We suggest that PSR B1509-58 and SNR MSH15-52 are the remnants
of the SN explosion of a massive star (Gvaramadze 1999b). In this case,
the structure of MSH15-52 could be determined by the interaction
of the SN blast wave with the ambient medium reprocessed by the
joint action of the ionizing emission and stellar wind of the SN
progenitor star (McKee et al. 1984; Shull et al. 1985; Ciotti & D'Ercole 1989; Chevalier & Liang 1989;
Franco et al. 1991; D'Ercole 1992; Gvaramadze 1999b,c, 2000a). The outer
shell of the SNR could arise due to the abrupt deceleration of the
SN blast wave after it encounters the density jump at the edge of
the bubble created by the fast stellar wind during the
main-sequence or the Wolf-Rayet (WR) stages. On the other hand,
some structures in the central part of the SNR
could be attributed to the interaction of the SN blast wave with
the circumstellar material lost during the late evolutionary
stages of the SN progenitor star [this is the material that
determines the appearance of young typeII SNRs, e.g. SN1987A
(e.g. McCray 1993) or CasA
(e.g. Garcia-Segura et al. 1996; Borkowski et al. 1996)]. During
the red supergiant (RSG) stage a massive star loses a significant
part (about two thirds) of its mass in the form of a slow, dense
wind. This matter occupies a compact region with a characteristic
radius of few parsecs (the high-pressure gas in the main-sequence
bubble significantly affects the spreading of this region,
Chevalier & Emmering 1989; D'Ercole 1992). Before the SN exploded, the progenitor
star (of mass >
)
becomes for a short time a
WR star (e.g. Vanbeveren et al. 1998). At this stage, the fast stellar wind
sweeps up the slow RSG wind and creates a low-density cavity
surrounded by a shell of swept-up circumstellar matter. The shell
expands with a nearly constant velocity
,
where
and
are, correspondingly, the mass-loss
rates and wind velocities during the WR and RSG stages (e.g.
Dyson 1981), until it catches up the shell separating the RSG wind
from the main-sequence bubble. For parameters typical for RSG and
WR winds, one has
.
The interaction of two circumstellar shells results in
Rayleigh-Taylor and other dynamical instabilities, whose
development is accompanied by the formation of dense clumps moving
with radial velocities of
(Garcia-Segura et al. 1996). The
dense clumps could originate much closer to the SN progenitor star
due to the stellar wind acceleration during the transition from
the RSG to the WR stage (Brighenti & D'Ercole 1997). The number density of clumps
is
10
provided they are not fully
ionized and were able to cool to a temperature of
102 K
(Brighenti & D'Ercole 1997). Direct evidence of the existence of high-density
clumps close to the SN explosion sites follows from observations
of young SNRs. For example, the optically emitting gas of
quasi-stationary flocculi in CasA is characterized by a
density of
10
and a temperature of
104 K (e.g. Lozinskaya 1992). Assuming that the optical emission
of a floccule comes from an ionized "atmosphere" around the
neutral core, one can estimate the density of the core to be
10
,
provided that the temperature of the core
is
102 K. Similar estimates could also be derived from
observations of the optical ring around SN1987A, the
inner ionized "skin" of which has nearly the same parameters
(e.g. Plait et al. 1995) as the optically emitting gas of flocculi in
CasA, or from observations of some other young SNRs (e.g.
Chugai 1993; Chugai & Danziger 1994). The radial velocity of flocculi in CasA
ranges from
80 to
(e.g. Lozinskaya 1992).
Initially, the new-born pulsar moves through the low-density
cavity created by the fast wind of the presupernova star until it
plunges into the first dense clump on its way. This happens at the
moment
,
where
-2 pc is the radius of the cavity,
,
and
are respectively the velocities
of the pulsar and the clump. For
and
,
one has
-
years
. Let us assume that all matter
captured inside the accretion radius
(
,
where
is the
sound speed in the cold, dense circumstellar clump) of the pulsar
moving through the clump penetrates into the region of the light
cylinder, where it is accelerated to relativistic velocities and
then leaves this region in the form of equatorial outflow (cf.
Istomin 1994; King & Cominsky 1994). The rate at which the ambient
medium arrives at the light cylinder could be estimated as
For accretion to occur, the standoff radius
of the bow
shock (formed by the outflow of relativistic particles) should be
less than
.
For the spherically symmetric outflow,
one has
,
where
we assume that only a fraction
of the spin-down
luminosity
is transferred to the ambient medium (cf.
Kochanek 1993). This condition can be re-written as
(cf. Kochanek 1993; Manchester et al. 1995), i.e.
should be much smaller than the usually adopted value of
1 (e.g. Kulkarni & Hester 1988; Cordes et al. 1993). Weak coupling (
)
of
the pulsar wind with the ambient medium is consistent with an
outflow composed of highly relativistic particles (e.g. Kochanek 1993
and references therein). Alternatively, if the outflow of
relativistic particles is confined to the vicinity of the
rotational equatorial plane, one can expect that the ambient
matter accretes onto the pulsar's magnetosphere along the polar
directions. Another possibility is that the ambient matter
penetrates in the pulsar wind bubble through instabilities in the
bow shock front. In the latter both cases
could be
larger than
,
and one can adopt
(see next section).
Copyright ESO 2001