Transient radio emission from the PSR B1259–63 system near periastron

Transient radio emission from the PSR B1259ÿ63 system near periastron

Simon Johnston,1 R. N. Manchester,2 D. McConnell2 and D. Campbell-Wilson3

1Research Centre for Theoretical Astrophysics, University of Sydney, NSW 2006, Australia2Australia Telescope National Facility, CSIRO, PO Box 76, Epping, NSW 2121, Australia3Astrophysics Department, School of Physics, University of Sydney, NSW 2006, Australia

Accepted 1998 September 4. Received 1998 September 4; in original form 1998 March 23

A B S T R A C T

Observations of the pulsar±Be-star binary system PSR B1259ÿ63 around the 1994 and 1997

periastron encounters, using the Australia Telescope Compact Array at four frequencies

between 1.4 and 8.4 GHz, and the Molonglo Observatory Synthesis Telescope at 843 MHz,

show a relatively strong and rapidly varying but unpulsed radio source. The source appears 22

d before periastron and lasts for more than 100 d. Over most of the event the spectral index is

, ÿ 0:5, indicating an optically thin synchrotron source. The peak ¯ux of ,50 mJy at 1 GHz

occurs about 20 d after periastron. There are strong similarities in the light curves of the 1994

and 1997 periastron passages. The radio light curve shows day-to-day variability and varies by

a factor of more than 100 over the observation interval. This is in contrast to the X-ray emission

which varies by only a factor of 2 close to periastron. The emission is interpreted as

synchrotron emission from relativistic electrons accelerated in shocks at the interface between

the pulsar wind and the Be-star emission-line disc.

Key words: radiation mechanisms: non-thermal ± binaries: close ± stars: emission-line, Be ±

pulsars: individual: PSR B1259±63 ± radio continuum: stars.

1 I N T R O D U C T I O N

PSR B1259±63 is a unique binary pulsar. It was discovered in a

Galactic plane survey at 20 cm (Johnston et al. 1992a) and is in a

highly eccentric (e , 0:9), 3.5-yr orbit around a 10th-magnitude Be

star, SS 2883 (Johnston et al. 1992b). It lies at a distance of 1.5 kpc

in the Sagittarius±Carina spiral arm. Optical observations (John-

ston et al. 1994) show that SS 2883 is of spectral type B2e, and thus

has a mass of ,10 M( and a radius R�,6 R(. Assuming a pulsar

mass of 1.4 M(, the implied inclination angle of the binary orbit to

the plane of the sky is 368. The Ha emission line shows that the

emission disc of the Be star extends to at least 20R�, just inside the

orbital radius of the pulsar at periastron. Timing measurements have

shown that the disc of the Be star is likely to be highly inclined with

respect to the orbital plane (Wex et al. 1998).

In this paper we will discuss the unpulsed radio emission from the

system based on observations made between 1994 January and

1997 July. This range spans two periastron passages of PSR

B1259ÿ63 with its companion star. We will denote the epoch of

the ®rst, on 1994 January 9, by T 1 and that of the second, on 1997

May 29, by T 2. The generic epoch of periastron is denoted by T .

Pulsed radio observations of the system were made for several

months on each side of the 1994 periastron event (Johnston et al.

1996). In brief, these results showed depolarization of the pulses at

T 1ÿ100, dispersion measure increase and scatter-broadening of

the pulses around T 1ÿ25, an eclipse of the pulsed emission at

frequencies up to 8.4 GHz from T 1ÿ18 to T 1+17, and extremely

large (negative) rotation measures varying on short time-scales near

T 1+30. Similar results were seen during the 1997 periastron

passage. These will be described in detail in a subsequent paper.

High-energy observations of the binary system were made close

to the 1994 periastron. The OSSE instrument on board the Compton

Gamma Ray Observatory (CGRO) made a 3-week observation

from T 1ÿ6 until T 1+14 (Grove et al. 1995). Following an earlier

detection of X-ray emission from the system around apastron with

ROSAT (Cominsky, Roberts & Johnston 1994), the ASCA satellite

observed PSR B1259ÿ63 on days T 1ÿ12, +1, +17 (Kaspi et al.

1995) and +50 (Hirayama et al. 1996). In summary, these results

show no evidence for pulsed emission, a power-law spectrum with

photon index ,1.6 between 10 and 200 keV, and evidence for ¯ux

variability by a factor of ,2. Tavani & Arons (1997) have modelled

the X-ray observations in considerable detail. They conclude that

the X-ray emission is synchrotron radiation from the pulsar wind

shock associated with the interaction between the pulsar wind and

both the wind and disc of the Be star.

The PSR B1259ÿ63 system bears some resemblance to the Be/

X-ray binary system LSI +61±303, which contains a B0e star and a

neutron star in an eccentric, 26-d orbit. Periodic non-thermal radio

emission is detected from the system (Taylor & Gregory 1982), and

X-ray outbursts are also detected (Taylor et al. 1996). LSI +61±303

is likely to be associated with a COS B g-ray source. Unlike other

high-mass X-ray binary systems, LSI +61±303 shows only low-

level X-ray emission (Lx , 1033 erg sÿ1) and little high-energy

radiation, both of which indicate that accretion on to the surface of

Mon. Not. R. Astron. Soc. 302, 277±287 (1999)

q 1999 RAS

the neutron star does not occur. The radio emission is explained as

either synchrotron emission from the boundary between the

(young) neutron star wind and the Be-star wind (Maraschi &

Treves 1981) or as emission from expanding jets resulting from

super-Eddington accretion on to the neutron star (Taylor & Gregory

1984; Paredes et al. 1991).

Non-thermal radio emission is also observed from massive early-

type stars (e.g. Bieging, Abbott & Churchwell 1989) and binary

systems (e.g. WR 140: Williams, van der Hucht & TheÂ 1987). In

both cases, the radio emission is interpreted as synchrotron emis-

sion from relativistic electrons accelerated in shocks in the wind; for

single stars the shocks are attributed to turbulence in the wind

(White 1985), whereas in the binary systems the shocks are at the

interface between the winds from the two stars (Eichler & Usov

1993, and references therein).

We report here on observations of the PSR B1259ÿ63 system in

the radio continuum at frequencies between 843 and 8640 MHz

from T 1ÿ1 to T 2+100. The original motivation for carrying out

these observations was to detect unpulsed emission from the pulsar

during the periastron passage, in order to distinguish between

different possible eclipse mechanisms. However, it was immedi-

ately clear that the source was much stronger than the pulsar, and so

the observations were extended to study the source and its time

variations.

278 S. Johnston et al.

q 1999 RAS, MNRAS 302, 277±287

Figure 1. Images of the PSR B1259ÿ63 ®eld obtained at the four ATCA frequencies (from the top: 1.38, 2.38, 4.80 and 8.40 GHz). The pulsar system is the left-

hand source in each frame, and the beamsize is shown in the lower right-hand corner. Each image has the contour levels set at ÿ3, 3, 5, 10, 20, 30, 50, 100, 150,

200 and 500j above the noise level, which is 0.2, 0.12, 0.09 and 0.17 mJy for the four frequencies, respectively.

2 O B S E RVAT I O N S

Observations at 843 MHz were made using the Molonglo Observa-

tory Synthesis Telescope (MOST) of the University of Sydney. The

Australia Telescope Compact Array (ATCA) was used for observa-

tions at frequencies near 1.4, 2.4, 4.8 and 8.4 GHz. On the ATCA,

observations are made simultaneously at either 1.4 and 2.4 GHz or

4.8 and 8.4 GHz, with a bandwidth of 128 MHz at each frequency, a

spectral resolution of 8 MHz and full Stokes parameters. Frequency

switching between pairs of bands was carried out every ,20 min.

One of the primary ATCA ¯ux calibrators, 0823ÿ500 or

1934ÿ638, was observed once during each session. The ¯ux

densities of 0823ÿ500 are 5.5, 5.6, 3.1 and 1.4 Jy, and those of

1934ÿ638 are 14.9, 11.6, 5.8 and 2.8 Jy at the four observing

frequencies, respectively (Reynolds, private communication). The

secondary (phase) calibrator, 1251ÿ713, was observed every ,20

min for ,3 min at each frequency. The ¯ux density of 1251ÿ713 is

close to 1 Jy at all four observing frequencies.

Images of the ®eld at the four ATCA observing frequencies are

shown in Fig. 1. These images result from the 12-h synthesis on

1994 February 10, and were formed in the standard fashion using

the imaging package MIRIAD. The position of the unpulsed source is

coincident with the timing position of the pulsar. This observation

and a second 12-h synthesis on 1994 March 9 allowed us to

determine the ¯ux density of the confusing source 54 arcsec west

of the pulsar. (The J2000 coordinates of the brightest region of the

source are right ascension 13h 02m 40s, declination ÿ638 509 160.)The source is extended (probably a compact double extragalactic

source) and has ¯ux densities of 19.7, 13.2, 5.8 and 3.0 mJy at our

four observing frequencies. These yield a spectral index of ÿ0.9.

We also recorded the clean components of the confusing source at

each frequency.

For most of the 1994 observations, we obtained only very short

integrations, typically 1 h or less. For these short integrations the

following technique was developed to determine the ¯ux density of

the PSR B1259ÿ63 system. Flux densities were calibrated using a

primary calibrator, and 1251ÿ713 was used for phase and bandpass

calibration. The clean-component model of the confusing source

was then removed from the uv visibilities with the routine UVMODEL,

and then UVFIT was used to calculate the ¯ux density of the source

associated with PSR B1259ÿ63 and the rms phase error during

each ,20-min integration. If the phase error exceeded 708, the

integration was rejected. For phase errors less than 708, the

computed ¯ux density was increased by a factor of 1/cos(rms

phase error). The effectiveness of this technique was determined

by subdividing the 12-h synthesis observation into 20 min integra-

tions. Applying the correction yielded a ¯at ¯ux-density curve over

the 12 h. This correction was particularly important for the high

frequencies, where observations during the daytime in the summer

often have high rms phase ¯uctuations. Occasionally, phases were

bad on the short baselines because of interference; in these cases

only the long (,6-km) baselines were used to determine the ¯ux

density.

During 1996, binning of data synchronously with the period of a

pulsar was implemented on the ATCA. This facility was used

throughout the 1997 observations, giving us eight time-bins

across the pulsar period. Since we also had 32 frequency channels

across each band, the pulsar signal could be dedispersed during off-

line analysis, and `on-pulse' and `off-pulse' images formed. A `total

¯ux density' image, summing across all bins, was also obtained. All

Transient radio emission from PSR B1259ÿ63 279

q 1999 RAS, MNRAS 302, 277±287

Figure 2. An image of the PSR B1259ÿ63 ®eld obtained at the MOSTat 843 MHz. The pulsar system and the extragalactic source are in the centre of the frame.

The rms level is 0.8 mJy; contours are at 3, 6, 9, 15, 20, 30, 45, 60, 75, 150 and 300j.

observations during 1997 were of at least 4-h duration. Conven-

tional techniques were used to produce images from these data.

Flux densities of the pulsar and unpulsed source were then calcu-

lated using the routine IMFIT.

The MOST is a synthesis telescope comprising two collinear

cylindrical±paraboloid re¯ectors, each 11.6 by 778 m2, and operates

at 843 MHz with a detection bandwidth of 3.25 MHz. The antenna

measures right circular polarization at transit as de®ned by the IEEE

standard. At the declination of the pulsar, the synthesized beam is

approximately 44 by 49 arcsec2. The smallest ®eld size available of

23 by 26 arcmin2 was used for all observations described here. This

gives a 5j noise level in one 12-h synthesis of 2.0 mJy (in 1994) and

1.5 mJy (in 1997).

The telescope forms a comb of 64 real-time fan beams, which

makes it dif®cult to calibrate incomplete syntheses. This was

overcome by making a reference image at a time far from perias-

tron. For each observation, the reference image was re-synthesized

with the same hour angle coverage as the pulsar, and subtracted

from the image. Variations in the telescope beam pro®le and

pointing over the observing period were estimated by short scans

of point sources made pre- and post-observation. The absolute ¯ux

density scale was established from observations of 1934ÿ638 and

0409ÿ75, which have ¯ux densities at 843 MHz of 13.65 and 19.8

Jy respectively.

Fig. 2 shows a 12-h synthesis image of the region made with the

MOST on 1994 January 16. The extragalactic source mentioned

above is the main source of confusion in the MOST images, as it is

just over one synthesised beamwidth from the pulsar. Its ¯ux

density at 843 MHz is 31.0 mJy, consistent with its spectral index

derived from the ATCA. The MOST made a total of 29 complete or

part observations during 1994, and a further 15 observations during

1997.

3 R E S U LT S

As expected, the source was unresolved at all epochs and all

frequencies, with the ATCA observations at 8.4 GHz putting an

upper limit of ,1 arcsec on the source size. At a distance of 1.5 kpc,

the entire orbit has an angular diameter of only , 4 mas. At times

when the pulsar emission is eclipsed, there is no evidence of

polarized emission in any of the observations in either 1994 or

1997. For the 1997 data, the approximate upper limit is 0.3 mJy (3j)

at all frequencies. This is less than 1 per cent of the ¯ux density at

maximum.

Measured ¯ux densities and their associated 2j errors are given

in Table 1 for the observations from 1994 to the end of 1996. The

quoted errors include some estimation of systematic (e.g. calibra-

tion) and residual confusion errors. Even away from periastron, the

¯ux densities vary considerably, especially at the two lower

frequencies. This may be partly a result of source confusion, but

is more likely to result from long-term refractive scintillation of the

pulsar signal (Johnston et al. 1996).

Table 2 lists the ¯ux densities measured during 1997. The

array designation is also listed for each observation. In all

observations the sixth antenna, located 3±6 km distant from

the remaining ®ve, was always used. For the ATCA, pulse

binning allowed separate determination of the pulsar and

unpulsed ¯ux densities, and these are given in the table along

with the total ¯ux density. At all four frequencies, the on-pulse

and off-pulse images each consisted of four time-bins. At the

lowest frequency, this is likely to underestimate the pulsed

fraction and overestimate the unpulsed fraction, as the pulsar

shows some emission over virtually the whole pulse period

(Manchester & Johnston 1995).

For observations taken with the 375-array (from T 2ÿ51 to ÿ36

inclusive) and the 122-array (T 2+42 and +47), we discarded the

short baselines and imaged with only the ,6-km baselines to avoid

confusion with the extragalactic source. Data taken on 1997 April

15 are badly affected by interference, and no results were obtained.

On 1997 April 28 and May 8 (T 2ÿ31 andÿ21), the pulsar was very

scattered at 1.4 GHz; it was not possible to make a separate on- and

off-pulse image. On 1997 May 2 (T 2ÿ27), observations made with

the Parkes telescope showed that the pulsar was very weak at all

four frequencies; and neither was there evidence of pulsed emission

seen in the ATCA data. From 1997 May 11 until June 9 (T 2 ÿ18 to

T 2+11), there is no evidence for any pulsed emission at any

frequency. On 1997 June 19 and 22 (T 2+21 and +24), observations

were made at the lower frequencies only and pulsar binning was not

in operation. On 1997 July 1 (T 2+33), no pulses were obvious in

the data at 4.8 and 8.4 GHz; we conclude that the pulsar binning

mechanism was not sampling synchronously with the pulsar period.

This could be due either to an ephemeris error or possibly to the

failure to add a leap-second to the observatory clocks on 1997 June

30 (Sault, private communication).

3.1 The 1994 light curves

Observed ¯ux densities at the ®ve frequencies in 1994 are shown in

Fig. 3. Observations of the unpulsed source began on the day of

periastron when, unexpectedly, the observed ¯ux densities were in

the range 10±30 mJy, much greater than the expected pulsar ¯ux

density of 2±3 mJy. Apart from a minor peak at about T 1+6, the

¯ux density remained relatively constant at all frequencies until T 1

+14. At this date, a steep rise (especially at the lower frequencies)

began, peaking at T 1+25. During this rise, the low-frequency ¯ux

density doubled, and at 4.8 and 8.4 GHz it rose by about 50 per cent.

The ¯ux density then dropped rapidly until about T 1+32.

Following this, the decay was much slower; the source

remained signi®cant in the ATCA data until about T 1+60 and

persisted in the MOST data until after T 1+100. Single-dish

observations show that the pulsar was absent at all frequencies in

data taken on T 1+14 and was present on T 1+24. The re-

appearance of the pulsar thus approximately coincides with the

maximum of the radio light curve.

3.2 The 1997 light curves

Observed ¯ux densities of the PSR B1259ÿ63 system (i.e. both the

pulsar and the unpulsed source) before and after the 1997 periastron

are shown in Fig. 4. (The system ¯ux density is plotted in preference

to the unpulsed source value for compatibility with Fig. 3, and

because the two components were not separately measured at some

epochs.) Until T 2ÿ27 there is no evidence for any unpulsed

radiation in the off-pulse bins down to a limit of ,1.0 mJy. At

T 2ÿ27 we see the ®rst sign of off-pulse emission with a ¯ux

density of ,2.5 mJy at 1.4 GHz. The ¯ux density then abruptly rose

to ,15 mJy at all four ATCA frequencies at T 2ÿ21, before

declining again to only ,5 mJy three days later. The ®rst MOST

observation showing signi®cantly increased emission was also on

day T 2ÿ21, when the 843-MHz ¯ux density was signi®cantly less

than that at higher frequencies. Following the short-lived dip at

T 2ÿ18, the light curves then rose somewhat irregularly to a plateau

at about T 2 ÿ8. Flux densities are lower at higher frequencies, and

,25 mJy at 1.4 GHz during the plateau phase. At T 2+8, the ¯ux


q 1999 RAS, MNRAS 302, 277±287

density rose again, peaking at 57 mJy at 1.4 GHz at T 2+20. There

was then a rapid decline at all frequencies to about T 2+30, when

the ¯ux densities were about 60 per cent of their plateau value.

Subsequently there was a more gradual decline which continued

until the last day of observations at T 2+47, when the 1.4-GHz ¯ux

density was still about 18 mJy.

The disappearance of the pulsed emission coincides with the

®rst peak in the radio light curve at T 2ÿ21. The pulsar was not


q 1999 RAS, MNRAS 302, 277±287

Table 1. Flux densities of the PSR B1259ÿ63 system: 1994±96.


q 1999 RAS, MNRAS 302, 277±287

Table 2. Flux densities of the PSR B1259ÿ63 system in 1997.


q 1999 RAS, MNRAS 302, 277±287

Table 2 ± continued

re-detected until T 2+16, which is somewhat before the prominent

post-periastron peak in the light curve.

3.3 Spectral index evolution

We ®tted a power-law spectrum (i.e. Sn ~ na; two free parameters)

to all observations made at three or more frequencies. In 1994 we

®tted all data from T 1 0 to +60, and in 1997 we ®tted all data from

T 2ÿ21 to +47. In almost all cases we ®nd that a simple power law

®ts the data well; the derived spectral indices are plotted in Fig. 5.

During the 1997 periastron, the spectral index was initially very

small, but changed quickly to a stable value near ÿ0:6 by T 2ÿ8.

There seems to be good evidence for a second rise in the spectral

index between T 2+18 and +30; this is seen in both the 1994 and

1997 data.

On some occasions when 843-MHz observations were simulta-

neous with ATCA observations, the ¯ux density at 843 MHz

appears somewhat lower than expected from the power-law ®t.

This may result from free±free absorption of at least part of the

synchroton emission by the stellar disc.


q 1999 RAS, MNRAS 302, 277±287

Figure 3. Flux densities of the PSR B1259ÿ63 system following the 1994 periastron.

Table 2 ± continued

3.4 Comparison between 1994 and 1997

Comparison of the 1994 and 1997 light curves shows that the gross

features are similar. At both epochs we see a relatively ¯at light

curve around periastron, a peak near 60 mJy at the lowest frequency

about 20 d after periastron follwed by a rapid decay for ,10 d, and

then a slow decay until at least 50 d after periastron. However, the

peak in the unpulsed emission occurs about 5 d earlier in 1997 than

in 1994, and both last for about 10 d. The spectral index as a

function of epoch is also similar for both periastrons. At both

epochs we see a steepening of the spectral index as the emission

fades towards T +30.

4 D I S C U S S I O N

Pulsar timing observations (Wex et al. 1998) suggest that the disc is

signi®cantly tilted with respect to the plane of the orbit, and that the

line of intersection of the disc plane and the orbital plane is almost

perpendicular to the orbital major axis (see Fig. 6). If this is the case,

then the pulsar crosses the disc twice, once pre-periastron and once


q 1999 RAS, MNRAS 302, 277±287

Figure 5. Spectral index evolution around the 1994 periastron (dashed line) and the 1997 periastron (solid line).

Figure 4. Flux densities of the PSR B1259ÿ63 system around the 1997 periastron.

post-periastron. We can estimate the disc-crossing epoch from the

disappearance and re-appearance of the pulsed ¯ux which occur

roughly at T ÿ20 and +16. This implies that the disc is rather thin,

with an opening angle of less than 108. The pulsar intersects the disc

plane at about T ÿ18 and T +13 (A and B in Fig. 6), and is behind

the disc (with respect to the observer) between these times. It takes

approximately 3 d for the pulsar to cross the disc.

The rapid time variations and spectral index of the observed radio

emission show that it is almost certainly optically thin synchrotron

emission from relativistic electrons accelerated in shocks associated

with the interface between the pulsar wind and the Be-star out¯ow.

There are several factors as to why the radio emission is likely to

come from the interaction of the pulsar wind with the Be-star disc,

rather than with the isotropic (polar) wind of the Be star. First, the

observed ¯ux density at periastron is a factor of 50 greater than at

T ÿ30, but the wind density for the rÿ2 wind differs by only a factor

of about 3. Secondly, at periastron when the pulsar is in the polar

wind, it is behind the Be-star disc which has a very high free±free

optical depth (*105), yet signi®cant emission is seen at this time.

Thirdly, the major peak in the radio emission occurs near the epoch

of the second disc crossing. We conclude, therefore, that the pulsar±

disc interaction is the main source of the unpulsed emission.

Since the observed X-ray emission from the PSR B1259ÿ63

system is also attributed to synchrotron emission from shock-

accelerated electrons and positrons (Tavani & Arons 1997), it is

of interest to compare the radio and X-ray observations. Kaspi et al.

(1995) quote an energy ¯ux in the 1±10 keV band of ,3 ´ 10ÿ11 erg

sÿ1 cmÿ2, with a photon number index of , ÿ 1:6 for T 1 615. This

translates to a ¯ux density at 1018 Hz of 1.4 mJy with a spectral

index of about ÿ0:6. Taking the radio ¯ux density to be 50 mJy at 1

GHz, we obtain a spectral index of ÿ0:51 from the radio to the X-

ray, remarkably similar to both the radio and X-ray spectral indices.

There are, however, signi®cant differences in the radio and X-ray

emission. As remarked above, the radio emission probably

originates from the pulsar±disc interaction, but the X-ray emission

almost certainly originates from both the disc and polar wind

interactions (Tavani & Arons 1997). Observational evidence for

this is the signi®cant X-ray emission around apastron (Cominsky et

al. 1994) and the absence of radio emission at these times.

Furthermore, very signi®cant variations are seen in the radio

emission on time-scales of a day or less (Figs 3 and 4), whereas

the X-ray emission at similar epochs varies by 10 per cent or less

over these time-scales (Kaspi et al. 1995).

We suggest the following interpretation for the observed radio

unpulsed emission and its variations. Synchrotron-emitting elec-

trons accelerate when the pulsar wind encounters the Be-star disc,

twice per orbit. If the disc has an ,rÿ4 density gradient, as

suggested by observed dispersion measure changes (Johnston

et al. 1996), it will be accelerating outward with an expansion

velocity increasing as ,r2 (assuming that the disc thickness is

proportional to r). At r , 50R�, the expansion velocity is likely to

be *1000 km sÿ1 (Waters, CoteÂ & Lamers 1987), much greater

than the orbital velocity of the pulsar (&85 km sÿ1). The wind

interface will therefore have a cometary form, expanding more or

less radially outward as illustrated in Fig. 6. Shocks will form on

both sides of this interface, and the particle acceleration could occur

on either side. As mentioned above, the different natures of the time

variations seen in the radio and X-ray emission suggest that

different populations of relativistic electrons are responsible for

the emission in the two bands. Indeed, Ball et al. (1998) show that

the radio-emitting electrons are accelerated in the stellar wind

shock, whereas Tavani & Arons (1997) argue that the X-ray

emission is generated in the pulsar wind shock.

For a source distance of 1.5 kpc (Johnston et al. 1994), an energy

index g � 2 for the relativistic electrons and a magnetic ®eld at the

shock of ,0:5 G, the source size must be ,2 ´ 1013 cm to be

optically thin to synchrotron self-absorption at 1 GHz. With these

parameters, there is approximate equipartition between the relati-

vistic particle and magnetic ®eld energy densities in the emission

region.

The size of the emission region is comparable to the stellar

separation at periastron (,1013 cm). We therefore suggest that

synchrotron electrons rapidly diffuse or are convected away from

the apex region of the bow shock to form an emission region with

size of this order. The light-crossing time of a source region of this

size is less than 1 h, consistent with the observed rise times. Another

important parameter is the acceleration time-scale for synchrotron

electrons, tacc , 10K=u2, where K , crB is the diffusion coef®cient,

u is the effective shock speed and rB is the electron gyro-radius

(Jones & Ellison 1991). For the stellar wind, u * 1000 km sÿ1 and

so the acceleration time-scale is also less than 1 h.

In our model, then, the pulsar hits the disc at about T ÿ22 and

remains within the disc for ,3 d. Initially, most of the emission

region is behind or within the disc as viewed by the observer. The

low amplitude and ¯at spectral index of the ®rst peak at T 2 can be

understood if most of the emission is absorbed by the disc. The drop

in amplitude to the minimum at T 2ÿ18 may result from either

increased free±free optical depth or a drop in synchrotron emissiv-

ity as the source expands out of the stellar disc. As the emission

region expands, an increasing fraction of it will have a line of sight

to the observer that does not pass through the thick region of the

stellar disc. Therefore the amplitude of the source will rise and its

spectral index will approach the optically thin value

ÿ�g ÿ 1�=2 < ÿ0:5, as observed (Figs 4 and 5).

After reaching a peak a few days before periastron, the observed

¯ux density decays slowly before the second encounter results in


q 1999 RAS, MNRAS 302, 277±287

Figure 6. A schematic diagram showing the interaction of the wind from

PSR B1259ÿ63 with the disc of the Be star SS 2883. The pulsar moves

around the orbit as indicated (the dashed portion of the orbit is behind the

plane of the sky), passing behind the stellar disc before periastron (P). The

stellar disc is believed to be inclined to the orbital plane, and so the pulsar

passes through it twice, before periastron (A) and after periastron (B). The

pulsar wind bubble is blown into a cometary shape by the ram pressure of the

wind in the stellar disc.

the rise to the main peak at about T +20. After expansion in the

stellar wind, the synchrotron lifetime is very long (*1 yr), and so

we attribute the slow decay to adiabatic losses in the expansion. At

the second encounter, the pulsar is closer to its companion than at

the ®rst, which implies higher densities and magnetic ®elds during

the encounter. This has the effect of producing stronger emission

and a more rapid decay (through synchrotron losses), as seen in the

data. The emitting material is in front of the absorbing disc with

respect to the observer, and hence the spectral index remains largely

constant. We interpret the slowly decaying emission after T +30 as

the tail of the decay from the ®rst encounter.

5 C O N C L U S I O N S

We have observed the pulsar±Be-star binary system PSR

B1259ÿ63 around the 1994 and 1997 periastron encounters at

®ve frequencies between 0.84 and 8.4 GHz. The observations reveal

a relatively strong and rapidly varying but unpulsed radio source

which persists from T ÿ22 to T +100. For most of this time, the

spectral index is , ÿ 0:5, consistent with optically thin synchrotron

emission. We associate the peaks of the emission near T ÿ10 and

+20 with relativistic electrons generated in the encounter of the

pulsar wind with the equatorial disc of the Be star. A model

involving adiabatic expansion of the emission region and variations

in the fraction of the emission absorbed by the stellar disc gives a

qualitative explanation of the observed radio ¯ux density and

spectral changes. There is a close relationship between the radio

and X-ray emission from the system, although they are probably

generated by different populations of relativistic electrons. Detailed

calculations on the emission mechanism and a more complete

investigation of the implications for the parameters of the PSR

B1259ÿ63 system will be presented in Ball et al. (1998).

AC K N OW L E D G M E N T S

We thank R. Ekers for granting `Target-of-Opportunity' time for the

observations of this pulsar during 1994. B. Piestrzynski was

instrumental in providing us with a uniform sample of images

from the MOST observations. We recognize the excellent Duty

Astronomer system at the ATCA, and particularly thank D. McKay

and B. Koribalski for help with the observations. We thank L. Ball

and O. Skjaraasen for valuable discussions, and E. Coughlin for

help with Fig. 6. The MOST is owned and operated by the

University of Sydney, with support from the Australian Research

Council and the Science Foundation within the School of Physics.

The ATCA is part of the Australia Telescope, which is funded by the

Commonwealth of Australia for operation as a National Facility

managed by CSIRO.

REFERENCES

Ball L. T., Melatos A., Johnston S., Skjaeraasen O., 1998, ApJ, submitted

Bieging J. H., Abbott D. C., Churchwell E. B., 1989, ApJ, 340, 518

Cominsky L., Roberts M., Johnston S., 1994, ApJ, 427, 978

Eichler D., Usov V. V., 1993, ApJ, 402, 271

Grove E., Tavani M., Purcell W. R., Kurfess J. D., Strickman M. S., Arons J.,

1995, ApJ, 447, L113

Hirayama M., Nagase F., Tavani M., Kaspi V. M., Kawia N., Arons J., 1996,

PASJ, 48, 833

Johnston S., Lyne A. G., Manchester R. N., Kniffen D. A., D'Amico N.,

Lim J., Ashworth M., 1992a, MNRAS, 255, 401

Johnston S., Manchester R. N., Lyne A. G., Bailes M., Kaspi V. M., Qiao G.,

D'Amico N., 1992b, ApJ, 387, L37

Johnston S., Manchester R. N., Lyne A. G., Nicastro L., Spyromilio J., 1994,

MNRAS, 268, 430

Johnston S., Manchester R. N., Lyne A. G., D'Amico N., Bailes M.,

Gaensler B. M., Nicastro L., 1996, MNRAS, 279, 1026

Jones F. C., Ellison D. C., 1991, Space Sci. Rev., 58, 259

Kaspi V. M., Tavani M., Nagase F., Hirayama M., Hoshino M., Aoki T.,

Kawai N., Arons J., 1995, ApJ, 453, 424

Manchester R. N., Johnston S., 1995, ApJ, 441, L65

Maraschi L., Treves A., 1981, MNRAS, 194, 1P

Paredes J. M., Marti J., Estalella R., Sarrate J., 1991, A&A, 248, 124

Tavani M., Arons J., 1997, ApJ, 477, 439

Taylor A. R., Gregory P. C., 1982, ApJ, 255, 210

Taylor A. R., Gregory P. C., 1984, ApJ, 283, 273

Taylor A. R., Young G., Peracula M., Kenny H. T., Gregory P. C., 1996,

A&A, 305, 817

Waters L. B. F. M., CoteÂ J., Lamers H. J. G. L. M., 1987, A&A, 185, 206

Wex N., Johnston S., Manchester R. N., Lyne A. G., Stappers B. W.,

Bailes M., 1998, MNRAS, 298, 997

White R. L., 1985, ApJ, 289, 698

Williams P. M., van der Hucht K. A., TheÂ P. S., 1987, A&A, 182, 91

This paper has been typeset from a TEX=LATEX ®le prepared by the author.


q 1999 RAS, MNRAS 302, 277±287

Documents

Transient radio emission from the PSR B1259–63 system near periastron