Southern planetary nebulae

B. LOUISE WEBSTER 19 66

A Thesis Submitted for the Degree of Doctor of Philosophy

The original work reported in this thesis is that of the candidate alone.

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Dr. B.E» Y/esterlund has been more than generous with encouragement, guidance and valuable discussion in all phases of

this thesis. I am greatly indebted to him both for his kindness

in this and for continued access to his extensive data on planetary nebulae before publication.

I wish to thank Professor B.J. Bok for his interest in the work and Professor S.C.B. Gascoigne and Dr. D.J. Faulkner

for the benefit of many stimulating discussions. I am very

grateful to Dr. A.W. Rodgers for obtaining scanner observations of one of the objects and for his helpful comments.

It is a pleasure to thank Dr. K.G. Henize for kindly permitting me to make use of his catalogue of planetary nebulae prior to publication.

The care and cheerfulness with which Mrs. E.R.Y/ilkie typed the thesis is much appreciated.

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Chapter 1. THE POTENTIAL OF SOUTHERN PLANETARY NEBULAE

1.1 Introduction 1

1 ,2 Physical processes and conditions 4

1.3 The evolution of a planetary nebula 6

1.4 The distances of planetary nebulae 13

1.5 Galactic distribution and kinematics 16

1.6 Groups of equidistant planetaries 19

1 ,7 The Magellanic Clouds 20

Chapter 2. THE OBSERVATIONS

2.1 Introduction 25

2-2 The photoelectric equipment 27

2.3 Calibration of the observations 29

2 r 4 The photoelectric observations 34

2.5 Qualitative observations 37

2.6 The spoctrographic equipment 39

2.7 Radial velocities 42

2.8 Relative line intensities 50

Chapter 3. DISTANCES FROM FLUXES AND DIMENSIONS

3.1 Distances from the assumption of constant mass 76

3.2 Hß fluxes in the Magellanic Clouds 79

3.3 A direct limit to the masses of Cloud planetaries 83

3.4 Masses by identification with galactic planetaries 86

3.5 Comparison with previous calibrations 90

3.6 Binary systems 92

3.7 The mass of a planetary nebula 93

3.8 Angular radii and filling factors 94

3.9 Absorption corrections and adopted distances 95

Chapter

4

4.1 The kinematics of planetaries in the Small

Magellanic Cloud 103

4.2 The kinematics of planetaries in the Large

Magellanic Cloud 104

4.3 The kinematics of the Large Magellanic Cloud 108

4.4 The kinematics of planetaries in the Galaxy 109

Cheap ter 5. PLANETARIES IN THE NORMA REGION OF THE GALAXY

5.1 The galactic distribution of planetaries in

Norma 114

6.2 The intensity of the nebular continuum 138

6.3 Star temperatures and luminosities 141

6.4 The evolution of the nuclei 145

CHAPTER 1

THE POTENTIAL OF SOUTHERN PLAHETARY NEBULAE

1.1 Introduction

lie physical state of a planetary nebula is, in general,

quite well understood, yet the phenomenon as a whole has not been

fitted completely satisfactorily into the framework of stellar evol­

ution. It is known that a planetary is a body of gas, about a tenth

of a solar mass, surrounding with some degree of symmetry and ionized

by a hot, subluminous star from which it is presumed to have been

ejected. The gas is expanding at about 20 km/sec and the nebular

radii range from a few hundredths of a parsec, where the object can

first be recognized as a planetary, to about 0.7 parsec, where the

nebula blends into the interstellar medium.

The large accumulation of knowledge on these objects has

been extensively reviewed in the last few years by Aller (1956), Seaton (i960), Gurzadyan (1962 a), Osterbrock (1964) and Minkowski (1966).

Perek and Kohoutek (1966) have collated the observational data into a

general catalogue. The conclusions from these reviews are that there

is a large quantity of observational material, the physical mechanisms

responsible for the spectrum are satisfactorily known and the physical

conditions in the nebula are known qualitatively, but the finer details

remain to be elucidated and there is an uncertainty in all quantities

for which the derived values depend on distance, since this cannot be

It is widely accepted that a planetary nebula is formed from a not very massive star after the red giant stage and immediately before the white dwarf state. The mass in the shell is a significant fraction of the mass of such a star and must play an important role in

the evolution of the star. The evaluation of this role is hindered

by a lack of fundamental data caused by the inadequate distances. The most important of these are:

(1) the masses of the nebulae and of the nuclei,

(2) the population characteristics of the planetary subsystem,

(

3

(H-R) diagram.

Means of improving this situation in the Galaxy are limited. However, in the closest external galaxies to the Sun, the Magellanic Clouds, over seventy, intrinsically luminous planetary

nebulae have been identified. An unparalleled opportunity to study

a large group of equidistant planetaries is thus presented.

The primary purpose of this thesis has been to use this advantage in studying the three, poorly known characteristics mentioned

above. The first of these is intimately connected with the cali­

bration of a distance scale for all objects. Information on the

second comes from the kinematics of the Cloud planetaries, the study of which is simplified by the knowledge of the relative positions of the

objects over the system. Finally, the temperature at which the central

3

One limitation of this work, at times an advantage, has been that only the brightest nebulae have been studied. Even

so the main difficulties associated with the program are observational ones. In particular, the central stars are fainter than 1 7th magni­ tude, yet narrow-band filters have to be employed in measuring them.

In some parts of the work it is important to have a comparison with galactic objects. The distribution of planetaries along the southern Milky Way has been studied by Henize (1 966) and shows a peak at new galactic longitude 1^'L = 330°. The region around this has been chosen for an intensive study, both as a comparison with the Clouds and to establish whether there is a true maximum in the space density in this direction or whether the interstellar absorption is abnormally low. Since the latter is partially true, we have been able to study the kinematics of objects distant from the Sun and close to, but not right at the galactic centre.

1.2 Physical processes and conditions

Ionization of the gas of the nebula by stellar ultra­

violet quanta more energetic than the ionization limit is followed

by recombination with the associated emission of lines as the electrons

cascade to the ground state (Zanstra 1927? Menzel 1926). References

to a series of papers discussing this and other physical processes in

gaseous nebulae are given by Aller and Menzel (1945)- In addition

a continuum is formed from, recombination, free-free and two-photon

emission (Seaton 1955)? the free-free continuum remaining strong at

radio wavelengths. Any series of line ratios is determined by the

energy distribution of the star and the radiation field within the

nebula and references to the calculated decrements are given by

Seaton (i960). In the event that all the radiation able to ionize a given ion is absorbed, the intensities of the permitted lines of that

ion give a direct measure of the amount of this ionizing radiation

which cannot otherwise be observed below the earth’s atmosphere (Zanstra

1927? 1961). Stellar temperatures can therefore be estimated if an energy distribution, such as one of those calculated by Gebbie and

Seaton (19^3) or Bohm and Deinzer (1965)? Is assumed.

Ions in the gas can be excited, but only rarely ionised,

by electron impact. At the electron temperatures common in gaseous

nebulae, excitation generally occurs only in complex atoms and ions and

to low-lying metastable levels. The electron returns to the ground

5

dependence of the intensity of this radiation on the electron temper­

ature and density permits these quantities to be evaluated from the

relative strengths of the forbidden lines (see e.g. Seaton i9 6 0).

Once the temperature and density have been found from lines of the

same ion, the chemical abundances may be estimated from the relative

intensities of both permitted and forbidden lines of different ions.

The enhancement of certain lines by the coincidental

fluorescent mechanism (Bowen 1934) has no application here.

The physical condition in the nebula can thus be deduced

from its spectrum. The ionization level within a nebula is stratified,

the atoms near the central star being highly ionized and the degree of

ionization decreasing outwards as photoionization is attenuating

radiation of progressively longer wavelength (Wilson 1930; Hummer

and Seaton 1964). Wilson (1958) has found that the shells are

expanding with velocities up to a few tens of kilometres per second.

The expansion velocities are greater for ions of lower excitation, from

which it is inferred that the outer parts are moving outwards more

rapidly than the inner regions. The general excitation level varies

considerably among the nebulae and the classification into ten groups

by Aller (1958) will be followed here.

The most recent abundance analyses by Aller (1964) show

that the helium to hydrogen ratio, 0.1

65

is higher than in the Sun but quite constant among the forty planetaries

the very underabundant planetary in M 15 (O'Dell, Peimbert and Kinman 1964) to the oxygen rich NGC 7027. There is some indication of a

variation in the neon to oxygen ratio. However, the heavy element

abundances are sufficiently close to normal to exclude the possibility that planetaries originate from metal-enriched peculiar stars (Aller

1959).

The electron temperatures are determined primarily by the balance between heating by photoionization and cooling by inelastic collisions, and lie between about 8000 °K and 20000 °K. The electron

7 -1

densities range from Nq ~ 10 cm in the densest planetaries known, to the low values at which the nebulae become indistinguishable from the interstellar medium (N < 1 0 cm" ). The values of both quantities depend on which line intensities are used for the determinations. The most plausible explanation for this is the presence of small scale inhomogeneities (Seaton and Osterbrock 19575 Osterbrock i 9 6 0 ) , which has been confirmed on large scale photographs of some objects (Zanstra

1955? Minkowski and Osterbrock i 9 6 0 ) .

These non-uniformities are superposed on the large scale density fluctuations responsible for the many curious but symmetrical shapes observed. In optically thick objects the outer edge is sharply defined at points where the ionizing radiation is exhausted (Stromgren 1939), but inside this edge, and throughout optically thin nebulae, the

distribution of the ionized hydrogen follows that of gas. (Unless

7

in the hydrogen Lyman continuum). The forms projected on the sky are,

in general, either ellip+ical, ring-like, biaxial or irregular. Evans

and Thackeray (1950) calculate that the distribution of axial ratios of the nebulae photographed by themselves and by Curtis (1918) is such that a considerable fraction of the planetaries must possess spherical

symmetry. Westerlund and Henize (

1966

conclusion in a more extensive and homogeneous photographic survey

of southern planetaries. The three-dimensional form of a ring-shaped

object has been suggested to be an ellipsoidal shell (Wilson 1950) or a dense toroidal form with faint extensions in other directions

(Khromov 1962). Each of these models is represented by certain of the

observed nebulae.

The form of the nebula has the potential to indicate the mechanism of its ejection and maintenance. Thus a flattened, equatorial nebula may result from rotationally forced ejection (Limber 1965)? whereas

spherical symmetry points to radial expulsion of gas after pulsation

or an explosion. Osterbrock (1

964

taining the shape of the nebular shell and concludes that a force is

needed to restrain its outer edge. Forces supplied by mass loss from

the star (Mathews

1966

1 .3 The evolution of a planetary nebula

A relatively small dispersion in expansion velocity and

mass among planetaries permits the average lifetime of an observable shell 4

to be determined as a few times 10 years. Furthermore, the linear

dimensions of the shell serve to indicate the time interval from the

beginning of the planetary stage, so the evolution of the shell and nucleus

can be traced. Hubble (1922) was the first to notice that the spectral

character of a nucleus is related to the size of its shell and since then

it has become increasingly clear that the nuclei change drastically over

the short lifetimes of the nebulae.

No sensitive luminosity criteria are applicable to the

spectra of the central stars but a large range in this parameter is indi­

cated. Features commonly seen in population I Wolf-Rayet (WR), 0 and

Of stars , which have extended atmospheres, are seen in some nuclei.

Others show broad, weak, low-luminosity features. The temperatures all

appear to be high. Aller and Wilson (1954) have measured the equivalent

widths and profiles of absorption lines in the cooler stars and derive

excitation temperatures greater than 30,000°K. The weakness or complete,

absence of observed lines in others indicates extremely high temperatures,

Greenstein and Minkowski (1964) have tentatively identified a high excit­

ation line of 0 VI in two faint nuclei which must therefore have excitation

temperatures of more than 660,000°K. Broad emission lines, indicating

instability and possibly some mass loss, occur in some of the nuclei over

9

The difficulties in obtaining spectra of the faint central stars, and in subsequently interpreting' them, make photometric

methods more suitable for studying large numbers of objects. The

first uniform program of this kind was carried out by O ’Dell (1963 a) with measurements of Hß fluxes, angular dimensions and star magnitudes. The distances were calculated from a model of an optically thin nebula while the temperatures of the central stars were computed by assuming that the nebulae are optically thick and applying corrections to the

derived values. (Nuclear temperatures found from the Hß flux and the

magnitude of the star are hereafter referred to as Zanstra temperatures). O ’Dell found that the nuclei are spread over a large range in temperature and luminosity and that the radii of the nebulae increase as the nuclei, presumably, evolve from bright, cool to faint, hot positions.

Harman and Seaton (1966) established definite criteria for estimating optical depth and demonstrated that a nebula starts its life optically thick to the far ultraviolet radiation, then becomes thin progressively to quanta capable of ionizing K, He and He+ , and finally, as the luminosity of the nucleus drops, increases again in optical

depth. The first change, caused by expansion, had been realized

earlier, by Shklovsky (1956 a), for example. Certain features observed in planetaries place some of them in each of the three stages of develop­

ment. Intrinsically small, dense nebulae are regular* in form and the

that neutral hydrogen is present. In rather larger objects the

Zanstra temperatures a.re lower than the star temperatures inferred

from their spectra, and faint, often irregular outer extensions to the

ionized gas confirm that ionizing radiation can escape. Minkowski

(1942)

LO

with He II in some large, low surface brightness planetaries, with

high Zanstra temperatures, points to their being optically thick.

With the correct optical depths, and temperatures

calculated from the H, He I and He II lines, Harman and Seaton (1964)

and Seaton (

1966

central star increases in temperature and luminosity. The luminosity

ohange soon halts,but the temperature becomes even higher until degeneracy

stops the star fcom contracting and a drop in luminosity and cooling

towards the white dwarf state follows. There are features of this in

common with theoretical tracks of collapsing stars of 0.6 and 0.4

solar masses calculated by Fayashi, Hoshi and Suginoto (1962).

The evolutionary track of a planetary can be extra­

polated back to investigate the nature of the object immediately pre­

ceding it. Such a procedure leads to a cool supergiant or giant star

with hydrogen emission lines appearing in its spectrum after some

expansion (Deeming

1965

stars also leads to the region of the H-R diagram above the main sequence,

but it is unlikely that the helium flash would supply enough energy to

throw off a shell at the top of the globular cluster red giant branch

11

The comparative ease with which planetaries are discovered

has resulted in much more information on their space density and distri­

bution than on those of subluminous blue stars with temperatures and

luminosities like those of planetary nuclei. However, it is known that

hot stars below the main sequence are found in all population types and

are one or two orders more numerous than planetary nuclei (Cox and

Salpeter 1961; Greenstein 1966). Two of the classes, in particular,

have features in common with the nebulae. Novae are known to eject

shells with kinetic energies comparable to those of the planetaries.

The masses in the former are about a thousand times smaller, however, the

expansion velocities being correspondingly larger, and a high proportion

of the novae are binary systems (Kraft 1964) . The link between the two

phenomena would thus appear to be slender (Minkowski 1948) although

the same mechanism may be responsible for both. The spectra of some

subdwarf 0 and B stars resemble those of planetary nuclei, and the

temperature and luminosity ranges of the two classes are not dissimilar

(Greenstein 1966). It seems plausible either that a star passes through

each group in turn or that some minor difference prevents a shell from

being ejected or from being observed (see Osterbrock 19 6 4)•

The available figures on space densities and lifetimes

suggest that about 10 per cent of the white dwarfs may have passed

through the planetary stage, but there is still quite a large uncertainty

of population type than planetaries (G-reenstein 1966) so it could not

be as high as 100 per cent. However, it is possible that all stars pass down Vorontsov-Velyaminov1s

(

)

or as subdwarf 0 or B s tars.

The stage immediately preceding the planetaries must

also be common to the evolution of a high proportion of the stars and

must be characterized by some sort of instability, either potential or

actual. This suggests an intrinsic variable, most of which arc also

connected with post-main-sequence evolution. The long period variables

have been suggested and are perhaps the most likely candidates. Their

kinematic and space properties cover a wide range but, like planetaries,

are mainly those of a disc population (Feast 1963)» Many of them

show hydrogen emission lines formed deep in the atmosphere and emission

lines of ionized iron and the light curves are not completely explained

by pure pulsation.

Merrill (1958) and others have suggested that planetaries

may pass through the symbiotic phase before the nebula reaches a more

conventional condition. The faint nebula around R Aqr, discovered

by Hubble, shows that gas may escape from such a system and with a type

of symmetry similar to that in some planetaries, but its expansion

velocity is probably higher than is typical for planetaries. If

13

characteristics of the cool component point to a connection of the latter

stars with the early stages of planetary formation.

1 »4 The distances of planetary nebulae

The planetary nebulae lie in the region of the H-R diagram

where spectroscopic and photometric luminosity and temperature criteria

are very insensitive. Since the nuclei cover a large luminosity range

and, furthermore, are evolving very rapidly over it, it is not possible

to derive their distances in ways similar to those for main-sequence

stars. As a consequence of its expansion, as well as of the change in

ionizing radiation, the nebular envelope also changes. Methods applicable

to all planetaries are therefore based on some assumptions regarding the

variation of the planetary with time and are calibrated by distances

found by the following means.

(a) Van Maanen (1923) measured trigonometrical parallaxes and

found a value significantly greater than the error only for NGC 7293»

(b) Proper motions have principally been measured by van Maanen

(1933) a^-d Anderson (1934). The data are inadequate to permit anything

but a mean parallax to be extracted as, for example, Barenago (1946) has

done.

(c) A few planetaries are binary systems and spectroscopic parallaxes

may be obtained for the companions. This has been done for the G8 to

KO dwarf associated with NGC 246 (Baum and Minkowski i9 6 0) and for the

super-posed on at least three other planetaries. The AO star in the direction of the centre of NGC 3132 has a velocity 17 km/sec higher than that of the nebula (Evans 1 9 6 5) and may not be connected with it. Vorontsov- Velyaminov (1 9 6 2) reports that the central star of W 68 is of spectral type B9? and Minkowski (1955) that there is an A5 star associated

with NGC 2346 (although there is no further reference to this in

There is some risk in generalizing the results for any of these binaries. Seaton (1966) regards NGC 246 as abnormal on the basis of its high He II \ 4 6 8 6 to Hß ratio. If the planetaries are highly evolved the masses of the stars from which they were formed must be greater than those of their unevolved companions which are assumed to have been formed at the same time. For the three objects with 339 to A O companions the masses would be greater than about three solar masses and the galactic distribution demonstrates that this is not true for the majority of the planetaries.

(d) The planetary in M 15 appears to be unusual in having an

abnormally bright central star and a subluminous nebula (^Minkowski 1966). It has a normal helium content but is deficient in oxygen (O’Dell,

Peimbert and Kinman 1964).

15

Osterbrock (1964). Liller’s work (19^5) on fourteen planetaries

demonstrates the difficulties in the approach, both observationally and

in the interpretation. The observed rate of expansion is affected by

the variation of expansion with ionization level, optical depth and

mass loss from the central star. It seems that the information is more

profitably used in studying these properties, as Liller has done, than

in deriving dubious distances.

(f) The upper limit to the luminosity has been measured in some

external galaxies and will be described later.

(g) Seaton (1966) assumes a homogeneous, spherical nebula and from

the observed value for the Hß brightness and an electron density from

forbidden line ratios derives the linear radii of fourteen planetaries.

The distances, which follow from the angular dimensions, are very

sensitive to the electron densities and systematic errors can follow if

density fluctuations axe present.

An observed decrease in surface brightness with increasing

radius has been the foundation of several general methods for determining

distances. Vorontsov-Yelyaminov (1954 a; b) explained the relation

in terms of constant absolute magnitude for all the nebulae. A similar

assumption had been made by Zanstra (1951) and. was developed by Berman

(1957) and Parenago (1946). The hydrogen line flux of an optically

thick nebula is determined by the radiation of the central star and is

nebula becomes optically thin, and the ionizing radiation can escape, the flux is proportional to the electron density so decreases with time and the assumption of constant magnitude breaks down.

In an optically thin nebula, in the absence of continued mass loss from the star, the mass of ionized gas remains constant and

this leads to a relation between surface brightness and linear dimensions. Shklovsky (1956 a$ b) used the fact that since the linear dimensions

calculated from this depend on the mass only to the two-fifths power, the assumption that the mass is constant can lead to quite accurate

distances. This method, employed by Minkowski and Aller (1954)?

has been recalibrated and adapted to the more meaningful Hß fluxes by O ’Dell (1962) and Seaton (1 966).

Kohoutek (196O5 1961) uses the star magnitudes and temperatures in an effort to take individual masses into account.

Minkowski (1964) outlines the principle on which the optical depth may be simply estimated and the correct method adopted for nebulae excited

by similar central stars. The initial rise in temperature and luminosity

of the nuclei found by Karman and Seaton (1964)5 makes the division less clear. Minkowski points out, however, that even if both these diffi­ culties are overcome, the irregularity of the gas distribution sets a limit to the attainable accuracy.

1 .5 Galactic distribution and kinematics

H

assumed that the majority of these are actually at the centre.

Minkowski (1966) points out that the excess of small diameter objects

(Minkowski 1951) relative to the distribution in the solar neighbour­

hood (O'Dell 1962) indicates that discoveries in the centre are incomplete

and that the concentration may be even more pronounced. On the other

hand, Perek

(1963

planetaries towards the centre by the general methods described in

section 1.4. The resulting galactic distribution is quite uniform,

with no preference for a particular distance. Although many of these

planetaries must obviously lie between the Sun and the galactic centre,

the majority of them are expected to be in the central bulge. The

cause of the apparently uniform distribution is, then, the inadequacy

either of a general absorption model or of the method of finding the

distances.

On the basis of this strong concentration to the centre

direction, coupled with a moderate concentration to the galactic plane

(Minkowski and Abell

1963

disc population (Oort 1959)« Recent estimates of the mean distance

from the plane range from 280 parsec (O’Dell 1962) to 500 parsec

(see Seaton

1966

equal to 1.2 solar masses. No planetary is known to be associated

with extreme population I, but those with A-type companions and the more

population I. Some membership in extreme population II is evidenced

by the planetary in M 15 and by the planetary near the north galactic

pole (Haro 1951) which, since it is faint and unresolved, must be a

considerable distance from the galactic plane. High velocity planetaries,

especially those near the anticentre, must belong to one of the oldest

populations.

A satisfactory kinematic model of the subsystem has not

been constructed because of the inaccurate distances. There are some

obvious tendencies, though. The radial velocities have been measured

principally by Campbell and Moore (1918) and Mayall and Minkowski

(unpublished). Between 8 and 12 kiloparsecs from the galactic centre

the system displays signs of differential rotation but with low angular

velocities and a higher dispersion than the extreme population I

(Minkowski 196 4). Minkowski (1966) notes, however, that since Berman

(1957) derives a value for Oort’s constant of A = 14 km/sec/kpc, the

kinematics of these two populations may be quite similar.

The most striking feature of the velocities is the

extremely large velocity dispersion of the planetaries in the centre

direction (Minkowski 1964). This discovery emphasizes the importance

of this subsystem, of which the members can be identified and studied

at large distances, in the investigation of the galactic distribution

19

1.6 Groups of equidistant planetaries

To investigate the validity of the existing ways of

finding the distances of planetaries and, hopefully, to advance to a

more satisfactory general method, accurate distances of many more

individual objects are needed. The galactic centre has frequently

been suggested as an excellent place to look for a distance dependent

parameter or combination of parameters. Perek’s work illustrates the

problems posed, in this direction, by the absorption corrections and by

contamination of the group of planetaries at the centre by objects

closer to the Sun. Other opportunities to study equidistant groups are

found in external galaxies. Here the complications of the galactic

centre are avoided, but the great advantage of knowing the angular

dimensions is lost for all but the very largest nebulae.

The most distant system in which planetaries have been

found is M 31 • Baade (1955) identified five suspected planetary

nebulae in a field rich in population II and Miss Swope (1963) measured

their mean photographic magnitude as -2.55* The similarity of M 31

to our Galaxy suggests that these objects are similar to galactic planet­

aries, but their large distance modulus limits further work.

The closest galaxies to the Sun, the Magellanic Clouds,

are sufficiently far away that their constituents may be regarded as

equidistant, their foreground obscuration is small and they have been

excellent for distance studies but they have the disadvantage that they

are dwarf, irregular systems, rich in extreme population I, and there

may be some intrinsic difference in planetaries in them that would

invalidate any generalizations. The data are sufficiently plentiful

for an evaluation of this possibility to be made and it will be discussed

in the next section.

It will be interesting, in the future, to look for

planetaries in the dwarf elliptical systems, Sculptor (m-M = 19»7) and

Draco (m-M = 20.o) being the closest. Their stars appear from colour

magnitude arrays to belong to population II. Sculptor (Hodge 1965)

bears some resemblance to M 5> a globular cluster of intermediate metal

richness, while Draco shows some evidence of being extremely metal poor

(Baade and Swope 1961). Planetaries as luminous as the brightest in

the Magellanic Clouds would be found at photographic magnitudes of

16/7-Data on planetary nebulae in these systems would complement the inform­

ation on the Magellanic irregulars.

1.7 The Magellanic Clouds

Before assuming similar properties for individual planet­

aries and the planetary subsystem in both the Galaxy and the Magellanic

Clouds the validity of this assumption should be examined. The most

important parameters that must be compared are the chemical abundance

and the age of any system of objects. Detailed reviews on the Large

21

published by Buscombe, Gascoigne and de Vaucouleurs (1954)? Thackeray

(1963)

(1966)

parameters will be discussed.

The subsystems within the Clouds have been outlined by

Westerlund (1965)» Each contains a central body of mixed population,

a complex of population I objects, an intermediate age disc and a halo

containing old globular clusters. That the relative proportions of

these populations differs between each Cloud and the Galaxy is most

graphically illustrated by the percentage of the mass that is in the

form of neutral hydrogen. In the LMC it is between 5 and. 9 per cent

(McGee and Milton 1966) and in the SMC it is about 30 per cent (Hindman

and Sinclair 1965). The corresponding value in the Galaxy is less than

5 per cent but the percentage rises to 17 per cent in the solar

neighbourhood (Oort

1966).

Some understanding of the progress of star formation in

galaxies has been gained in the last few years. The Galaxy is thought

to have collapsed from a spherical cloud over a short time span, while

the halo population was forming (Eggen, Lynden-Bell and Sandage 1962).

Star formation was most active in the initial stages (Dixon 1965)» In

the Magellanic Clouds the comparative scarcity of objects as old as those

in the galactic halo suggests that the stars were being formed less

rapidly in the Clouds than in the Galaxy in the early lives of these

systems (Tifft

1964).

Westerlund (19615 1 9 6 4) finds that Shapley*s Constel­ lations in the Large Cloud are of the same age. Westerlund and

Mathewson (1966) suggest that the reason for bursts of star birth could by type I super-supernovae, the surface distribution of young stars

and hydrogen supporting this interpretation. The distribution of the

periods of cepheids is also consistent with discontinuous star formation

(Gascoigne 1965). Longer period cepheids (more massive) are common in

the main body of the SMC, but the cepheids in the bar of the LMC are predominantly of short period.

In the older populations, which are more relevant to planetaries, a similar phenomenon is found among the globular clusters. Gascoigne (1966) divides the colour-magnitude diagrams of these into

four main classes. Old globular clusters, similar to galactic halo

ones, are found in both Clouds. Five of the seven bright Small Cloud

globular clusters resemble galactic disc clusters and were formed,

9

practically simultaneously, about 10 years ago. NGC 17^3? in the

Large Cloud, has some features in common with these. The last group

comprises three similar clusters in the LMC unlike any found previously. If star formation occurs in bursts, as these investi­ gations suggest, the distribution of the ages and masses of the stars from which planetaries are formed cannot be expected to be the same in

eaeh Cloud and in the Galaxy. However, the subsystem will bear the

same relation to the other populations in all three galaxies.

23

Inherent differences, such as could be attributed to

different chemical compositions, may also be present and the available

data is summarized by Roberts

(1963).

arise in the period-luminosity relation for the cepheids which differs

in all three galaxies. Indications of metal deficiency in open

cluster colour-magnitude diagrams (see Arp 1959

)

(Feast

i960).

investigations, mostly on population I objects, which find no abundance

or other anomalies in the Clouds. To these can be added the high

dispersion analysis of an LMC supergiant by Przybylski

(1965)

shows a fairly normal metal abundance. Dickel, Aller and Faulkner

(

)

(1965)

ances in gaseous nebulae in both Clouds to be only very slightly, about

30 per cent, lower than normal. Thus no inherent difference between

the Clouds and the GaRaxy that will affect the interpretation of the

planetaries has been established.

Neither has any difference between galactic planetaries

and those in the Clouds been found with the available information. All

the brighter planetary nebulae in both Clouds are thought to have been

discovered and have been discussed by Henize

(

),

(

),

Westerland and Rodgers

(

)?

(1961),

(1962),

and Mullan

(1963),

nebulae primarily on the basis of their mass. It was shown that the

planetary nebulae are unresolved, fainter than M = -3.0 and generally PS

with moderately high excitation. The number per unit mass of the

galaxy is of the same order in both Clouds as in the Galaxy. Feast

(1964 a) has obtained spectrograms of nine planetaries and detects the

lines normally present in galactic objects.

Certainly, then, to a first approximation, and possibly

to quite a high degree of accuracy, the Cloud planetaries resemble their

galactic counterparts. However, when any generalization on population

characteristics, except differential ones, derived from investigations

on the Clouds is made, one should bear in mind the effect of the

CHAPTER 2

THE OBSERVATIONS

2.1 Introduction

The distance moduli of the Large and Small Magellanic

Clouds adopted in this work are (m - M)q = 18.7 and 19.0, respectively

(Bok 1966), and the absorption is taken to be A^ = 0.2 magnitudes in

front of each system (Feast, Thackeray and Wesselink i960). Since

planetary nebulae in the Clouds are fainter than M ^ = -3.0, the data

that can be collected on them in a reasonable observing time is limited.

Accordingly, the observing program has been planned to measure absolute

values of the fluxes of only the strong lines, N ^ , Hß, HY and LO IIJ

k 3727» by the accurate photoelectric method. The interference filter

technique has been used, rather than that of spectral scanning, as this

enables more, and fainter objects to be observed. The intensities of

the weaker lines, particularly those of helium, are estimated relative

to the strong ones from supplementary, medium-dispersion spectrograms.

Radial velocities are also determined.

In unresolved objects the magnitude of the central star can

only be found from the flux at wavelengths between strong nebular emission,

corrected for the contribution from the nebular continuum. Narrow-band

o

filters have been chosen to isolate the wavelength regions around 3500 A,

Johnson U, B, V effective wavelengths, which facilitates comparison

with the widely used two-colour diagram.

The monochromatic magnitudes of the continuum at these

three wavelengths will be referred to as m_,-_-, m.«-- and m r7AA, and the

° 3 5 0 0’ 4200 5 3 0 0’

derived monochromatic magnitude at 4861 A as • For convenience

in tables, some of the more common nebular emission lines will be denoted

by their wavelengths only. These are

Co

[S II] X 4068, X 6730, [0 III] X 4363, He I X 4471» He II X 4 6 8 6, [Ar IV]

X 4740, He I X 5875 and [n II] X 6 5 8 3.

Planetaries in the Small Cloud were selected from the

nebulae that are unresolved on the large-scale photographs of Henize and

Westerlund (1963)« In the Large Cloud the source was the catalogue of

Westerlund and Smith (1 96 3) and the objects were chosen to include most

of the brightest planetaries, but with sufficient of the fainter ones to

cover a considerable range in magnitude and surface distribution over the

Cloud, Some planetaries in common with those of Feast's work (1964 a)

were studied.

In the Galaxy, all objects, except H 168 (NGC 6164—65)y

in a region in Norma between galactic longitudes 1 ^ = 320° and 340°

that Henize (19 6 6) has identified as planetary nebulae from an Ha survey

have been observed. In order that space densities may be compared at

different galactic longitudes, N^ and Hß have been measured for planetaries

27

few objects closer to the plane were excluded because of the very uneven

interstellar absorption at low galactic latitude.

The catalogue of Perek and Kohoutek (1966) will define the

future notation for planetaries, but in the interim the objects are

referred to by their numbers in the catalogues from which they were

selected. N prefixes planetaries in the SMC, P those in the LMC and H

those in the Galaxy.

2.2 The -photoelectric equipment

The photoelectric observations were made with a single­

channel photoelectric photometer attached to the Cassegrain focus of the

40-inch reflector at Siding Spring Observatory. The scale at the focus

of the f/18 bean is 11.5 seconds of arc per millimetre. The photometer

is similar to that described by Johnson (1962) and incorporates an R.C.A.

1P21 photomultiplier, always used refrigerated by dry ice. The output

current from the cell is amplified by a General Radio amplifier, type

number 12$0-A, and recorded on a Brown recorder. The steps in the volt­

age range of the amplifier were calibrated at approximately yearly inter­

vals and found to remain constant. The resistant steps were calibrated

during each observing run on star signals or on a constant light source,

and variations of up to 2 per cent between any two steps were found over

the eighteen months in which measurements were made.

The transmission characteristics of the interference

28

TABLE 2.1

The transmission characteristics of the filters

Peak Half­

wavelength width

Feature ( A ) ( A )

continuum 5300 140

N 1 5006 30

Hß 4864 25

H y 4346 ₂₅

continuum 4200 70

LO II] 3727 25

continuum 3500 90

Peak transmission

( * )

60 Spoctrolab

57 Spectrolab

65 Spectrolab

32 Grubb Parsons

50 Spectrolab

30 Spectrolab

40 Schott

These were measured with a Beckman spectrophotometer both

at the beginning and towards the end of the observing period and no changes

could be detected over this interval. The accuracy in the wavelength

measures is a few angstroms, adequate for the wider-band filters and as

a check on the stability of the others, but the transmission curves of

the very narrow-band filters, , Hß and Hy, were remeasured using the

coude spectrograph. A beam converging from a tungsten filament lamp

was passed through a filter into the spectrograph and photographed at a

dispersion of 2.7 /-/mm. An argon comparison spectrum and a wedge cali­

bration were produced on the plate as in the normal operation of the

instrument. Intensity tracings of the plates were calibrated for the

percentage transmission with the values from the spectrophotometer.

le

d

fi

lt

e

r

tr

ans

mis

sio

n

c

u

rv

e

29

a spectral scan of a typical planetary in Figure 2.1. In Figure 2.2

the wavelength regions occupied by the Doppler-shifted lines are shown

with respect to the very narrow filters. No correction for velocity is

necessary to the N . or Ilß fluxes but the Hy fluxes need to be corrected

both for wavelength shifts and for the contribution from LO III] \ 43^3 •

Irregularities in transmission over the filters were found

to be as high as 30 per cent. This could be expected to give rise to

appreciable errors in the measurement of planetaries of large angular

extent. In the present program the objects measured generally had

small angular dimensions similar* to those of the standards and the error

introduced was small.

2.3 Calibration of the observations

The wavelength bands chosen are narrow, and any change of

effective wavelength with energy distribution is expected to be too small

to affect the extinction coefficient. This was verified by deriving

extinction coefficients from observations of two stars of spectral types

AO and KO.

Multilayer interference filters are an effective means of

achieving such bands, but their transmission characteristics are subject

to long term changes with time. Intercomparison of intensities should

therefore only be made between objects with similar spectral features,

so that these changes will affect the transmission of light from both

selected, at convenient positions in the sky, one set applicable to the

emission line intensities and. the other to the strength of the continuum.

For the continuum measurements stars of «pectral type AO

or earlier were used. The continuum was measured in magnitudes per

unit frequency interval, adjusted to equal the Johnson V magnitude at

o

its effective wavelength, 5460 A. Since the pass-bands of the filters

avoid the hydrogen and helium lines, the secondary standards could be

calibrated by direct comparison with spectrophotometric standard stars.

The secondary standards themselves form a system consistent to 0.01

magnitudes, but the four primary standards gave different calibrations

of the colour indices in this system. Differences in the n u _ ^ band

were as great as 0.1 magnitudes. The pass-band of this filter is quite

wide. Consequently, to minimize any inaccuracy introduced by the

application of the different band-widths of scanning and filter techniques

to various energy distributions, the greatest weight in the calibration

of the southern standards was given to the primary standards closest in

spectral energy distribution to planetary nuclei. In Table 2.2, part (a),

the primary standards and their monochromatic magnitudes as measured by

Code (i960) and Oke (1964) are given. In part (b), the secondary standards set up for the present v/ork. are listed.

Secondary standards for the line fluxes were chosen,

wherever possible, from the bright planetaries with angular diameters

as well as continuum standards in each of the Magellanic Clouds, since there were no suitable unresolved emission line objects and the frequent use of very small diaphragms made a standard of stellar appearance

necessary. This situation is very sensitive to drifts in the filter

characteristics, since the intensity of the continuum source, which depends on the whole pass-band of the filter, will change in a different

way from that of the line source, at a specific wavelength. Standard

planetaries were therefore measured on several nights during each run. The calibration of the secondary standard stars changed by only 10 per cent over the observing period.

The difficulties associated with photographic photometry of emission lines are avoided by the use of photoelectric methods and the absolute fluxes and relative intensities of the strong lines are

now known for many planetaries. Whereas the absolute fluxes refer to

the emission over the whole surface of the nebula, frequently only a section of the planetary is selected by the slit of the spectrograph or scanner ./hen relative line strengths are determined. Since the distri­ bution of the ions around the central star depends on their excitation, the ratios of lines from a small area are not necessarily the same as

++ +

from the whole nebula. The distribution of the 0 and H ions is

similar in most planetaries but differs from that of the 0"r ions.

The published to Hß ratios are not always the same

TABLE 2 . 2

( a ) P r im a r y s t a n d a r d s f o r m onochrom atic m a g n itu d e s

O b je c t Sp V

rn5300 m4200 m3500 R e f e r e n c e s

r 2

§ Get AO 4 . 2 7 4 . 2 5 4 . 0 6 5 .1 5 0 , JM

_3

n _{O r i F6 3 .1 6 3 . 2 0 3 . 5 6 4 . 5 3 C, JM}

V O r i BO _{4 . 5 9} 4 . 5 6 4 .1 8 _{4 . 0 7} C, JM

58 Aql AO 5 .5 7 5 .5 7 5 .5 3 6 . 8 5 0

U ) S e c o n d a r y s t a n d a r d s f o r m onochrom atic m a g n itu d e s

O b j e c t Sp V

m5300 m4200 m3500 R e f e r e n c e s

ED 3719 AO 6 .8 4 6 . 8 5 6 . 7 9 8 . 0 7 A, W

ED 25938 AO 6 . 5 7 6 . 5 6 6 . 4 6 7 . 7 0 ¥

ED 147152 B6 5 . 3 6 5 .2 9 5 .1 2 5 .7 4 M

HE) 86659* B4 6 .2 5 6 . 1 5 5 . 9 4 6 . 5 0 M

ED 158186* _B3 6.81 6 . 9 5 6 .8 7 6 . 8 7 M

* Lovr w e i g h t . Used on o n l y two n i g h t s

R e f e r e n c e s

0 _{O ke, J . B . , Ap. J •} j ,440., 689, 1964

JM _{J o h n s o n , H,, L . , an d Morgan, W.W., Ap. J •> U 2 . 313, 1953}

A A rp, H ., A,, J . , 6 j , 118, 1958

W W e s s e l i n k , A. J . , M.N.R.A . S . , 12£, 359 , 1962

M M o r r i s , P.M. , M .N .R .A .S ., .122, 325, 1961

33

TABLE 2 .3

( a ) P r i m a r y s t a n d a r d s f o r l i n e f l u x e s

O b j e c t l o g F(1T1 ) l o g F(Hß) l o g F(Hy) l o g f(^3 7 2 7 ) R e f e r e n c e s

IG 418 > 9 . 4 0 - 9 . 5 4 - 9 . 9 2 - 9 .3 7 L,CD,LA,OCB,0

NGC 3242 8 . 6 4 9.81 L , CDO

NGC 6572 8 . 7 0 _{9 .7 7} 10.21 L,CD,LA,OCB

ngc _{7009 8 . 6 9 9 .7 6 1 0 .1 4 L,Li,OCB}

0>) S e c o n d a r y s t a n d a r d s f o r l i n e f l u x e s

O b j e c t _{l o g F (ll^j ) l o g F(Hß) l o g F(Hy) l o g F ( \ 3727) R e f e r e n c e s}

NGC 2857 _{- 9 . 4 3} - 1 0 . 5 6 _{- 1 0 . 9 5} AF

NGC 3918 8 . 8 0 10.01 AF

NGC 5882 9 . 3 2 1 0 .3 3 1 0 .7 4

NGC 6326 9 . 9 0 11.0 5 11 .41 AF

HD 3719* 9 . 7 0 9 . 8 9 9 .4 5 9 .7 0

HD :25538* - 9 .6 0 - 9 .7 4 - 9 .3 5 - 9 . 5 3

* The f l u x e s g i v e n a r e f u n c t i o n s o f t h e f i l t e r s and t h e v a l u e s a r e l i s t e d o n ly f o r c o m p le te n e s s

mean

R e f e r e n c e s

L L i l l e r , W. , Ap. J . , 122, 240, 1955

CD _{C a p r i o t t i , E . R . , and Daub, C . T . , Ap.. J . , V£2, 677, 1960}

AF A l l e r , L . H . , an d F a u l k n e r , D . J . , IAU/URSI Symposium N o .20, 45 ( A u s t r a l i a n Acad. S e i , , C a n b e r r a ) , 1964

CDO C o l l i n s , Q .Yi., Daub, C

4 7 1 , I960

. T . , and O’D e l l , C .R ., Ap ♦ j . , 3^33.»

LA _{L i l l e r , W. , and A l l e r , L . H . , Ap. J . , L20, 48, 1954} OCB O s t e r b r o c k , D . E . , C a p r i o t t i , E .R . and B a u tz ,

Ap. J . , jj>ß, 62, 1963

L . P . ,

quoted as about 5 per cent. The system of southern secondary standards

set up for the present work is consistent to 1 per cent and the line

strengths of the previously measured planetaries used in its calibration

were therefore taken from both types of measurement, as specified in

Table 2.3? to give the best agreement among themselves for the southern

system. LO III was measured only in the Magellanic Clouds and IC 418

was the sole standard. The emission of both [o II] and H comes princi­

pally from the outer shell of this nebula (Aller 1956), so the absolute

flux of LO II] \ 3727 can be calculated from its intensity relative to H ß .

Line fluxes derived by comparison with NGC 6572 'were found

to be systematically 10 per cent brighter than those from the other

nebulae. Its declination is +6° while the other three lie between -10°

and -20°, and it is evident that a declination effect is responsible for

the difference. NGC 6572 is small, bright, and a northern standard,

whereas the more southern primary objects are more difficult to measure

from northern latitudes. However, the calibration from the three latter

objects has been taken here, as they agree well with each other and are

widely distributed in right ascension. The values adopted for the -2 -1 logarithms of the fluxes, log F(\), are measured in ergs cm” sec” and

are listed in Table 2.3, part (b).

2.4 The photoelectric observations

The Hß fluxes were corrected for the Doppler shift in

35

curve measured on the coude spectrograph and the velocities and line

intensities given in Tables 2.11, 2.15 and 2.14 were used. In cases

for which the intensity of the oxygen line relative to that of Hy was not known, a ratio of one to five was assumed. This ratio was taken as the mean of the line ratios of planetaries that Aller (1956) uses to illustrate the medium excitation classes 5 to 8,

The continuum contributed a significant amount to the observed line fluxes in about one quarter of the objects, although it was less than 5 per cent in all but ten peculiar ones. To correct for this, four Of stars were selected on the basis that, on spectrograms from the Mount Stromlo plate files, the Hß and Hy absorption lines were

filled in with emission, giving an effectively continuous spectrum. Their energy distribution is comparable to that expected in the central

star of a planetary. Each of these stars was observed through both

the line and tne continuum filters and the contribution to the observed line intensity from the continuum was determined relative to the magnitude of the continuum. Corrections to the line fluxes were then applied.

For two planetaries with bright stars of spectral type AO in the direction of their centres, corrections were derived from similar measurements

of the AO secondary standards.

The corrected line fluxes and the continuum magnitudes

are listed in Tables 2.4 and 2.5. These, and other tables giving the

end of this chapter. The weight to be assigned to the observations is

denoted by n, and is, in general, the number of nights on which the

object was measured. Some delay in the acquisition of the Hy and

Lo

X

3727

these lines. Where three numbers are specified for the lines, instead

of one, they refer to and Hß, to Iiy and to

X

The probable error in a single observation, n = 1, is

a function of the magnitude or the logarithm of the flux and is given

in Table 2.6.

TABLE 2.6

The accuracy of the photoelectric observations

- log flux p.e. observed

magnitude _P

observed • £ •

calculated

10 0.005 13 0.03

11 0.008 14 0.04

12 0.016 15 0.06

13 0.03 16 0.14 0.08

17 0.28 0.20

Observations in the lines could be pursued to the faint, small nebulae

with moderate results, but in planetaries with large angular dimensions

and low surface-brightness the accuracy was limited by increased noise

from the sky background and interference from field stars. Except in

these cases, no dependence of the errors on the angular diameter of a

C lo u d s w ere s m a l l. C row ding o f th e f i e l d s an d i r r e g u l a r i t i e s i n th e

f a i n t i o n i z e d g a s o f t h e b ack g ro u n d in c r e a s e d th e e r r o r s , w hich were

a b o u t 0 .1 0 i n th e l o g a r i th m o f th e f l u x .

I n th e co n tin u u m th e a c c u r a c y was l i m i t e d by p h o to n

f l u c t u a t i o n s f o r s t a r s f a i n t e r th a n 15t h m a g n itu d e . To a s s e s s t h i s

e f f e c t , t h e o r e t i c a l e r r o r s w ere c a l c u l a t e d f o r th e f i l t e r s u s e d , an

i n t e g r a t i o n tim e o f 100 se c o n d s and a quantum e f f i c i e n c y o f 10 p e r c e n t .

I t was assum ed t h a t th e sk y h a s a m a g n itu d e o f V = 1 6 .0 th r o u g h a 10 seco n d

o f a r c d ia m e te r d ia p h ra g m ( A lle n 1 9 6 2 ), t h i s b e in g c o n s i s t e n t w ith th e

d e f l e c t i o n s t h a t w ere o b ta in e d on a good n i g h t . The a c c u r a c y was lo w e re d

by th e p r e s e n c e o f th e moon, by th e u s e o f l a r g e r a p e r t u r e s and by d i f f i ­

c u l t i e s i n r e a d i n g s m a ll d e f l e c t i o n s on th e r e c o r d e r , b u t lo n g e r i n t e g ­

r a t i o n tim e s w ere u s e d to c o u n t e r a c t t h e s e . I n view o f t h i s , th e

e r r o r s fo u n d i n p r a c t i c e , w hich a r e s e e n i n T a b le 2 .6 to be o n ly a l i t t l e

g r e a t e r th a n th e t h e o r e t i c a l oncL f o r 100 seco n d i n t e g r a t i o n s , a r e

r e g a r d e d a s v e r y s a t i s f a c t o r y .

2 .5 Q u a l i t a t i v e m easu rem en ts

On s e v e r a l n i g h t s when a c c u r a t e p h o to m e try was im p o s s ib le

th e r e l a t i v e i n t e n s i t i e s o f l i n e s c o u ld s t i l l be e s t i m a t e d . The l a r g e

d i f f e r e n c e s i n th e r a t i o o f t o Hß fo u n d among p l a n e t a r i e s and e m is s io n ­

l i n e s t a r s a llo w e v en ro u g h v a lu e s t o g iv e u s e f u l i n f o r m a tio n . The

r a t i o was i n v e s t i g a t e d i n a b o u t s i x t y u n r e s o lv e d o b j e c t s , p r o v i s i o n a l l y

Way. Similar percentages of the light from the lines were transmitted so the relative intensity was estimated directly from the amplified

output. The sky readings were effectively the same for each filter

and had negligible effect on the result.

Identification charts were made from direct plates in the Ha + Cn II] region while the brightness of the object in visual light was determined by the green lines of LO IIIJ and H ß . Variations in the ratio of the red to green lines caused changes in the brightness of the suspected planetary relative to surrounding field stars on the

charts and visually at the telescope. In some objects for which the

red light was strong neither nor Hß could be detected photoelectrically,

and it is suspected, because of their high Ha/Hß ratios, that these might be Be stars.

The sixty stellar objects have been divided into three classes according to the line ratios and are listed in Table 2.7« Hebulae with 11^ = Hß are regarded as being probably, though not con­

clusively, planetaries and are placed in class A. Class B contains

objects with Hß > 1T^ which may be either low excitation planetaries or

emission-line stars. Objects in Class C should be regarded as very

doubtful planetaries until further information is obtained. The first