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    {\small\rm\zihao{6}\songti\thepage\hfill
    \quad\quad ACTA ASTRONOMICA SINICA \hfill
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    {\small\rm \zihao{6}\songti suppl.\hfill
    K. Ferri$\grave{\rm e}$re: Magnetic fields in galaxies
    %: their origin and their impact on the ISM %%（作者：题目名称）
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    Vol.44~ Suppl.
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\vs\vs\vs \bc {\Large\bf\zihao{2}\heiti Magnetic Fields in
Galaxies: Their Origin
and \\[2mm]Their Impact on the Interstellar Medium    %%（文题）
%\\[1ex]
%$\!\!^*$
}



\vs\vs{\normalsize\zihao{4}\fangsong Katia Ferri$\grave{\rm e}$re}  %%（本文作者）

\vs
(\footnotesize{\it Laboratoire d'Astrophysique de Toulouse, Observatoire
Midi-Pyr$\acute{e}$n$\acute{e}$es, 14 avenue Ed. Belin, 31400 Toulouse, France}) \\%%（本文作者单位）

\vs\vs \ec

\normalsize
\begin{quotation}

\no{\bf A{\scriptsize\bf BSTRACT }}~~~\small
In the first part of the paper, I discuss the origin of galactic
magnetic fields.
After showing that the primordial field theory is in disagreement
with observations, I present the dynamo theory
in which large-scale magnetic fields are amplified under the combined
action of the large-scale galactic differential rotation and small-scale
cyclonic turbulent motions.
I write down the dynamo equation and interpret it physically.
%%I describe the threefold role played by
%%turbulent motions in the dynamo process: generation of magnetic field
%%in the direction perpendicular to the prevailing field (alpha-effect),
%%vertical transport of magnetic field lines, and turbulent magnetic diffusion.
I then present numerical solutions of the dynamo equation for our own
Galaxy, assuming that the required turbulence is driven by supernova
explosions.

In the second part of the paper, I discuss the impact of galactic
magnetic fields on the interstellar medium.
Through the Lorentz force, they affect both the dynamics and
the spatial distribution of the interstellar matter at all scales.
At large scales, they help to support it against
its own weight, while they confine cosmic rays to the galaxy.
At smaller scales, they oppose the expanding gas motions driven
by supernova explosions, they constrain the random motions
of interstellar clouds, and they control the star formation process.
In addition to their dynamical role, they provide
a heat source for the interstellar gas through magnetic reconnection.

\end{quotation}

\vs
\para{\large\bf{1}~~ Introduction}

\vs\vs
Numerous observations, based on starlight polarisation,
Zeeman splitting, Faraday rotation, and synchrotron emission,
reveal the presence of magnetics fields in our Galaxy as well as
in external spiral galaxies. These magnetic fields appear to have
a pressure comparable to the galactic cosmic-ray pressure and
somewhat higher than the interstellar gas pressure.

Two important issues will be addressed in this paper.
The first issue (discussed in \S2) concerns the origin
of galactic magnetic fields:
Are they of primordial origin or the result of a dynamo process?
What is the nature of the turbulent motions necessary for dynamo action?
And how are large-scale magnetic fields predicted to behave
(both temporally and spatially) in the dynamo scenario?
The second issue (discussed in \S3) concerns the impact of galactic
magnetic fields on the interstellar medium (ISM), in particular,
on the dynamics and spatial distribution of interstellar matter
at large and small scales, and on their energetics.


\para{\large\bf{2}~~ The Origin of Galactic Magnetic Fields}

\vs\vn{\bf{2.1~~Primordial Field {\bf\it versus} Dynamo Field}}

Two main theories have been put forward to explain the origin
of magnetic fields in galaxies: the primordial field theory and
the dynamo theory.

In the primordial field theory, present-day galactic magnetic fields
would simply be the relics of a coherent magnetic field
existing in the early Universe prior to galaxy formation.
Presumably, the gas motions associated with the collapse of
a protogalaxy would have compressed the lines of force of the ambient
magnetic field (which tend to be frozen into the highly conductive gas),
and from then on, the differential rotation of the galaxy would have
wrapped them up about its center. In the absence of any other process,
the galactic magnetic field would be wound up 50 to 100 times
at the present time, i.e., much more than indicated by observations.

The traditional way of resolving the discrepancy is to invoke magnetic
diffusion (supposedly of the turbulent kind). But if magnetic diffusion
is sufficiently efficient parallel to the galactic plane to avoid
a tight wind-up of magnetic field lines,
it must also be sufficiently efficient perpendicular to the plane
for field lines to diffuse out of the disk.
%%one may not ignore its effect perpendicular to the plane,
%%which is to let field lines diffuse out of the disk.
Under these conditions, galactic magnetic fields
would have completely decayed away by now,\footnote{
This conclusion could possibly become invalid
in the case of ambipolar diffusion, which allows magnetic field lines,
tied to the charged particles, to slip through the neutral gas.}
\up{1} without ever reaching the observed strengths of a few $\mu$G
(Rosner \& DeLuca 1989).

\begin{figure}[h]
\vspace{-28mm}
\center {\includegraphics [bb=54 360 558 720,scale=0.7]{figure1.eps}}
   \parbox[tb]{12cm}
   %{\hspace{3.5cm}
   {%\center
 \vspace{-45mm}
   \footnotesize
      Fig.\,1~~ Face-on view of the galactic disk showing the azimuthal stretching
of magnetic field lines by the large-scale differential rotation}

    \end{figure}
    \begin{figure}[h]
\vspace{-48mm}
\center {\includegraphics [bb=54 360 558 720,scale=0.7]{figure2.eps}}
   \parbox[tb]{12cm}
   %{\hspace{3.5cm}
   {%\center
   \vspace{-56mm}\footnotesize
      Fig.\,2~~ Oblique view of the galactic disk illustrating the alpha-effect
arising from small-scale cyclonic turbulent motions}

    \end{figure}

  \vspace{-25mm}
  This short reasoning suggests that an additional mechanism
partakes in the generation of galactic magnetic fields.
In the dynamo theory, this additional mechanism is due to
small-scale turbulent motions which are cyclonic, i.e.,
which have acquired a preferred sense of rotation
under the action of the Coriolis force.
As these turbulent motions stretch and twist magnetic field lines,
they impart to them a net rotation, whereby magnetic field is created
in the direction perpendicular to the prevailing field.
This process has been termed the ``alpha-effect''.

    Thus, in the dynamo theory, the large-scale differential rotation
stretches magnetic field lines in the azimuthal direction
about the galactic center (see Figure~1),
while small-scale cyclonic turbulent motions regenerate,
via the alpha-effect, the meridional component of the field
from its azimuthal component (see Figure~2).
It is the combination of these two complementary mechanisms
that leads to magnetic field amplification in galaxies
(Parker 1971; Vainshtein \& Ruzmaikin 1971).

Obviously, the operation of a galactic dynamo
requires a seed magnetic field to initiate the amplification process.
Several possibilities have been advanced regarding the nature
of this seed field (see Rees 1987).
Briefly, the seed field could have a pregalactic origin
(as in the primordial field theory), or it could arise
in the protogalaxy as a result of charge separation due to electrons
interacting with the microwave background photons (battery effect),
or else it could originate in the first generation of stars
and be expelled into the ISM by their winds and/or supernova explosions.


\vs\vn{\bf{2.2~~The Dynamo Equation}}

In the dynamo theory, the time evolution of the large-scale galactic
magnetic field is governed by the dynamo equation,

\begin{equation}
{\partial \langle {\bm B} \rangle \over \partial t}
\, = \, {\bm \nabla } \times \Big( \langle {\bm v} \rangle
\times \langle {\bm B} \rangle \Big)
\, + \, {\bf \nabla } \times \mathcal{E} \ ,
\end{equation}
where
${\bm B}$ is the magnetic field, ${\bm v}$ is the velocity field,

\begin{equation}
\mathcal{E} \equiv \langle \delta {\bm v} \times \delta {\bm B} \rangle
\end{equation}
is the electromotive force due to turbulent motions,
angle brackets denote large-scale (or ensemble-averaged) quantities,
and the symbol $\delta$ denotes small-scale turbulent quantities
(e.g., Steenbeck, Krause, \& R$\ddot{\rm a}$dler 1966).
The first term on the right-hand side of equation~(1) represents
the effect of the large-scale velocity field -- including chiefly
galactic rotation -- on $\langle {\bm B} \rangle$ and the second term
represents the effect of small-scale turbulent motions.

In general, $\cal E$ can be expressed as a linear function of
$\langle {\bm B} \rangle$ and its first-order spatial derivatives:

\begin{equation}
{\cal E}_{i} =
\alpha _{ij} \, \langle B_{j} \rangle
+ \beta _{ijk} \, {\partial \langle B_{j} \rangle \over \partial x_{k}} \ \cdot
\end{equation}
The so-called alpha-tensor, $\alpha_{ij}$, embodies not only
the alpha-effect but also the effective advection of magnetic field
by turbulent motions,
whereas the tensor $\beta_{ijk}$ describes turbulent magnetic diffusion
(e.g., Moffatt 1978).

In the case of a galactic disk, where the parameters of the ISM
and the sources of turbulence vary essentially in the vertical direction,
the alpha-tensor takes on the form
\begin{equation}
\alpha _{ij} =
\left(
\begin{array} {ccc}
\alpha_{R}   & - v_{\rm esc}     & 0 \\
v_{\rm esc}  & \alpha_{\Phi}     & 0 \\
0            & 0                 & \alpha_{Z} \\
\end{array}
\right)
\end{equation}
in a cylindrical reference frame ($\displaystyle \widehat{e}_{R},
\widehat{e}_{\Phi}, \widehat{e}_{Z}$)
with origin at the galactic center and $\displaystyle \widehat{e}_{Z}$
perpendicular to the galactic plane.
The diagonal components,
$\alpha_{R}$, $\alpha_{\Phi}$, and $\alpha_{Z}$, give the effective
rotational velocity associated with the alpha-effect when
$\langle {\bm B} \rangle$ is radial, azimuthal, and vertical, respectively,
and the off-diagonal component, $v_{\rm esc}$, represents the effective
vertical velocity at which $\langle {\bm B} \rangle$
is advected by turbulent motions (Ferri\`ere 1993a).
The diffusivity tensor, for its part, can be written as
\begin{equation}
\beta_{ijk} =
\beta_{h} \,
(\epsilon_{ijR} \, \delta_{kR} + \epsilon_{ij\Phi} \, \delta_{k\Phi})
+ \beta_{v} \, \epsilon_{ijZ} \, \delta_{kZ} \ ,
\end{equation}
where $\beta_{h}$ and $\beta_{v}$ are the horizontal and vertical turbulent
magnetic diffusivities, respectively, $\delta_{ij}$ is the unit tensor,
and $\epsilon_{ijk}$ is the three-dimensional permutation tensor
(Ferri\`ere 1993b).

The rotation curves of spiral galaxies are reasonably well established
observationally, mainly thanks to H~{\sc i} and CO velocity measurements.
In contrast, the dynamo tensors, $\alpha_{ij}$ and $\beta_{ijk}$,
are poorly constrained observationally, and their determination mostly
relies on theoretical models of interstellar turbulence.
For example, Ferri\`ere (1998) calculated $\alpha_{ij}$ and
$\beta_{ijk}$ in our Galaxy on the assumption that the turbulent motions
responsible for dynamo action are driven by supernova explosions.


\vs\vn{\bf{2.3 Numerical Solutions for our Galaxy}}

Several authors have set out to solve the galactic dynamo equation
numerically (e.g., Donner \& Brandenburg 1990; Brandenburg et al. 1992;
Panesar \& Nelson 1992; Schultz, Elstner, \& R\"udiger 1994;
Elstner, R\"udiger, \& Schultz 1996).
Here we present the numerical solutions obtained for our own Galaxy
by Ferri\`ere \& Schmitt (2000),
who adopted the Galactic rotation curve of Fich, Blitz, \& Stark (1989)
together with the supernova-driven dynamo tensors of Ferri\`ere (1998).

Ferri\`ere \& Schmitt (2000) first studied linear modes,
which were allowed to be either axisymmetric
(azimuthal wavenumber $m = 0$) or bisymmetric ($m = 1$) with respect
to the rotation axis, and either symmetric (symbol S) or
antisymmetric (symbol A) with respect to the equatorial plane.
For each mode, they started by computing a reference solution corresponding
to the case when all the input parameters take on their default values.
Then, they systematically varied each of the input parameters, so as to
determine its impact on the temporal evolution and on the spatial structure
of the large-scale magnetic field.

\begin{figure}[h]
%\vspace{-12mm}
\center {\includegraphics [bb=56 417 566 713,scale=0.7]{figure3.eps}}
   \parbox[tb]{12cm}
   %{\hspace{3.5cm}
   {%\center
   \footnotesize
      Fig.\,3~~ Contour lines of the three cylindrical components
of the large-scale magnetic field and lines of force
of its poloidal component, in a given meridional plane of our Galaxy,
for the S0 mode under the reference conditions.}

    \end{figure}

They found that axisymmetric modes are always easier to excite than
bisymmetric modes.
%%Amongst the former, the S0 mode often has the larger growth rate,
%%while the A0 mode generally oscillates more readily.
Under the reference conditions, the S0 and A0 modes have very similar
properties. Both grow monotonously with time at a slow exponential rate
$\simeq 0.45$\,Gyr$^{-1}$, which suggests that the Galactic magnetic field
has presently reached a state close to saturation.
The azimuthal field component dominates by more than one order of magnitude;
as indicated by Figure~3, it exhibits, on each side of the midplane,
two well-defined polarities peaking at $R \simeq 5$\,kpc and separated
by a node at $\vert Z \vert \simeq 4$\,kpc.
The poloidal field peaks at a slightly smaller radius, and its lines
of force rotate about a point located close to the maximum of $\alpha_\Phi$.

When the effective escape velocity, $v_{\rm esc}$, is set to zero,
the exponential growth rate decreases and even turns slightly negative
for the S0 mode.
The magnetic field becomes oscillatory, and its overall pattern
propagates toward the midplane. During most of a dynamo cycle,
$B_\Phi$, which is again the dominant field component, possesses
four alternating polarities along the vertical (see Figure~4).
Oscillatory behaviors are also obtained, albeit with an increased
growth rate, when the diagonal components of $\alpha_{ij}$ are enhanced
by a factor $> 3$ or when the magnetic diffusivities are reduced
by a factor $> 1.7$ with respect to their reference values.

\begin{figure}[h]
%\vspace{-12mm}
\center {\includegraphics [bb=56 417 566 713,scale=0.7]{figure4.eps}}
   %\parbox[tb]{12cm}
   %{\hspace{3.5cm}
   {\center
   \footnotesize
      Fig.\,4~~ Same as Figure~3, with $v_{\rm esc} = 0$}

    \end{figure}

The large-scale rotation rate, $\Omega$, too, has
a considerable influence on the solutions. In particular, when
in addition to being a decreasing function of $R$,
$\Omega$ is assigned a vertical dependence in
$\exp \left[ - (Z / 5 \, {\rm kpc})^2 \right]$,
both the S0 and A0 modes decay oscillatorily with time.
As before, the magnetic field is predominantly azimuthal, but now,
the alternating polarities of $B_\Phi$ (displayed in Figure~5)
migrate along an oblique axis, which can be verified as roughly parallel
to lines of constant $\Omega$.

\begin{figure}[tbph]
%\vspace{-12mm}
\center {\includegraphics [bb=56 417 566 713,scale=0.7]{figure5.eps}}
   %\parbox[tb]{12cm}
   %{\hspace{3.5cm}
   {\center
   \footnotesize
      Fig.\,5~~ Same as Figure~3, with
$\Omega \propto \exp \left[ - (Z / 5 \, {\rm kpc})^2 \right]$}

    \end{figure}
For bisymmetric modes, the azimuthal stretching of magnetic field lines
by the large-scale differential rotation has a destructive, rather than
amplifying, effect. As a result, the magnetic field rapidly vanishes from
the differentially-rotating parts of the Galaxy (outside $R \simeq 4$~kpc),
while it undergoes a much slower exponential decay, accompanied by
an azimuthal propagation, in the rigidly-rotating innermost region.
Note that there exist significant differences between S1 and A1 modes,
not only in the decay rate, but also in the magnetic configuration.

The main limitation of the linear results described above resides in
the fact that the adopted expressions for $\alpha_{ij}$ and $\beta_{ijk}$
are proper to the present-day Galaxy, whereas the true dynamo parameters
have undoubtedly evolved in the course of time, being presumably greater
at the beginning, when the large-scale magnetic field was weaker.
For this reason,
 Ferri\`ere \& Schmitt (2000) carried out a number
of nonlinear calculations
in which $\alpha_{ij}$ and $\beta_{ijk}$ were multiplied by a decreasing
function of $\vert \langle {\bm B} \rangle \vert$, meant to represent
the back-reaction of the Lorentz force on the turbulent motions responsible
for the alpha-effect and for magnetic diffusion.
They found that, for the growing solutions, the magnetic field
amplification saturates
when its maximum intensity reaches $\sim 20 \ \mu$G, corresponding to
a magnetic pressure roughly equal to four times the local gas pressure.
The time to saturation, which depends on the seed field strength adopted,
is typically on the order of a few 10~Gyr.
%%In some cases, nonlinear mode interactions produce long-term changes
%%in the symmetric {\it vs.} antisymmetric vertical parity
%%or in the monotonous {\it vs.} oscillatory temporal behavior.


\para{\large\bf{3}~~ The Impact of Galactic Magnetic Fields on the ISM}

\vs\vn{\bf{3.1~~Dynamic Role}}

The magnetic field of a galaxy acts on the interstellar matter
through the Lorentz force.
Of course, the field acts directly on the charged particles only,
but its effect is then transmitted to the neutrals by ion-neutral
collisions.
Apart from the densest parts of molecular clouds, whose ionization
degree is exceedingly low,
virtually all interstellar regions are sufficiently ionized
for their neutral component to remain tightly
coupled to the charged component and, hence, to the magnetic field.

At large scales, the galactic magnetic field helps to support
the ordinary matter
against its own weight in the galactic gravitational potential,
and it confines cosmic rays to the galactic disk.
In this manner, both magnetic fields and cosmic rays partake in the overall
hydrostatic balance of the ISM and influence its stability.
Boulares \& Cox (1990) were the first authors to fully appreciate
the importance of magnetic fields and cosmic rays in the hydrostatic balance.
By the same token, they managed to solve the long-standing problem
of apparent mismatch between the total interstellar pressure at a given point
and the integrated weight of overlying interstellar material: by adopting
higher magnetic and cosmic-ray pressures than previously estimated, they were
able to bring the total pressure at low $\vert Z \vert$ into agreement
with the integrated weight, and by including the magnetic tension force
at high $\vert Z \vert$, they could explain why the weight integral
falls off faster than the total pressure.

As magnetic and cosmic-ray pressures inflate the gaseous disk, they tend
to make it unstable to a generalized Rayleigh-Taylor instability,
now known in the astrophysical community as the Parker instability
(Parker 1966).
When this instability develops, magnetic field lines ripple,
and the interstellar matter
slides down along them toward the magnetic troughs, where it accumulates.
This whole process, it has been suggested, could give birth to new
molecular-cloud complexes and ultimately trigger star formation
(Mouschovias, Shu, \& Woodward 1974; Elmegreen 1982).

At smaller scales, the galactic magnetic field affects all kinds
of turbulent motions in the ISM.
Of special importance is its impact on supernova remnants
and superbubbles (see Tomisaka 1990; Ferri\`ere, Mac Low, \& Zweibel 1991;
Slavin \& Cox 1992).
First, the background magnetic pressure acting on their surrounding
shells directly opposes their expansion.
Second, the magnetic tension existing in the field lines swept into
the shells gives rise to an inward restoring force,
while the associated magnetic pressure prevents the shells
from fully collapsing and, therefore, keeps them relatively thick.
Third, the enhanced external ``signal speed" causes the shells to merge
earlier than they would in an unmagnetized medium.
All three effects conspire to lower the filling factor of hot
interstellar gas.

The galactic magnetic field also constrains the random motions of
interstellar clouds. The latter are magnetically connected
to their environment,
namely, to the intercloud medium and possibly to neighboring clouds,
through the magnetic field lines that thread them. When a given cloud moves
relative to its environment, these field lines get deformed,
and the resulting magnetic tension force modifies the cloud's motion,
transferring part of its momentum to its environment (Elmegreen 1981).
Likewise, angular momentum can be transferred from a rotating cloud to
its environment by magnetic torques (Mouschovias \& Paleologou 1979).
This mechanism is particularly relevant to the star formation process,
as it allows the contracting protostellar cores to get rid of angular
momentum (e.g., Mouschovias \& Morton 1985).

Finally, the galactic magnetic field plays a crucial role
in the support of molecular clouds against their self-gravity
and in the eventual gravitational collapse of protostellar cores.
The magnetic support of molecular clouds is essentially provided by
magnetic pressure gradients in the directions perpendicular to
the mean field and, presumably, by nonlinear Alfv\'en waves
in the parallel direction (Shu, Adams, \& Lizano 1987).
In the case of protostellar cores, magnetic support is insufficient
to prevent their ultimate gravitational collapse.
This is generally because their ionization degree is so low that
neutrals are not perfectly tied to magnetic field lines,
which enables them to drift inwards under the pull of their self-gravity
and eventually form stars (Nakano 1979; Mestel 1985).


\vs\vn{\bf{3.2~~Energetic Role}}

In most regions of interstellar space, the electrical conductivity,
$\sigma$, is extremely high, with the twofold consequence that magnetic
field lines are frozen into the plasma and that ohmic dissipation is
completely negligible.
There exist, however, localized regions where magnetic field gradients
are very steep, so that the electric current density,
${\bm J} = {c \over 4 \pi} \, {\bm \nabla} \times {\bm B}$, is very strong,
and where the effective electrical conductivity is greatly reduced with
respect to its classical Coulomb value (for instance, an anomalous
conductivity may result from microscopic wave fluctuations which increase
the effective collision frequency through wave-particle interactions).
The combination of strong ${\bm J}$ and reduced $\sigma$ entails both
a partial loss of the magnetic field frozen-in property and a resistive
dissipation of magnetic energy.

In particular, when field lines of opposite polarities are forced
to approach one another, a thin transition layer develops, in which
the magnetic field gradient steepens and, accordingly, the current
density rises, to the point where field lines decouple from the plasma,
reconnect with field lines of the opposite polarity, and leave
the transition layer sideways.
The reconnection zone is subject to intense ohmic dissipation, with
the dissipated magnetic energy being essentially converted into plasma
thermal energy.

Thus, magnetic reconnection constitutes a source of heating
-- and, at the same time, ionization -- for the ISM.
This heating/ionization source could even be dominant in the galactic
halo, which is not easily accessible to the powerful radiation from
the luminous O and B stars.
Based on this idea, Birk, Lesch, \& Neukirch (1998) invoked magnetic
reconnection within thin current filaments assumed present in the halo
to solve the ionization/heating problem of the warm ionized medium at
high $Z$.
It was also suggested by Zimmer, Lesch, \& Birk (1997) that interaction
regions between high-velocity clouds traveling through the halo and
the surrounding halo gas could be the sites of intense magnetic reconnection
and, hence, plasma heating.

\bc
{\bf References}
\ec

\vspace{-4mm}
\no \fns

 \nref Birk G T, Lesch H, Neukirch T.  MNRAS, 1998, 296: 165
\nref Boulares A, Cox D P. ApJ, 1990, 365, 544 \nref Brandenburg
A, Donner K J, Moss D, et al. A\&A, 1992, 259: 453 \nref Donner K
J, Brandenburg A. A\&A, 1990, 240: 289 \nref Elmegreen B G.  ApJ,
1981, 243: 512 \nref Elmegreen B G.  ApJ, 1982, 253: 655 \nref
Elstner D, R\"udiger  G, Schultz M. A\&A, 1996, 306: 740 \nref
Ferri\`ere K M.  ApJ, 1993a, 404: 162 \nref Ferri\`ere K M.  ApJ,
1993b, 409: 248 \nref Ferri\`ere K M.  A\&A, 1998, 335: 488 \nref
Ferri\`ere K M, Mac Low M-M, Zweibel E G. ApJ, 1991,375: 239 \nref
Ferri\`ere K M, Schmitt D. A\&A, 2000, 358: 125 \nref Fich M,
Blitz L, Stark A A. ApJ, 1989, 342: 272 \nref Mestel L. in
Protostars and Planets II. In: Black D C \& Matthews M S eds.
Tucson: University of Arizona, 1985, 320 \nref Moffatt H K.
Magnetic Field Generation in Electrically Conducting Fluids
(Cambridge: Cambridge University Press), 1978, 145 \nref
Mouschovias T C, Morton S A.  ApJ, 1985, 298: 190 \nref
Mouschovias T C, Paleologou E V.  ApJ, 1979, 230: 204 \nref
Mouschovias T C, Shu F H,  Woodward P R. A\&A, 1974, 33: 73 \nref
Nakano T.  PASJ, 1979, 31, 697 \nref Panesar J S, Nelson A H.
A\&A, 1992, 264: 77 \nref Parker E N.  ApJ, 1966, 145: 811 \nref
Parker E N.  ApJ, 1971, 163: 255 \nref Rees M J.  QJRAS, 1987, 28:
197 \nref Rosner R,  DeLuca E.  in IAU Symp. 136, The Center of
the Galaxy, ed. M. Morris (Dordrecht: Kluwer), 1989, 319 \nref Shu
F H, Adams F C, Lizano S.  ARA\&A, 1987, 25: 23 \nref Schultz M,
Elstner D, R\"udiger G. , A\&A, 1994, 286: 72 \nref Slavin J D,
Cox D P.  ApJ, 1992, 392: 131 \nref Steenbeck M, Krause F,
R\"adler K-H. Z. Naturforsch., 1966, A21: 369 \nref Tomisaka K.
ApJ, 1990, 361: L5 \nref Vainshtein  S I, Ruzmaikin A A.  Astron
Zh, 1971, 48: 902 [Soviet Ast., 15, 714 (1972)] \nref Zimmer  F,
Lesch H, Birk G T.  A\&A, 1997, 320: 746


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