Rotation curves of high luminosity spiral galaxies are flat, to distances as great as r=49 kpc. This implies a significant mass at large r. Rotational velocities increase about 20 km/s across a spiral arm, as predicted by the density wave theory. By analogy, it is suggested that our Galaxy has a flat rotation curve out to r∼60 kpc, with V ∼ constant at near the solar rotational velocity, and m ∼7×1011 m⊙. Values of A and B imply that the sun is not located in a spiral arm.

Studies of the rotation curve of our Galaxy at galactocentric radii, R, greater than the solar distance, Ro, from the center require the use of conventional optical techniques since the distances to as well as the radial velocities of Population I objects are needed.

The galactic rotation velocity at the Sun, , can be derived several ways, none of them direct and unambiguous - (1) the solar velocity can be found relative to the halo population (the RR Lyrae stars, globular clusters etc.), but may contain an unknown contribution from possible systematic rotation of the halo system (2) the product Ro ω(Ro) = Ro (A-B) can be calculated but is uncertain because of large uncertainties in each of these three quantities (3) the motion of the Sun with respect to the center of the Local Group can be found but includes the motion of the galactic center of mass and (4) the velocity-longitude dependence of the outer HI boundary can be examined to deduce the most likely value of . The incorporation of new data into analyses using methods (1) and (3) gives essentially the same answers as older studies. Examination of the accumulated current evidence suggests that the best values for the solar rotation velocity and the galactocentric distance Ro are 220 km s−1 and 8.5 kpc respectively.

The conventional value of the galactic circular velocity near the Sun is Vo = 250 km s−1 (Schmidt 1965). Recently both lower (Mathewson) et al., 1974; Einasto et al., 1976) and higher (Lynden-Bell and Lin, 1977) values have been suggested. The following summarizes our dynamical determinations of Vo. Altogether six independent methods have been used.

It is well known that the velocity dispersion of nearby stars increases systematically with age, and this fact is conventionally explained by postulating that stars are randomly accelerated by encounters with massive gas clouds (Spitzer & Schwarzschild 1951, 1953; Wielen 1977). However, a strong constraint on the possible role of random accelerations is provided by the apparent existence of fairly old moving groups of stars (Eggen 1969, Boyle & McClure 1975). An alternative explanation (Tinsley & Larson 1978) is that the age dependence of the velocity dispersion of stars older than 109yr is produced by a gradual decay with time of interstellar turbulent motions, as predicted by plausible collapse models. This effect cannot account directly for the variation of velocity dispersion with age observed for stars younger than 109yr, but Tinsley & Larson suggested that this could be explained if the velocity dispersion of the youngest stars reflects only the local turbulent motions of the gas, while the velocity dispersion of older stars reflects in addition larger-scale non-circular motions in the galactic gas layer. If the interstellar medium possesses a hierarchy of motions whose velocity dispersion increases with the size of the region considered, older stars which have traveled farther since their formation will sample the gas motions over a larger volume of space, and thus will have a larger velocity dispersion than the younger stars.

The mean angular momentum of stars in the solar vicinity with respect to the galactic center shows a significant change as a function of age. In figure la the variation is given for B stars using the Strömgren index [u-b] as abscissa since it depends nearly linear on the logarithm of the age (i.e. the interval corresponds approximately to 7.7 < log(age) < 8.2). This can be explained by a change in the birthrate as a function of place, velocity and time and/or an azimuthal variation in the galactic potential. To investigate if the presence of a density wave potential could cause the observed distribution a number of models were made. In these models stars were formed close to the spiral arms according to the density wave theory and traced back to the present epoch. The mean angular momentum of stars being within 300 pc from the sun is shown in figure lb where the amplitude of the wave was 5 percent, the inclination angle and one arm coincided with the Sagittarius arm. For a pattern speed Ωp= 13.5 km/s/kpc the observed values could be reproduced quite well which indicats the existence of a density wave in our Galaxy.

…“there exist many galaxies where most of the light originates in an apparently flat rotating disk, and where the random stellar motions appear to be small compared to the systematic circular motion i.e. they are apparently “cold.” Our own Galaxy is such a system, and it does not seem to suffer any large scale, large amplitude, short-time-scale instability”… Ostriker and Peebles, Ap. J. 186, 467, 1973.

Space motions are calculated for 145 dK2-M2 stars with radial velocities and with parallaxes and proper motions determined and published at the Van Vleck Observatory. The stars are divided into young and old disk components kinematically and also according to the age-sensitive CaII emission intensities. Rigorous solutions for the solar motion and velocity ellipsoid were calculated for each population group using three methods of weighting errors in parallax, proper motion and radial velocity. All methods show a mean motion of the young stars outward away from the galactic center of about ten km/sec when referred to the old stars. Details are presented by Upgren (1978). The conclusion appears to confirm an outward motion suggested by Kerr (1962) from 21-cm observations. It is reasonable to inquire whether a similar outward motion can be seen in the space motions of other stars. Unlike the Van Vleck parallaxes combined with the most rigorous of the weighting methods used here, earlier results may be too insensitive to measure such a motion. Nonetheless a search of some of the existing literature shows that a small outward motion is consistent with sources so far examined. Of the motion solutions listed by Delhaye (1965) only two involve space motions of stars covering most of the main sequence. For stars with velocity dispersions similar to ours, both studies show a small outward motion of the order of five km/sec (for the A dwarfs relative to the generally older F and G dwarfs). The A dwarfs brighter than 5m.5 (Eggen 1965) have a planar motion distribution very similar to our young dK-M stars and both groups possess a mean motion close to the basic and standard solar motions. These last are mostly based on young stars and the solar motion should be reexamined relative to stars of all ages. It might be worthwhile to redetermine the mean motions of stars for which age-dependent parameters are now available, such as the F-stars with uvby measures or the giants with DDO photometry, since the local standard of rest appears to be a function of stellar age. This study was supported by grant AST77-26554 of the National Science Foundation.

Five years ago, at the previous IAU Symposium (No. 60) on the galaxy, most of us pictured the interstellar medium (ISM) as a mainly quiescent medium which evolved by orderly processes–the “steady state” ISM. Radhakrishnan reviewed the “intercloud medium” which was warm (some few thousand degrees), rarefied (some few tenths cm−3), uniformly - distributed, partly ionized, and responsible for the smoothly-distributed pulsar dispersion measures and Faraday rotations. Heiles reviewed the “clouds,” which had densities of tens cm−3, temperatures of tens K, and nonspherical shapes. These two components were in approximate pressure equilibrium; Grewing, in his review of heating mechanisms, emphasized pervasive, nonvarying processes.

A large number (approximately 7) of different components or phases are needed to describe the interstellar medium. The neutral intercloud medium is probably a composite of (a) “lukewarm, substandard” clouds (heated by grain photoeffect and shockwaves), (b) the interfaces between clouds and coronal gas and (c) some “phase 2” gas heated by soft X-rays. Ionizing UV photons are mainly produced by OB-stars and are responsible for most of the average electron density. Bulk kinetic energy for “stirring” the medium and soft X-rays are mainly produced by supernova remnants, less by O-star stellar winds.

The survey of 21-cm galactic absorption towards 819 extragalactic sources, recently carried out with the Nancay radio telescope (Crovisier et al. 1978) provides a sample of absorbing neutral hydrogen clouds which is an order of magnitude larger than those resulting from other available surveys (Hughes et al. 1971, Radhakrishnan et al. 1972, Lazareff 1975, and Dickey et al. 1978). A statistical study of this sample offers insight into the properties of local interstellar matter.

The interstellar medium may be characterized by several physically rather distinct regimes: coronal gas, intercloud gas, diffuse clouds, isolated dark clouds and globules (of small to modest mass), more massive molecular clouds containing OB (and later) stars, and giant molecular clouds. Molecules first appear in the denser diffuse clouds, and occur everywhere that AV ≳ 1m. Values of temperature, density, ionization fraction, mass, size, and velocity field are discussed for each regime. Heating and cooling mechanisms are reviewed. Nearly all molecular clouds exceed the Jeans criteria for gravitational instability, yet detailed models reveal no cases where observations can be interpreted unambiguously in terms of rapid collapse. The possibility that clouds are supported by turbulence, rotation, or magnetic fields is discussed, and it is concluded that none of these agencies suffice. Comments are made about fragmentation and star formation in molecular clouds, with possible explanations for why only low mass stars form in low mass clouds, why early-type stars form only in clouds with masses ≳ 103 M⊙, and why O-stars seem to form near edges of clouds. Finally, large-scale interactions between molecular clouds and the galactic disk stellar population are discussed.

We summarize here some of the distribution properties of CO clouds in our galaxy as determined by observation and by computer simulations of these observations. The simulations adequately account for the number of clouds, their spatial separation, and their relative velocities with respect to one another. The principal shortcoming of the axisymmetric simulations is their failure to describe adequately the clustering of clouds seen in the velocity-longitude domain.

Observations of CO emission at ℓ=0 to 70°, |b| ≤ 1° are analyzed to give a map of the molecular cloud distribution in the galaxy as viewed from the galactic pole. From the fact that this distribution shows no obvious spiral pattern we conclude that the giant molecular clouds sampled in the CO line are situated in both arm and interarm regions and they must last more than 108 years. A similar age estimate is deduced from the large mass fraction of H2 in the interstellar medium in the interior of the galaxy. An implication of this longevity is that the great masses of these clouds may be accumulated through cloud-cloud collisions of originally smaller clouds.

I want to make three points about star formation in spiral galaxies that follow from consideration of the internal structure of giant molecular cloud complexes (GMCC). The first point comes from pressure considerations. The total pressure inside the star-forming core of a GMCC may be written 106k)v/3kms−1)4(17pc/D)2 for virial theorem line width v and cloud diameter D; the pressure from a spiral density wave shock (SDWS) is 105 k(ns/1cm−3)(vs/20kms−1)2 and the thermal pressure in the cloud is 104 k(n/103cm−3) (T/10K) for Boltzmann constant k. These three pressures differ by factors of 10. An SDWS has too low a pressure to affect a cloud core; the only way an SDWS could influence a GMCC is if it interacted with the low thermal pressure in the cloud, i.e., the SDWS could propagate into a cloud along the direction of a magnetic field which may be the source of large scale pressure in a transverse dimension. The second point is that the density and mass of a GMCC are so large that the cloud will enter an SDWS like a cannon ball and will not be readily deflected. GMCC in other galaxies would then look like spurs on the spiral pattern and not like dust lanes. The alternative to these two points is that an as yet undiscovered (or uncommon) population of low density (100cm−3) clouds exists involving GMCC-type masses, or that smaller clouds coalesce at the SDWS. This implies that the star-forming clouds studied by molecular observers would be post-SDWS and post-gravitational collapse objects. Finally, the maximum age of a GMCC in the solar neighborhood is probably less than 50 million years. Its destruction is a result of pressure forces from the stars which it creates. Destruction in this sense does not necessarily imply that the molecules are converted into atoms – only that the cloud is pushed around. In the solar neighborhood, some clouds may, in fact, turn into 21-cm features; e.g., an HI half shell with a radius of 100 pc and a visual extinction through the shell of 0.2 mag. contains 3×105 M⊙, the mass of a GMCC. However, in the 5-kpc ring of the Galaxy, there is too much H2 relative to HI to allow any cycling between H2 and HI that is in phase with an SDWS unless the cloud remains molecular for 80% of the cycle. More likely, the cloud will be “destroyed” before that time. The implication is that cloud destruction at 5 kpc must produce molecular shells in addition to some atomic shells. This could be observed.

Galactic CO emission surveys such as those carried out by Gordon and Burton (Ap. J., 208, 346, 1976) and references cited therein are sensitive only to CO with J = 1 → 0 excitation temperature ≳ 5 k “Cold” molecular (H2) gas with CO excitation temperature ≳ 5 k would go undetected in such surveys. Low Tex could be due to either low kinetic temperature or low H2 density. Evidence for the possible existence of substantial amounts of cold gas (Tex ∼ 3 K) towards the galactic center is apparent in spectra of Liszt et al. (Ap. J. 213, 38, 1977 and 198, 537, 1975) in the form of narrow gaps, suggestive of self-absorption, at velocities of 0, −30 and −55 km/s with respect to the LSR.

The interstellar HI is best discussed in temperature categories because the requirement of pressure equilibrium leads to gross differences between hot and cold gas. The temperature of the hot (6000K) neutral gas has been measured by five techniques while the cold (60K) gas is visible mainly in 21cm absorbtion. There is evidence for warm (200K to 1000K) gas. Because of emission line blending, the small scale structural morphology is undetermined; however, there is evidence it is filamentary. With few exceptions, spatial sizes observe a lower limit of about 5 pc. On a larger scale, the gas is organized into sheets which reflect the recent history of cloud formation.

The distribution of the column density of HI over the sky shows strongly organized features of large angular extent. In particular, there is (1) a conspicuous set of arching concentric filaments extending more than 80° on either side of ℓ = 330°, and (2) a prominent extended region of low column density above ℓ ∼ +40° and centered on ℓ ∼ 130°. The changes is column density in the b-direction in this region of low NH is almost step-like at ℓ ∼ +40°.

I consider an ordinary shell to be one which can be produced by ordinary means such as a supernova. Good examples are the HI shells associated with the prominent radio loops, discussed today by Prof. Weaver, and the large HI shell in Eridanus (Heiles, 1976). Figure 1 shows the shell which is associated with radio Loop I, made from a consolidation of data from Martha Cleary of Australia, Raul Colomb and Wolfgang Poppel of Argentina, and myself. It changes size with velocity in the manner expected for an expanding shell. The distance is only 100 pc or so. Similar shells should be easily resolvable in the galactic plane, even with a modest telescope, because differential galactic rotation eliminates the confusing effects of foreground and background gas.

Recent abundance determinations of galactic H II regions and planetary nebulae are reviewed. The presence of O/H and N/H abundance gradients is well established; there is observational evidence indicating the presence of N/S, He/H and C/H abundance gradients. Some implications of these results are discussed.