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\title{The Galactic Center Environment\footnote{Paper to appear in
  ARA\&A 1996, Vol. 34 - this Preprint was retrieved electronically
  through GCNEWS and does not include the original figures.}}

\vskip 1cm
\author{Mark Morris\\
University of California at Los Angeles, Los Angeles, California 90095-1562
\\ \\and\\ \\ Eugene Serabyn\\
California Institute of Technology, Pasadena, California 91125}

\begin{document}
\maketitle

\begin{abstract}

The central half kiloparsec region of our Galaxy harbors a variety of phenomena
unique to the central environment. This review discusses the observed structure
and activity of the interstellar medium in this region in terms of its
inevitable inflow toward the center of the Galactic gravitational potential
well. A number of dissipative processes lead to a strong concentration of gas
into a ``Central Molecular Zone'' of about 200-pc radius, in which the molecular
medium is characterized by large densities, large velocity dispersions, high
temperatures, and apparently strong magnetic fields.  The physical state of the
gas and the resultant star formation processes occuring in this environment are
therefore quite unlike those occuring in the large-scale disk. Gas not consumed
by star formation either enters a hot x-ray emitting halo and is lost as a
thermally-driven galactic wind, or continues moving inward, probably
discontinuously, through the domain of the few parsec-sized circumnuclear disk,
and eventually into the central parsec. There, the central radio source SgrA$^*$
currently accepts only a tiny fraction of the inflowing material, likely as a 
result of a limit cycle wherein the continual inflow of matter provokes star
formation, which in turn can temporarily halt the inflow via mass-outflow winds.
\end{abstract}

%\vskip 1cm
\noindent
\section{Introduction}
   
In its most inclusive sense, the Galactic center encompasses a wide variety of 
phenomena occuring on stellar to galactic scales.  Because of this, and because 
research in this area has advanced rapidly for almost three decades, volumes
could be written to summarize our current state of knowledge of this region. 
Indeed, four conference volumes dedicated largely to this subject can now be
consulted (Riegler \& Blandford 1982, Backer 1987, Morris 1989, Genzel \& Harris
1994), and the pages of this review have already seen three major summaries on
Galactic center research: those of Oort (1977), Brown \& Liszt (1984), and
Genzel \& Townes (1987). In related contributions, Frogel
(1988) discusses the stellar content of the Galactic bulge, and Reid (1993)
summarizes research on the distance to the Galactic center; here, we adopt his
suggested distance of 8.0 $\pm$ 0.5 kpc. The most recent reviews on the Galactic
center are those by Blitz et al (1993) and Genzel et al (1994). 
  
Given this plethora of both information and publications, we have chosen to
concentrate this review on the interstellar environment, and in particular, on
the causes and ramifications of radial mass flow toward the center.  This
encompasses the nature of the central molecular layer, the hot medium 
surrounding it, and the
strong magnetic field in which they are found (\S 2), the central stellar bar
and its effect on the gas (\S 3), star formation (\S 4), the nature and fate of
the circumnuclear gas disk (\S 5), and the central nuclear radio source,
Sagittarius A$^*$ (\S 6). In the final section we address the possibility of
recurrent, highly energetic activity in the Galactic nucleus. Timely topics
notably missing from this review include the compact sources of
high-energy radiation near the Galactic center, large scale jets, stellar
populations and distributions within the Galactic bulge, stellar kinematics,
collisions and mass segregation in the central stellar core, and determinations
of the radial mass distribution beyond the central few parsecs.  Recent 
information on these subjects can be found in previous reviews, and in Sofue et al 
(1989), Sofue (1994), Lee (1994, 1995) and Blitz (1996).  
 
\section{Contents of the Interstellar Medium}

\subsection{The Arena}

The emissivity of our Galaxy's core reflects a unique environment. Line emission
from both molecular CO and atomic C peak sharply in our Galaxy's central few
hundred parsecs, and a somewhat smaller
region is quite bright in radio and infrared continuum emission as well (Fig. 1;
Altenhoff et al 1978, Mezger \& Pauls 1979, Odenwald \& Fazio 1984, Dame et al 
1987, Handa et al 1987, Cox \& Laureijs 1989, Bennett et al 1994). This compact 
and luminous nuclear region, hereafter
designated the ``central molecular zone", or CMZ (to distinguish this largely
molecular region from the more extensive ``HI nuclear disk" in which it is
ensconced), produces $\sim$ 5 - 10\% of our Galaxy's infrared and Lyman
continuum luminosities, and accounts for roughly 10\% of our Galaxy's molecular
gas content (Mezger 1978, Mezger and Pauls 1979, Hauser et al 1984, Scoville \&
Sanders 1987, Cox \& Laureijs 1989, G\"usten 1989, Scoville \& Good 1989,
Bloemen et al 1990, Wright et al 1991, Bennett et al 1994). The 
Galactic center is obscured at optical and UV wavelengths by a line-of-sight
interstellar extinction of roughly 30 visual magnitudes, but becomes accessible
again at energies above $\sim$ 1 keV, where X-ray line emission again peaks 
strongly on the Galactic center (Koyama et al 1989, Yamauchi \& Koyama 1993). On
the other hand, the CMZ does not stand out in the continuum at high energies
(Blitz et al 1985, Bloemen et al 1986, Cook et al 1991, Yamauchi \& Koyama 1993,
Diehl 1994, Goldwurm et al 1994). The focus of this review is on processes
occuring in this bright central region of our Galaxy. In this section we outline
the unique properties of the ISM in the CMZ, and examine the links between its
different components. 

\begin{figure}
%\vskip 13.5cm
\caption{Our Galaxy's Central Molecular Zone. Top: integrated intensity of \twelco, 
J = 1-0 emission, from the AT\&T Bell Labs survey (Uchida et al 1996). Bottom:
60-$\mu$m emission map, showing the distribution of relatively warm
dust, as measured by IRAS.}
\end{figure}

\subsection{The Central Molecular Zone}
 
The structure of our Galaxy's CMZ is best delineated by high resolution
observations of the CO molecule (Bania 1977, Liszt \& Burton 1978, Heiligman
1987, Bally et al 1987, 1988, Oka et al 1996). At galactocentric radii in excess
of a few hundred pc, the emissivity of CO is low (Scoville \& Sanders 1987), 
reflecting an average molecular surface density $\sim$ 5 M$_\odot$ pc$^{-2}$ in
the ``HI nuclear disk'' (Burton 1992, Burton \& Liszt 1993, Boyce \& Cohen 1994,
Liszt 1995, and references in these), a mixed molecular/atomic (and markedly
tilted) layer which occupies the region inside of our Galaxy's ``4-kpc molecular
ring''. However, at a radius of roughly 200 pc, a transition occurs to a largely
molecular, high density (n $\wig>$ 10$^4$ cm$^{-3}$), high volume filling factor (f
$\wig>$ 0.1) medium containing  5 - 10 $\times$ 10$^7$ M$_\odot$ of gas
(Armstrong \& Barrett 1985, Walmsley et al 1986, Bally
et al 1987, 1988, G\"usten 1989, Stark 1989, Tsuboi et al 1989).  Such densities
are usually found only in molecular cloud cores, with small net filling
factor, but in the CMZ, stable clouds require such densities to withstand tidal
shearing (G\"usten \& Downes 1980). However, while high densities are required,
the high surface density (several hundred M$_\odot$ pc$^{-2}$) and total mass
content are not, suggesting mass concentration due to inflow from larger radii. 

The clouds in the CMZ also show significantly elevated temperatures (30 - 200 K,
typically $\sim$ 70 K, G\"usten et al 1981, Morris et al 1983, Armstrong \&
Barrett 1985, Mauersberger et al 1986, H\"uttemeister et al 1993). Such gas
temperatures are high in comparison both to outer Galaxy clouds and to
coextensive dust, calling for a direct and widespread gas heating mechanism.
Highly supersonic internal velocity dispersions ($\sim$ 15 - 50 km s$^{-1}$),
comparable to the inter-cloud velocity dispersion, are also the rule. 
The elevated temperatures and linewidths may be linked, as
dissipation of turbulent energy is a prime candidate for direct heating of the
gas (Wilson et al 1982). This remains true if the cloud velocity dispersions are
due to magnetosonic disturbances (\S 2.5). Energy extraction from the tidal field
had been proposed as the origin of the large linewidths (Fleck 1980); Das \& Jog
(1995) conclude that this mechanism is efficient only for a central bar
potential, in which tidal fields vary markedly along elongated orbits. Viscous
magnetic heating resulting from the drift of ions and grains through the neutral 
gas (ambipolar diffusion) also warrants strong
consideration in this highly magnetized environment (Roberge et al 1995). 

The distribution and kinematics of the gas in the CMZ are consistent neither
with axial symmetry nor with uniform circular rotation (Bania 1977, Liszt \& 
Burton 1978, Morris et al 1983, Heiligman 1987, Bally et al 1988, Jackson et al
1996). The distribution is such that roughly 3/4 of the dense molecular gas is
located at positive longitudes, 3/4 is at positive velocities, and large radial
and vertical motions are present in a significant fraction (30\%) of the gas
(Bally et al 1988). On the basis of molecular kinematics, it is possible to 
divide the CMZ gas into two
components. The first is a high velocity (130 - 200 \kms), quasi-continuous ring
structure surrounding the nucleus and having a radius ($\sim$ 180 pc) suggesting
a location at the outer boundary of the CMZ. The kinematics of this ring (hereafter
referred to as the ``180-pc molecular ring''), and its
tilt relative to the Galactic plane, are continuous with those of the exterior HI
nuclear disk, indicating that this structure likely marks the location of an
HI/H$_2$ transition (Binney et al 1991). 

Inside of this ``boundary'' lies the second, mass-dominant molecular component:
a lower-velocity (typically $\wig<$ 100 km s$^{-1}$) population of dense and
massive molecular clouds. This cloud population (which includes the
familiar clouds associated with the Sgr A - Sgr E HII regions) lies quite close
to the true Galactic plane, and will be referred to as the Galactic center
``disk population" (Bally et al 1988, Heiligman 1987). This disk population
contains filament-like clouds with coherent velocity gradients over scales of
30-100 pc, suggestive of dust lanes and tidally stretched arcs or arms of gas
(Stark and Bania 1986, Serabyn \& G\"usten 1987, Bally et al 1988, Sofue 1995a, 
b). Several of these extended clouds contain a large fraction 
of the total molecular mass (Bally et al 1988, G\"usten 1989). Sgr B2, the most massive
molecular cloud in the Galaxy, contains about 5\% of the gaseous mass present in
the CMZ, while a handful of large clouds contribute up to 1/3 of the CMZ's CO
line flux. Due to our vantage point in the disk of the Galaxy, the
true geometric arrangement of these clouds and structures remains ambiguous,
but a framework for interpretation appears possible in the context of gas flows
in a barred potential (\S 3). 

One of these massive molecular clouds, the ``50 \kms cloud'', is adjacent to the 
Galactic center in projection.  It is associated with the bright, central Sgr A 
complex (Ekers et al 1983, Yusef-Zadeh \& Morris 1987a, Pedlar et al 1989), which
itself consists of the Sagittarius A West HII region, centered more or less on 
the central stellar cluster, and the somewhat larger-scale 
nonthermal shell source Sgr A East, which is offset from the center, 
but which encloses Sgr A West in projection (\S 5 provides further detail), and 
finally, the compact, central radio source Sgr A$^*$(\S 6).

The distributions of both ionized gas and diffuse infrared emission
in the CMZ  (Altenhoff et al 1978, Odenwald \& Fazio 1984, Handa et al 1987, 
Cox \& Laureijs 1989) are more symmetric than the CO distribution (Fig.\ 1), 
suggesting heating
sources more evenly distributed than the discrete clouds. In particular, the
dust temperature drops regularly as $r^{-0.3}$ from its central peak to its
asymptotic value of 23 K near a radius of 200 pc. The Lyman continuum production
rate and FIR luminosity of the CMZ are $\sim 1$ - $3 \times 10^{52}$ s$^{-1}$
and 10$^9$ L$_\odot$, respectively, and infrared excesses of roughly 10 and 30 
are deduced for individual HII regions and the extended emission, 
respectively (Mezger \& Pauls 1979, Reich et al 1987, Cox \& Laureijs 1989). 
The former is typical of Galactic H II regions, while the latter requires
a later population of stars. Star formation with an underabundance of O stars
had been considered (Odenwald \& Fazio 1984), but the abundant population of K
and M giants in the Galactic center suffices to heat the dust in the extended
region to the observed levels (Cox \& Laureijs 1989) without adjusting the O/B 
ratio. A soft interstellar
radiation field also receives support from the underabundance of C$^+$ emission
from the CMZ (Nakagawa et al 1996). 


\subsection{The Hot Component} 

Extended X-ray emission centered on the Galactic nucleus has also been detected,
with a size (1.8\deg\ x 1.0\deg) roughly half that of the CMZ. This region
stands out most clearly in the intense 6.7-keV K$\alpha$ transition of
He-like Fe (Koyama et al 1989, Yamauchi et al 1990), although it is also
seen in the continuum (Kawai et al 1988, Sunyaev et al 1993, Markevitch et al
1993), and most recently, in the 6.4-keV K$\alpha$ transition of neutral Fe
(Koyama 1995, Koyama et al 1996). The continuum emission drops rapidly with
frequency, and this central region is not seen at energies above 35 keV
(Goldwurm et al 1994). Skinner (1989) provides a complete review of earlier
X-ray observations; here we concentrate on recent results from the ASCA, Ginga,
and GRANAT spacecraft. 

The detection of the 6.7-keV line from He-like Fe clearly reveals the
existence of a high temperature plasma in the central hundred or so pc.
Estimates for its temperature vary, from 10-15 keV (Koyama et al 1989,
Yamauchi et al 1990, Nottingham et al 1993) down to 1-3 keV (Markevitch et al
1993). The latter estimate is lower because of the finding that the
X-ray emission changes character for energies above $\sim$ 10 keV, the lower
energy emission showing a roughly elliptical distribution, but the higher energy
emission essentially mimicking the CMZ's flattened distribution. 
This, and the image of the neutral Fe K$\alpha$ transition, which shows a clear
correlation with
the dense clouds (Koyama 1995, Koyama et al 1996), strongly suggest that
scattering of X-rays by the high-column-density medium in the CMZ is of
importance, both in terms of the 6.4-keV line, and of the continuum above 10 keV.
X-ray data
thus also provide information on the cold molecular medium, and by inference, on
the discrete X-ray sources which illuminate it
(Sunyaev et al 1993, Koyama et al 1996). 

Although estimates vary, the hot plasma has a temperature of some 10$^7$ -
10$^8$ K. The upper end of this range is quite high, more akin to the
temperatures found for intergalactic gas bound to clusters of galaxies than to
the temperatures found for individual galactic nuclei or supernova remnants
(Holt \& McRay 1982, Pietsch 1994, Koyama et al 1996). At temperatures near
10$^8$ K, the gas would not be bound to the Galaxy at all (Yamauchi et al
1990, Sunyaev et al 1993), and a wind would expand from the center at a few
thousand \kms, implying a lifetime of $\sim$ 10$^5$ yr.  The energy and mass
requirements would then be severe, with the plasma containing some 10$^{53}$
ergs and $\sim$3000 M$_{\odot}$ (Yamauchi et al 1990, Koyama et al 1996), at 
an average density of
0.3-0.4 cm$^{-3}$. To account for the extreme inferred temperatures, these
authors propose an energetic explosion in the Galactic
center at a past epoch. However, the necessary prior luminosity, $10^{41-42}$
erg s$^{-1}$, very dramatically exceeds anything observed today (\S 6). On the
other hand, the lower temperatures implied by removal of the high-energy
scattering component from the temperature determination yields plasma
temperatures more in line with supernova remnant temperatures, but the
emission measure then requires a rather large number of supernovae ($\sim$
1000). It has thus been difficult to settle on a single scenario for this plasma
which does not involve past releases of energy in the Galactic center far
exceeding the current rate. Perhaps the simplest explanation would be a past
starburst, with the ensuing supernovae accounting for the high temperature
plasma. 

Scattering of high-energy photons also provides a tracer for past activity through
the source-scatterer time delay; Sunyaev et al (1993) conclude from the 
observed distribution 
of high-energy X-rays that the X-ray luminosity of the
central $\sim$100 pc has not exceeded the current luminosity of $\sim10^{37}$ 
erg s$^{-1}$ by more than a factor of 10-100 in the past several
hundred years, while scattering in the K$\alpha$ line of
Fe$^{\rm o}$ leads Koyama et al (1996) to infer that a single source must 
have brightened
to $> 10^{39}$ erg s$^{-1}$ in roughly this same interval. The 1.8 MeV line of 
$^{26}$Al, which traces past massive star formation on a few-million year
timescale, shows a peak near, but offset from, the Galactic Center (Diehl et al
1995), and so may be relevant in tracing past activity in the cental region 
also. However, a coherent picture of the recent past of our Galactic nucleus has 
not yet emerged. 

\subsection{The Galactic Center Magnetosphere}

The magnetic field within the CMZ probably plays a significant role in the
physical interactions occuring there.  Its presence is revealed in several 
ways, the most striking
of which is the observation of 7 or 8 systems of nonthermal radio filaments (NTF's)
within $\sim$150 pc of the Galactic center (Yusef-Zadeh 1989, Morris 1990, 1994,
1995).  These polarized structures are tens of
parsecs long, and only a fraction of a parsec wide.  They may occur in isolation
(in which case they have been called ``threads'', Morris \& Yusef-Zadeh 1985) or
in bundles, such as those comprising the linear portion of the prominent Radio 
Arc (Yusef-Zadeh et al 1984). 
Morphologically, the NTF's are strikingly uniform in brightness and curvature 
(Fig.\ 2), and therefore quite different from the meandering, contorted  
filamentary structures one finds in supernova remnants and emission-line nebulae.
The morphology of the NTF's suggests that they reflect the
local magnetic field direction, and radio polarization evidence supports that
notion.  Tsuboi et al (1986) corrected the observed polarization angles
within the NTF's of the Radio Arc for Faraday rotation, and deduced that
the intrinsic field is parallel to the filaments, a result confirmed by Reich
(1994) and Tsuboi et al (1995) at high frequencies (32 \& 45 GHz, 
respectively), where Faraday rotation is small. 

\begin{figure} 
%\vskip 6cm 
%\picture 61.3mm by 31.4mm (threads)
\caption{Radiograph of 20-cm emission arising from the Northern and Southern
Threads, made with the VLA (Echevarria \& Morris 1996).  The vertical scale is 16.5
arcminutes, or $\sim$40 pc. The Sgr A complex is visible at 
bottom center.  The radio Arc (upper left) appears weak because it is beyond the telescope's
primary beam.}
\end{figure} 

A clue to the strength of the magnetic field in the NTF's is provided by the
near-absence of deformation or bending along their lengths. Every NTF
which has been sufficiently well studied has been found to be associated with,
and is probably interacting with, at least one molecular cloud (detailed below).
 However, the magnetic filaments are not subject to large distortions at the
interaction sites, in spite of the large velocity dispersion within Galactic
center molecular clouds, and in spite of the likelihood that, given the large
inter-cloud velocity dispersion at the Galactic center, most clouds have a
typical velocity of at least a few tens of \kms\ with respect to the ambient
magnetic field. By equating the apparent turbulent pressures within clouds (or
the ram pressure associated with presumed cloud motion relative to the field) to
the magnetic pressure, as a minimum condition on the strength of the magnetic
field, Yusef-Zadeh \& Morris (1987b,c, 1988) have determined that the magnetic
field within the NTF's has milligauss strength.  (The one known exception is
G359.1-0.2, the ``Snake" filament [Gray et al 1991, 1995], which is endowed with
a few ``kinks", suggesting that turbulence in its environment has partially
overcome the rigidity of the field. However, this filament is the one
with the greatest projected distance from the Galactic nucleus [125 pc], so it
may imply a declining field strength within filaments at larger distances.) 
This ``rigidity" method of constraining the magnetic field
strength in NTF's needs to be recast if the clouds consist of a number of
independently-moving clumps (\cf, Serabyn \& Morris 1994), because the field
lines can then be deflected around and between the clumps as the ensemble of
clumps moves through the field. A modified picture of this sort does not
dramatically change the constraint on the magnetic field strength, however: the
time scale to transport clump-induced distortions of the interclump field lines 
out of the cloud by Alfv\'en waves, $2\times10^4$ B$_{\rm mG}^{-1}$ (n/10
cm$^{-3}$)$^{0.5}$ (D/10 pc) yr for cloud size D and interclump density n, must
be less than a clump crossing time (typically 2 $\times$ 10$^4$ years), still implying
$\wig>$ mG field strengths. 

The fact that all of the known NTF's are perpendicular to the Galactic plane
to within about 20\deg\ (Morris \& Yusef-Zadeh 1985, Anantharamaiah et al 1991) 
suggests that these structures trace an 
ubiquitous dipole magnetic field occuring on a scale comparable to that of the
CMZ.  In this view, the observed filaments are the magnetic flux tubes which 
happen to be illuminated by the local injection of relativistic particles.  A
few alternative hypotheses have been considered, however.  One -- that the site
of each NTF or NTF bundle represents a locally generated field
enhancement -- encounters difficulties because of the strength of the magnetic
field in the NTF's. If the milligauss fields are present only in the
NTF's, then a strong confinement mechanism operating along their full
$\sim$30-pc length is needed to prevent them from expanding at
the Alfv\'en speed, and thus disappearing on a time scale of $\sim$10$^3$ years.
This is much faster than is needed to establish the currents necessary to
generate the structure in the first place (comparable to a dynamical crossing
time of $\wig>$ a few times 10$^4$ yr), and shorter than the energy-loss time
scale for relativistic electrons in the filaments (Anantharamaiah et al 1991,
Sofue et al 1992). Another alternative is that the NTF's lie upon the
surface of a cylindrical wall of compressed magnetic field surrounding the
Galactic center (Uchida et al 1985, Heyvaerts et al 1988), and that they are again 
illuminated by some local mechanism for the generation of relativistic
particles.   Such a geometry could result from expansive gas motions from the
center (Umemura et al 1988), if the gas momentum were sufficiently large,
although the expansive events would have to be continuous, quite recent, or even
in progress (such as would be manifested by the hypothetical, expanding
molecular shell described in \S3.1).  This geometry cannot be ruled out,
although it would imply a remarkable pattern of currents which might be
difficult to maintain. 

A twisted poloidal magnetic field geometry has been considered for
the Galactic center, on the assumption that this would result from flux-freezing
of field lines into the differentially rotating disk gas (Uchida et al 1985,
Sofue et al 1987, Shibata \& Uchida 1987).  However, the evidence for such a
geometry is based on rotation measures, which are difficult to localize to the
Galactic center, and which in any case probe the presumably relatively weak,
line-of-sight magnetic field component.  The more direct indicators of
the dominant field geometry -- the NTF's -- pass through the gas layer
in the Galactic disk without showing any appreciable large-scale twist or
distortion, other than perhaps a slow divergence. A large-scale twist might be
hidden by a fortuitous projection of one or two filaments, but it is
unlikely that all 7 or 8 known NTF systems would be similarly projected.
The apparent immunity of the NTF's to the inertia of the disk gas
could be understood if the fields within molecular clouds are 
only weakly coupled to the external field (\ie, separated by magnetopauses).

\noindent{\sc mid \& far-infrared polarization.} The second most informative probe of
the magnetic field near the Galactic center has been the polarization of thermal
infrared emission from dust.  Various mechanisms have been proposed for
orienting the emitting dust grains in such a way that their spin axes are
preferentially aligned along the magnetic field (Hildebrand 1988, Lazarian 1994,
Roberge et al 1995). All potentially viable mechanisms imply the same 
relationship -- orthogonality -- between the polarization position angle and 
the projected magnetic
field direction. A map of the distribution of polarization position
angles is thus readily transformed into a map of the projection of the magnetic
field onto the plane of the sky, weighted by the volume emissivity of the dust 
along the line of sight. The dust emissivity is maximized in dense
clouds or at warm cloud surfaces, so in contrast to the NTF's, which depict 
the field geometry in the intercloud medium, far-IR
polarization tends to probe the magnetic field within dense clouds.  The gross
uncertainties about the specifics of the grain alignment mechanisms make 
quantitative links between the field
strength and the fractional polarization difficult, although the magnitude of 
field direction fluctuations can be indicative of the
field strength (Morris et al 1992, 1996a,b). 

Polarization measurements were first applied to the Galactic center by Aitken et al (1986, 
1991), who, observing at 10 $\mu$m, found that the magnetic field within the
northern arm of Sgr A West (\cf, \S5) is parallel to that arm, a result which
would arise naturally from the action of shear within this orbiting gas stream.
Shear is expected to affect much of the gas present within the strong tidal
field of the Galactic center, so this result may have widespread applicability. 
Strong far-IR polarization has now been mapped in four clouds near the Galactic
center, as reviewed by Davidson (1996). In essentially all of the
clouds, the inferred magnetic field direction is predominantly parallel to the
Galactic plane (Sgr B2 being a complex mix of absorption and emission which is
still ambiguous in this regard), which stands in marked contrast to the
perpendicular field inferred for the intercloud medium. 
Figure 3 shows the 60-$\mu$m polarization vectors toward G0.18-0.04 (Morris et al 1996b)
superimposed upon a 6-cm radiograph (Yusef-Zadeh \& Morris 1987b) and
the underlying cloud (Serabyn \& G\"usten 1991).  
The implied field direction within the cloud is parallel to the ridge of molecular
and thermal
radio emission, which is largely parallel to the Galactic plane. 
The implication that the magnetic field follows the 
ridge of the molecular cloud in this and in the nearby ``arched filament" cloud
suggests that
the field within clouds is shaped by the stresses which shape the clouds
themselves (Serabyn \& G\"usten 1987, Morris et al 1992, 1995). Furthermore, the
general othogonality of fields inside and outside of clouds suggests that the two
are relatively independent, consistent with the apparent immunity of NTF's to the 
inertia of disk gas.

\begin{figure}
%\vskip 12cm
\caption{100-$\mu$m polarization measurements across the HII region G0.18-0.04 
(from Morris et al 1996b), superimposed upon a $\lambda$6-cm continuum radiograph
(from Yusef-Zadeh \& Morris 1987b). 
The orientations of the line segments represent the position angles of
the electric vectors, while their length is proportional to the percent
polarization.  A horizontal segment representing 4\% polarization is shown at
bottom left.  The position angle of the Galactic plane is $\sim$30\deg, so the 
polarization vectors are all close to being perpendicular to the Galactic plane.  The
superimposed contours show the intensity of CS J = 3-2 emission, from Serabyn \&
G\"usten (1991).} 
\end{figure}

The shaping of the fields within clouds by the stresses which shape clouds 
themselves appears to apply to the circumnuclear disk (CND) as well (\S 5, Werner et
al 1988, Hildebrand et al 1990, 1993). With the exception of a presumably
infalling stream of dust and gas which coincides with the Northern Arm of Sgr A
West, all of the far-IR emission from dust in the CND is polarized in a manner
consistent with a dominant magnetic field component parallel to the
plane of the disk.  Self-similar models for a poloidal field 
well outside the disk (\ie, perpendicular to the disk plane) which has
been deformed and sheared within the CND by a combination of
differential rotation and radial infall, were proposed by Wardle \& K\"onigl
(1990). Variations of these models have generally been rather successful in
accounting for the polarimetric mapping results on this source (Hildebrand et al
1993). Another model based on numerical magnetohydrodynamic calculations
carried out by Meglicki et al (1994) is also able to reproduce the gross
geometric and magnetic characteristics of the CND, although this model disk is
undergoing rapid evolution, with inward radial motions ($\sim$50 \kms) well in
excess of the observational limits ($\sim$20 \kms, \cf, \S5). 

\noindent{\sc the zeeman effect} has been used to probe the magnetic field in
Galactic center clouds directly, although the results are mixed. The large
velocity breadth of the lines from essentially all Galactic center clouds
have made Zeeman measurements difficult for all but the
strongest fields.   Using HI absorption, Schwarz and Lasenby (1990) derived a
modest line-of-sight field, B$_{||}$ = 0.5 mG toward the northern
side of the CND.  With 1667-MHz OH absorption, Killeen et al
(1992) measured B$_{||}$ = -2.0 mG, averaged separately
over broad portions of the northern and southern sides of the CND, with
considerable spatial variation in B$_{||}$
on scales of 5 to 10". Most recently, Plante et al (1995) used
HI absorption to detect fields ranging up to about -3 mG on the northern side of
the CND.  They attribute the strongest fields to a feature which they interpret
as a stream of infalling gas rather than a part of the disk. These results, all 
with the VLA, 
and the negative HI Zeeman results of Marshall et al (1995), who obtained
B$_{||}~<$ 0.5 mG in 45x45 arsec$^2$ regions of the CND, suggest that the
average line-of-sight component of the field within the CND is $\sim$1 mG,
but the likelihood that B$_{||}$ changes considerably and even undergoes
reversals along the line of sight (for example, in the model of Wardle \&
K\"onigl 1990) implies that the Zeeman signals are diminished by projection
effects. 

Elsewhere in the Galactic center, Zeeman measurements have yielded negative
results. Uchida \& G\"usten (1995) observed OH absorption in a number of clouds
with a relatively large beam, and obtained typical upper limits to
B$_{_{||}}$ of a few tenths of a milliGauss. While their observations are
subject to dilution of the Zeeman signal by averaging over spatial variations
within the beam, as well as along the line of sight, these results raise
questions about the conclusion that milligauss fields pervade the Galactic
center region, unless either the line of sight component is small everywhere
(\ie, a field which is everywhere within $\sim$ 10\deg\ from the vertical to the
Galactic plane), or the field has strong spatial variations. 

\noindent{\sc faraday rotation measures} provide yet another probe of the field
surrounding nonthermally-emitting Galactic center radio structures.  Vast areas
of the Galactic center display polarized emission (Sofue et al 1987, Haynes et
al 1992), so in principle the magnetic field within a large volume of the
intervening medium can be studied, although this probe has not yet been
extensively utilized.  The
rotation measures (RM's) can be quite large at some sites within the Galactic
center. Using 6-cm radio observations, Yusef-Zadeh \& Morris (1987c, 1988) 
report RM's up to 5500 radians
m$^{-2}$ toward the Radio Arc, with somewhat smaller values
(up to 1450 radians m$^{-2}$), at higher Galactic latitudes (see also Tsuboi et al
1995).  In a
2-cm study of the Radio Arc (Inoue et al 1989), complex structure was
seen in the foreground Faraday screen, even to the extent that depolarizing
filaments were identified, but with no obvious counterpart in total
intensity images.  When this kind of study is done more thoroughly and
quantitatively, with independent measures of electron density, it will
provide much-needed information on the field strength and geometry of
magnetic structures. 

\noindent{\sc the production of radio filaments.} Several
hypotheses about the origin of the Galactic center NTF's have been offered
(Chudnovsky et al 1986, Heyvaerts et al 1988, Benford 1988,
Morris \& Yusef-Zadeh 1989, Serabyn \& G\"usten 1991, Lesch \& Reich 1992, Rosso
\& Pelletier 1993, Serabyn \& Morris 1994), the merits and drawbacks of which
are summarized by Morris (1995). Here, we describe only the
hypothesis which, in our judgement, conforms best to the observational evidence. 
The key to this hypothesis is that all sufficiently well studied NTF's
appear to be interacting with molecular structures. It was originally suggested
following examination of molecular CS emission from the cloud underlying the
G0.18-0.04 HII region (Serabyn \& G\"usten 1991, Serabyn \& Morris 1994, see
also Morris 1995, Serabyn 1995). Yusef-Zadeh \& Morris (1987b) had pointed out
that most of the nonthermal filaments of the Galactic Center Radio Arc undergo
abrupt discontinuitites where they encounter the ionized ridge of G0.18-0.04,
and Serabyn and G\"usten (1991) demonstrated that this HII region lies at the
surface of a relatively massive molecular cloud.  With a partial
interferometric view of the molecular cloud, we found that dense clumps are
present at a number of the locations where the NTF's undergo
brightness discontinuities, or small deflections (Serabyn \& Morris 1994). This
led us to propose that the NTF's in the radio Arc, and perhaps all Galactic
center threads, originate at the surfaces of such dense molecular clumps.  The
hypothesis can be stated as follows: when certain conditions are satisfied by a
molecular cloud moving through a more diffuse ambient medium, magnetic field line
reconnection taking place at leading clump surfaces leads to the acceleration of
charged particles, and the resultant relativistic particles then stream away from
their point of origin along the field lines.  The flux lines to which these 
particles are attached are then illuminated by synchrotron emission. 

The hypothesized conditions for particle acceleration are: i) the cloud surface
must be ionized, presumably by a fortuitously placed local source of ionizing
radiation. The ionized gas provides a source of free electrons for the
acceleration mechanism to act upon, and the turbulence associated with the
ionization front ensures that the cloud and intercloud magnetic fields are
sufficiently mixed for reconnection to be efficient. ii) The cloud surface at
which the acceleration takes place must be moving at a relatively large velocity
into the ambient intercloud medium, in order that the magnetic fields in the two
media are forced into contact.  iii) The orientations of the fields in the cloud
and intercloud media must be quite different for rapid energy extraction, a
condition which has very recently been verified in the case of G0.18-0.04, where
far-IR polarization measurements show that the internal cloud field is
perpendicular to the linear radio filaments with which the cloud is interacting
(Fig.\ 3).  The same may be true for a substantial
fraction of all Galactic center clouds, if the trends from far-IR polarimetry
are any guide. However, the conditions listed here are
collectively stringent enough to explain why not every molecular clump in the
Galactic center magnetosphere has an associated NTF. 

On the other hand, it is possible that these conditions are satisfied by every
filament or thread in the galactic center.  The NTF's
associated with Sgr C arise (in projection) at the interface
between that bright HII region and its associated molecular cloud (Liszt \&
Spiker 1995); the filament G359.5+0.18, located to the north of Sgr C, is
apparently interacting with at least one, and possibly two, molecular clouds
(Bally et al 1989, Staguhn et al 1995); the Northern thread (G0.08+0.15, Morris
\& Yusef-Zadeh 1985) undergoes an abrupt intensity discontinuity where it is
superimposed upon the thermal arched filament HII region and associated cloud
(Fig. 2); and finally, an HII region/molecular cloud complex is located at the 
northern tip of the ³Snake² (G359.1-0.2; Gray et al 1995, Uchida
et al 1996).  However, in none of these cases is a cause-and-effect relationship
between an NTF and a molecular cloud/HII region as convincing as the
Radio Arc case because high-resolution observations are lacking.


\noindent{\sc origin of the strong poloidal field.}
If the Galactic center magnetosphere is pervasive on a scale of $\sim$100 pc,
then its total energy content is $\sim 10^{54}\langle B(mG)\rangle$ ergs,
comparable to the energy content of the hot, x-ray emitting gas and only a few
times larger than the kinetic energy associated with non-circular motions in the
180-pc molecular ring (\S 3). The maintenance of such a strong, dynamically
important field requires a rather powerful ring current circulating about the
center, presumably somewhere within, or at the boundary of, the CMZ. The
maintenance of such a current, and the confinement of the field, might be
accomplished by the Lorentz forces accompanying the quasi-steady radial inflow
or outflow (or both) of gas through a vertical magnetic field (\eg, Lesch et al
1989), although the physical elements of such a dynamo mechanism still require
elaboration. 

One hypothesis for the vertical field at the Galactic center which is appealing
both for its simplicity and its seeming inevitability, is that of Sofue
\& Fujimoto (1987). Noting that Galactic evolution is characterized by the
inexorable radial inflow of matter, these authors hypothesized that the vertical
component of the early Galactic field is dragged inwards by the gas accreting to
the Galactic center region over the lifetime of the galaxy; its rate of outward
radial diffusion with respect to the gas is small compared to the rate of 
inward gas flow, so the vertical field becomes concentrated at the nucleus.
In contrast, the component of the field parallel to the disk can be transported
vertically out of the Galaxy by diffusion or by the Parker instability on a time
scale short compared to a Hubble time, so that the azimuthal field represents 
an equilibrium between the flux lost by vertical transport
and that regenerated by a Galactic dynamo. 

According to this hypothesis, the strength of the vertical field concentrated at
the nucleus reflects the strength of the primordial magnetic field.  However, it
also depends on the spectrum of spatial fluctuations, since the only
fluctuations which could survive to form a uniform field when concentrated at
the nucleus are those present initially on Galactic scales or larger
(Morris 1994).  Smaller-scale field fluctuations would have suffered
reconnection and annihilation in the Galactic center mix.  This, plus the
conversion of gravitational potential energy of the inflowing gas, may have been
important sources of heating for Galactic center gas throughout Galactic
history. 
 
\subsection{Pressure Balance}

The physical conditions present in the CMZ are consistent with
thermal, nonthermal, and magnetic pressures several orders of magnitude higher
than those present in the large-scale Galactic disk. For molecular densities of
10$^4$ cm$^{-3}$ at the inferred temperatures (\S 2.2), $P_{thermal} \sim
10^{-10}$ erg cm$^{-3}$. In the hot plasma, $P_{hot} \sim 4 \times 10^{-10}$ erg
cm$^{-3}$, close enough to consider pressure balance (Spergel \& Blitz 1992).
However, turbulent pressures in the clouds greatly exceed this value,
approaching $P_{turb} \sim 10^{-8}$ erg cm$^{-3}$.  For a field strength of 
0.1 - 1 mG, the magnetic pressure is comparably large: $ P_{mag} \sim 4
\times 10^{-10}$ - $4 \times10^{-8}$ erg cm$^{-3}$.  Internal cloud turbulence 
is therefore likely to be linked to magnetic pressures.  In particular,
the observed velocity dispersions in Galactic center clouds likely
reflect hydromagnetic waves (Arons \& Max 1975), which could be generated by the
angular momentum loss processes affecting Galactic center clouds (described in
\S3), including shocks, cloud collisions, and the irregular viscous interactions
suffered by clouds moving through a magnetized, low-density medium. 

\noindent
\section{The Central Bar, and Gas Dynamics near the Galactic Center}

The presence of a substantial stellar bar in our Galaxy, first invoked to 
account for the noncircular motions of HI near the Galactic center 
(de Vaucouleurs 1964, Peters 1975), is now well established by a variety of methods,
including techniques based on photometry (Blitz \& Spergel 1991b, Sellwood 1993,
Weiland et al 1994, Dwek et al 1995), kinematic studies of gas (Liszt \& Burton
1980, Mulder \& Liem 1986, Binney et al 1991, Wada et al 1994) and stars (Zhao
et al 1994, Blum 1995), and counts of luminous stars (Nakada et al 1991,
Weinberg 1992a,b, Whitelock \& Catchpole 1992, Stanek et al 1994). The picture
which has emerged is that the bar is quite pronounced, with axial ratios of
about 3 : 1 : 1 (though Binney et al 1991 and Wada et al 1994 prefer 1.5:1 for
the axial ratio in the Galactic plane), it extends at least to the co-rotation
radius of about 2.4 kpc, it has a total mass of 1 - 3 $\times$ 10$^{10}~\Msun$,
and its long axis is oriented at a modest angle with respect to our line of
sight (15 to 45$\deg$ toward positive longitude, depending on the model).
%choose not to include Calbet et al 1996, ApJL, 457,L27-30 because their data
%extend well beyond 2.4 kpc.  They differentiate between bar & bulge. 
The excess of microlensing events toward the Galactic bulge provides 
yet another indication that the
Galactic bulge is strongly barred, and that the bar is somewhat along our line of
sight (Paczy\'nski et al 1994, Han \& Gould 1995, Zhao et al 1995). 
Leaving aside the details of these determinations, we turn to the issue of the
interstellar material which co-exists with the stellar bar.  A strong bar is an
essential element for understanding gas dynamics near the
Galactic center, and especially for appreciating the potential importance of
bar-induced migration of interstellar gas toward the Galactic nucleus. 

%\vskip 0.5cm
\subsection{Response of Interstellar Gas to a Bar Potential}

The  response of orbiting gas to such a pronounced bar is quite strong, and
leads both to large deviations from circular motion and to strong shocks, so it is
essential when studying gas dynamics to properly account for the $m=2$
(quadrupole) deviation from axial symmetry represented by the bar (where the 
gravitational potential depends on
azimuthal angle, $\phi$, via the factor $e^{im\phi}$). The velocity field of the
gas provides one of the best ways of probing the shape of the Galactic
gravitational potential.  The primary data to be explained by any bar model is
the distribution of molecular emission (or absorption) throughout the 
longitude-latitude-velocity (l,b,v) cube, and
the most remarkable feature of this distribution in our Galaxy is the 
180-pc molecular ring (\eg, Bania 1977; Bally et al 1987), seen as a parallelogram, or
ellipse, in projections of this cube onto the l-v plane (Figure 4). 
Models for the response of orbiting gas to a symmetric m=2 bar potential
(described more fully below) have met with considerable success in accounting
qualitatively for the gas kinematics measured in both HI and CO surveys (Binney
et al 1991), notably including this parallelogram, and it is now becoming the
standard paradigm for Galactic center kinematics, although a quantitatively
faithful model is still being sought (\eg, Binney 1994; Jenkins \& Binney 1994).
Most importantly, existing models readily account for the abundance of gas 
near the Galactic center having ``forbidden'' radial
velocities (\ie, with sign opposite that of Galactic rotation).

\begin{figure}
\indent
%\vskip 6cm
\caption{\twelco\ J = 1-0 emission contours in a position-velocity plane in 
which position extends
along the line from l,b = 2.5\deg, -0.333\deg\ to -2.0\deg, 0.1333\deg,
from the AT\&T Bell Labs survey (Uchida et al 1996). This line, tilted by 
$\sim$6\deg\ with respect to the Galactic plane, was chosen to reflect the tilt
of the plane containing the 180-pc molecular ring (Uchida et al 1994a). The
horizontal axis shows only the dominant galactic longitude component of the position.  Contour units are 1.4, 2.8, 4.3, 6.0, 8.0, 10.3, 12.8 ... K.}
\end{figure}

\noindent{\sc orbits of gas clouds in the presence of a strong bar.}
Gas moving in response to a bar potential tends to settle into closed, elongated orbits
because cloud collisions and consequent energy dissipation act to enforce
conformity. Angular momentum loss by the orbiting
gas resulting from the processes described more fully in \S3.2 causes it to
drift inward along a family of nested, closed orbits.  When the gas is orbiting
at radii between that of corotation and the inner Lindblad resonance (ILR) of
the bar pattern, it moves on the so-called ``X$_1$'' orbits (Contopoulos \&
Mertzanides 1977): elongated oval orbits having their major axes aligned with
the bar (Fig. 5). 
   
\begin{figure}
\indent 
%%\vskip 6cm
\caption{Schematic diagram illustrating the shapes and relative orientations 
of the X$_1$ and X$_2$ orbits, and the locations of the shocks resulting 
from the interaction between the innermost X$_1$ orbit and the outermost 
X$_2$ orbit.  The arms in the CMZ hypothesized by Sofue (1995a) are also shown.}
\end{figure}

As the ILR is approached, however, there is an innermost stable X$_1$ orbit
inside of which these orbits become self-intersecting, or cusped. Any angular
momentum loss by gas in the innermost stable X$_1$ orbit leads to
orbit-crossings and shocks. The shocks, in turn, imply a further, abrupt loss of
angular momentum which causes the gas to fall inward until it settles onto a new
family of closed, elongated orbits lying deeper in the potential well: the
``X$_2$'' orbits. These oval orbits have their long axes oriented perpendicular
to the bar.  At their apocenter, the outermost  X$_2$ orbits are thought to
graze the pericenter of the innermost X$_1$ orbits, so angular momentum loss
by gas in the innermost X$_1$ orbit, for whatever reason, be it via viscous
transport of angular momentum outward or by orbit crossings near the cusps,
would lead it to collide with gas in the X$_2$ orbits, presumably creating a
spray which leads to a shock at the far side of the innermost X$_1$ orbit
(Binney et al 1991, Athanassoula 1992, Jenkins \& Binney 1994, Gerhard 1995). 

Binney et al (1991) hypothesize that the 180-pc molecular ring corresponds to
the innermost stable X$_1$ orbit.  The small radial extent of this feature is
responsible for the narrowness of the trace of emission defining the
parallelogram and for its well-defined, sharp vertices.  The shocks along the
inside edge of the X$_1$ orbits are held responsible for compressing the gas
into molecular form, thereby explaining why predominantly molecular gas occupies
only the innermost X$_1$ orbits.  Indeed, gas located outside the 180-pc molecular
ring is evident in HI data (Burton \& Liszt 1978), and the dynamics of this HI
nuclear disk can be modelled in terms of a hierarchy of X$_1$ bar orbits (Liszt
\& Burton 1980, Binney et al 1991). 
One potential problem with this interpretation of the
parallelogram is the velocities of the vertices 
(Uchida 1993, Binney 1994): the velocity of the negative-longitude,
intermediate-velocity vertex is near zero in the models, but is measured to be
large ($\sim$150 \kms) and in the forbidden quadrant (Fig. 4).

Recall that the 180-pc molecular ring surrounds a complex, asymmetric
concentration of disk population molecular clouds.  Binney et al (1991) propose
that these clouds are distributed on X$_2$ orbits. An attempt to identify some
spatial order in this complex distribution was recently made by Sofue (1995a, 
see also Scoville et al 1974), who defines two ``arms'' of molecular material,
seen as continuous streams in the (l,b,v) data cube of $^{13}$CO emission 
(fig. 3 of Liszt 1992 nicely shows the intersection of these streams
with the b = -3' plane), as the repository of most of the molecular mass
interior to the 180-pc ring. These hypothetical arms have small pitch angles,
and are roughly symmetrically placed about the galactic center at radii of
$\sim$120 pc.  Their proposed morphology is similar to the
innermost, spiral density enhancements resulting from shocks in a numerical
model of gas response to a bar (Gerhard 1995, Athanassoula 1992), and it is similar to the
molecular distribution seen in several nearby galaxies (Kenney et al 1992; Turner
et al 1993), so the molecular arms may have a natural exlanation
in the context of bar-induced kinematics.  These shocks result from the
transition from X$_1$ to X$_2$ orbits, as discussed above, and they should
presumably extend continuously out to the 
180-pc molecular ring, although l-v plots show no clear continuity between the arms
and the ring. 

%%\vskip 0.3cm
\noindent{\sc the emr versus the bar.}
Perceived in early molecular data as an ellipse in the l-v
plane, the 180-pc molecular ring was originally interpreted as a radially
expanding molecular ring surrounding the Galactic center, (Scoville 1972, Kaifu
et al 1972) and was subsequently referred to as the ``EMR''. It was hypothesized
to be the result of an explosive event at the center $\sim$10$^6$
years ago which gave a radial momentum impulse to $\sim$10$^7$ \Msun\ of nearby
gas (Bania 1977). This hypothesis requires a rather extreme energy ($>$10$^{55}$ 
ergs when dissipation, work against the gravitational potential and
the non-spherical geometry of the ring are taken into account; Sanders 1989, Saito 1990),
although such energies can, in principle, be generated by episodic accretion
events at the Galactic center (\eg, Bottema \& Sanders 1986).  Another challenge
for the EMR hypothesis has been the dearth of evidence that non-EMR
clouds in the CMZ have been affected by the passage of 
%the shock accompanying 
this putative flash flood of radially-moving material.
However, Uchida et al (1994a,b) recently pointed out that a molecular structure
associated with the far-infrared source AFGL 5376 (at l,b = 359.5,+0.43) is
apparently the site of a large-scale (100-pc) shock oriented perpendicular to
the Galactic plane.  They argue that it and another vertical complex at
l=1\clap{\deg}{.}2 could be the signatures of the interaction of the EMR with
the ambient medium. The projection of these shock structures out of the Galactic
plane is attributable in part to the tilt ($\sim$20\deg) of the EMR with respect
to the Galactic plane. Uchida et al (1994a) also cautioned against abandoning
the EMR hypothesis by noting that, when CO emission within 0\clap{\deg}{.}15 of
the Galactic plane is excluded from the average over latitudes, the
resultant l-v plot reveals an elliptical envelope, as forseen by the
original EMR hypotheses. 

It would be curious indeed if an energetic explosion at the Galactic center left
behind no kinematic signature outside of the single plane defined by the EMR.
In another examination of the latitudinal structure of the EMR, Sofue (1995b)
presents evidence that the EMR is really an expanding molecular shell, (EMS).
Using the $^{13}$CO survey data of Bally et al (1987), Sofue finds that the EMS
is an oblate spheroid which is pinched around its equatorial waist because the
expansion has suffered more impedance there by the higher density in
the midplane of the CMZ. He also concludes that the EMS has gaps in it along
radial rays which intersect the most prominent interstellar structures in the
CMZ, as would be expected if these structures were ponderous and thick enough to
block the passage of the EMS.  While this interpretation of the data is not
completely unequivocal, it definitely warrants further investigation. 

Which of the two competing hypotheses for the 180-pc molecular ring is 
more correct is not yet clear.  While the kinematic response of gas to the
Galactic bar should inevitably
produce a signature in the (l,b,v) cube looking roughly like the data, episodic
explosions at the nucleus might also be both in evidence and inevitable (\cf, \S
7). Thus, although it appears that the dominant characteristics of the gas
kinematics are naturally determined by the bar potential, the final analysis of
Galactic center gas dynamics may well require an occasional injection of radial
momentum. 

%%\vskip 0.5cm
\subsection{The Inward Transport of Gas to the Nucleus}

The mass of CMZ material said to be on X$_2$ orbits constitutes
85 - 90\% of the total molecular mass in the Galactic center arena, so the
residence time on these orbits is clearly much longer than in the innermost
X$_1$ orbit.  Nevertheless, this reservoir of gas can only be temporary. 
Angular momentum loss by orbiting disk gas is inevitable, in the face of the 
many processes acting near the Galactic center. Clouds on X$_2$
orbits circulate with velocities greater than the pattern speed of the bar
(estimated at 19 \kms\ kpc$^{-1}$ by Wada et al 1994); they thereby lose
angular momentum to the stellar bar by gravitational torques.
These clouds are also subject to tidal friction from stars in
the bulge (Stark et al 1991), and the most massive of them are doomed 
to spiral into the Galactic
center on time scales of a few times 10$^8$ years.  The magnetic field is
another contributor to angular momentum loss, particularly for the less massive
clouds; if clouds continually move through a strong (mG), pervasive vertical magnetic field
in the inner 100 pc or so, with a relative velocity equal to a substantial
fraction of the cloud's orbital velocity, then the
magnetic viscous force would cause angular momentum loss
on a time scale of $\sim$10$^8$ years (Morris 1995). 

In addition to 1) bar-induced torques, 2) dynamical friction and 3) magnetic
viscosity, several other angular momentum loss processes contribute to the
inexorable inward transport of matter: 4) shocks associated with the X$_1$-X$_2$
transition, 5) viscous drag in the differentially rotating Galactic disk,
including that resulting from cloud-cloud collisions (von
Linden et al 1993, Biermann et al 1993), and 6) dilution of the gaseous disk's
specific angular momentum by stellar mass loss material raining down out of 
the slowly-rotating
Galactic bulge (Jenkins \& Binney 1994).  While assessing
the relative importance of these mechanisms is rather difficult at present, 
several of them are individually quite important, so gas
brought into this arena is destined to migrate into the center on time scales
much smaller than a Hubble time. 
 
The rate of mass flow through the ILR, \Mdot$_{ILR}$, can be estimated using the
mass of molecular gas in the 180-pc ring, 8 $\times$ 10$^6$
\Msun\ (Bania 1977, Sofue 1995a), and, following Gerhard (1992), noting that the gas cannot
stay on this innermost (cusped) X$_1$ orbit for more than about one orbital
period, 2 $\times$ 10$^7$ years.  Thus, \Mdot$_{ILR}$ = $\sim$0.4 \Msun\
yr$^{-1}$.  This estimate, which is larger than that of Gerhard because of the
larger mass assumed for the 180-pc ring, is, in any case, quite uncertain, and
we regard a range of \Mdot$_{ILR}$ $\approx$ 0.1 - 1 \Msun\ yr$^{-1}$ as more
appropriate. This range can be compared to the 0.07 \Msun\ yr$^{-1}$ estimated
by Jenkins \& Binney (1994) to be the rate at which the bulge stars within 2 kpc
shed mass. The actual contribution from bulge stars is likely to be larger than
this, since this estimate is based on the assumption that the bulge is similar
to an elliptical galaxy, whereas many bulge stars are younger than this would
imply (Lindqvist et al 1992, Rich 1993, Rieke 1993). Of course, if it is 
to contribute to the mass budget of the
galactic center gas reservoir, the matter shed by bulge stars does indeed have
to migrate down to the Galactic plane without being lifted off by a bulge wind.
In addition, some fraction of the material moving inward at the ILR may come
from further out in the disk, as a result of some of the same angular momentum
loss processes, although the torque exerted by the bar
outside corotation has the opposite sign, and may partially counteract the other
loss mechanisms. In any case, there
appears to be no fundamental problem with finding enough material to maintain
the gas inflow at its present rate. 

The ratio of the total gas mass inside the 180-pc molecular ring, 4 - 9 $\times$
10$^7$ \Msun, to \Mdot$_{ILR}$ provides an estimate of 
0.4 - 1 $\times~10^9$ years for the mean residence time for gas in
X$_2$ orbits, if we assume a steady state inflow.  This range of time scales is
similar to that for dynamical friction to extract angular momentum from massive
clouds.  The gas moves inward from $\sim$150 pc at a mean rate of 0.2 - 2 \kms,
which is comparable to the estimate of 0.3 \kms\ for inward radial motion near
the solar circle in spiral structure models (Lacey \& Fall 1985), and to the
azimuthally averaged inflow velocities found in the bar models of Athanassoula
(1992). This inwardly migrating gas meets one of three fates: star formation
(0.3 - 0.6 \Msun\ yr$^{-1}$; G\"usten 1989, \S 4), a thermally driven Galactic
wind or fountain (0.03 - 0.1 \Msun\ yr$^{-1}$; \S2.3), and
accretion into and through the domain of the much smaller-scale circumnuclear
disk (0.01 - 0.1 \Msun\ yr$^{-1}$; \S5). While none of these routes is 
negligible, star formation evidently dominates.

In principle, the CMZ could undergo a global gravitational instability, creating
a ``bar within a bar,'' leading to yet further angular momentum loss, and
thus to a greatly enhanced inflow rate, as well as to a starburst (Shlosman et
al 1989).  However, the mass fraction of the gas inside the ILR, $\sim$5 - 10\%,
is apparently at or below the limiting value for such an instability.  We
therefore presume that the list of angular momentum loss processes given above
is complete, and that the inward flow of gas on 100-pc scales is relatively
steady, although the possibility that the Galactic center has undergone
substantial convulsions in the past must be kept in mind. 

%%\vskip 0.5cm
\subsection{Evidence for an m=1 Wave}
  
In addition to the $m=2$ bar mode, there is evidence that an $m=1$ wave 
with substantial amplitude is present near the Galactic center.  Such a wave
would reveal itself as a displacement of the central stellar cluster from the
centroid of the bar, of the bar from the center of mass of the Galaxy, and/or as
a global tilt of the inner disk of the Galaxy, and would have implications
for the stellar velocity field as well. The pronounced longitudinal asymmetry of
molecular line emission both from the Galactic center disk population (\cf, \S2.2) 
and from the molecular parallelogram (\S 3.1, Uchida et al 1994a, Blitz 1994)
is a long-standing clue that the
gas may be responding to an asymmetry in the Galactic potential. 
A pure $m=2$ bar would show some displacement from l $=$ 0\deg\ because of
projection effects, but the predicted offset is considerably less than that
observed.  Blitz (1994) finds that more distant gas -- HI at radii out to 750 pc
-- shares this rotation center offset from the assumed galactic center: Sgr
A$^*$ and its surrounding stellar cluster. 

A second clue for an m=1 wave in our Galaxy is the velocity displacement, 
with respect to the local standard of rest, of
30 or 40 \kms\ of the appropriate velocity centroids of gas in the 180-pc
molecular ring and the HI nuclear disk from the 0 \kms\ expected in an
axisymmetric galaxy (Blitz 1994).  A portion of this offset (15 \kms) can
perhaps be attributed to motion of the LSR in response to a large-scale triaxial
potential (\eg, Blitz \& Spergel 1991a), but the remainder suggests that the
entire central gas layer, and perhaps the bulge as well, are in motion,
presumably oscillatory, with respect to the Galactic disk. This is reminiscent
of observations of lopsided galaxies, \ie, barred spirals in which one spiral
arm is much longer than the other (Baldwin et al 1980). In these galaxies, the velocity
profiles near the center are invariably found to be highly anomalous, and
attempts to pinpoint the rotation center always find it at a large distance from
the center of the bar.  It is curious in this regard that the gas in the
circumnuclear disk, lying within several parsecs of the dynamical center of the
Galaxy, shows no LSR velocity displacement to within about 20 \kms\ (\S 5), so
the central stellar core and its immediate environment appear to be
kinematically independent of the gas motions on scales of a few hundred parsecs.

Theoretically, m=1 asymmetries in galaxies may arise in response to
gravitational interaction with a passing companion, wherein the time scale for
subsequent relaxation is different for the bulge and disk
components.  However, interaction with another galaxy is not required.  Tagger
\& Athanassoula (1990), for example, point out that an m=1 mode, or wave, can
emerge in a galactic disk via nonlinear coupling to m=2 spiral waves.  From a
different perspective, Miller \& Smith (1992) find in numerical
experiments that a test particle initially located at the galactic center is
subject to overstable oscillations about the galaxy's mass centroid.  This
result would imply that a mass concentration such as a galaxy's central stellar
cluster should typically appear to be offset from the galaxy center as defined
by nearby isophotes.  These authors summarize the literature on off-center and
lopsided galaxies which can perhaps be understood in this light. 

Another form of an $m=1$ wave at the Galactic center is the pronounced tilt of
the plane defined by gas in X$_1$ orbits with respect to the large-scale
Galactic plane (Liszt \& Burton 1980; Uchida et al 1996).  A tilt of the
presumably triaxial bulge has been suggested (Blitz \& Spergel 1991a, Izumiura
et al 1995), but COBE near-IR maps rule out a large tilt (Dwek et al 1995).  
The tilt of the gas layer, considered as a coherent, propagating warp, can perhaps be
understood as a response to the central triaxial bar, even if the major axis of
the bar is not itself tilted with respect to the Galactic plane (Binney 1978,
Sparke 1984). Long-lived warps can also be fed by coupling to m=2 spiral density
waves (Masset \& Tagger 1995, 1996).  For our Galactic center, it remains to be
seen whether the tilt of the gas layer and the longitudinal asymmetries near 
the Galactic
center bear any relationship to each other.  The understanding of both of these
effects is still in its earliest stages. 

\section{Star Formation}

The initial conditions for star formation in the Galactic center environment 
differ dramatically from those found elsewhere in the Galaxy: the temperature,
pressure, velocity dispersion, and estimated magnetic field strength are all
much larger in the CMZ than in the Galactic disk, in some cases by several 
orders of magnitude (\cf, \S2).
Furthermore, potentially collapsing gas clouds are subject to the unusually
strong tidal field near the nucleus, capable of overcoming a cloud's self-gravity
for cloud densities $<$ 10$^7$ cm$^{-3}$ [1.6 pc/r]$^{1.8}$, where r is
the galactocentric distance (G\"usten \& Downes 1980).  Consequently,
self-gravity can initiate collapse only in the densest clouds.  The conditions 
prevailing within the CMZ thus imply an extremely large Jeans mass 
($\sim$10$^5$ \Msun) and the inhibition of star formation via slow, quasi-static
contraction of cloud cores. 
Nevertheless, star formation may proceed in this manner in the occasional
exceptional cloud.  The unusually dense and massive Sgr B2 cloud, located
$\sim$100 pc in projection from the nucleus, is furiously forming stars 
(\eg, Mehringer et al 1993, Gaume et al 1995),
albeit very probably with an initial mass function (IMF) quite
different from that of the galactic disk.  Even in this
cloud, however, the possibility has been raised that the star formation was
provoked by a cloud collision (Hasagawa et al 1994). 

In spite of all of the impediments to star formation within the CMZ, stars are
clearly forming there, at many locations. It seems likely that much of the
star formation is induced by events external to the clouds, notably by shocks
associated with cloud collisions, supernova remnants, and perhaps violent gas
outflows from the nucleus.  This mode of star formation is likely to lead to
an IMF skewed toward relatively massive stars.  The lower mass
cutoff may be elevated as well (Morris 1993 and references therein).  Also, the
enhanced metallicity of gas in the CMZ (up to twice solar; Lester et al 1981, 1987, 
Lacy et al 1989, Wannier 1989, Shields \& Ferland 1994) implies a correspondingly 
large opacity per gram, which affects
the IMF by prolonging protostellar collapse, thus allowing more time for matter
to accrete. 

The global star formation rate in the CMZ, $\Phi$, has been estimated at 0.3 - 0.6 \Msun 
yr$^{-1}$ by G\"usten (1989), assuming a relatively normal IMF, and using the global
production rate of Lyman continuum photons derived from the radio continuum
flux (\S2). This estimate is sensitive to the
IMF, however; if massive stars are favored, relative to the 
Galactic disk, then the global rate of star formation should be correspondingly reduced.
G\"usten (1989) gives a lower limit of 0.05 \Msun yr$^{-1}$ from the luminosity of the
discrete far-infrared sources measured by Odenwald \& Fazio (1984), but this again
assumes a normal IMF.  A proper accounting of the mass budget of interstellar gas in
the galactic center will clearly require new information on the IMF, as well as an improved
census of the sites of star formation. 

Attempts to understand the spatial distribution of star formation near the
Galactic center, and whether, for example, it reflects large-scale shocks in
the CMZ, are still in their earliest stages.  Infrared imaging has revealed several new
star formation sites (Moneti et al 1992, 1994).  Another promising approach has been to
survey H$_2$O masers in IRAS sources (Taylor et al 1993; 
Levine 1995).  After the masers associated with evolved stars are differentiated
from those associated with sites of star formation, about a dozen sites of star
formation in the Galactic center have been identified. This ongoing survey
should eventually reveal large-scale star formation patterns. 

Main sequence stars have been difficult to detect near the Galactic center
because of confusion with the more luminous giants, so the IMF has not yet been
directly probed. However, access to the upper main sequence is becoming
possible with near-infrared observations. Three spectacular clusters of young
stars are now known to be located near the Galactic center, each of which 
contains a sizeable number of exceptionally luminous (10$^{6\pm0.5}$
\Lsun) stars: 1) the well-known cluster occupying the central parsec of the Galaxy,
centered roughly on Sgr A$^*$, 2) the ``Quintuplet'' cluster, or AFGL 2004,
located near the G0.15-0.05 HII region and the NTF's in the 
Galactic Center Arc (Nagata et al 1990, Okuda et al 1990, Glass et al 1990,
Cotera et al 1995a, Figer et al 1995) , and 3) G0.121+0.017, or Object 17,
located near the thermal arched filaments of the radio Arc (Fig.\ 6, Cotera et 
al 1994,1996, Nagata et al 1995, Serabyn \& Shupe 1996).  All of these clusters 
have a population of massive and luminous
emission-line stars, including Wolf-Rayet stars (late-type WN and WC stars,
notably WN9/Ofpe and WC9) and blue supergiants (luminous blue variables, B[e]
stars).  The near-IR emission-line spectra of these stars indicate that they
have high-velocity winds with speeds 500 to 1000 \kms, and that many of these
stars are helium rich, consistent with their presumed post-main-sequence status.

\begin{figure}
%\vskip 8cm
\caption{The G0.121+0.017 cluster, imaged at 2.2-$\mu$m with the NIRC camera on
the Keck Telescope (Serabyn and Shupe 1996)} 
\end{figure}

\noindent{\sc star formation in the central parsec.} The evidence for recent massive
star formation within the central parsec has grown over the years
(Rieke \& Lebofsky 1982, Lacy et al 1982, Allen 1987, Rieke \& Rieke 1989, Allen
et al 1990), and is now rather widely accepted.   The cluster of emission line
stars centered on the core of the central stellar cluster has been the most
thoroughly studied, and in many respects is the most remarkable (Allen 1994, 
Genzel et al 1994, Blum et al 1995a, b, Eckart et al 1995, Krabbe et al 1995, 
Libonate et al 1995, and Tamblyn et al 1996).  The presence of at least two dozen 
HeI/HI emission line
stars and four M or K supergiants ascribed to this cluster has been interpreted
by Krabbe et al (1995) in terms of a modest burst of star formation between 3
and 7 million years ago.  The young stars created in this burst are
intermingled with the older population of the central stellar core, of which
only the giant stars have so far been observed. The young stars dominate the
luminosity of the central parsec, but the total mass of the stars formed in the
burst is only $\sim$10$^4$ \Msun, far less than the mass of old stars in the
central parsec of the Galaxy's central cluster ($\sim$10$^6$ \Msun).  The newly 
formed stars have thus likely not yet equilibrated with their elder brethren.
This scenario places the newly formed cluster in a relatively brief, windy
phase, which should subside in a few million years.  The picture is complicated
somewhat by the presence of about 10 medium-luminosity, late-type stars in the
central 8 arcseconds, objects which appear to be intermediate-mass AGB stars,
signalling that at least one other star formation event took place there within
the past $\sim$10$^8$ years. The implication that the central parsec is the site
of repetitive bursts of star formation is perhaps not surprising, given the
inevitability of radial inflow of gaseous matter (\S3), although the factors
inhibiting star formation -- tidal forces, cloud turbulence, magnetic fields --
are probably nowhere more extreme than in the central parsec. The
star formation events taking place in the deep gravitational potential well of
the Galactic nucleus must be catastrophic events unique to this environment. 

\noindent{\sc the other major, young clusters.} The possibility that the
luminous, blue, windy objects in the central parsec are something more exotic
than young stars (Morris 1993, Eckart et al 1993) was rendered unlikely by the
finding of the same kinds of stars in two other clusters located well outside
the Galaxy's central stellar cluster: the Quintuplet and G0.121+0.017.  The
luminous post-main-sequence stars in these clusters, again identified by their
emission line spectra in the near-IR, constitute a large fraction of the
brightest cluster members (Figer 1995, Cotera 1995), especially in G0.121+0.017
(Fig.\ 6), where a dozen stars show HeI and HI emission lines (Cotera et al
1996).  If these are WN stars, as Cotera et al suggest, then this one cluster
contains 14\% of all known Galactic WN stars.  The brightest members of the
Quintuplet cluster -- the original five -- are featureless in spectra measured
so far, so their nature is unclear; they may be cocoon-like protostellar
objects (Okuda et al 1990), or dust-enshrouded WC9 stars (Figer, pers. 
comm.).  The
bright stars which dominate the Quintuplet and G0.121+0.017 are likely to be
quite massive (50 - 100 \Msun), so both of these clusters have masses of
at least several thousand \Msun, and possibly much larger if their IMF is not
highly unusual. 

The ages of these two remarkable clusters are comparable to that of the young
cluster occupying the central parsec.  A possible scenario for their formation
is that the burst of energy accompanying the formation of stars in the central
parsec strongly shocked nearby, massive clouds, inducing a
relatively catastrophic gravitational collapse over a region much more
widespread than the central parsec. On the other hand, the star formation which
gave rise to these clusters may resemble what is now going on in Sgr B2, where the
star formation is presumably determined by local events. 
The two known, major, non-central
clusters are likely responsible for ionizing the surfaces of molecular clouds in
their immediate environments (Figer 1995, Cotera 1995).  The cloud near the
Quintuplet does not share the velocity of the cluster however (Figer 1995); the
parent clouds of these clusters may no longer be in evidence. 

There are several other sites of star formation near the galactic center (\eg,
Moneti et al 1992, 1994), but no young clusters having the status of the
Quintuplet or G0.121+0.017 are known.  A survey carried out at 2 $\mu$m by Figer
(1995) over a 60 by 30-pc region near the Galactic center revealed no new
examples of clusters of emission-line stars, so if there are further
instances of such clusters, they are either highly extincted or are located further
from the center, in projection, than the known clusters. 

One particularly interesting string of HII regions known as G-0.02-0.07, or the
Sgr A East HII regions, lies quite close (10 pc) to the Galactic center in
projection (Ekers et al 1983, Goss et al 1985).  The HII regions, embedded in a
ridge within the 50 \kms cloud (Mezger et al 1989, Ho et al 1991,
Serabyn et al 1992b) are neatly aligned adjacent to, and along the edge of, the
Sgr A East nonthermal shell source, giving the strong impression that the
expansion of this shell has provoked the formation of the stars which ionize the
HII regions. However, there seems to be a mismatch between the expected
expansion time of Sgr A East and the substantially longer time required for
stars to form in response to the shock from the expansion of the shell, and then
evolve to the main sequence, as these stars have apparently done.  Therefore,
the origin of this cluster is not yet understood.   Furthermore, although
the cluster
appears to contain at least one evolved emission line star (Cotera
et al 1994), star formation may still be adding to it, as evidenced by the
rather high densities inferred for the molecular ridge 
(Serabyn et al 1992b), where
an \water\ maser is located (Yusef-Zadeh \& Mehringer 1995). 

\noindent{\sc the environmental effects of stellar winds.} The powerful winds emanating 
from the massive, post-main-sequence
stars at the Galactic center have a strong effect upon their surroundings,
especially when they act collectively, as in the three major clusters described
above.  The cluster within the central parsec is the only one where the effect
of the winds has been demonstrated, and there the result is profound.  The
1.5-pc-radius cavity inside the circumnuclear disk (\cf, \S5) may be largely
swept clear by the cumulative mass outflow winds (Gatley et al 1984, 1986), 
aside from a few streamers of gas that appear to
have enough inertia to maintain their integrity as they move through this
region. Consequently, the matter currently accreting onto the central object,
Sgr A$^*$, may be almost entirely dominated by the material in the wind (\S6).
Other effects of the wind on gas at small galactocentric radii are described 
in \S5. 

However, the most spectacular effect of the Galactic center wind is the ablation
of the envelope of the red supergiant, IRS7.  The envelope of this well-studied
star (\eg, Sellgren et al 1987) is not only ionized by hot stars near the center
(Serabyn 1984, Rieke \& Rieke 1989, Yusef-Zadeh et al 1989); it also has an
extended, cometary ``tail'' of ionized gas pointing away from the sources of the
wind and showing a pronounced velocity gradient (Yusef-Zadeh \& Morris 1991,
Serabyn et al 1991). This apparently is a case of colliding winds: the
supergiant has its own radiation-pressure-driven wind, although this modest
wind is quickly overcome by the ram pressure of the Galactic wind, and since 
a bow shock is evident on the side of the star facing the wind source
(Yusef-Zadeh \& Melia 1992). The hope that the shape and size of this bow shock
can allow one to derive the momentum in the wind is diminished by the conclusion
of Dyson \& Hartquist (1994) that the ram pressure of the Galactic wind plays no
role in determining the size of the large-scale bow-shock; the long tail implies
that the stellar envelope is clumpy and permeable, so that the wind is
decelerated primarily in bowshocks around the individual clumps.   

The same phenomenon should be happening at some level to the atmospheres of all
red giants and supergiants in the central parsec, although no further cometary
stellar wind tails have yet been observed. Coming years should see more examples of
this phenomenon, as sensitive telescopes with improved spatial resolution are
brought to bear. 

\section{Circumnuclear Material - Disk or Debris?}

Our Galaxy's innermost molecular feature is a relatively compact ($<$ 7 pc
maximum radius) and dense disk- or torus-like structure which orbits about
the center (Genzel et al 1994 and references therein). Although normally
referred to as the circumnuclear disk (CND), this nomenclature may actually hide
a multitude of sins (\S5.2). The molecular medium in general extends only to
within about 1.5 pc of the center, its inner edge presenting a rather sharply
defined ionized boundary layer basking in the radiation from the encircled
stellar cluster. Much of the filamentary Sgr A West HII region (Fig. 7; Killeen
\& Lo 1989,
Lacy et al 1991, Roberts \& Goss 1993, Yusef-Zadeh \& Wardle 1993) 
located in the innermost 1.5 pc
radius region, can then be attributed to photoionization of the 
molecular medium's inward-facing surfaces. In order for centrally originating
photons to propagate out to the CND's inner boundary, the wind-evacuated 
``central cavity'' must be comparatively transparent (Becklin et al 1982). 
However, gas
kinematics indicate that several ionized filaments and clumps are found in
closer proximity to the center (Serabyn \& Lacy 1985, Serabyn et al 1988,
Serabyn 1989, Lacy et al 1991, Herbst et al 1993b, Lacy
1994), as is neutral and molecular gas (Davidson et al 1992, Marr et al 1992,
Jackson et al 1993, Pauls et al 1993, Yusef-Zadeh et al 1993, Zhao et al 1995,
Telesco et al 1996), both likely the result of dense gas plunging into the
central cavity from the CND or beyond. However, the location of some of the
molecular material remains controversial (Liszt \& Burton 1993, Marshall \&
Lasenby 1994b). In this section, we discuss recent developments which
bear upon this general scenario, and also address evidence regarding its
shortcomings. As this is a well studied topic, we limit ourselves to a quick
summary. 

\begin{figure}
%\vskip 12cm
\caption{12.8 $\mu$m Ne$^+$ line emission from the Sgr A West HII region (Lacy
et al 1991). The nomenclature for the ionized filaments is as follows: the
Western Arc is the diffuse, slightly curved filament crossing the frame from
top
to bottom to the right (west) of center, the Northern and Eastern Arms are the
well-defined filaments extending from the vicinity of the center of the image
to
the north and east, respectively, and the Bar is the bright, clumpy east-west
structure near the center of the image.}
\end{figure}

\subsection{The standard model}

The established picture of a clumpy, centrally-illuminated torus or disk of
largely molecular gas provides a good first-order description of the
excitation, distribution, and kinematics of the gas making up the CND. Although
a one-to-one identification of ionized structures with individual molecular
clumps is far from complete, a large and increasing fraction of the ionized gas
in the central few parsecs can indeed be accounted for in this manner (Telesco
et al 1996). Thus, the intricate distribution of ionized gas evident in the
Ne$^+$ 12.8 $\mu$m image of Sgr A West in Fig. 7 (Lacy et al 1991) can likely
be attributed largely to photoillumination of the clumpy molecular medium
surrounding the center by stars in the central cluster. Local heating by
individual stars is also observed (Smith et al 1990, Gezari 1992). 

Copious measurements have by now probed the molecular medium lurking behind the
outermost filaments in Fig. 7 (most recently, by Sutton et al 1990, Jackson et
al 1993, Marr et al 1993, Marshall \& Lasenby 1994a, Serabyn et al 1994). The
HCN 1-0 map of the inner part of the CND, along with the 90-$\mu$m dust continuum 
emission,
is shown in Figure 8. The gas near the CND's inner rim is both hot (few hundred
K) and dense (10$^4$ - 10$^7$ cm$^{-3}$), with line emission consistent with a
dense photodissociation region (Genzel et al 1994). The column abundances and
excitation states of the ionized and neutral gas components, as well as of the
dust component, provide constraints on the incident radiation field's effective
temperature and total luminosity (\S6). The molecular medium is extremely clumpy
(with clumps most likely taking the form of tidally sheared streamers), with
low filling factor, allowing for penetration of radiation well beyond the CND's
innermost edge. Large linewidths ($\wig>$40 km s$^{-1}$) at the CND's inner edge
(Marshall \& Lasenby 1994a) likely reflect a large inter-clump velocity
dispersion, the magnitude of which suggests frequent inter-clump collisions and
shocks. Toward the outer edge of the CND, the linewidths decrease (Serabyn et
al 1994), probably due to decreasing clump overlap. 

\begin{figure}
%\vskip 11cm
\caption{The circumnuclear disk, viewed in the emission of three probes.
 a) 90-$\mu$m continuum emission (Davidson et al 1992), b) HCN J=1-0 emission (Wright 
 et al 1987), c) 2-$\mu$m molecular hydrogen emission 
(Gatley \& Merrill 1995, I. Gatley, personal communication) 
The last panel (d) shows the velocity field of the H$_2$
 emission, indicating that the maximum velocities do not lie at the ends 
 of the apparent major
 axis of the elliptical emission trace, and the nodes in the velocity distribution do
 not lie along the apparent minor axis.}
\end{figure}


Circular rotation provides a good first-order fit to extant kinematic data
on the CND. Most kinematic models yield low values for radially-directed
velocities, although other interpretations remain possible. The differences in
interpretation appear to stem largely from the method of selection of a disk
major axis orientation: a major-axis selected to coincide with maximal
velocities necessarily yields low radial velocities ($\wig<$ 20 km s$^{-1}$;
Jackson et al 1993, Marshall \& Lasenby 1994a, Serabyn et al 1994), while an
orientation based on the morphology of the intensity distribution (or any other
criterion) can yield a larger radial velocity component (up to 50 km s$^{-1}$;
Gatley et al 1986, Gatley \& Merrill 1994; see Fig. 8). A common major-axis
orientation may not apply at all radii (or for all molecular species), but even
so, speculations about disk warp are likely premature. The most general
approach would fit all disk parameters, but few data sets have been complete
enough (beyond the CND's inner rim) to warrant such a treatment. 

At smaller radii, gas kinematic modelling of the dense, ionized gas filaments in
the Sgr A West HII region is possible. The ionized gas shows velocities as high
as $\pm 300$ km s$^{-1}$, with pronounced velocity gradients in the central 10
arcsec. For several of the inner filaments, eccentric orbits in the field of a
point mass reproduce the gas velocities well (Serabyn and Lacy 1985, Serabyn et
al 1988, Serabyn 1989, Herbst et al 1993a), under the assumption of gas flow
along the filaments. However, flow primarily along the filaments need not be
taking place. The alternatives include tidally stretched cloud models
(Quinn and Sussman 1985, Lacy et al 1991, Zhao et al 1995), which call for
similarly elongated orbits with large radial motions, and the contrasting
mini-spiral model of Lacy et al (1991), which allows an interpretation of the
data (along a unified ``Western Arc" - ``Northern Arm" feature; see Fig.\ 7 for
nomenclature) in terms of circular orbits. However, growing evidence indicates
distinct natures for the Western Arc and Northern Arm, the former being the
inner photoionized boundary of the CND in near circular rotation, and the latter
the inner photoionized boundary of a neutral gas feature presumably falling
towards the center (the northern intruder, Davidson et al 1992, Jackson et al
1993, Telesco et al 1996).  As even the ``eastern arm" and a portion of the 
innermost ``bar'' feature (Fig.\ 7) have velocity patterns consistent with
orbital trajectories transiting the central 0.5 pc (Serabyn 1989), it is
difficult to avoid the conclusion that the kinematics of the innermost filaments
call for an infall component. 

However, while large radial motions are indicated for the ionized filaments
interior to the CND (and the underlying molecular material, such as the northern
intruder), the CND exhibits relatively small radial motions in most models.
Thus, net mass inflow from the CND is probably not the result of small-scale
viscosity in a homogeneous disk, but is perhaps more accurately described either
in terms of sporadic infall of individual clumps resulting from dissipative
inter-clump collisions, or as tongues of gas splitting off the inside edge of
the CND as a result of an instability at this interface. Recent estimates for
the inflow rate for the CND gas and the material in the central cavity are
$\sim$ 3 - 5 $\times 10^{-2}$ M$_\odot$ yr$^{-1}$ (G\"usten et al 1987, Jackson
et al 1993). 

The inner rim of the CND also exhibits emission from vibrationally excited 
H$_2$, which may result from a strong shock caused by the impinging central wind
(Gatley et al 1984, 1986, Fig. 8), although UV fluorescence has not been
completely excluded (Pak et al 1996). This wind, arising from stars in the
central stellar cluster (\S4), probably helps to explain the abrupt density
discontinuity at the CND's inner edge, where a quasi-steady equilibrium may be
established between the pressures of the wind and the turbulent disk. The impact
of the wind upon the disk may also contribute to the large linewidths at the
disk's inner edge, and it has been held responsible for creating the radio
continuum ``streamers" that originate at the interface (Yusef-Zadeh \& Morris
1987a, Yusef-Zadeh et al 1995). However, we note that the sharp inner edge of
the CND might also be the result of a discontinuity in the radial gradient of
the Galactic mass distribution at that point (Duschl 1988), a density wave in
the field of a central point mass (Lacy et al 1991), a hydromagnetic instability
(Fridman et al 1994), or as we discuss in the next section, the possible
short-lived nature of the CND. The wind also appears to affect the distribution
of the ionized filaments in the centralmost few arcseconds, causing compressed
and possibly rippled surfaces on the infalling gas streams (Yusef-Zadeh \&
Wardle 1993), and it may be responsible for the formation of the ``mini-cavity''
in the gas distribution within the bar (Morris \& Yusef-Zadeh 1987, Yusef-Zadeh
et al 1989, 1990, Wardle \& Yusef-Zadeh 1992, Eckart et al 1992, Lutz et al
1993, 1994, Yusef-Zadeh 1994). 

\subsection{A Symmetric Disk?}

The major objection to the simple model of the CND described above is that the 
material beyond the CND's inner edge is distributed highly asymmetrically about
the center (Serabyn et al 1986, 1996b, Fukui \& Churchwell 1987, Sutton et al
1990, Marshall \& Lasenby 1994a): at negative Galactic longitudes the disk
extends up to 7 pc from the center, whereas at positive longitudes the disk's
radial extent is much smaller, $\wig<$ 3 pc. In galactic latitude, it is the
ionized gas which is asymmetrically distributed, with a bright ionized filament
(the Western Arc) outlining the positive latitude side of the CND (Fig. 7), but
only a weak counterpart in the southeastern quadrant. Although this may result
from selective shielding by interior features, the molecular gas distribution
along the CND's inner rim is also not entirely symmetric (Jackson et al 1993,
Marr et al 1993). In addition, molecular gas extends asymmetrically into the
CND's central cavity from the northeast. Kinematically, the velocities are also
not quite azimuthally symmetric, although it is not yet clear whether this is
the result of non-circular streamlines or merely clumpiness (\ie, regions of
``missing" gas). Finally, largely ignored and still unaccounted for is the 
positional coincidence between the southwestern rim of the nonthermal Sgr A East
radio shell source, and the thermal Western Arc filament in the Sgr A West HII
region (Ekers et al 1983, Yusef-Zadeh \& Morris 1987a, Pedlar et al 1989). 


\begin{figure}
%\vskip 10cm
\caption{The Sgr A complex, showing the Sgr A East
nonthermal radio shell source (outermost white contours), the Sgr A West
thermal
HII region (the fourth white contour and above), the compressed 50
\kms\ molecular cloud layers surrounding the expanding radio shell 
(grey scale and
dashed contours), and the circumnuclear disk (solid black and dotted
contours). The (0,0) position marks the location of Sgr A$^*$. All maps except
the radio continuum map derive from observations of
the 492 GHz fine-structure transition of CI (Serabyn et al
1996b). The velocity intervals are --- 50 \kms cloud: 30 to 80 \kms;
blueshifted CND gas (solid black contours): -130 to -70 \kms; redshifted CND
gas (dotted contours): 80 to 130 \kms.}    
\end{figure}

Symmetric models thus likely oversimplify matters. As symmetry implies
longevity, asymmetry suggests time dependence, \ie, a short-lived disk. 
The asymmetric distribution of the CND beyond its inner
rim suggests that the current appearance of the CND may result either from
the gravitational capture and tidal disruption of a passing molecular cloud, or
from an energetic disruption to a stable disk.  In either case, the relative 
symmetry of the CND's inner rim is a natural result of the shorter rotational 
timescales there.

The origin of either of these alternatives may lie in the expanding Sgr A East
shell, as the atomic carbon map presented in Fig.\ 9 shows a very high degree of
anticorrelation between the southern lobe of the CND and the ring of neutral
carbon which evidently encloses the Sgr A East radio shell (Serabyn et al 1996b)
--- the carbon ring fades in intensity precisely at the location of the southern
lobe of the CND. Thus, either the expansion of the Sgr A East shell has
compressed and swept up molecular gas, pushing it toward the center (Mezger et al 1989), 
where it is now settling into a disklike configuration
(in which case the notion of a wind-evacuated ``central cavity" may be
unnecessary), or the expansion has disrupted a pre-existing and originally more
symmetric disk. Alternatively, the geometric relationship could be radiatively 
highlighted; that is,the central 
cluster could be illuminating the inner edge of the Sgr A East shell except where it
is shielded by the nearer CND material,  but even in this case the expansion of
the radio shell must be considered in the context of mechanical interaction with
the CND. While the most likely scenario is not manifest at this point
(except that Sgr A East is largely behind Sgr A West; Yusef-Zadeh \& Morris
1987a, Pedlar et al 1989), what does emerge clearly is that in all of these
possibilities the expanding Sgr A East shell, and more specifically the
molecular material which the shell is compressing (Mezger et al 1989, Genzel et
al 1990, Zylka et al 1990, Ho et al 1991, Poglitsch et al 1991, Serabyn et al 
1992b, 1996b, Ho 1994), is
indeed intimately associated with the CND. As a result of this interaction, the
long-lived nature of the CND must be seriously questioned. 

\section{The Compact Central Object}

In common with several categories of external galaxies (Owen et al 1980), our
own Galaxy harbors a compact radio source at its center (Lo 1989, 1994,
Backer 1994, 1995). Referred to as Sgr A$^*$ to differentiate it from the
extended Sgr A complex in which it is ensconced, our Galaxy's central radio
source shares a unique set of characteristics with its siblings: central 
location, compact size, high brightness temperature, and a relatively flat radio
spectrum. However, its luminosity falls near the low end of the observed range
for compact nuclear radio sources, and this is almost certainly the 
result of selection, as
proximity makes our own nuclear source comparatively easy to detect. Thus, the
standard argument calling for a massive central black hole on the basis of a
prodigious energy output from a small volume is inapplicable in the case of our
own Galactic nucleus. Indeed, independent of wavelength, the flux from the
vicinity of Sgr A$^*$ is quite low compared both to observed active galactic
nuclei, and to theoretical expectations for the radiative output from an
accreting compact object. 

\subsection{Source Characteristics}

Although Sgr A$^*$ is frequently referred to as a ``point source", in fact VLBI
observations at centimeter wavelengths routinely resolve its emission, showing
an apparent source size very closely proportional to $\lambda^2$, for
wavelength $\lambda$. An axial size ratio of 0.5 is typically
observed (Alberdi et al 1993, Lo et al 1993), with a major axis position angle
roughly EW, and compilations of existing measurements yield a major-axis source
size in $\mu$as of $\theta_{\mu as} = 14\, \lambda^2_{mm}$ (Jauncey et al 1989,
Marcaide et al 1992, Backer 1994, Yusef-Zadeh et al 1994). This behavior is 
consistent with anisotropic scattering by the intervening interstellar medium,
and as nearby maser sources show similar characteristics, the highly magnetized
medium in the central few hundred parsecs is a prime candidate for the
scattering agent (Frail et al 1994, Yusef-Zadeh et al 1994). Unfortunately, the
resultant image broadening prevents direct observation of source structure in
Sgr A$^*$ at cm wavelengths (and also masks any possible intrinsic size
variation with $\lambda$). 

VLBI observations have recently been extended to millimeter wavelengths,
yielding tantalizing suggestions of intrinsic source structure (Krichbaum et al
1993, 1994), but as yet no confirmed evidence for a source size exceeding that
set by interstellar scattering (Backer et al 1993, Rogers et al 1994). The most
stringent observational limit gives an intrinsic source diameter $<$ 130
$\mu$as, or 1.1 AU, at $\lambda = $ 3.5 mm (Rogers et al 1994), and an intrinsic
brightness temperature $> 1.4 \times 10^{10}$ K, implicating synchrotron
emission. As the inverse Compton limit and the absence of short-time-scale 
scintillations at near-millimeter wavelengths both imply a source
diameter $\wig >$ 10 $\mu$as, or 0.1 AU (Kellermann \& Pauliny-Toth 1981, Gwinn
et al 1991, Serabyn et al 1992a, Zylka et al 1992), the intrinsic source size at
near-millimeter wavelengths is evidently constrained to within an order of
magnitude. This scale size is quite interesting from a theoretical perspective, 
as the smallest stable circular orbit about a
massive black hole (with a diameter of 6 Schwarzschild radii, or $0.12\, M_6$
AU, where $M_6$ is the black hole mass in units of $10^6$ M$_\odot$), falls near
the lower end of the range, for masses suggested by kinematic data ($M_6
\sim$ 2 - 3; \S 6.3). VLBI observations are thus beginning to probe a very 
telling regime, although it is likely that the shortest millimeter
wavelengths, and perhaps even submillimeter wavelengths, will be needed to
overcome the scattering handicap and reveal the intrinsic source
structure. 

The spectrum of Sgr A$^*$ is also key to understanding its nature, and
the radio source has now been detected at frequencies up to 670 GHz
(Zylka et al 1992, 1995, Dent et al 1993). Most observations at frequencies below
100 GHz show a slowly rising, approximately power-law spectrum (Lo 1989), 
but as moderate
variability is established at cm and mm wavelengths 
(Zhao et al 1989, 1992), care is 
required in deriving the spectral slope.
In particular, as Sgr A$^*$ shows short time-scale
($\wig<$ month-long), weak ($\sim$ a factor of 2), ``outbursts'' at cm and mm
wavelengths several times per year, with
concommitant spectral changes (Wright \& Backer 1993), a concatenation of
numerous flux density measurements taken over the years could be misleading. To
avoid these concerns, we instead show in Fig.\ 10 nearly simultaneous measurements
of the radio/millimeter spectrum of Sgr A$^*$ taken inside a span of 12 days in
March 1993 with the VLA, OVRO, and CSO-JCMT interferometers (Serabyn et al
1996a). These ``simultaneous'' data cover nearly two decades in frequency, from 5
to 354 GHz, and the measurements at frequencies below 100 GHz are best 
fit by a $\nu^{0.25}$ power law (Fig.\ 10), consistent with
earlier spectra. At frequencies above 200 GHz, an excess
of 1 - 2 Jy above the extrapolated power law (up to a factor of 2) is evident, and
the higher frequency ``simultaneous'' data points are in good
agreement with prior interferometric and bolometric measurements 
(Zylka \& Mezger 1988, Zylka et al 1992, 1995, Serabyn et al 1992a,
Dent et al 1993). 

\begin{figure}
%\vspace{8cm}
\caption{The radio/infrared spectrum of Sgr A$^*$. The solid squares show the
near-simultaneous measurements of Serabyn et al (1996a), the empty squares the
submillimeter data of Zylka et al (1995), the downward-pointing arrows the
measured upper limits to the source flux at shorter wavelengths, the inverted
triangles the corresponding extinction-corrected limits, and the asterisks the
measured 8.7 $\mu$ flux (25 mJy; Stolovy et al 1996) and its
extinction-corrected value. The measured upper limits are: 8 Jy at
350 $\mu$ and 3 Jy at 24.3 $\mu$ (Serabyn et al 1996a), 20 Jy at 30 $\mu$
(Telesco et al 1996),  1 Jy at 20 $\mu$ (Gezari 1994), 0.3 Jy at 18.1
$\mu$ and 0.1 Jy at 12.4 $\mu$ (Gezari 1992), 6.5 mJy at 4.8 $\mu$ and
5.5 mJy at 3.8 $\mu$ (Herbst et al 1993b), and 1.6 mJy at 2.2 $\mu$ (Eckart
et al 1995). The upper limits are corrected using the extinction curve of
Mathis (1990) with A$_{\rm V}$ = 30.  The
conservative 1 $\mu$ limit of 75 $\mu$Jy (Liu et al 1993) is too low to plot on
this scale. }
\end{figure}

At frequencies beyond those measured interferometrically, the shape of the
spectrum remains unsettled. Single-dish measurements between 300 and 670 GHz
yield flux densities of roughly 3-4 Jy for Sgr A$^*$ across this range, although
an initial report of a 1.5 Jy upper limit at 670 GHz (Dent et al 1993)
conflicts. Large error bars leave spectral slopes indeterminate at present, 
and the extant database is insufficient to definitively address the question
of source variability at submillimeter wavelengths. In any case, the fact that
Sgr A$^*$'s flux density does not exceed the 4 Jy level even for frequencies
approaching 700 GHz suggests that the meager near-mm excess is due either to a
slight rise in an undulating spectrum, as is seen in external nuclear radio
sources at lower frequencies (Owen et al 1980), or perhaps to a final ``hump''
just before turnover (Falcke 1995). 

Because of its positive spectral index at radio frequencies, the luminosity of
Sgr A$^*$ is dominated by its highest frequency emission, making submillimeter
and shorter wavelength observations critical for a determination of this
quantity. The power law spectrum to 100 GHz has a total luminosity of
$\approx 10^{34}$ erg s$^{-1}$, and inclusion of the emission to 700 GHz increases the
total luminosity by roughly an order of magnitude, to $1.4 \times 10^{35}$ erg
s$^{-1}$. At just under 40 L$_{\odot}$, this is well below the Eddington limit
of even a solar mass object ($3.4 \times 10^4$ L$_{\odot}$). 

Searches have also been carried out for counterparts to Sgr A$^*$ at many
shorter wavelengths, with success only recently becoming apparent. Several far-
and mid-infrared upper limits (Fig.\ 10), as well as the mid-IR detection of a
source toward Sgr A$^*$ (Stolovy et al 1996), constrain the high frequency
rolloff of Sgr A$^*$'s synchrotron emission. It is clear from the
extinction-corrected data at wavelengths shortward of 20 $\mu$m (Fig.\ 10) that
the synchrotron spectrum must cut off between 1 and 10 THz, with a falloff at least
as steep as $\nu^{-1}$. The maximum
synchrotron luminosity then cannot exceed several hundred L$_{\odot}$. We expect
that measurements of the submillimeter and mid-IR spectra will soon be dramatically 
improved by the new generation of high-resolution cameras.

In the near-IR, where angular resolution and sensitivity have
increased markedly in the past decade, no definitive counterpart to Sgr A$^*$
has yet been identified, owing primarily to confusion by surrounding stellar 
sources. Increased pointing accuracy has eliminated several candidate
detections, and most recently, Eckart et al (1995) have found a compact cluster of
half a dozen near-IR sources located within a few tenths of an arcsecond of the
position of Sgr A$^*$ (Fig.\ 11), possibly representing a central cusp in the
stellar distribution.  While intriguing in its own right, this result makes the
identification of any non-stellar near-IR source which may be associated with
Sgr A$^*$ itself problematic. At the same time, the extremely low limits (m$_K
\wig>$ 14, Fig.\ 10) to the possible near-IR flux attributable to Sgr A$^*$
severely constrain accretion disk models (Sect.\ 6.2). 

\begin{figure}
%\vskip 8cm
\caption{High-resolution K-band image of the vicinity of Sgr A$^*$, 
from Eckart et al (1995).}
\end{figure}

At visual and UV wavelengths, interstellar extinction obscures the
galactic center completely, but constraints on the radiation field at these
wavelengths can nevertheless be derived from the excitation of the ISM in and
surrounding the central parsec (Becklin et al 1982, Lacy et al 1982, Serabyn et
al 1985, Davidson et al 1992, Maloney et al 1992, Genzel et al 1994, 
Shields \& Ferland 1994, Zylka et al 1995, Telesco et al
1996). The result for a single-component cloud model is a radiation temperature 
of about 35,000 K, an ionizing photon flux of 2 $\times$ 10$^{50}$ 
s$^{-1}$, and a total
luminosity of $\approx$ 2 $\times$ 10$^7$ L$_\odot$. However, as stellar sources
apparently dominate the production of Lyman continuum photons in the central
parsec (Rieke et al 1989, Krabbe et al 1991, 1994, Tamblyn \& Rieke 1993), the
optical/UV flux in the central parsec which arises from Sgr A$^*$ itself is
likely only a small fraction of the total ($\wig< 10^6$ L$_\odot$) and is hence
difficult to determine reliably. 

At an energy of roughly 1 keV, it again becomes possible to search for emission
from the direction of Sgr A$^*$. A 5-arcmin resolution GRANAT
ART-P map at 4-20 keV shows emission from the general direction of the
extended Sgr A source (Syunyaev et al 1991, Pavlinsky et al 1994), but high
resolution 2-10 keV ASCA data (Koyama et al 1996) show the main emission to be
associated not with Sgr A$^*$, but with the Sgr A East shell source. 
The 2-10 keV flux from Sgr A$^*$ itself is then limited by the ASCA 
observations to $\sim 10^{35}$ erg s$^{-1}$ (Koyama et al 1996),
an order of magnitude below the large-beam GRANAT estimate, and more consistent with the
1.5 $\times 10^{35}$ erg s$^{-1}$ seen below 4 keV by Watson et al (1981). The
lower ASCA flux is also more consistent with data in the softer 1.2-2.5 keV
ROSAT band (Predehl \& Tr\"umper 1994), where a source coincident with Sgr A$^*$
was also found. Comparison of the ROSAT and GRANAT data had earlier suggested
the need for an anomalously high extinction at long X-ray wavelengths, but this
need is removed by the the lower ASCA flux actually attributable to the compact
source. Of course, source variability may affect this discussion.
At higher energies still (35-150 keV), no source was detected toward Sgr A$^*$
by the Sigma/GRANAT telescope (Goldwurm et al 1994). Thus in total, the X-ray
flux which arises in a small enough region to be attributable to Sgr
A$^*$ itself is no more than a meager $2.5 \times 10^{35}$ erg s$^{-1}$. 

\subsection{Source Models}

The preceeding section establishes that Sgr A$^*$ is a rather weak source
across the electromagnetic spectrum, with at most a few hundred L$_\odot$ at
radio and X-ray wavelengths. The intermediate-frequency situation is not as
clear, but allowable near-IR/optical/UV luminosities have decreased markedly
over the past decade (by a factor of 100 at K band), ruling out luminosities for
Sgr A$^*$ in excess of 10$^6$ L$_\odot$. Because of the low radio and X-ray
luminosities definitively attributable to Sgr A$^*$, accretion onto a stellar
source cannot be excluded solely on energetic grounds. However, the radio
spectrum of Sgr A$^*$ is much flatter than that of a pulsar, too stable for a
mass-transfer binary, and much more energetic than typical examples of either
(Lo 1989). Furthermore, as proper motion studies (Backer 1995) establish a lower
limit to the mass of Sgr A$^*$ of at least 100 M$_\odot$, and possibly as high
as 1000 M$_\odot$ (depending on whether one assumes kinetic equilibrium with low
or high mass stars, respectively), stellar sources are effectively ruled out.
Thus, models advanced recently to account for the emission from this unique
radio source are all variants on the general scenario of accretion onto a
massive compact object. 

Proposed models differ widely in their assumptions,
including the size of the central mass, the accretion rate and geometry, the
radiative mechanisms and efficiency, the source opacity, and the disposition of
infalling material. Central masses of $500$ - $3 \times 10^6$ M$_\odot$
appear in different models (Falcke et al 1993, Ozernoy 1993, 1994, Melia 1994a,
1994b, Falcke 1995, Narayan et al 1995), and accretion rates vary even more widely,
from 10$^{-10}$ to $10^{-4}$ M$_\odot$ yr$^{-1}$, the high end of which is
coincidentally the maximum possible average accretion rate over a Hubble time.
Emission geometries range from spherical to disk-like to jet-like, with the
curious result that the same radio spectrum can apparently be fitted (to some
level of accuracy) with either a
jet-like outflow or an accretion inflow. The radio spectrum as measured to 
date thus may be less than definitive, and indeed, models running the gamut from
optically thick to thin have been proposed (de Bruyn 1976, Reynolds \& McKee 
1980, Duschl \& Lesch 1994). Higher frequency continuum flux measurements and
VLBI imaging will no doubt aid in discriminating
between models, as a definitive measure of source geometry, size, and cutoff
frequency would provide tangible constraints. 

Potentially one of the most telling probes of SgrA$^*$ may well be its
spectrum beyond the radio (\ie, synchrotron) regime. It is too soon to tell
whether mid-IR emission might arise from the tail of the synchrotron spectrum
(which is at least as steep as $\nu^{-1}$), a
dust component, inverse-Compton scattered photons, or from an accretion disk
photosphere. 
With only upper limits available from the near-IR through the UV, definitive
model fitting is unfortunately not yet possible at these wavelengths either. It
has however long been clear that standard accretion disk photospheres are in
conflict with stringent near-IR/UV source constraints (Rieke \& Lebofsky 1982,
Lacy et al 1982, Rieke et al 1989). The only potential means of evading the
near-IR constraints, for a classical thin-disk, is to assume a disk
orientation very close to edge-on, certainly a distinct possibility (we are,
after all, located in the plane of the Galaxy), but the UV constraint arising
from the excitation of the surrounding medium (after removal of the stellar UV
contribution) is relatively independent of disk orientation. Furthermore,
general-relativistic light bending tends to isotropize emergent radiation,
rendering even an edge-on disk visible (Falcke \& Heinrich 1994, Hollywood \&
Melia 1995) so that disk inclination may not be a panacea after all. Thus,
near-IR/UV limits remain extremely restrictive. Finally, high resolution
X-ray data (Watson et al 1981, Koyama et al 1996) also provide very tight
luminosity constraints. 

How do existing models fare? Four schools of thought have recently emerged in
rather distinct corners of parameter space, and we briefly outline the salient
points. The first two models are related, both being based on the idea that a
massive black hole would capture and accrete a sizable fraction of the observed
stellar mass-outflow wind in its vicinity (Ozernoy 1989), a scenario which has
the advantage of properly accounting for the specific environment of Sgr A$^*$.
Expanding on this idea, Ozernoy (1993, 1994) and Mastichiadis \& Ozernoy
(1994) claim that the observed radio and X-ray
luminosities of Sgr A$^*$ determine upper limits to its mass of $\sim$ 500
M$_\odot$ and $\sim$ 6000 M$_\odot$, respectively. As the Bondi-Hoyle 
accretion rate is proportional to the square of the
central mass, such a low mass results in very slow accretion, $10^{-10}$
M$_\odot$ yr$^{-1}$, and so little near-IR radiation is produced. However, the
low mass also introduces a small length scale, implying optically thick
synchrotron emission at high frequencies, and so a further emission component
not currently in the model is needed to account for the flat low-frequency
spectrum observed (Falcke 1995). Beyond this, a detailed prediction for the
source spectrum would certainly aid in evaluation. 

A similar Bondi-Hoyle accretion scenario is considered by Melia (1994a,
1994b), but diametrically opposite conclusions are reached: a central mass of
1-3 $\times 10^6$ M$_\odot$ is called for, and an extremely high accretion rate
emerges, 1.6 $\times 10^{-4}$ M$_\odot$ yr$^{-1}$, necessitating an edge-on disk
unless a general-relativistic treatment is applied (Hollywood \& Melia 1995).
However, while the low-frequency radio spectrum (below 100 GHz) can be fitted, 
the turnover frequency predicted tends to be too low, no excess emission
at submillimeter wavelengths is predicted, and the most recent near-IR/UV and
X-ray constraints to the flux from Sgr A$^*$ are likely violated. 

Of course, the near-IR and UV luminosity problems are obviated if the accreting
material radiates with low efficiency. This can occur if electrons and ions are
decoupled in temperature in the inner parts of the disk (Rees et al 1982,
Narayan et al 1995), a situation which can arise because the viscously heated
heavy ions cannot
easily transfer energy to the much lighter
electrons via two-body collisions. Thus, the more efficient radiators do not 
have ready access to the thermal energy, and the energy is advected with the ions'
flow (Abramowicz et al 1995,
Narayan \& Yi 1995).  To remove angular momentum on a timescale short
compared to the ion-electron equilibration timescale, a reasonably high
viscosity is also needed, implying radial infall rates higher than in a
standard thin-disk (10$^{-7}$ - 10$^{-5}$ M$_\odot$ yr$^{-1}$ in Narayan
et al 1995). Accreting material is then able to fall through the
Schwarzschild radius before radiating away its thermal energy content, yielding
a ``dim disk" in the near-IR/optical/UV range. The resulting spectrum 
(Narayan et al 1995) consists of a series of roughly equal energy
humps, due to (in increasing frequency) synchrotron, inverse-Compton, and
thermal bremsstrahlung emission. A central mass of rougly 10$^6$ M$_\odot$ 
provides a good match to the submillimeter luminosity, but not to the radio
spectrum (without the addition of an {\it ad hoc} component), and there may also
be difficulties in fitting the X-ray and submillimeter fluxes in a single model,
especially when the data of Koyama et al (1996) are considered. However, most
worrisome, as plasma instabilities may serve to couple the ions and electrons
much more effectively than single-particle collisions, this scenario is based on an
untested assumption. 

Finally, in analogy with external galaxies, a core-jet model has been suggested
(Blandford \& K\"onigl 1979, Falcke et al 1993, Falcke \& Heinrich 1994, Falcke
1995), in which the observed radio emission arises not in the accreting
material, but in an ensuing jet-like outflow. This model thus has the specific
prediction that at high enough angular resolution, a jet should be seen, and it
is also the only model which predicts a slight excess of emission at
submillimeter wavelengths. The latter results from the emission from the jet
nozzle, the smallest scale structure (Falcke 1995). The predicted accretion rate
is $\sim 10^{-7}$ M$_\odot$, a mid-range value, and is thought to be set by long
viscous timescales in a remnant steady state accretion disk. As the accretion
rate is governed not by the capture rate, but by viscosity, the central mass is
not well constrained. Possible drawbacks to this model are the large discrepancy
between the Bondi capture rate and the low steady-state accretion rate required,
the fact that no jet has yet been conclusively detected, and the
lack of a prediction for the X-ray emission. 

While several of these models show promise, the contrast between
them also serves to highlight the continuing scarcity of definitive
source-structure information. However, with continually improving observational
capabilities throughout the electromagnetic spectrum, this situation will no doubt
soon improve. 

\subsection{The Central Mass}

Clearly one of the most vital of Sgr A$^*$'s unknown parameters is its mass,
which cannot be determined uniquely from models of the emerging radiation.
On the other hand, observational evidence has been mounting for over a decade
which consistently shows that line-of-sight velocities increase toward the
center (inside the central 0.5 pc or so), for both the gas and the stars (Lacy
et al 1980, 1982, 1991, Serabyn \& Lacy 1985, Serabyn et al 1988, McGinn et al
1989, Sellgren et al 1990, Herbst et al 1993a, Krabbe et al 1995, Haller et al
1996). While the conversion of these findings to an enclosed mass distribution
is fraught with uncertainties, for both the stellar and gaseous components,
different techniques yield remarkably good agreement both for the size (2 - 3 
$\times$ 10$^6$ \Msun) of the central dark mass
concentration (observationally constrained to lie within the central few tenths
of a parsec), and for the distributed stellar mass, 
$\sim$3 $\times$ 10$^6$ r$_{pc}$ \Msun, just beyond the 
core radius of the central cluster   
(\cf, Rieke \& Rieke 1988, McGinn et al
1989, Lindqvist et al 1992, Serabyn et al 1994, Haller et al
1996). While the radial size constraint on the dark mass concentration cannot
be used to tie the mass excess directly to Sgr A$^*$, new imaging techniques
now being employed are pushing observations ever closer to the center, where
the discrimination will no doubt improve. In particular, the possible stellar
cusp around Sgr A$^*$ may soon provide important proper motion data (Eckart et
al 1995). 

\section{The Episodic Nature of Activity at the Galactic Center}

The inexorable inward flow of matter in our Galaxy makes nuclear activity
inevitable; if the accretion rate onto a presumed central black hole of mass
M$_6$ were only 0.02 M$_6$(0.1/$\epsilon$) \Msun\ yr$^{-1}$, where $\epsilon$ is
the efficiency with which the accretion energy is radiated away (often estimated
at $\sim$0.1), then the nucleus would emit at the Eddington rate, with an energy
output of Seyfert proportion, $\wig>$10$^{43}$ ergs s$^{-1}$.  Given the mass
flow budget outlined in \S3.2 and \S5.1, a time-averaged mass flow rate into the
central parsec equal to this value is plausible.  However, the current accretion
rate and luminosity of Sgr A$^*$ are many orders of magnitude smaller than this
(\S6). It therefore appears that the Galactic center may now be in a lull
between brief accretion events. The massive young stars clustered within the
inner parsec are indicative of a substantial mass accumulation event within the
past 10$^7$ years, while the probably inward-moving circumnuclear disk likely
heralds the succeeding accretion epoch, possibly within the coming 10$^5$ -
10$^6$ years. 

The implied limit cycle is one that begins with the relatively sudden ($\wig<$
10$^4$ years) appearance of gas within the gravitational potential well defined
by the core of the central stellar cluster (0.2 - 1.2 pc, Rieke \& Lebofsky
1987, Allen 1994, Rieke \& Rieke 1994, Eckart et al 1995).  This would occur,
for example, if the inward radial velocity of the CND were to be as high as 50
\kms\ during its last stages of inward migration, or if the CND underwent a
strong disk instability during its descent, causing unusually rapid transfer of
angular momentum outwards.  Alternatively, dissipative inter-cloud or
inter-clump collisions might be responsible for sending material rapidly
inwards. The gas arriving in the inner few tenths of a parsec has no exit other
than to accrete onto a compact object (either a supermassive black hole or one
of a collection of smaller ones; Morris 1993, Lee 1995), or to form stars.
However, as discussed in \S4, star formation is inhibited until either the
density grows above a few times 10$^7$ cm$^{-3}$, or the gas
is subjected to a violent shock. Such a shock may be provided by fast-moving
cloud collisions in the small central volume, or by the energy emerging from 
an accretion disk around the supermassive black hole.  The combination of
accretion energy, in the form of radially-propagating winds and radiation,
coupled with the energy released in the formation of dozens, if not hundreds, of
massive stars, reverses the infall of gas, and begins a period of winds
(initially from the massive O stars, and later from the post-main-sequence
WR-type stars), aided considerably by radiation pressure on dust.  This would be
the moment at which an expanding molecular ring, if such exists in our Galaxy,
could be launched. However, gas clouds initially in the near vicinity of the
nucleus can only be pushed to a relatively large radius under the most extreme
circumstances; for a 0.5-pc diameter cloud of density 10$^5$n$_5$ cm$^{-3}$, for
example, the ratio of the outward radiative to the inward gravitational
accelerations at galactocentric distance r$_{pc}$, even assuming a Seyfert-like
central luminosity of 10$^{43}$ ergs, is only
$\sim$0.1n$_5^{-1}$r$_{pc}^{-1}$.  And the ratio of the ram pressure
acceleration of that same cloud caused by a steady wind carrying \.M \Msun\
yr$^{-1}$ at a velocity of 1000v$_3$ \kms\ to the gravitational acceleration is
4\.Mv$_3$n$_5^{-1}$r$_{pc}^{-1}$, so an exceptionally large mass outflow rate
would have to be sustained for an extended time period
($\sim$10$^4$(\.Mv$_3$)$^{-1}$n$_5$r$_{pc}$ yr) in order to give a
significant radial impulse to even a modest cloud starting near the center. 
Alternatively, in place of a steady wind, this could be accomplished with a
succession of at least 10$^2$ - 10$^3$ supernovae occuring within the central
parsec on a similar time scale. If and when these conditions are met, the
radially moving gas mixes with orbiting gas so that its specific angular
momentum rises and it cannot simply fall back in on a free-fall time scale. 

At this stage, the Galactic nucleus would resemble that of an emission-line
galaxy, a starburst galaxy, or even an AGN, depending on how much mass had
accumulated before the reversal took place, and whether the central accretion
disk generates more energy than the luminous, newborn stars. Also, those stars
would resemble the kinds of stellar objects which have been considered as the
source of activity in emission line galaxies (\eg, ``warmers'' Terlevich \&
Melnick 1985, or the stars in Wolf-Rayet galaxies, Conti 1991).

As these massive stars evolve on a time scale of a few million years, their
population declines as some of them become supernovae, and the total stellar
luminosity decreases.  The luminosity of the accretion disk also declines, as it
exhausts the supply of gas with which it was impulsively endowed during the
relatively brief accumulation phase. The surrounding gas transfers its angular
momentum outwards by the processes described in \S3.4 and migrates back in,
against the resistance of the winds from the central stars, for as long as those
stars stay in a windy stage of evolution.  In broad terms, this is the state in
which we now apparently find the Galactic center.  However, the gas is not
smoothly distributed, so the timing of the recurrent accretion episodes and the
manner in which they occur can only be characterized stochastically.  In a
picture proposed by Sanders (1981), activity at the Galactic center is initiated
when a cloud passes close enough to be accreted. We adopt that picture here,
noting that a cloud which is ``captured'' near the Galactic nucleus is likely to
first form a sheared, lumpy and asymmetric circumnuclear disk much like the
observed CND.  The task before us is to unravel the final stages of this
recurrent process by understanding how rapidly the CND is now moving inward and
then how the gas will behave as it descends upon the denizens of the center. 

Interstellar gas is not the only possible source of matter for the accretion
disk around the presumed supermassive black hole at the center, as stars
whose orbits carry them within the central object's Roche limit are
susceptible to tidal disruption (Hills 1975, Ozernoy 1979, Lacy et al 1982, Rees
1988, and references therein).  Some fraction of the resulting debris settles
into a disk which can have a large accretion rate, and a consequent luminosity
approaching the Eddington luminosity.  Phinney (1989) estimates that, in the
Galactic center, this happens about once every 10$^4$ years, with the
accompanying luminosity impulse lasting 10 - 100 yr.  This may affect the x-ray
appearance of the Galactic center by producing, for example, a large flux which
can be scattered in the 6.4-keV iron line (Koyama et al 1996), thereby giving
rise to an iron-line glow which persists for hundreds of years after the
initial, brief, production event.  Another potential consequence of stellar
disruption arises directly from the ejecta created during the tidal disruption
event. Khokhlov \& Melia (1995) note that half the stellar mass is 
asymmetrically ejected with an energy well in excess of that of a supernova, 
and so propose that the energetic and asymmetric Sgr A East shell source was
created in this manner.  This stellar snacking will not typically have much
effect on the limit cycle in which Galactic center gas is engaged, but if it is
timed to occur within the $\wig<$10$^4$-year interval during which gas is
descending at a large rate upon the nucleus, then the pulse of radiation (and
perhaps matter, in the form of a jet or a wind) can play a pivotal role in
inducing star formation within the central parsec and controlling the rate at
which gas accumulates onto the central black hole.

Finally, we remark that a profound shift of emphasis has taken place in Galactic
Center research since the review by Oort (1977).  That review stressed outflows
and violent ejections from the Galactic nucleus, whereas we have traced the
thread of inflow and accretion.  The recognition of the bar-like nature of the
central potential has dramatically reduced the need for explosive origin
models for observed kinematic structures.  Still, winds and jets remain an
important part of the story of Galactic center activity (\eg, Sofue 1994),
although the ultimate driver for that activity is the far larger amount of mass
flowing toward the Galactic center.

%\vskip 0.4cm
\noindent
{\it Acknowledgements}

We thank our many colleagues who avidly contributed with discussions, unpublished data,
preparation of figures,
and/or readings of the preliminary manuscript: R Blandford, J Carlstrom, JA Davidson, 
D Dowell, I Gatley, RH Hildebrand, JH Lacy, DC Lis, N Mastrodemos, S Phinney, DL Shupe, 
S Stolovy, K Uchida, and MCH Wright. 
 We also appreciate little Nicholas Serabyn's unrelenting interest 
in the flavor of the manuscript.

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% check published version of Meglicki et al.
% look up Contopoulos & Grosbol reference, & see if ILR is defined 
%	there, and thus whethre we can use that where we mention ILR here.
% REFERENCE on ages of bulge stars:  rich or rieke, haller
% Tagger message -- find an appropriate reference for asymmetric galaxies
%     'last year review by Sancisi (I believe) and somebody else, in A&A I believe'
% see Gusten & Ungerechts 1985, A&A, 145, 241 for rate of mass loss from bulge stars
\end{document}