galaxies. Radiation. The existence of sources which emit energy in many regions of the electromagnetic spectrum and also the shapes of their energy curves indicate that a variety of physical processes may be at work in galactic nuclei. This area of research is certainly a most speculative one since adequate theory does not exist either to determine uniquely the ultimate source of energy in nuclei or the exact mechanisms by which this is transferred into the observable spectrum. The only exception to this is for nonactive nuclei in the visual region where the bulk of the radiation is due to stars. 10 One well-known theory Is that of the synchrotron process. Observations of the radio continuum show complex nuclear activity (Burbidge et al. 1963) and indicate that large particle energies and magnetic fields are needed to explain the observed flux (Burbidge 1956, Burbidge and Burbidge 1965, Kellerman 1971). The basic data consist of the flux density of radiation as a function of frequency. The shape of the energy curve along with a peak or trough may be used to calculate physical parameters when the spectrum of particle energies is provided (Kellerman 1971). Obser vations of variability and polarization provide additional con straints for the problem, namely, concerning the number of active regions, homogeneity of the magnetic fields and limits to the electron density (Aller 1970). Similarly, angular size and a peak in the radio spectrum have been used to put limits on the physical mechanisms and the sizes of the emitting regions in volved (Wills 1971). Observations of the far infrared continuum, variability and angular size have also enabled limits to be put on the emitting volumes and on the type of mechanism responsible for this radiation (Jones and Kellogg 1972, Bisnovatyi-Kogan and Syunyeav 1972). Compact Mass Component. Gravitational wave observations (Weber 1969) as well as more general arguments about the evolution of matter (Hoyle 1969) sussest the existence of highly compact matter. Such matter is also postulated since nuclear activity seen in galaxies could be related to large gravitational fields and, 11 hence, to large masses in small volumes. Since recent developments in work on the later evolutionary stages of stars (Wheeler 1966) and other recent observations (Pulsating Stars 1968) show that very compact matter probably exists (see also Zwicky 1967), the assumption that it exists in the nucleus where star formation was most likely abundant at one time is not difficult to make. The real difficulty is in deciding how such matter might be observed. Collapsed sources have been proposed as the origin of the energy seen in galactic nuclei and quasi-stellar sources (Burbidge 1962, Hoyle and Fowler 1963, Hoyle, Fowler, Burbidge and Burbidge 1964). Several more exact theories have been proposed to explain the nonthermal radio and infrared emission by gravitational collapse (Lynden-Bell 1969, 1971) or through matter accretion onto neutron stars (Bisnovatyi-Kogan and Syunyaev 1972). An explanation in volving supermassive rotators has also been proposed (Woltjer 1970). All theories are as yet highly speculative. The adjustable para meters in them are many, the assumptions overwhelming and the observations limited to the basic ones described above. These theories will not be discussed in detail in this thesis, as they deal with no observables directly pertinent to this part of the problem. A relatively complete summary is given in "Nuclei of Galaxies" (O'Connell 1970). Clearly a large number of observed quantities can be Isolated from measurements of the radiatici from nuclei and many more can be derived when these are combined with a suitable theory. As the accuracy of most of the observables depends upon many things (mostly 12 dispersion, seeing or resolution and the available energy flux) and theory is still quite poor, derivable results range from obviously incorrect to merely good approximations. Table 1 lists a number of the observed and also the derived quantities related to them. Methods of Population Synthesis There are a number of reasons for attempting to determine the stellar content of galactic nuclei. The age of a galaxy is a function of the main sequence turn-off point of its H-R diagram. Physical conditions in the early stages of star formation are implied by the relative numbers of stars. Dynamical calculations in the nuclear regions use the relative number of stars of various masses. The radiation field depends upon the number of hot stars. Evolutionary corrections in cosmology are determined by comparing present and past luminosity functions. The only hope of obtaining any of this information for the majority of galaxies lies in studying the integrated light from them. A basic technique is population synthesis wherein quantitative measures of galaxy populations are matched with combinations of simi lar measures from standard spectral types. The error in the model fitting process as well as the wavelength baseline over which the measurements are made determines the accuracy of the synthesizing process. Whipple (1935) first outlined quantitatively the methods for combining galaxy colors and absorption line equivalent widths (along 13 Table 1 OBSERVED AND DERIVED QUANTITIES Derived Observed $(M, Sp, Radius) <4- Continuum Stellar Abundances ^ Absorption line strengths Nuclear mass 4 Absorption line widths Dynamics of Stars\ Absorption line z*s Dynamics of neutral gas 21 cm emission Abundances in gas Interstellar absorption Source of ionization ^ Emission line strengths Source of excitation Emission line widths Physical conditions Emission line z's Radio flux, spectrum Radio flux, angular size Radio flux, polarization Radio flux, variability Infrared flux, spectrum Infrared flux, angular size Infrared flux, variability 14 with similar measures for standard spectral types) to build models of galaxy populations. All recent work is based on these principles. Humason (1936) and Humason, Mayall and Sandage (1956) were the first observers to give spectral types and correlate the appear ance of the spectra of galaxies with their morphological types. The first researchers to recognize the idea that the composite blue spectral types in a galactic nucleus were related to the precise mixture of stars and the peak periods of star formation were Morgan and Mayall (1957), Morgan (1959), and Morgan and Osterbrock (1969). These observers worked out in detail the spectral type, galaxy morphological type relation first found by Humason, Mayall and Sandage. Recent quantitative approaches have been introduced by Wood (1966), Spinrad (1966), McClure and van den Bergh (1968), and Splnrad and Taylor (1971). These authors set up photoelectric narrow-band color and line index measurement systems sensitive in varying degrees to stellar temperature, luminosity and abundance differences. These systems were then used to construct models of the stellar content of galaxies, hoping that their photoelectric precision would offset the lack of spectral resolution. An approach to this problem used by other workers (de Vaucouleurs 1959; Moore 1968; Alloin, Andrillat, and Souffrin 1971) is to return to Whipple's original idea by combining both equivalent widths and continuum colors for producing models. Difficulties with extreme blending of lines and lack of adequate continuum points 15 makes this method less precise, yet rough quantitative measures of the nuclear populations have been obtained. The method followed by Alloln et al. (1971) is the one which will be used for population studies In this thesis and will be des cribed in detail in Chapter 4. Most of the inherent difficulties cannot be overcome, but the use of an image tube spectrograph reduces exposure times and allows spectra to be taken of fainter objects. Many more spectra can be taken over several different spectral regions with relative ease, since the plate sensitivity is not the limiting factor. Isolation of the nuclear regions themselves through the use of the sky suppressor has not been previously attempted, nor would it have been possible before the advent of the image tube. Most previous work using equivalent widths widened the whole nuclear region along the slit or attempted to isolate the inner nucleus later with the densitometer slit. The former method loses spatial re solution and results in blending of lines, while the latter method does photographic photometry on plates with large variations in density, sometimes even burned out regions. The method used in this thesis ameliorates the first problem and eliminates the second. Emission Line Studies in the Nuclei of Galaxies Burbidge's summary (1970) deals with much of the work on the nuclei of galaxies and a chapter in the unpublished volume 9 of Stars and Stellar Systems by Spinrad and Peimbert deals specifically with the emission line work up to 1970. Since the latter is not in print and widely available and since this history is quite important 16 to the following chapters of the thesis, the major points will be Bunmar i zed he re. Very little quantitative work has been done on the emission regions in the nuclei of normal galaxies, a fact which Is mostly due to the long exposures needed. One of the first qualitative studies was due to Mayall (1958), who investigated the occurrence of X3727 in morphological types. This study did not very accurately reflect the percentage of elliptical galaxies with X3727 as was shown by Oster- brock (I960). This line was also studied in detail for several galaxies, namely, NGC 1052 (Minkowski and Osterbrock 1959), M81 (Munch 1959) and M51 (Boesgaard 1967). Values of the electron -3 densities derived were < 300 cm for NGC 1052, ^ 1000 for M81, and ^ 500 for M51. Masses derived from the X3726/3729 ratios and Ho intensities showed that HII region masses are much smaller than the stellar masses in the nucleus. One o* he most important and controversial results was by Burbidge and Burbidge (1962, 1965) where it was found that the ratio of forbidden nitrogen to Hot intensities in some nuclei were different from normal HII regions. On spectra taken for rotation curves it was seen that this ratio changed from less than one to more than one as the center of the nucleus was approached. Quantitative foundation was given to this conclusion by emission line intensities measured by Peimbert (1968), and Peimbert and Spinrad (1970a, 1970b), who showed that in M51 and M81 this ratio (averaged over most of the 17 nuclear region) ranges from 3 to 10 times the values in normal Hll regions. In several galaxies the [01] and [SII] lines have also been seen to be stronger than in galactic HII regions and planetary nebulae. Burbidge and Burbidge (1962) report the presence of emis sion lines in NGC 4258 in the order [Oil], [Oil], [SII], [Nil],Ha . For M81 and M51 the strengths of the [SII] lines are about an order of magnitude larger than in galactic HII regions of similar ioni zation (Peimbert 1968). Since forbidden line ratios are a function of ionization, excitation and ionic abundance, several proposals have been made to explain this phenomenon. Burbidge and Burbidge (1962) and Burbidge, Gould, and Pottasch (1963) suggest collisional ionization as well as high electron temperatures to explain the [Nil]/Ha ratios. Burbidge and Burbidge (1965) emphasized the possibility that abundances might be different in nuclear and galactic HII regions. Morgan and Osterbrock (1969) suggested differences in ionization, while Peimbert (1968) showed that collisions by thermal electrons are not respon sible for the ionization in M81 and M51. He also suggested radiative ionization and nonsolar abundances to explain the line ratios. Recently Searle (1971) showed that HII regions in galaxies probably showed composition gradients across their disks. In all cases, however, the values of [Nil]/Ha were less than unity and the author did not survey very far into the nuclear region. A recent paper by Alloin (1971) suggests that the nucleus of M81 is stratified 18 into different ionization regions and that the ratio of nitrogen to oxygen is approximately ten times normal. Only Peimbert (1968) made any attempt to get emission line strengths within a small region of the nucleus. In M51, [NII]/Ha line ratios measured using two diaphragm sizes (7" and 20" diameter) showed significant differences indicating a rather steep gradient within the nucleus. It is not very clear from spectra taken with the nucleus fixed on the slit just how this intensity ratio changes within the normally burned out part of the nuclear region. Seeing eventually limits the spatial resolution, and most spectrograms of normal galaxies are not easy to interpret on short exposures since the weak lines just barely show above the continuum. The method used in this thesis is designed to partially re lieve this problem. Using several diaphragm sizes (on the smallest of which the resolution is essentially set by the seeing) one can compare emission line strengths from increasingly larger volumes of the nuclear region. A strong gradient In the intensity ratio of [Nil]/Ha shows up well and density criteria become available in the form of the line ratio jj|sil A6731) * also shows that there exist changes in the oxygen line (Oil, OIII) intensities. 19 CHAPTER 2 OBSERVATIONS AND REDUCTION OF DATA Choice of Galaxies The galaxies M51, M81 and NGC 5195 each have a bright well- defined nucleus and hence were obvious choices for nuclei which could be conveniently observed. In addition, both M51 and M81 pro vide sufficient emission line data to enable physical conditions to be studied. Neither M51 nor NGC 5195 have population models avail able for them, so these galaxies were chosen on this basis. A number of other galaxies fill some but not all of the above require ments, but their study will be left until a later time. Observations Observations were made with the Department of Terrestrial Magnetism (Carnegie Institution of Washington) image tube spectrograph attached to the 72" reflector of the Ohio State and Ohio Wesleyan Universities at Lowell Observatory. The Perkin-Elmer transfer lens and 5" camera were used with the spectrograph. Table 2 lists grat ings, settings, number of lines/mm, dispersions and spectral regions used for the work in this thesis. Table 2 SPECTROGRAPH PARAMETERS Grating Dial Setting Order Slit Width Dispersion Range (lines /tan) GO (X/nnn) 600 15.5 1st 300 110 U3300-6000 600 15.2 1st 300 110 4600-7000 1200 11.4 1st 300 44 5700-7000 1200 10.5 2nd 300 22 3500-4000 ro o 21 All objects were trailed using a skylight suppressor. This device consists of two slots 5 mm apart in a piece of metal attached to a threaded mount. The mount travels on a screw driven by a small motor driving at a rate of approximately 1 mm/minute. To widen spectra the telescope tracking rate was adjusted so that the desired region of the galaxy was kept in the upper (eastern) aperture de fined by the spectrograph slit and the crossed slot. The second slot, off the nucleus, normally gave information on the amount of night-sky emission on the plate. The guiding error using this equipment was estimated to be +_ 1". Seeing also spread out the size of the area actually being sampled. Average seeing over the four runs was about 2"-3M , so that a slot size corresponding to 2M on the sky was actually sampling light from 6" to 8M of the galactic nucleus. Since seeing is approximately a Gaussian distribution the majority of the light would have been from a smaller area than this. Two slot sizes were used, an 0.4 mm slot corresponding to an area of 275 x 178 = 5 square arc seconds and an 0.8 mm slot corresponding to an area of 6" x 178 “ 11 square arc seconds on the sky. Several advantages were obtained by using the skylight suppres sor. It enabled widened spectra of Isolated regions in the galaxy comparable to the seeing disk to be produced, and contamination by night sky and moonlight was suppressed to a large extent. The ratio of such contamination on a plate using this device to one not using it is Ax (mm)/5mm where Ax is the slot size in nm. These values ranged from 1/15 to 1/5 for the slots used. 22 It must be emphasized that for the absorption line work done in this thesis, no spectrograms were used which were exposed while the moon was above the horizon. Plates of comparison (Hayes 1970) stars and standard spectral types were taken using the same equipment, photographic plates and slit sizes. They were normally trailed 12 mm at the slit. Comparison stars were normally taken with each set of galaxy plates developed together. Various difficulties prevented this from being done always. For observations without a standard star mean sensitivity curves were drawn up and used. Such observations were given lower weight. Calibration plates giving intensity versus plate density were made for each set of galaxy and standard star plates developed together. These were made with a calibrated filter Inserted into the path between plate and a light source close ap proximating the color of the phosphur on the image tube. All plates were IIa-0 and were baked from AS to 6A hours at 50® centigrade within ten days before their use. Exposed plates were developed in D19 for 5 minutes at 68® Farenheight, emersed in stop bath for 30 seconds, then in hypo (F-5) for 5 minutes. A one minute water wash, hypo-clearing agent for 5-10 minutes and a final wash and cleaning in distilled water completed the development pro cess. Table 3 lists the galaxy, plate number, spectral region, dispersion , and slot size of all galaxy plates used in this thesis (B - XA3300-6000, B-R = AAA600-7000, R - AX5700-7000) . 23 Table 3 GALAXY PLATES USED GALAXY PLATE NUMBER SPECTRAL REGION DISPERSION SLOT SIZE (A /mm) (arc secs) M 51 24a B 110 2.4 34 B 110 2.4 48a R 44 2.4 137c B-R 110 2.0 137d B-R 110 no slot 138b B-R 110 6.0 141b B-R 110 2.0 M 81 3 B 110 2.4 23a B 110 2.4 23b B 110 2.4 27a B 110 no slot 41 R 44 2.4 55 R 44 2.4 47b R 44 2.4 140d B-R 110 6.0 140e B-R 110 no slot 143d B-R 110 no slot 144b B-R 110 6.0 144c B-R 110 2.0 NGC 5195 24b B 110 2.4 33a B 110 2.4 33b B 110 2.4 45a R 110 2.4 139a B-R 110 6.0 141a B-R 110 6.0 24 Plate 2 illustrates "blue" and "red" spectra of the nuclei of M51, NGC 5195, M81, and M64, taken with the skylight suppressor. Working Data Transmission traces were made of all galaxies used for absorption line work at a scale of 110 X/inch with the Grant comparator-densitometer using a slit corresponding to approximately three times the resolution of the plate. The purpose of this was to suppress noise and to get a good idea of the basic real absorption features in the spectra. The calibration plates were then traced in transmission mode and H and D curves were drawn. The approximately linear portion was determined, and the curves were redrawn on a scale suitable for use with the curve-follower attached to the Grant equipment. All plates, galaxies, star iar*’ stars, and standard spectral types were then traced in intensity mode using the Grant equipment. Slit widths were determined by the resolution of the plate and the noise in the spectrum. In practice a slit width of 34u was used for an effective resolution of 55m on the 100 A/mm plates. This re solution was determined from widths at 1/2 Intensity of comparison lines. Slit heights were determined from the width of the spectrum, wide enough that slight errors in lining up the spectrum would not introduce clear plate into the slit. Galaxy plates were traced at two scales, 55 A/inch for measuring equivalent widths and 110 A/lnch for determining the continuum level. Fogged plate level at blue and 25 CAPTION FOR PLATE Plate 2. "Blue" and "red" spectra of the nuclei of M51, NGC 5195, M81 and M64 taken with the skylight suppressor. 3500 4000 4500 5000 5500 — I— T — I M51 NGC 5195 M 81 M 64 K H G band 5000 5500 6000 6500 r T ------j M 51 NGC 5195 M 01 M 6 4 [Nil] [SII] ro 27 red ends of Che plates were also traced. Fogged plate level at a given point on the spectrum was found by assuming linearity along the plate and interpolating between these two points. Line Identification Several standard spectral type traces were used as a pre liminary guide for line identification. These consisted of traces of an F, G, and a K star whose lines had been identified from the Bonner Atlas (Seiter 1970). The galaxy traces were then compared to these master traces to identify strong lines and to gain a feel for the complexities of composite galaxy spectra. Figure 1 shows two intensity traces of NGC 5195 from spectra taken on the same plate and an intensity trace of a G2V standard. The main things to notice are (1) the strongest spectral features reproduce quite well in cluding the shapes of the lines and their blends, (2) occasional spikes of noise due to ion spots from the photocathode (and some times from the upper nonlinear part of the H and D curve) hinder the determination of "near continuum" levels, (3) weak lines vary in shape, strength and presence at all due to superposed noise caused by ion spots, grain, and plate imperfections, (4) strong as well as weak features in the standard star can be seen in the galaxy spectrum. On an average spectrum over 100 features were definite pos sibilities for real lines. The identification of all of them for each plate would have been laborious and would not have contributed measurably to the goals of this thesis. Thus only the strongest BS7569 (GZV) Figure 1. Intensity traces of NGC 5195 and a G2V standard. 29 lines were measured and their wavelengths obtained. The measured wavelengths were then used to isolate the features to be used In Chapters 3 and 4. Drawing the Continuum Since galaxy spectra are composites of various spectral types with a large velocity dispersion, it is safe to say that no point in the blue to visual regions truly represents the continuum. Still, significant information on the luminosity as a function of wavelength can be obtained by drawing a pseudo-continuum through regions with few lines and at points of near continuum on the intensity traces. The spectra of several standard spectral types were checked for such areas. Galaxy spectra were intercompared to visually isolate noisy areas and ion spikes which would give grossly inaccurate levels. These traces were also compared with transmission traces made with a large slit to suppress noise. Pseudo-continuum levels were then drawn in as smooth curves on traces of small scale. These curves were measured from fogged plate (due both to normal fogging of the plate and to the background noise of the image tube), and all numbers were normalized to the total area under the curve between X4000 and X5500 for the "blue*' region and X5007 and X6S83 for the "red" region. Usable comparison star data was not available for all of the plates taken. Such individual data is necessary to account for the nightly variations in the response of the systems Involved. These 30 factors are mainly due to atmospheric extinction and to the response of the telescope-spectrograph system. The true Intensity distribution from an object is modified by the combined response function R(X) into blackening on the plate. Assuming a perfect translation of density into intensity using the * H & D curve, R(X)F^ is simply the intensity deflection as a function of wavelength as measured on the galaxy trace (where F* is the in- 2 tensity of the galaxy outside the atmosphere in ergs/cm /sec/AA/ste- radian). Normalization by dividing by the area under the intensity curve with wavelength limits X^ and A^ yields the relative energy distribution of the galaxy as seen by the system, I^(gal). Using observations of standard stars with well-known energy distributions (Hayesl970), the relative intensity outside the atmos phere, F^, can be recovered: I.(gal) -kX sec z X _ FX<*al> ■ I."(star) V scar) e-kX sec s star ( ~ > {1> A G where F,(gal) - relative intensity outside the atmosphere I*(gal) = observed intensity distribution of the galaxy 1,(star)- observed intensity distribution of the standard star Fi(star)- known relative energy distribution of standard star - standard wavelength (Hayes (1970) - atmospheric absorption coefficient - zenith angle. 31 When a comparison star for a given galaxy spectrum is not available, then an averaged response curve must be used to derive the relative flux. In equation (1) the factor F. (star) X - k.sec z star c! ------c -r > e <2> 1^(star) is the correction for one comparison star. By setting sec *» 1.0 and I(gal)~1000, these factors for several stars can be computed and averaged: X „ . , gal star. R"1 l°3 F. (star) ( )2 e-k^ sec 1 -“ c * > . X I (star) X 103 e_k C (3) then, The values of Continua for the galaxies M51 and NGC 5195 were calculated from the blue and the red plates. Regions of overlap which were particularly sensitive to nightly variations or to focus were thrown out. Blue and red curves were matched in the overlapping regions by 32 34 36 38 40 42 4 4 46 48 50 52 54 56 58 Wavelength (AxICf2) Figure 2. Inverse Averaged Response of the System. 33 visually comparing curves. It was assumed that this process in creased the error bars of the red region by + 101. Errors are then estimated to be + 10X for the region X < 5000 X and + 20X for X > 5000 Table A and Figure 3 give the calculated average con- tinua for these galaxies normalized to 1000 at 5360 £ and corrected for galactic absorption with a Whitford law and a constant C(H£) - 0.08. The error bars take into account errors due to intensity calculations, continuum estimation, and averaged correction factors. Also in Figure 3 are normalized continua of an AOV star and an MOV star for comparison. Equivalent Widths Determining the continuum essentially determines the values of the equivalent widths, . For the lines of interest (see Chapter A) these were measured with a planimeter. In each case sufficient revolutions were made to reduce the measurement error to less than 5%. The major source of error is the le^el of the con tinuum. In each case the reasonable limits to this level were estimated, and this spread determined a systematic percentage error in W^. The shapes of the lines, extrapolation to continuum levels for depressed regions, and noise also affected the values. An estimate of this error can be found by averaging the values from several traces. Table 5 lists equivalent widths and their estimated errors for M51 and NGC 5195. 34 Table 4 2.5 LOG F” 1 FOR H51 AND NGC 5195 A x ( X ) A-1 M51 NGC 5195 3700 2.70 7.18 6.49 3800 2.63 7.18 6.66 3900 2.56 7.30 6.86 4000 2.50 7.37 6.93 4100 2.44 7.42 6.99 4200 2.38 7.44 7.02 4300 2.32 7.48 7.08 4400 2.27 7.58 7.13 4500 2.22 7.71 7.23 4600 2.17 7.70 7.30 4700 2.13 7.68 7.32 4800 2.08 7.68 7.36 4900 2.04 7.70 7.37 5000 2.00 7.58 7.39 5100 1.96 7.54 7.42 5200 1.92 7.50 7.44 5300 1.89 7.50 7.48 5400 1.85 7.50 7.52 5500 1.82 7.51 7.55 5600 1.79 7.49 7.58 5700 1.75 7.50 7.65 5800 1.72 7.48 7.65 5900 1.69 7.48 7.68 6000 1.67 7.47 7.68 6100 1.64 7.47 7.69 6200 1.61 7.43 7.68 6300 1.59 7.43 7.72 6400 1.56 7.46 7.72 6500 1.54 7.49 7.74 6600 1.52 7.49 7.74 6700 1.49 7.48 7.79 iue . . lg o 5, G 59 n w tnad Stars. Standard Two and 5195 NGC M51, for log 2.5 3. Figure - 16 - 20 22 4 26 2- 0 3 -8 2 6 2 -4 2 2-2 0 2 1-8 1-6 1-4 25 Log F\,+C (relative units) M 51 M 1) V x G 5195 NGC O star MOV O ar ta s AOV Table 5 OBSERVED GALAXY EQUIVALENT WIDTHS M51 NGC 5195 Wavelength Identification Obs’d Corrected AW Obs'd Corrected AW 3933 Call (K) 11.6 10.4 1.0 16.4 14.8 1.4 3968+ Call (H), H 14.0 11.9 1.2 17.6 15.0 1.5 4101 H 7.8 7.0 0.7 12.0 10.8 1.1 4226 Cal 2.3 2.1 0.4 2.5 2.2 0.4 4300 G band (CH) 6.8 6.8 2.0 7,4 7.4 2.2 4340 H 5.0 4.0 0.4 9.9 8.0 0.8 5184 Mgl 2.8 2.8 0.6 2.5 2.5 0.5 5890 Nal"D" 7.8 7.8 1.6 9.0 9.0 1.8 w O' 37 Equivalent widths of standard spectral types are also needed for the population synthesis process. A number of these are avail able in the literature (de Vaucouleurs 1959, Spinrad 1962, Rense and Kynek 1937), the best compilation being due to Alloin et al. (1971). It was decided that because of the inherent crudeness in the model building process, the curves of versus spectral types of these authors would be adopted after checking for consistency with my method of determining the equivalent widths. Accordingly, the W^'s were measured for a number of standard spectral types. Plate 3 shows the results of this along with the curves of Alloin et al. (1971) where the spread in is a measure of the scatter in the original data. It is seen that a systematic shift of from 10% to 20% is found for the various lines (10% for Call 13933 and Cal 14226, 15% for Call 13968 + He, 20% for HS14101 and Hy 14340) no doubt due to systematic differences in drawing the continuum. The equivalent widths and continua of the galaxies were corrected by appropriate amounts to out them on the same system as that of Alloin. For the G band the scatter is seen to be somewhat large and not systematic. For this reason no correction was made, but the error bars of this feature were Increased by ±10% in the galaxy measure ments. Similar corrections were made to the error bars of the Mgl feature and the sodium "D" lines since no measures of these were available. Little systematic change is expected for the "D" lines, since the continuum is so well defined. I ■ i 4 M^ » I . CaII K | 18 .. Call H*Ht Hi»Hr IB 16 J 16 14 14 12 12 W*X> 1°V^ e 8 6 * 6 A I 4 2 t 2 i I Ca I M 226 .. G band 18 Mgl dwarfs # 4 - No 10 dwarfs* 18 giants o glcnts o 16 ♦ 16 14 14 12 12 * v ° I 10 e t 8 6 6 A * .. . ■■ 4 2 ♦a *" I* 2 i t f ■*? OBAFGKMOBA F G KMOBA FGKMOBA FG KM SPECTRAL TYPE Plate 3. Equivalent-Width - Spectral Type Relationship for Standard Stars. Emission Line Intensities The process of reducing line intensities is the same as that of the continuum described in the previous section "Drawing the Continuum". The averaged sensitivity curve of Figure 2 were used for plates with poor or no standard star-plates, while the normal reduction process was estimated to be 101 for strong lines and ^ 25% for weak lines. Relative line intensities were normalized to Ha - 1000. The X3727 line was ratioed to the X5007 line which overlapped both blue and red 110 X/mm plates. This process in creases the error bars of X3727 to + 35%. Reddening corrections were made using the relation lo* hnt. - loS ^bs + C(V f(X) (7> with C(Ha) * 0.08 and the Whitford reddening function (Seaton 1960). p This is the value used by Peimbert (1968) and Alio in (1971) and was adopted here for consistency. This correction takes account of the average galactic absorption only since reddening internal to the galaxy cannot be computed. This latter must be considered the largest potential source of error in the line measurements. Several plates were taken with Ha in first order and X3727 in second order on the same plate. Similar plates taken of standard stars showed no large changes in continuum in spite of contamination by overlapping orders. Hence, the ratios of the Nil lines X6548 and 40 A6584 to Ha were used from these plates unreduced. Similarly the ratio of the lines of SII A6717 and A6731 were used unreduced. Such ratios are probably known to better than 10Z. Table 6 lists relative line Intensities for different slot sizes for M81 and M51 along with measurements from other authors. Reasonable agreement exists between my data and the photoelectric data of Peimbert (1968). Measurements by Alloin (1971) were made with the nucleus fixed on the photographic slit, a method which is subject to uneven exposure and poor reduction of plate densities to intensities. Such data should be given lower weight. The A6300 line of {01) in all my measurements is primarily due to night sky as is shown by a check of its wavelength. The velocity shifts are small and the resolution not sufficient to separate them. Velocity Dispersions Velocity widths of both absorption and anission lines are desirable for studies of the energetics of nuclei. Observed emission line widths are affected by the instrumental resolution (given by the comparison line widths) and the velocity dispersion along the line of sight in the nucleus. Absorption line widths are affected by the instrumental resolution, stellar velocity dispersion and in trinsic widths of the lines of the stars in the nucleus. It was determined that the resolution of the image tube spectro graph at the dispersions used was not good enough to give useful Table 6 M 81 H h id 2'.'5 6 " 7"(P) 7"(A) No Slot Orion III 3727 [Oil] 360 - 690 1700 - 760 4363 [OIII] ? - - 50 - 5007 [OlIIJ 240 - 450 500 - 467 6300 [01] -- 235 700 - 4.5 6563 Ha 1000 1000 1000 1000 1000 1000 6584 [Nil] 2300 1740 1400 2900 3750 300 6717 + 31 [SII] 260 800 460 - 1260 300 TABLE 6 (Concluded) M 51 x(X) id 2V5 6" 7"(P) 7"-20"(P) 20"(P) No Slot Off .30" 3727 [Oil] 380 - 980 645 710 - 4363 [OIII] 40 ------ 5007 [OIIIJ 650 430 1100 107 350 - - 6300 [01] - - 170 245 245 - - 6563 Ha 1000 1000 1000 1000 1000 1000 1000 6584 [Nil] 4100 2600 2500 930 1350 1320 400 6717 + 31 [SII] 440 320 776 — ** 540 * Caption: Table 6. M81 and M51 Emission Line Strengths for Different Aperture Sizes Including Values from Peimbert 1968 (P) and Alloin 1971 (A). 2V5, 6uf 7" and 20" denote increasingly larger apertures. 7n-20" denotes intensities in a ring defined by two circular apertures. "Off 30" denotes intensities in a 2V5 aperture 30" from the nucleus. "No Slot" denotes widening without a slot. A3 Information concerning the stellar velocity dispersions* It is clear that much better means are now available to obtain these numbers (Morton 1972). With sufficient resolution, intrinsic emission line widths may be separated out from the effects of the instrumental width. Com parison lines measured on several red plates taken at AA A/mm give an Instrumental width of approximately 1,5 A * 60 km/sec at A6563. A check of the instrumental profile showed that it approximated a Gaussian to a high degree. Assuming the intrinsic gas velocity dis tribution is Gaussian, the observed, instrumental, and intrinsic gas velocity widths are related by 2 2 v v , gas obsd (8) The results of these calculations for M51 show an average intrinsic dispersion of 123 + A km/sec and for M81, 155 + 3A km/sec, where the error denotes the dispersion around the mean. These numbers are not very accurate since the intrinsic width is approximately the same as the Instrumental width. They perhaps should be considered upper limits CHAPTER 3 PHYSICAL CONDITIONS AS DERIVABLE FROM THE EMISSION LINES Electron Temperatures and Densities Forbidden line ratios can be used to calculate electron tem peratures and densities once the required atomic parameters are known. These lines originate in transitions from low lying metastable levels populated by electron collisions in a low density medium. Figure 4 gives the energy level diagrams for Oil, OIII, SII, and Nil with the wavelengths of the important transitions marked. The lowest level is normally designated as (1), with (2) and (3) being the next higher levels in terms of energy. (2 1) transitions are called nebular, (3 -*■ 2) auroral and (3 -*■ 1) transauroral. Seaton (1954) gives the general expression for the Intensity ratio of (2 -> 1) to (3,m) transitions by solving a three level atom: K(m) exp(e32/tfi) (?) 1 + d2 x 45 2 p ‘ O 1S 2 1D OH 2pr Olll 2p‘ 0 n Sll 3p' 3/2 2p 1/2 ^ .5/2 2n .3/2 U ID (XI U» In u> O O O r^IT' cn * 'A .3/2 45 Figure 4. Energy level diagrams for Ions of Interest Giving Rise to Forbidden Lines. 46 where, T - 104/T e e -4 -42 -3 -1/2 x “ 10 NeTe (cm deg ' ) K(m) = E21 0(1,2) A 31 + A 32 E3m A3m e32 " 10 E32/,k and, d - 8.54 do'4) (10) 2 (ti 2 0(1,3) Collision strengths, 0(1,m) were taken from Czyzak and Krueger (1963) and Saraph, Seaton and Shemmlng (1969). Transition pro babilities Aj ^ were taken from Garstang (1968), E^ ^ is the energy difference between levels in electron volts, k is the well-known Boltzmann constant. Values for K(m) , d2 anc* ^32 were calculated and are given in Table 7. The above expression is valid under the conditions that a) « q13 q12 b) >> A31 q31 » c) A32 q32 Regarding condition a) we know that _1JL „ >A) expT - E /k Te 1 (11) q12 0(1,2) 32 16 J and since 0(1,3) « 0(1,2), this condition is satisfied for all T^. Table 7 CONSTANTS FOR TEMPERATURE CALCULATIONS (Unltless) Ion m K(m) X 21 3m d2 e32 Nil 2 6584 + 6548 5755 8.272 0.2811 2.50 OIII 2 4959 + 5007 4363 7.129 0.0281 3.30 Oil 2 3727 7319 + 7330 8.422 6.3770 1.96 SII 1 6717 + 6731 4069 + 4076 2.928 1.9960 1.39 "Vi 48 Conditions b) and c) require that x « 013/8.54 IHm,3) , <12> where (1)3 is the statistical weight of level (3). This is true if x « 200, i.e., if x ^ 10. This condition is satisfied when N < 10 T . (13) e 3 e1/2 It will be seen that most possible combinations of temperature and density in the objects under study satisfy this inequality. For each of the ions equation (9) is a slow function of tem perature. In the limit of low density the temperature approaches a constant value given by i o \ T - ---- — ----- (14) loge[R/K(m)] Hence we must have R > K(m) (15) to get finite values of T . e Densities may be obtained by comparing intensities of lines produced by transitions from the same level to levels separated by a small fraction of a volt. The strongest such doublets are [Oil]: A3726, A3729 and [SII] A6717, A6731. The latter pair is usually resolved in galaxies due to its 4.3 A separation while the former pair is not. 49 The ratio x [sil(6731) j— as a ^uncti°n of temperature has been worked out by Weedman (1968) and is given by 16717 1.01 +5.61 e1,40/Te + x(1.97 + 2.84 e1,A0/Te) 16731 = 1.20 + 3.74 e1*40^ + x (7.9o + 7.42 elt4°/Te) (16) where T^ and x have the same meanings as previously used. The low and high density limits for this ratio are 1.50 and 0.38 respect ively. Table 8 lists this ratio for the galaxies as a function of slot size. 13727 The ratio as a function of T^ and Ne is given in Kaplan and Pikelner (1970): 2 3N 1.0 + 0.33(10“19600/Tc)+ — (i.o + 0. 75(10" 19600/Te ) 13729 . . , ______10 Te ______13726 " 9.9 Ne 1.0 4- 0.40(10"19600/Te) + _ 1/2 (1.0 + 0.84(10 19600/Te 10 T e + O.UflO-39200/V , + 0.17 If this doublet is blended» a method worked out by Minkowski and Osterbrock (1959) can be used to calculate the ratio. For a blended line one must assume that the wavelength of the blend is given by X . (13729) (3728.9) + (13726) (3726.2) ^ 13729 + 13726 50 Table 8 THE RATIO I (6731) FOR M81 AND M51 FOR VARIOUS SLOT SIZES Slot (arc sec) GALAXY 2.5 6 no slot M81 0.87 1.15 M51 1.2 - 1.45 51 Figure 5 shows how the mean wavelength of this blend changes as a 13729 function of the -r- - ratio. This wavelength must be measured re- I j / £0 lative to other forbidden lines in the same spectrum. For H81 a plate was taken with Ha at 44 S/nm and X3727 blend at 22 S/m on the same plate in second order. The wavelength of 13727 was calculated using that of Ha and [SII] 16717 as standards. Their wavelengths are given in Table 9 along with the computed value of the ratio 13729 13726 * Not all lines necessary to form the ratios used in equation (1) were available for each galaxy. However, as long as a limit can be put on a line's strength and one of the lines in a (2 -► 1) trans ition can be measured, one can calculate a limit to the line ratio. Then, as long as R> K(m), an upper limit to the temperature as a function of density may be calculated. Due to the strong underlying absorption spectrum in the galaxies under consideration the upper limits on these lines are difficult to estimate, however a rough estimate may be made. The smallest measureable feature which has definitely been recognized in tracings of these galaxies has a ratio to H a° 1000 of 40. Choosing 50 to be the upper limit to a line in general, the limits to the temperature as a function of density of the regions in the galaxies were calculated. The auroral line X4363 [Oil] is recognizable on one of 3 traces of the spectrum of M51 and has a value of 40 (Ha - 1000). The wavelength of the observed feature agrees well with the laboratory 2-5 2-0 H M 81 1-5 1.0 0 5 3 7 2 7 .0 0 25 5 0 .75 3 7 2 5 -0 0 25 .5 0 75 Observed Wavelength (A) Figure 5. The Intensity Ratio 1(3726.2)/l(3728.9) vs Observed Wavelength, 53 Table 9 OBSERVED WAVELENGTHS OF Oil, Ha, SIX LINES IN M81 X 1(3729) xo i i Ha SII 1(3726) 3728.53 6581.01 6714.8 0.59 +0.18 +0.2 +0.2 value of the line. Alloin (1971) finds a similar case for M81 with a similar strength for the line. However, the feature is not seen on 3 galaxy traces of plates taken by this author. It is possibly seen on one transmission trace of an overexposed plate of this galaxy. The most reasonable conclusion is that there Is only slight evidence for the existence of 14363 in these two galaxies and that its strength is at best only poorly known. Still, to see its influence on the discussion of physical conditions, temperature- density curves for each of these cases were calculated giving upper limits to the temperature of the OIII region. The [Neill] X3868 line has been reportedly detected by Munch (1959) and Alloin (1971). It occurs in a deep absorption feature and is possibly apparent on only one of three widened blue plates of M81 taken by this author. Its existence is also certainly question able . Figures 6 through 9 display the results of the temperature- density calculations discussed above (to the left of the dashed line equation (3) is satisfied). The limits to the electron temperartures indicated by the [NIII], [SII], and [Oil] lines appear to be In the range found in galactic HII regions, i.e., approxi mately 12,000°K. The two possible measurements of the [OIII] ratio 1(4363)/I(4959 + 5007) would indicate a somewhat higher temperature. 4 -3 For low densities ( < 10 cm ) temperatures of ^ 30,000°K are in dicated for the inner 2**5 nucleus of M51 and ^ 60,000°K for the 7" Log Te(°K) A 5 3 3 5 4 3 2 iue . o ev o efrM1 (275). M51 for Ne Log vs Te Log 6. Figure L o g N e f c n r r 3 ) Lum J ..55) flimlt) ..K5755) . 1(6548*6584) 1(6548*6584) 1(4959*51 WW}li Log Te(*K) 4 5 iue . o ev o efrM1 (670). M51 for Ne Log vs Te Log 7. Figure 1(4363) (limit) l/l ON Log Te CK) 4 3 5 2 Figure 8. Log Te vs Log Ne for M81 M81 for Ne Log vs Te Log 8. Figure m m si m 3 1(6731) 1(6717) 2. 5). (2'.f5 ) 4 5 Log Te(°K) 4 3 5 2 Figure Figure 1(7319*7330) 9. 1(3727) Log Te Te Log vs 3 Log Ne for M81 M81 for Ne Log o Ne(cm‘3) Log 1(4069*4076) (6V0). 4 1(5755) (99 ) 7 0 0 5 1(4959* 1(4363? (limit) 5 59 region in M81 from Alloin’s data (Alloin 1971). To produce lower temperatures much higher densities would be required in each case. Densities indicated by the [SII] and [Oil] ratios are 3 -3 -3 ^ 10 cm for the inner nucleus of M51 and < 500cm for an area larger than this [here the value given by Peimbert (1968) was used and an approximate log Te - log Ne curve drawn in to show the de pendence]. The [SII] ratio for this region was 1.5 which is the low density limit indicating that the [SII] density is at most a few -3 3 hundred cm . For the inner region of M81 densities of 2 to 5 x 10 _3 cm are Indicated whereas approximately one-half this is found for the larger region around the nucleus. Since each of these numbers is an average over the whole reglon_seen by the slit, a comparison of densities derived for in creasingly smaller slot sizes indicates that the average density Increases by at least a factor of two from outside the semi-stellar nucleus to the nucleus itself. The actual density no doubt increases much more than this. Peimbert (1968) has shown that a considerable fraction of the oxygen in H51 and M81 is in neutral form. Boeshaar (1972) has demonstrated that in the low-temperature, hlgh-denslty filamentary regions in planetary nebulae a predominance of singly ionized and neutral atoms exists, and that the strongest emission comes from the forbidden nebular transitions. Both [Oil] and [Nil] emission reflect the conditions in the filamentary regions more strongly than in the ambient gas, while the opposite is true of the [OIII] emission. 60 Peimbert (1971) has also shown that partial ionization regions probably play an important role in the physical properties of HII regions and especially of galactic nuclei. These consist of tran sition zones between HI and HII regions, volumes partially ionized by X-rays or cosmic rays, and recombination regions where the ionizing source no longer plays a role. Since both galaxies are radio sources, large amounts of high energy particles are expected to be produced and to occupy at least part of the volume of the nucleus and are thus available as a source of ionization in this volume. Peimbert (1968) has also calculated filling factors for M51 -5 -4 and M81 of approximately 10 and 10 for these galaxies re spectively, and concludes that density fluctuations in these galaxies are probably more pronounced than in the Orion nebula. What is clear from this is that several types of regions could exist in galactic nuclei. These are: 1) partial ionization regions 2) filamentary structure (condensations) 3) high-temperature, low density regions 4) low temperature, high density regions. In the following we will assume that only two regions exist in nuclei: a) High density low temperature regions with HI, HII, 01, Oil, NI, Nil, SI, SII, SIII. b) Low density high temperature regions with OIII, NIII, N$ III. The ionizing source will be assumed to be radiative since the available evidence indicates that collisions by thermal electrons are probably not responsible for ionizing oxygen in these galaxies (Peimbert 1968). Abundances Seaton (1954), gives the expression for the intensity of a forbidden line produced by a (2 -*■ 1) transition: Ix (2,l) - h v ^ N l(X) Ne [ 1 < M > where N^(X) is the abundance of the ion X producing this line and the other terms have been defined previously. Now a . ^2 ci e £21/t:e -1.14e /t 12 id 2 1 e - 8.63(10 ) ^ e 21 (n T 1 e (20) Hence -1.14 e51/t C_ N.(X) e ^ IY (2,1) - (21) A ^3+1 x where . -2 „ _ -1/2 x = 10 Ne T 62 3 The emission of hydrogen/cm /sec is given by I(HB) - N(H+ ) Ne E° 2 10~25 (24) In Table 10 are given for case B and as a function of temper ature E° 2 CHS), I(HB)/I(H8) and E^j 2(Ha). The starred values were interpolated from values given by Aller and Liller (1968) for E° _ and Aller (1956) for I(Ha)/I(HB). Hence, 9 —lto 1*14 e_ /t lx (2,l) C2 10 Ne Tfi ' N(X) e (25) ^ ( H a ) o _i/9 i n oc (C3 + 10 Ne T£ ' ) N(H ) Ne E° 2 10 or _ oc 1<1A E*./t o 1/9 N m E 10 e e [102 C« T + Ne ] I(X) N£22_ _ 3------(26) N(H ) C2 I(Ha) The constants C2> and 1.14 e2^ are given in Table 11, for the transitions of interest. In every case it was assumed that 1(5007) 1 31(4959) and 1(6584) - 31(6548). (Seaton 1960). The ionic abundances were derived as a function of temperature for [OIII], [Nil], and [SII]. The ratio [Oil]/[OIII] as a function of temperature was derived from the ratio of the lines 13727(011)/ 1(4959 + 5007). Using equation 26 twice, the result is: .» Table 10 HYDROGEN EMISSION COEFFICIENTS (Case B) I (Ha) Te(°K) EA ,2 I(H(3) 3,2 1250 6.05 2,37 14.34 2500 3.72 2.39 8.89 4000 2.82 2.41 6.80 5000 2.22 2.43 5.39 6000 1.868 2.44 4.56 8000 1.63 2.47 4.03 10000 1.241 2.50 3.10 12000 1.125 2.52 2.84 14000 1.008 2.54 2.56 16000 0.897 2.55 2.27 18000 0.776 2.57 1.99 20000 0.6556 2.59 1.71 25000 0.5784 2.62 1.51 30000 0.4962 2.65 1.31 35000 0.414 2.68 1,11 40000 0.3327 2.71 0.90 64 Table 11 CONSTANTS IN EQUATION 26 Ion 1.14e21 C2 C3 OIII 6.19 (10-14) 67.8 2.828 Nil 6.74 (10-15) 7.39 2.149 SII 7.40 (10‘15) 3.77 2.101 65 The ionization potential of 01 and HI are similar, hence the nitrogen abundance is . Step. , ! + S ^ _ ] (28) H W N(H+) N(H+j N(0+ + ) The ionization potentials of nitrogen are approximately that of oxygen, hence the nitrogen abundance is SIS1 . "(S* + / * > . sia^i.,! + i s a p . } (29) U ' N(H ) N(H ) NCO ) Sulphur is more difficult since the ionization potentials of SI and SII do not resemble oxygen. Since 01, Oil and OIII are approximately equal in ionic abundance a good approximation is N(S+) - lUS**) (30) Then, since Xg ++ ~Xq+++ , we have for the sulphur abundance: sgl - SfiJ. , 2 + Sl^J_ , (31) N(H ) N(0 ) Again It is assumed, that since [01], [Oil], [SI], [SII], HI, HII are formed in the same region, the ratios N(0°) a N(S°) N(H°) (32) N(0+) N(S+) N(H+) hold. 66 Tables 12 through 15 list the derived abundances relative to solar abundances for various apertures and for Orion III, a typical HII region studied by Peimbert and Costero (1969). Solar values of oxygen and sulphur were taken from Aller (1963) (H ■ 12.00, 0 ■ 8.83, S ■ 7.22) while the nitrogen value is from the new determination by Withbroe (1969) (N - 8.48). Figures 10 and 11 display the abundances for the inner nucleus of K81, and for two regions of M51, the 275 nucleus and the 7M-20" region measured by Peimbert (1968). The dashed line in each case shows the. relative nitrogen abundance which would have been derived If the old value of 7.96 had been used (Goldberg et al. 1960). M51 From Figure 11 nitrogen is solar at a temperature of ^ 11000°K in the inner nucleus and at 5000°K in the outer region. On the other hand oxygen is solar at 6000°K in both cases. Sulphur is seen to be solar at the same temperature as nitrogen. Clearly in no case is there a single temperature at which both N and 0 can be near the solar values. Using the new nitrogen value for the sun changes these temperatures from what they would be if the GMA (Goldberg, Muller, Allen) value were used (30,000°K for 275 and 8000°K for 7"-20''). The value of 30,000°K is extreme and is ruled out by the upper limits to the electron temperature set by the [SII] and [Nil] lines. This implies an apparent over abundance of nitrogen in the nucleus. 67 Table 12 [N(X)/N(H)] R M81 (2-5) ABUNDANCES, LOG ^ > 7^ ) 1 i> UN 0+ Te(°] O/H S/H N/H______+ H Withbroe GMA ™ 2500 3.87 2.59 2.48 3.00 0.97 4000 1.41 1.21 1.60 2.12 0.90 5000 0.56 0.71 0.60 1.12 0.84 60C0 0.00 0.38 0.28 0.80 0.80 8000 -0.63 0.027 -0.06 0.46 0.72 10000 -1.08 -0.26 -0-34 0.18 0.67 12000 -1.33 -0.39 -0.46 0.06 0.64 14000 -1.52 -0.51 -0.57 -0.05 0.61 16000 -1.67 -0.60 -0.66 -0.14 0.59 18000 -1.81 -0.69 -0.75 -0.23 0.57 20000 -1.93 -0.78 -0.83 -0.31 0.56 25000 -2.08 -0.87 -0.91 -0.39 0.53 30000 -2.20 -0.95 -0.99 -0.47 0.52 35000 -2.31 -1.02 -1.06 -0.54 0.51 40000 -2.42 -1.12 -1.15 -0.63 0.50 68 Table 13 MSI NITROGEN ABUNDANCES (RELATIVE TO WITHBROE'S SOLAR VALUES) fN(N)/N(H)]M51 log ------[n (n )/n (h )Js u n Te(°K) 2V5 7" 7"-20" 20" 30" Off 2500 2,72 2.50 2.04 2.23 1.55 4000 1.38 1.14 0.63 0.84 0.17 5000 0,91 0.66 0.12 0.34 -0.32 6000 0,01 0.36 -0.22 0.02 -0.65 8000 0.31 0.04 -0.59 0.33 -0.98 10000 0.06 -0.22 -0.87 -0.60 -1.25 12000 -0.06 -0.34 -1.02 -0.74 -1.38 14000 -0.05 -0.44 -1.14 -0.84 -1.49 16000 -0.18 -0.47 -1.24 -0.88 -1.58 18000 -0.31 -0.60 -1.33 -1.02 -1.66 20000 -0.39 -0.69 -1.42 -1.14 -1.75 25000 -0.40 -0.76 -1.50 -1.19 -1.83 30000 -0.53 -0.83 -1.59 -1.26 -1.90 0 o> o 1 35000 -0.60 • -1.67 -1.34 -1.98 40000 -0.68 -0.99 -1.76 -1.43 -2.06 69 Table 14 M51 OXYGEN ABUNDANCE [n(0)/n(h)]M51 log — ------[n (o )/n (h )]s u n Te(*K) 2V5 7" 7"-20" 20" 30" Off 2500 3.94 4.15 4.07 4.21 4.19 4000 1.53 1.6? 1.63 1.75 1.73 5000 0.706 0.766 0.762 0.892 0.381 6000 0.170 0.183 0.183 0.330 0.317 8000 -0.433 -0.28 -0.484 -0.312 -0.322 10000 -0.859 -0.776 -0.986 -0.764 -0.767 12000 -1,09 -0.954 -1.21 -1.02 -1.02 14000 -1.27 -1.14 -1.42 -1.21 -1.21 16000 -1.32 -1.25 -1.60 -1.27 -1.37 18000 -1.55 -1.48 -1.72 -1.50 -1.50 20000 -1.67 -1.60 -1.86 -1.63 -1.62 25000 -1.81 -1.74 -2.02 -1.78 -1.78 30000 -1.92 -1.85 -2.14 -1.90 -1.90 35000 -2.03 -1.96 -2.26 -2.01 -2.00 40000 -2.14 -2.06 -2.38 -2.12 -2.12 Table 15 M51 SULFUR ABUNDANCES [n(s)/n(h)]M51 log ------[N(S)/N(H)]SUN Te(®l 2V5 7" 2500 2.76 2.85 4000 1.42 1.50 5000 0.938 1.01 6000 0.628 0.684 8000 0.301 0.357 10000 0.0426 0.0911 12000 -0.082 -0.041 14000 -0.182 -0.141 16000 -0.272 -0.241 18000 -0.352 -0.319 20000 -0.435 -0.403 25000 -0.512 -0.484 30000 -0.581 -0.554 35000 -0.654 -0.630 40000 -0.784 -0.772 -J o -J RATIOCM8D iue 0 M1 25 Abundance. C2V5) M81 10. Figure 10 20 Te (103 o K) 040 30 NCH) ) P N N(H) N(H)- N ( N ) , N ( S ) - 71 1-0 72 O 00 to ^ CM CM CM M U SOI CM T e<103'K) T e<103'K) y ie p K ) CM Figure 11. M51 Abundance. 73 The change from the GMA values also changes the relationship between nitrogen and oxygen in the region of 7"-20" from the nucleus. Previously, if oxygen was assumed to be solar at ^ 6000°K then nitro gen is over abundant by a factor of 2 whereas use of the n w value puts nitrogen underabundant by about the same factor. This part of the curve is more sensitive to small changes in temperature. Since changes in atomic parameters can shift the curves slightly, perhaps this difference is not as significant. Since the nitrogen abundance of the sun is unknown by perhaps close to this amount (as shown by the large difference in GMA and Withbroe values), the differences seen in the abundances 7M-20" from the nucleus do not present a clearly defined problem. The situation is quite different for the nucleus of M51. If, as Peimbert assumed, we take the oxygen abundance to be solar at 'v 6000°K, then nitrogen is overabundant by a factor of 3.5 (Withbroe value) or 12.5 (GMA value). Assuming nitrogen to be solar would put oxygen down from the solar values by a factor of 14 or 1401 A de finite problem clearly exists in this case, although the Withbroe nitrogen value.reduces the problem to one no worse than that en countered by Peimbert. From Table 6 it is seen that the ratio I[Nil]/I(Ha) increases by a factor of 10 from 30" away from the nucleus to the inner part of the nucleus. Such a large change can be accounted for by changes in 1) the N degree of ionization (Morgan and Osterbrock 1967), 2) the electron temperature (Burbidge et al. 1963), 3) the nitrogen 74 abundance (Burbidge and Burbidge 1965* Peimbert 1968), or 4) the relative importance of partial ionization regions (Peimbert 1971) or some combination of these four. Peimbert (1968) has argued that since the degree of ionization in H51 (and M81) is similar to that in galactic HI1 regions, a change in the degree of ionization alone could not explain the change in the [Nil]/Ha line ratio and hence in the observed abundance. From Table 6 and Figure 11 it is clear that rather than increase, the relative proportion of OII/OIII decreases going into the nucleus. However, this change could also be due to absorption internal to the nucleus region of M51. An Increase in absorption in the direction of the nucleus would affect X3727 [Oil] to a greater extent than X5007 [OIII]. Similarly absorption would affect the strengths of A3727 and '5007 relative to Ha much more than it would affect [Nil], Of the combination of 1-4 above, all but 3) would affect [Oil] and [Nil] similarly since these are presumed to co-exist in the same volume and have similar ionization and excitation parameters. Hence one would expect a change in the [Oil] line A3727 in the same direction as that of [Nil] A6584 if the effects 1), 2), or 4) were present. Such effects plus an increase in absorption internal to the nucleus would be expected to change to ratios of [Nil], [Oil], and [OIII] in a manner similar to what is seen in Table 6. If [SII] co-exists with [Nil] it would perhaps be expected to exhibit similar changes. Inspection of Table 6 shows that it roughly remains the same (within the error bars) or possibly decreases 75 slightly Into the nucleus. Superficially one would consider that the reddening argument should be abandoned, yet since sulphur is solar at approximately the same temperature as nitrogen, this argues for a partially wavelength dependent process. The effect of a specific amount of reddening has been com puted and is displayed in Figure 12 assuming 2*?9 or absorption at A3727 exists in the 2V5 nucleus. All lines were corrected with a Whitford law and the abundances of 0, N, and S were calculated. The temperature at which oxygen is solar has changed to ^ 11,000°K and the degree of ionization has changed to values closer to the 7"-20u area. The abundances of nitrogen and sulphur remain essentially unchanged. Thus absorption internal to the galaxy could play an important role in resolving the abundance problem. The effects of temperature change were also calculated for a AT ■ 12000°K - 7000*K assumed to exist between the 7"-20" area and the nucleus. The emissive properties of oxygen and nitrogen change slightly differently with temperature due to excitation potential: e(0/N, T2) -1.14(Ae21) ( |---- ) (33) e(0/N, Tx " e e2 el where, again, tg ■ Te/lo\ and e2^ is the excitation potential of the level of the ion. For oxygen ([OIII] A5007), t2^ = 2.481 and for nitrogen ([Nil] 6584), t2 » 1.885. Since we are not certain of the exact temperature relationship between [OIII] and [Oil], we assumed for simplicity that the electron temperature for both regions is the same. LOG WIQ(M51) LA-'° RATIO(SUN) iue 2 M1 25 Audne orce fr Reddening. for Corrected Abundances (2V5) M51 12. Figure 10 0 4 NMLNSl N f l H ) N C H ) “ O N 76 1-0 77 For the 2V5 region the line strengths of [0X1] and [OIII] in the 7M-20" region were assumed to be the true strengths and were corrected for the assumed temperature gradient and the abundances calculated. They are displayed in Figure 13. The temperatures at which both elements are solar has been changed. Nitrogen is solar at 8000°K and oxygen at 11000°K and the abundance differences have been reversed. Nitrogen is underabundant with respect to oxygen at 11000'K. Since a number of temperature differences can produce the same gradient AT, it appears possible to fit the observed obser vations with a temperature gradient for which the upper value is the temperature at which both oxygen and nitrogen are solar in the nucleus. The GMA. value of the abundance would necessitate larger temperature gradients and higher nuclear temperatures. There is some possibility that [OIII] may exist at this tempera ture, but not [Oil]. This possibility cannot be ruled out, but the evidence for it is poor. It can be concluded that it should be possible to fit the observations with a single model including both a temperature and an absorption increase as the inner nuclear regions are approached and the question becomes whether this suggestion is astrophyslcally reasonable. Some information is known about absorption in the nuclei of galaxies. Spinrad and Peimbert (1970) derived a figure of 4?32 78 M 51*5 AT =(12000 °K-7000eK) ‘ 10 5 i/> S q s s -2 10 20 30 4 0 Figure 13. M51 (2V5) Abundances Corrected for Temperature Gradient. 79 for M82, but this exceptional galaxy is an Irregular possibly seen edge on. Evidence does exist for massive absorbing clouds within 10" of the nucleus of H31 (Johnson and Hauna 1972, Rubin and Ford 1972) which clearly remove most of the blue light. Wampler (1968) has also shown that up to 3?3 of absorbing material exists in the nuclei of Seyfert galaxies (see also Osterbrock 1971). M81 From Table 6 it is clear that the situation for M81 is similar in some ways and different in others. A decrease in the ratio of [Nil]/Ha is seen with increasing aperture until the largest region is sampled where the ratio increases significantly. This must have been due to a large increase in the ratio at a distance of 10"-50" from the nucleus since the absolute brightness in this region is less than in the nucleus. Since Peimbert (1968) sampled a region of only 20" in diameter this effect did not show up in his data. Alloln (1971) perhaps saw this region, but it is not clear exactly which region she sampled. Considering the above named possible mechanisms for producing such an effect, an actual change in abundance is the most difficult to believe. Assuming the intensity change is axially symmetric (it could be due to an isolated region on only one side of the nucleus) the observations would imply concentric rings of differing nitrogen abundance. This is somewhat more difficult to accept than a simple Increase of nitrogen into the nucleus. Absorption, temperature or 80 the influence of partial ionization regions could be expected to vary over such distances (1" * 15 pc) in an extended nucleus (M81 is Sb). Either explanation is still possible but the abundance change is perhaps least plausible. Figure 14 shows that nitrogen and oxygen are solar at temperatures of 6000°K and 7500°K respectively. Nitrogen would then be 2.5 times overabundant using the new value of Withbroe and 8 times using the GMA value. Sulphur Is again similar to nitrogen. Figure 22 shows the results of using 2™1 of absorption at X3727. The relative abundances are shifted Indicating that less than this amount should produce solar values at ^ 8000“K. A discrete point of intensity is not available off the nucleus so that a temperature gradient model was not calculated. However, the temperature dif ferences needed to produce solar values are small if the new nitrogen value is used. Since values for the intensities of the oxygen lines are not available for all slot sizes it is difficult to compare the changes with the [Nil] lines. The two values which are available imply that the [Oil] line A3727 decreases into the nucleus whereas [Oil]/[OIII] increases slightly. The lines [SII] XX6717 + 6731 are seen to in crease when [Nil] increases outside the nucleus whereas, they apparently decrease on the average into the nucleus. Since these results differ in general from M51, simple comparisons are not easily made. Yet, some of the same conclusions are certainly valid. 81 -T it o- N(N) N(S) ~ NCH)'N(H) - NO. NOH) 10 20 3 0 4 0 Te(103oK) Figure 14. M81 (2'.'5) Abundances Corrected for Reddening. 82 Conclusions It has been observed that the [OIII] line Increases with respect to Hot going Into the nucleus of M51 although not as much as the nitrogen line X6584. The [Oil] line is seen to decrease going into the nucleus. These facts coupled with a similarity between the abundances of sulfur and nitrogen with respect to the solar values implies the possibility that a wavelength dependent process is at work in the nucleus of M51, namely, reddening internal to the galaxy itself. This could be the explanation for the apparent large abundance change going into the nucleus. If a change in a physical parameter such as electron temperature or ionization were re sponsible for the large change in the [Nil]/Ha ratio with distance from the nucleus, then similar changes would be expected to occur in the strengths of the [Oil] and [SII] lines. Such a change is not observed. The [01] line was not observable, hence any change in this line with distance is unknown. For M81 a large apparent overabundance is also observed and is due to a similar increase in the [Nil]/Ha ratio. The Increase does not appear to be smooth with distance in this galaxy, however. The choices for the explanation for the large change in the [Nil]/Ha line with distance from the nucleus are then: i) internal reddening in the coupled with an increase in the electron temperature, 83 li) a change in the abundance of nitrogen, ill) a change in the influence of partial ionization regions. Presumably a combination of all three effects coul<^also be at work. The problem could be solved in principle by an independent measurement of the reddening using the methods discussed by Wampler (1968). Information about the change in Influence of the partial ionization regions could be obtained by further obser vations of the neutral oxygen lines with distance from the nucleus. Strong support to the overabundance of nitrogen as the explanation is given by the observations of Searle (1971), 84 CHAPTER 4 MODELS OF STELLAR POPULATION The Methods As described in Chapter 2 I have adopted the method used by Alloin et al. (1971) for determining the relative proportions of the various spectral types in the nuclei. It was chosen since it lends itself to the type and quality of data available from the 110 A/mm image tube plates. The disadvantage of the method is that it Is not very sensitive to some of the parameters Involved, a fact due ultimately to the low accuracy of the continuum level. Several factors influence the observed measures of continuum and equivalent width. These are: 1) the luminosity function 2) abundances 3) reddening internal to the galaxy observed 4) interstellar line absorption in the nucleus 5) thermal continuous emission 6) nonthermal continuous emission. Both M31 and NGC 5195 are radio sources, so there is the possibility that 6) might influence the continuum. Each also has emission lines in its spectrum, although they are considerably 85 weaker in NGC 5195 than in M51, hence, 5) might be Involved. As discussed in the previous chapter there is evidence that reddening internal to the galaxy plays an important role; hence, 3) might be involved. Previous discussions (Splnrad 1962) have led to the con clusions that interstellar line absorption would add only minute amounts to the calcium and sodium equivalent widths. This con clusion was based on an upper limit for the number of neutral sodium atoms available, a number dependent upon the electron temperature and density. From Chapter 3 it is seen that the electron 3 -3 density in the nuclei of normal galaxies can reach 2-4 x 10 cm 4 and perhaps higher in the condensation regions. This is 2-4 x 10 higher than was used by Spinrad in his calculations, and since »Ne — oe " X / T e (34) o from the Saha equation, the number of atoms available to produce interstellar sodium could clearly be greater by several orders of magnitude than the amount suggested by Splnrad. Since the ion ization potentials of calcium and sodium are approximately the same and since much neutral material and a range of ionization levels exist in the nucleus, it is quite possible that either of these elements could contribute significantly to the absorption of these elements. The presence of super metal rich (SMR) stars in the nucleus of M31 was found to be necessary by Splnrad and Taylor (1972). 86 Such stars would affect the calcium and sodium line strengths used in the model making by perhaps factors of 3 (Andrillat et al. 1972). Increased strengths are not expected for magnesium, however (Andrillat et al. 1972, Cayrel and Vaziage 1972). Hence, the effect on the line strengths would be similar to interstellar absorption, and some Independent means would be necessary to decide which in fluence was dominant. The method for calculating stellar models consists of match ing synthetic equivalent widths and continuum with the observed quantities. The initial parameters consist of the fraction of visual light contributed by eleven stellar groups. Values for the standard star continue were adopted from curves given by Alloin et al. (1971) and are listed in Table 16 (normalized to 1*00 at 15360). Table 17 lists values of the equivalent widths for the eight spectral lines used in the calculations. These were taken from the relations published by Andrillat et al. (1972). Figure 15 shows a schematic H-R diagram for the eleven stellar groups. Calculations The following calculations are based on similar ones by Splnrad and Taylor (1971). Define, ■ number of stars of the ith type (i=l NTOT) l^j * relative luminosity of stars of the ith type at wavelength j (J**l, ..., N) 1^ “ relative luminosity of stars of the ith type at 15360 Table 16 ADOPTED STANDARD STELLAR CONTINUA \ BO AOV A5V FOV GOV KOI II KOV K5III K5V MOIII 3700 1.43 0-75 0.64 0.54 0.34 0.11 0.23 0.06 0.15 0,015 3933 1.35 1.20 1.00 0.79 0.45 0.24 .33 0.16 0.22 0.09 3968 1.34 1.20 1.00 0.79 0.47 0.26 .35 0.18 0.24 0.11 4050 1.31 1,18 0.99 0.80 0.50 0.30 ,37 0.22 0.26 0.13 4101 1.30 1.17 1.00 0.82 0.52 0.33 .39 0.24 0.28 .14 4226 1.25 1.14 1.00 0.85 0.57 0.40 .44 0.29 0.32 .18 4300 1.23 1,12 1.00 0.87 0.60 0.44 .46 0.32 0.35 .21 4340 1.22 1.11 1.00 0.88 0.64 0.47 .48 0.36 0.36 .24 4500 1.19 1.10 1.00 0.89 0.68 0.57 .60 0.44 0.46 .32 4900 1.09 1.05 1.00 0.94 0.84 0.78 .73 0.69 0.66 .60 5184 1.02 1.02 1.00 0.97 0,93 0.98 .86 0.r,0 0.84 .82 5360 1.00 1.00 1.00 1.00 1.00 1.00 1.00 1.00 1.00 1.00 5700 0,93 0.96 1.01 1.06 1.14 1.14 1.13 1.25 1.22 1.36 5890 0.89 0-83 1.01 1.01 1.17 1.23 1.21 1.46 1.35 1.68 6100 0.86 0.91 1.00 1.10 1.23 1.31 1.30 1.63 1.48 1.95 6500 0.81 0.87 1.00 1.12 1.32 1.43 1.40 1.82 1.68 2.41 6700 0.78 0.84 0.98 1.11 1.3' 1.48 1.44 2.03 1.76 2.57 Table 17 ADOPTED STANDARD STELLAR EQUIVALENT WIDTHS line Spectral K H + He Cal G Hy Mgl Nal "D" Type BO 0 5.0 4.0 0 0 4.0 0 0 AOV 1.0 11.0 12.0 0 0 12.0 0 0 A5V 4.0 17.0 15.0 0 0 15.0 0 0 FOV 6.0 13.0 9.0 0 1.0 9.0 0 0 GOV 14.0 9.0 3.0 1.0 7.0 3.0 1.0 1.0 KOI II 17.5 12.0 1.2 3.0 13.0 1.2 2.0 2.0 KOV 17.5 12.0 1.2 3.0 13.0 1.2 3.0 4.0 K5III 18 15 1 5 12 1.0 2.5 3.5 K5V 18.0 15.0 1.0 5.0 12.0 1.0 3.5 9.0 MOV 19.0 13.0 0.5 7.0 11.0 0,5 4.0 16.0 MOIII 19.0 13.0 0.5 7.0 11.0 0.5 3.0 6.0 OD 00 89 -6 - 5 0 7 - 0 9 V - 4 BOV 3 2 1 KOft* K S t o AOV 1 V 2 A 5 V FOV 3 GOV 4 5 6 KO V 7 8 K 5 V 9 MOV Spectral Type Figure 15. Schematic H-R Diagram. 90 Lj “ synthetic galaxy luminosity (relative to 1.00 at *5360) = observed galaxy luminosity (relative to 1.00 at J A5360) P.. - fraction of light at wavelength j due to stars of of the ith type * fraction of light at *5360 due to stars of the ith type w*al » observed galaxy equivalent widths (X) *3 ySy11 = synthetic galaxy equivalent widths (X) J Then, NTOT t P.. - 1.0 (35) i-1 J Ki » n (36) ij NTOT I N± I i-1 1 iJ NTOT Ni*jj L « X N I - (37) i«l 1 J ij therefore, N±N s- TPii b. “ P! 0 8 ) and NTOT T - <39) 91 A program to calculate these quantities was written for the IBM 360/75 Conversational Programming System with which direct online access to the computer is possible. Fractions for each group were entered on a trial and error basis until satisfactory results were obtained. This method of calculation enabled results to be seen lmmedlately so that the actual time spent calculating models was minimized. Results and Discussion NGC 5195 Figures 16 and 17 display the results of model calculations for NGC 5195. Table 17 describes the characteristics of in dividual models. It was concluded that it was not possible to fit both the hydrogen lines and the K line of Call in a single model. Model number 35 (see Figures 24 and 25) fits lines and continuum quite well with the exception of the K line. The difficulty with this model is that the fraction of late type dwarfs is too large. Spinrad (1962) classifies this galaxy as G*, meaning mostly giants with a early spectrum in the blue. The Cal line at X6162, a luminosity indicator, is not seen on two good plates which 1 took Implying a predominance of giants at 6100 X. Model number 53 is the best giant dominated model which fits the red continuum reasonably well, yet the late type star indicators, the K line and 92 79 • Model No. 35 o NGC 5195 observed o o 0 °0 0 • o o ° o o o o o Oo 73 ■ ° o * ' * 0 o> • o* • • * .2 71 o IO o CM ° o 6 9 o 6 5 - 1-5 2 0 2 5 A W 1) Figure 16. Comparison of Observed and Calculated Continua for NGC 5195. Percentage Visual Light N G C 5195 80°A MOB KSV KSV A5V A5V A5V MOV A5V K5V.M0V KOV k o v A 5V K O I MOV MOV KOV Calculated W^'s 18 16 ■ 14 H y Mg I Nal 39 46 35 50 52 53 Model Number Figure 17. Stellar Models for NGC 5195. 94 Table 18 CHARACTERISTICS OF NGC 5195 MODELS Model $ Characteristics 39 Good fit for all but K, poor continuum 48 Good fit for all but Na"D", red end fits 35 Good fit for all but K, fair continuum fit 50 Best giant only model, late type indices down 52 Best giant dominated model, G, K, down, poor continuum 53 Giant dominated model, fits red continuum, late type indices down 95 the sodium MD" line are both much lower than observed. It was not possible to fit both blue and red continua with a giant dominated model. Model number 53 is consistent with an amount of interstellar reddening and with either SMR stars or interstellar line absorption for the sodium and calcium. An unidentified band at 16284 normally correlates well with the presence of interstellar sodium in stellar spectra. On two plates of this galaxy (tf 139a and # 14lc) a dif fuse feature is seen partly redshifted into the 16300 [01] night sky feature. Similar a diffuse feature is found on intensity traces of this galaxy at the position of the 14430 interstellar feature. Both pieces of information argue for a partly inter stellar origin for the sodium and calcium lines. The location of this absorbing material is another problem. A dust lane is seen crossing NGC 5195 which could give rise to the absorption. This lane does not appear in front of the nucleus, but matter could easily be there without being seen. Spinrad's plates (1962) were widened over the whole galaxy and yet no indicators of large amounts of interstellar absorption were found. Since this author's plates were taken in the nuclear region only, these facts argue for the placement of the absorption in the nucleus. 96 M51 Figures 18 and 19 display the results of the models of M51, while Table 18 describes their characteristics* It is clearly not possible to fit the observed continuum with a combination of black - body curves. Introduction of both interstellar reddening and thermal continuum in appropriate amounts could possibly fit the curve, but any increase in blue light would wash out the calcium and hydrogen lines and make the line strengths difficult to fit. Hence, it is concluded that no model will fit the available measures of line strengths and continuum. Model number 18 fits lines well with poor continuum (see Figure 19), while number 15 fits the con tinuum as well as possible with a poor line fit. Number 40 is the best giant dominated model (Splnrad also classifies this galaxy as G ) which a;-proximately fits the blue continuum. Clearly no detailed information can be obtained about M51 with the present data. The continuum indicates an early type while the line strengths indicate many late type stars. Further observations which may help to resolve the problem are: 1) an independent determination of the reddening 2) a redetermination of the continuum by photoelectric means. Percentage Visual Light M 51 MOUt MOB MOB 6 0 ®/< MOM FOV AOV.GOV FOV MOV MOV MOV A5V A5V A5V FOB K5* BO BO CalculatedW -v 's 18 16 Figure Figure 18. Models M51. for Stellar 10 I la I Nil C at Mg I _ Cat O b se rve d 18 15 35 37 40 Model Number 98 3 3 CM CM K Figure 19. Comparison of Observed and Calculated Continua for M51. 99 Table 19 CHARACTERISTICS OF M 51 MODELS Model # Characteristics 15 Fits continuum fair, lines poor 18 Fits lines well, poor continuum 35 Best giant model, NaMD" down, poor continuum 37 Best giant dominated model, poor continuum 40 Giant dominated model, fits blue continuum fairly well 100 CHAPTER 5 PROPERTIES OF GALACTIC NUCLEI Sunsnary of Conclusions in This Thesis Several conclusions concerning the nuclei of galaxies have been derived: i) Electron densities have been shown to be in the range -3 1000 to 5000 cm which is larger than previously thought, ii) An electron density gradient clearly exists in the nuclei of galaxies, iii) Concerning the abundance problem, either nitrogen and sulphur are overabundant to a larger extent than previously thought (see for instance Peimbert 1968) or large amounts of absorption exist internal to the nuclei of galaxies and a tem perature gradient exists or partial ionization regions play a very important part in determining the emission line strengths in the nucleus. iv) Ionization changes probably do not cause the observed change in the |NII] /Ha ratio. v) Large amounts of interstellar line absorption are possibly present in the spectra of the nuclear regions of some galaxies or supermetal rich stars play an Important role. 101 Summary of Known Data Table 20 lists all currently known data about the nuclei of M31, M51, M81, M321 NGC 1052, NGC 5195, the Galactic Center and the Seyfert galaxy NGC 1068.A few conments would be useful about the dependability of this data. The distance are reason ably well known for all of these galaxies. The nuclear sizes are in general upper limits, since recent work has shown that angular sizes for M31 and NGC 1068 are 0.9 arcseconds (Morton 1972). Models of stellar populations are known at least roughly for all but our own galactic nucleus. The stellar velocity dispersions are mostly likely upper limits except for M31. The gas velocities are more accurate. Mass-lumlnosity ratios are not well known. The abundance problem has not been solved, so that the physical conditions are only poorly known. Nevertheless, tem peratures and densities are not abnormally high in the nuclei, although they are possibly higher than in many normal galactic HII regions. From the objects with sufficiently well-known radio spectra a few properties are derivable. A synchrotron spectrum plus self absorption is apparently responsible for the shape of the spectrum in NGC 1052 (Wills 1968). If it is responsible then, for magnetic -4 fields of ^ 1 gauss, sizes of the order of 7 x 10 arc seconds are necessary. If B <0.1 gauss, then some other mechanism is at least partly responsible for the observed shape. Table 20 SUMMARY OF THE PROPERTIES OF THE NUCLEI OF SEVERAL GALAXIES OBJECTS Galactic Properties M31 M51 M81 M32 NGC 1052 NGC 5195 NGC 1068 Center Type Sb Sc Sb E2 E4 EO-Irr Sb Distance (mpc) 0.7 4.0 3. 0.7 19 4.0 13. 0.010 Nuclear Size 0.9 2.7 5 0.8 6" 7" 2 60 (arc sec) Spectrum KO-MOIII F8-G0III KV Kill F-KIII 6-Kill V Stars 120 15 V Gas 225 125 100 600-900 3000 (km/s) Abundances S tar s SMR? SMR? SMR? SMR? Gas nonsolar nonsolar Optical nonthermal no no? possibly yes radiation Nuclear ejection possibly possibly yes yes 102 Mass-luminosity 50 2 27 2 ratio Table 20 (Concluded) Galactic Properties M31 M51 M81 M32 NGC 1052 NGC 5195 NGC 1068 Center Infrared flux yes limit yes yes yes yes yes Electron 500- 1000- 100, 300 density (cm ) 1000 3000 40000 Electron 6000- 6000- temperatures (°K) 30000 40000 Radio component yes yes yes limit yes yes yes yes Radio structure complex complex complex complex complex complex Radio variability yes 44 44 Bolametrlc 0.7X 2x10 0.5x10 luminosity (ergs/sec) 10AA 104 For the galactic center the assumption of self-absorption to explain the observed spectrum would imply improbably high magnetic fields of the order of 1 0 ^ gauss (Lequeux 1971) . Most likely the explanation is free-free absorption due to weakly ionized clouds in front of the nonthermal source. Electron densities of a few hundred cm ^ are Implied. A Composite Model The general properties of a galactic nucleus may be assessed and inferred by referring to Table 21. These are given in Table 18 as a composite idealized model. The basic components consist of a star cluster 10 parsecs in diameter with a core of about 0.1 parsec, an HII region several hundred parsecs in size filled 4 with a low density plasma with electron temperatures of *\# 10 °K _3 and densities of a few hundred cm , condensations of higher density and perhaps lower temperature ( ^ 10 3 cm -3 and 'v 5000°K) and a radio source with two components ( ^ 0.01 pc and 'v 5 pc) . 6 3 Stellar densities are 'v 10 stars/pc . Gas velocities are 100- 400 km/sec, while stellar velocities are 100 km/sec. The stellar component consists of mostly G, K and M stars with an enriched lower main sequence, although anomalies do show up. The nucleus is overabundant in the metallic elements, a fact which shows up in both the stars and the interstellar matter. Much neutral material and dust exists in the nucleus, both playing an important role in the dynamics and energetics of this region. 105 Table 21 PROPERTIES OF COMPOSITE GALACTIC NUCLEUS Components: Star Cluster a. 10 pc diameter, core 'u 0*1 pc - G,K,M stars HII region 10-100 pc Te 5000-10000 °K Ne 200-5000 cm-3 Radio Source '< 0,,01 pc and ^ 5 pc - double concentric sources Other Properties: Stellar Mass ^ 107-108 M_ Mass of Gas ■>, lO6 ^ Star Density ^ 106/pc3 a Gas 100-400 km/s o Stars ^ 100 km/s AA ^ 10 ergs/sec. ^SoL 106 The bulk of the radiation in the ultraviolet region is due to hot young stars and to thermal radiation from the HII region. In the visible and near infrared stars emit most of the radiation, although forbidden emission lines provide a fraction of what is observed. In the far Infrared a large amount of energy comes from a source probably connected with the presence of dust in the nucleus and the energetic particles producing the radio source. In the radio region a synchrotron source with a low frequency cut off provides the energy. The total energy given off by the nucleus is approximately 10 ergs/sec. Comparison With Active Objects Properties of Seyfert nuclei have been sunmarized by several authors (Rees and Sargent 1972, Sargent 1971). They have nearly stellar nuclei ( 'v OVl * 5 pc for NGC 4151), a rich forbidden line spectrum in which both high and low excitation lines coexist. This latter feature is found only in strong radio sources and Super Nova Remnants and is correlated with the presence of a nonthermal source for the excitation. Line widths are greater than a few hundred km/sec, reaching 6000-8000 km/sec in some galaxies. Many Seferts show wide Balmer lines with narrow cores and forbidden lines with only narrow cores. Physical conditions Include electron temper- A 6 7 3 atures of ^ 10 °K and electron densities of up to 10 -10 cm” . 44 Bolometric luminosities range from 1-100 x 10 ergs/sec. 107 Most Seyferts are strong radio sources, some such as OQ208 have peculiar spectra indicating the presence of radio activity. Some objects such as NGC 1068 are variable over time scales of days to months in the radio, optical and infrared regions. 0Q208 which ought to have much activity on the basis of its radio spectrum shows little activity in the radio and optical regions, although it is a radio and optical variable (Craine and Warner 1972). One object (3C120) is thought to contain a very small double source which is moving apart at faster than the speed of light (Shaffer and Cohen 1972). There are other indications of activity. Most objects of this type have evidence for mass ejection from their nuclei, and several have random gas velocities of several hundred km/sec. Redshifts for Seyfert galaxies have not shown any deviation from the Hubble relation. The properties of QSO's have been summarized by several authors (Burbldge and Burbidge 1967). They are defined as a class of objects which have starlike appearance on direct plates and have very large redshifts in their spectra. Broad emission lines are present which are weaker relative to the continuum than in Seyfert nuclei. Sharp absorption lines are sometimes present. Most of these objects show strong and often variable nonthermal continua and are powerful radio sources and infrared emitters. Their bolometrlc luminosities are 44 50-5000 x 10 ergs/sec. 108 Multiple absorption and emission redshifts are found, and if the redshift is entirely due to the expansion of the universe, then these objects are moving away from us at speeds approaching 90% of the speed of light. Many of these sources contain a double radio source. In two objects (3C279 and 3C273) their components are observed to be moving away from each other at possibly greater than the speed of light. QSO spectra are peculiar in that several components are sometimes seen. They show rapid variability over time scales of years to days. The optical variations are similar in kind but are not necessarily correlated with the radio variations. From variability and long baseline in- terferometry observations, the existence of several small com ponents can be inferred in these objects. The basic conclusion that can be drawn from a comparison of properties of galactic nuclei with those of Seyfert nuclei and QSO's is that a sequence of power levels and activity exists in going from the normal nuclei such as M31, M51, M81,and NGC 5195 to NGC 1052 to NGC 1068, NGC 4151, 3C120, and OQ208 to 3C273 and 3C279. Each has some of the properties of the others. The so- called normal galaxies have some indications of activity if good enough resolution and enough data is available. Some Seyferts show more activity than others, and some showalmost as much as the QSO’s without the excessive redshift. 109 The link between these objects is either one of evolution in the sense that QSO's become Seyferts which become normal galaxies or vice-versa, or it is one of simple activity in the sense that there exists a QSO-like object in the nucleus of every galaxy and the differences become one of the relative activity of this object (Sandage 1971). Resolution of this problem be comes dependent upon the distances assumed for these objects. If they are present in every nucleus one would assume that the QSO stage would be found somewhat randomly distributed in time and space. If there is an evolutionary sequence, then this phase would be found most likely in the early stages of the universe. The problem then depends upon some independent proof of the existence of non-velocity redshifts. The problem is only partly dependent upon the ultimate source of energy in all of these objects. Gravitation as the primary source has been shown to be efficient enough in the rotating (ergosphere) black-hole model to provide this energy (Bardeen 1970), although the existence of collapsed objects is not yet certain. One fundamental problem connected with the nuclei of these objects is the direction of the evolution of matter. The two possibilities which are apparent are that (1) matter is evolving toward singularities on a small scale, meaning that galaxies form out of the general medium pervading the universe at some epoch, or that (2) matter is coming from some source such as a singular ity so that galaxies are formed around highly condensed regions which may have geometrical connections with the rest of the matter in the universe. 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