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Fic. 5.—Galactic distribution of accreting pulsars. We plot the known Be transients as open circles, the low-mass objects as pluses, the disk-fed OB supergiants as asterisks, and the wind-fed OB supergiants as squares. This contains all pulsars listed in Table 1. Triangles denote sources with unknown companions and orbit. This symbol convention is the same as in Fig. 4, except here the symbols are not filled.

Table 1 Fic. 5.—Galactic distribution of accreting pulsars. We plot the known Be transients as open circles, the low-mass objects as pluses, the disk-fed OB supergiants as asterisks, and the wind-fed OB supergiants as squares. This contains all pulsars listed in Table 1. Triangles denote sources with unknown companions and orbit. This symbol convention is the same as in Fig. 4, except here the symbols are not filled.

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Accreting X-ray pulsars were discovered over 25 years ago (Giacconi et al. 1971), and a qualitative description of both the accretion process and the origin of the pulsed emission was understood almost immediately (Pringle & Rees 1972; Davidson & Ostriker 1973; Lamb, Pethick, & Pines 1973). X-ray pulsars are rotating and strongly magne- tized (B = 10’! G) neutron stars that accrete gas from a stellar companion. As the accreting material approaches the neutron star, the plasma is channeled to the magnetic polar caps, where it releases its gravitational energy as X- and y-radiation; these rotating hot spots are the sources of the pulsed emission. Despite more than two decades of study, however, many details of this scenario remain poorly under- stood. Subject headings: accretion, accretion disks — binaries: close — pulsars: general - stars: magnetic fields — stars: neutron — X-rays: stars
Accreting X-ray pulsars were discovered over 25 years ago (Giacconi et al. 1971), and a qualitative description of both the accretion process and the origin of the pulsed emission was understood almost immediately (Pringle & Rees 1972; Davidson & Ostriker 1973; Lamb, Pethick, & Pines 1973). X-ray pulsars are rotating and strongly magne- tized (B = 10’! G) neutron stars that accrete gas from a stellar companion. As the accreting material approaches the neutron star, the plasma is channeled to the magnetic polar caps, where it releases its gravitational energy as X- and y-radiation; these rotating hot spots are the sources of the pulsed emission. Despite more than two decades of study, however, many details of this scenario remain poorly under- stood. Subject headings: accretion, accretion disks — binaries: close — pulsars: general - stars: magnetic fields — stars: neutron — X-rays: stars
KNOWN ACCRETION-POWERED PULSARS (AS OF 1997 FEBRUARY)
KNOWN ACCRETION-POWERED PULSARS (AS OF 1997 FEBRUARY)
Fic. 1—DISCLA data from a single BATSE detector for a full space- craft orbit. The rates are from the 20-50 keV channel of Large Area Detector 2, on 1996 August 21 when Vela X-1 was in a high state. Shown are 8.2 s average rates, with errors <16 counts s~!. The 800 counts s~! modulation is due to the difference between the cosmic diffuse background and the diffuse background from Earth’s atmosphere. The lowest rates are when the detector mainly sees Earth, while at the highest rates, it sees the sky. The data gaps are due to a loss of telemetry. Vela X-1 rises above the horizon at 71,566 s and sets at 75,517 s, as noted by the arrows. Accreting pulsars are typically detected in the lowest DISCLA channel, covering 20—50 keV, and in CONT chan- nels 1-4, typically covering 20-70 keV, with detections sometimes extending to energies as high as 160 keV. In this sense, BATSE is measuring only the high-energy spectral tails of accreting pulsars. Often, most of the flux from these objects is at lower energies, so that our flux measurements are subject to a large (and often unknown) bolometric cor- rection. The pulsed flux histories we provide thus reflect the history of the bolometric luminosity only when the overall
Fic. 1—DISCLA data from a single BATSE detector for a full space- craft orbit. The rates are from the 20-50 keV channel of Large Area Detector 2, on 1996 August 21 when Vela X-1 was in a high state. Shown are 8.2 s average rates, with errors <16 counts s~!. The 800 counts s~! modulation is due to the difference between the cosmic diffuse background and the diffuse background from Earth’s atmosphere. The lowest rates are when the detector mainly sees Earth, while at the highest rates, it sees the sky. The data gaps are due to a loss of telemetry. Vela X-1 rises above the horizon at 71,566 s and sets at 75,517 s, as noted by the arrows. Accreting pulsars are typically detected in the lowest DISCLA channel, covering 20—50 keV, and in CONT chan- nels 1-4, typically covering 20-70 keV, with detections sometimes extending to energies as high as 160 keV. In this sense, BATSE is measuring only the high-energy spectral tails of accreting pulsars. Often, most of the flux from these objects is at lower energies, so that our flux measurements are subject to a large (and often unknown) bolometric cor- rection. The pulsed flux histories we provide thus reflect the history of the bolometric luminosity only when the overall
Fic. 2.—The 1 day BATSE sensitivity for a single large-area detector to an unocculted pulsed source at three pulse frequencies. We define the threshold pulsed intensity (upper panel) as the count rate where the error in the measured rate is 20%. The degradation in sensitivity at low frequencies is due to residual background not fully accounted for by the background subtraction process. The lower panel shows the energy flux corresponding to these threshold count rates, assuming that the source has a power-law spectrum with photon number index « = 4.0 and is viewed at normal incidence. A typical live time of ~ 42,000 s is obtained per day. a, ee CO The resulting 1 day pulsed- -intensity sensitivities of the CONT data as a function of energy are shown in Figure 2 for three representative pulse frequencies. The upper panel
Fic. 2.—The 1 day BATSE sensitivity for a single large-area detector to an unocculted pulsed source at three pulse frequencies. We define the threshold pulsed intensity (upper panel) as the count rate where the error in the measured rate is 20%. The degradation in sensitivity at low frequencies is due to residual background not fully accounted for by the background subtraction process. The lower panel shows the energy flux corresponding to these threshold count rates, assuming that the source has a power-law spectrum with photon number index « = 4.0 and is viewed at normal incidence. A typical live time of ~ 42,000 s is obtained per day. a, ee CO The resulting 1 day pulsed- -intensity sensitivities of the CONT data as a function of energy are shown in Figure 2 for three representative pulse frequencies. The upper panel
Fic. 3—The 1 day BATSE 5 o detection threshold for an unocculted source (CONT channels 1-4 summed) following background subtraction. The thresholds plotted were obtained by averaging the thresholds for each of the eight BATSE detectors from MJD 49,081 to 49,093. A typical day is characterized by a mean background rate (CONT channels 1-4 summed) of +2000 counts s~/ and ~ 42,000 s of useful data. At frequencies v > 0.02 Hz, the background noise is essentially Poisson, and the threshold is simply that given by eq. (1). The thresholds at lower frequencies are found by assuming local Gaussian statistics for the measured noise strength. The anomalous rise in the thresholds at v ~ 0.007 Hz is attributable to the second harmonic of Vela X-1, which was not accounted for in the back- ground subtraction model.
Fic. 3—The 1 day BATSE 5 o detection threshold for an unocculted source (CONT channels 1-4 summed) following background subtraction. The thresholds plotted were obtained by averaging the thresholds for each of the eight BATSE detectors from MJD 49,081 to 49,093. A typical day is characterized by a mean background rate (CONT channels 1-4 summed) of +2000 counts s~/ and ~ 42,000 s of useful data. At frequencies v > 0.02 Hz, the background noise is essentially Poisson, and the threshold is simply that given by eq. (1). The thresholds at lower frequencies are found by assuming local Gaussian statistics for the measured noise strength. The anomalous rise in the thresholds at v ~ 0.007 Hz is attributable to the second harmonic of Vela X-1, which was not accounted for in the back- ground subtraction model.
CONVERSION TO PULSED FLUX FROM COUNTS NotTe—These conversions are for normal incidence. For other incident photon angles, see Appendix A and Fig. 42 * Spectral model used is dN/dE oc E~*. > Energy flux in units of 107 1° ergs cm~? s~' corresponding to a pulsed intensity of 1 count s~ 1 in DISCLA Channel 1 (20-50 keV). ° Energy flux in units of 107 1° ergs cm~? s~! corresponding to a pulsed intensity of 1 count s~' in CONT Channels 1-4 (20-70 keV) summed. 4 Percentage of error in energy flux due to variations in energy edges for different detectors. ° Spectral model used is dN/dE oc (1/E) exp (— E/kT)g,,(E, kT). 2
CONVERSION TO PULSED FLUX FROM COUNTS NotTe—These conversions are for normal incidence. For other incident photon angles, see Appendix A and Fig. 42 * Spectral model used is dN/dE oc E~*. > Energy flux in units of 107 1° ergs cm~? s~' corresponding to a pulsed intensity of 1 count s~ 1 in DISCLA Channel 1 (20-50 keV). ° Energy flux in units of 107 1° ergs cm~? s~! corresponding to a pulsed intensity of 1 count s~' in CONT Channels 1-4 (20-70 keV) summed. 4 Percentage of error in energy flux due to variations in energy edges for different detectors. ° Spectral model used is dN/dE oc (1/E) exp (— E/kT)g,,(E, kT). 2
A convenient organization of the high-mass systems emerges by plotting those with known orbital periods on a Pyyin-P op diagram (Corbet 1986; Waters & van Kerkwijk 1989), where P,,;, is the neutron star spin period. As is evident in Figure 4, the neutron stars orbiting Roche lobe— filling supergiants (asterisks) have short spin periods (Pspin S 10s) and short orbital periods (P,,, < 4 days). They are quite luminous (L = 10°” ergs s~‘) and tend to show Fic. 4.—Spin period of the accreting neutron stars vs. the binary period (the Corbet diagram). The different symbols refer to the type of binary the neutron star resides in. Asterisks are supergiant companions that are Roche lobe filling, squares are supergiant companions that underfill their Roche lobe, circles are confirmed Be-transient binaries, and plus signs are those with low-mass ($2 M.) companions (Her X-1, GRO J1744— 28, and 4U 1626—67). Triangles refer to sources for which the companion type is unknown, though the orbital period has been determined. (GS 0834—430, OAO 1657—415, GRO J1750—27, 2S 1553—54, and GRO J2058 + 42.)
A convenient organization of the high-mass systems emerges by plotting those with known orbital periods on a Pyyin-P op diagram (Corbet 1986; Waters & van Kerkwijk 1989), where P,,;, is the neutron star spin period. As is evident in Figure 4, the neutron stars orbiting Roche lobe— filling supergiants (asterisks) have short spin periods (Pspin S 10s) and short orbital periods (P,,, < 4 days). They are quite luminous (L = 10°” ergs s~‘) and tend to show Fic. 4.—Spin period of the accreting neutron stars vs. the binary period (the Corbet diagram). The different symbols refer to the type of binary the neutron star resides in. Asterisks are supergiant companions that are Roche lobe filling, squares are supergiant companions that underfill their Roche lobe, circles are confirmed Be-transient binaries, and plus signs are those with low-mass ($2 M.) companions (Her X-1, GRO J1744— 28, and 4U 1626—67). Triangles refer to sources for which the companion type is unknown, though the orbital period has been determined. (GS 0834—430, OAO 1657—415, GRO J1750—27, 2S 1553—54, and GRO J2058 + 42.)
ORBITAL PARAMETERS OF ACCRETION-POWERED PULSAR SYSTEMS?
ORBITAL PARAMETERS OF ACCRETION-POWERED PULSAR SYSTEMS?
where I ~ 0.4M,, R2 is the neutron star’s moment of inertia (Ravenhall & Pethick 1994). The timescale for spinning up the neutron star is then tion rate. The fiducial torque sets a scale, and the actual torque may by significantly smaller. A pulsar subject to the torque in equation (3) will spin up at a rate
where I ~ 0.4M,, R2 is the neutron star’s moment of inertia (Ravenhall & Pethick 1994). The timescale for spinning up the neutron star is then tion rate. The fiducial torque sets a scale, and the actual torque may by significantly smaller. A pulsar subject to the torque in equation (3) will spin up at a rate
SPECTRAL PARAMETERS MEASURED WITH BATSE
SPECTRAL PARAMETERS MEASURED WITH BATSE
Fic. 6—Long-term frequency history for all pulsars detected by BATSE that were previously known. The squares show the pre-BATSE data taken frot Nagase (1989) and additional references. The line is the BATSE data, which we discuss later in great detail. The long-term frequency history for X-ray pulsa: observed by BATSE that were known prior to the Compton Observatory launch commences 1991 April. For Her X-1, Cen X-3, Vela X-1, 4U 1538—52, G 301—2, 4U 0115+634, and EXO 2030+375, all frequencies have been orbitally corrected. For OAO 1657—415, GS 0834—430, 2S 1417—62, an A0535 + 262, orbital corrections have been applied only to the BATSE observations. No orbital corrections have been applied for 4U 1626 —67, GX 1+4, 41 1145—619, or A1118—615, which have unknown, or incompletely known, orbital elements. The BATSE frequencies for OAO 1657—415, GS 0834-431 EXO 2030+375, 4U 1145—619, and A1118—615 are from the daily frequency history files we have deposited at the Compton Observatory archive. For th remainder of the objects, the frequencies are from source-specific studies: Her X-1 (Wilson et al. 1997), 4U 1626—67 (Chakrabarty et al. 1997a), GX 1+ (Chakrabarty et al. 1997b), Cen X-3 (Finger et al. 1997), Vela X-1 (Finger et al. 1997), 4U 1538—52 (Rubin et al. 1997), GX 301—2 (Koh et al. 1997), 41 0115 +634 (Cominsky et al. 1997), 2S 1417 — 624 (Finger et al. 1996b), and A0535 + 262 (Finger et al. 1996c).
Fic. 6—Long-term frequency history for all pulsars detected by BATSE that were previously known. The squares show the pre-BATSE data taken frot Nagase (1989) and additional references. The line is the BATSE data, which we discuss later in great detail. The long-term frequency history for X-ray pulsa: observed by BATSE that were known prior to the Compton Observatory launch commences 1991 April. For Her X-1, Cen X-3, Vela X-1, 4U 1538—52, G 301—2, 4U 0115+634, and EXO 2030+375, all frequencies have been orbitally corrected. For OAO 1657—415, GS 0834—430, 2S 1417—62, an A0535 + 262, orbital corrections have been applied only to the BATSE observations. No orbital corrections have been applied for 4U 1626 —67, GX 1+4, 41 1145—619, or A1118—615, which have unknown, or incompletely known, orbital elements. The BATSE frequencies for OAO 1657—415, GS 0834-431 EXO 2030+375, 4U 1145—619, and A1118—615 are from the daily frequency history files we have deposited at the Compton Observatory archive. For th remainder of the objects, the frequencies are from source-specific studies: Her X-1 (Wilson et al. 1997), 4U 1626—67 (Chakrabarty et al. 1997a), GX 1+ (Chakrabarty et al. 1997b), Cen X-3 (Finger et al. 1997), Vela X-1 (Finger et al. 1997), 4U 1538—52 (Rubin et al. 1997), GX 301—2 (Koh et al. 1997), 41 0115 +634 (Cominsky et al. 1997), 2S 1417 — 624 (Finger et al. 1996b), and A0535 + 262 (Finger et al. 1996c).
dwarf (4U 1626—67), and two red giants (GX 1+4 and GRO J1744— 28). The absence of an observable companion in very deep optical and IR searches and the lack of orbital detections also suggest low-mass companions for 4U 0142+61, 1E 1048.1—5937, RX J1838.4—0301, and 1E 2259+589 (see Mereghetti & Stella 1995 and references therein). We now discuss the BATSE observations of Her X-1, 4U 1626—67, GX 1+4, and GRO J1744— 28. Only four accreting pulsars are definitely known to be orbiting low-mass (M < 2 Mo) stars: Her X-1 and GRO J1744—28 on the basis of timing-based measurements of their companions’ mass functions (Tananbaum et al. 1972; Finger et al. 1996a), 4U 1626 —67 from optical photometry (Middleditch et al. 1981; Chakrabarty 1997), and GX 1+4 from spectroscopy (Davidsen, Malina, & Bowyer 1977; Chakrabarty & Roche 1997). This is a very heterogeneous class of objects: the mass donors are a main-sequence A star (Her X-1), a <0.1 Mo helium or carbon-oxygen degenerate
dwarf (4U 1626—67), and two red giants (GX 1+4 and GRO J1744— 28). The absence of an observable companion in very deep optical and IR searches and the lack of orbital detections also suggest low-mass companions for 4U 0142+61, 1E 1048.1—5937, RX J1838.4—0301, and 1E 2259+589 (see Mereghetti & Stella 1995 and references therein). We now discuss the BATSE observations of Her X-1, 4U 1626—67, GX 1+4, and GRO J1744— 28. Only four accreting pulsars are definitely known to be orbiting low-mass (M < 2 Mo) stars: Her X-1 and GRO J1744—28 on the basis of timing-based measurements of their companions’ mass functions (Tananbaum et al. 1972; Finger et al. 1996a), 4U 1626 —67 from optical photometry (Middleditch et al. 1981; Chakrabarty 1997), and GX 1+4 from spectroscopy (Davidsen, Malina, & Bowyer 1977; Chakrabarty & Roche 1997). This is a very heterogeneous class of objects: the mass donors are a main-sequence A star (Her X-1), a <0.1 Mo helium or carbon-oxygen degenerate
circular, 1.7 day orbit (Deeter, Boynton, & Pravdo 1981) around the low-mass companion HZ Her (Doxsey et al. 1973; Gottwald et al. 1991). This disk-fed system exhibits “supercycles” of intensity modulated with a period of +35 days (Giacconi et al. 1973; Soong et al. 1990). At energies of 1-10 keV, the source is observed during both a “Main High” and “Short High” interval of the 35 day cycle. Detailed discussion of BATSE observations have appeared elsewhere (Wilson et al. 1994d, 1994c). _ 4U 1626—67W—SAS-3 discovered 7.68 s pulsations from 4U 1626—67 in 1977 (Rappaport et al. 1977), and the optical counterpart was later identified to be KZ TrA (McClintock et al. 1977). See Figure 9. X-ray timing limits imply an ultracompact binary with an extremely low-mass companion (Levine et al. 1988; Chakrabarty et al. 1997a). There is optical photometric evidence for a 42 minute orbital period (Middleditch et al. 1981; Chakrabarty 1998), which suggests a mass of <0.1 M, for the mass donor.
circular, 1.7 day orbit (Deeter, Boynton, & Pravdo 1981) around the low-mass companion HZ Her (Doxsey et al. 1973; Gottwald et al. 1991). This disk-fed system exhibits “supercycles” of intensity modulated with a period of +35 days (Giacconi et al. 1973; Soong et al. 1990). At energies of 1-10 keV, the source is observed during both a “Main High” and “Short High” interval of the 35 day cycle. Detailed discussion of BATSE observations have appeared elsewhere (Wilson et al. 1994d, 1994c). _ 4U 1626—67W—SAS-3 discovered 7.68 s pulsations from 4U 1626—67 in 1977 (Rappaport et al. 1977), and the optical counterpart was later identified to be KZ TrA (McClintock et al. 1977). See Figure 9. X-ray timing limits imply an ultracompact binary with an extremely low-mass companion (Levine et al. 1988; Chakrabarty et al. 1997a). There is optical photometric evidence for a 42 minute orbital period (Middleditch et al. 1981; Chakrabarty 1998), which suggests a mass of <0.1 M, for the mass donor.
Fic. 7.—Pulse profiles of accreting pulsars from BATSE, in 20-35 keV except as noted. The 1 o error bar in each phase bin is also shown. The range of days (in MJD) summed to construct the profiles for each of the sources (and energy range if different from above) are (a) OAO 1657—415 (48,450—48,715), A1118 — 62 (48,621—48,633), GRO J1948 + 32 (49,448-49,482), GRO J1008 — 57 (49,182—49,215), GS 0834 — 430 (48,518—48,522), GRO J1750—27 (49,960-49, 967; 20-70 keV), GX 1+ 4 (48,450-48,715), 4U 1626 — 67 (48,450-—48,715), GRO J1744 — 28 (50,092-50,098), 4U 0115 + 634 (50049-50057; 20-40 keV) and (b) 4U 1145-619 (49,428-49,439), EXO 2030+ 375 (48,450-48,715), 4U 1538 —52 (49,350-—49,421; 20-50 keV), A0535 + 262 (49,379-49,430), GRO J2058 + 42 (49,987—-49,993; 20-40 keV), GX 301—2 (48,450-—48,715), Cen X-3 (48,985-48,992), Vela X-1 (48,450-48,715), Her X-1 (49,586-49,593; 20-40 keV), 2S 1417 — 62 (49,623—49,698).
Fic. 7.—Pulse profiles of accreting pulsars from BATSE, in 20-35 keV except as noted. The 1 o error bar in each phase bin is also shown. The range of days (in MJD) summed to construct the profiles for each of the sources (and energy range if different from above) are (a) OAO 1657—415 (48,450—48,715), A1118 — 62 (48,621—48,633), GRO J1948 + 32 (49,448-49,482), GRO J1008 — 57 (49,182—49,215), GS 0834 — 430 (48,518—48,522), GRO J1750—27 (49,960-49, 967; 20-70 keV), GX 1+ 4 (48,450-48,715), 4U 1626 — 67 (48,450-—48,715), GRO J1744 — 28 (50,092-50,098), 4U 0115 + 634 (50049-50057; 20-40 keV) and (b) 4U 1145-619 (49,428-49,439), EXO 2030+ 375 (48,450-48,715), 4U 1538 —52 (49,350-—49,421; 20-50 keV), A0535 + 262 (49,379-49,430), GRO J2058 + 42 (49,987—-49,993; 20-40 keV), GX 301—2 (48,450-—48,715), Cen X-3 (48,985-48,992), Vela X-1 (48,450-48,715), Her X-1 (49,586-49,593; 20-40 keV), 2S 1417 — 62 (49,623—49,698).
by BATSE and subsequent pulse timing unambiguously established GRO J1744—28 to be a neutron star in a circu- lar 11.8 day orbit around a low-mass companion and indi- cated that the neutron star was spun up by an accretion disk during the outburst (Finger et al. 1996a). These are the first persistent pulsations seen in a bursting X-ray source. BATSE observed a number of bright flares in the hard X-ray (20-100 keV) band that were accompanied by epi- sodes of enhanced spin-down. A smooth torque reversal to spin-up accompanied an extended bright state during late 1994 and early 1995 (Chakrabarty et al. 1994b), followed by a return to spin-down and a lower average hard X-ray flux (Chakrabarty et al. 1995b). During spin-down, the torque fluctuations exhibit a 1/f power density spectrum, similar to that seen in Cen X-3.
by BATSE and subsequent pulse timing unambiguously established GRO J1744—28 to be a neutron star in a circu- lar 11.8 day orbit around a low-mass companion and indi- cated that the neutron star was spun up by an accretion disk during the outburst (Finger et al. 1996a). These are the first persistent pulsations seen in a bursting X-ray source. BATSE observed a number of bright flares in the hard X-ray (20-100 keV) band that were accompanied by epi- sodes of enhanced spin-down. A smooth torque reversal to spin-up accompanied an extended bright state during late 1994 and early 1995 (Chakrabarty et al. 1994b), followed by a return to spin-down and a lower average hard X-ray flux (Chakrabarty et al. 1995b). During spin-down, the torque fluctuations exhibit a 1/f power density spectrum, similar to that seen in Cen X-3.
Fic. 9—4U 1626-67 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have not been orbitally corrected since the orbit is unknown, were determined at 5 day intervals by searching the Fourier power spectrum of the 20-50 keV DISCLA data for the strongest signal in a small range around 7.7 s (see Appendix B, § 1.1). The pulsed fluxes were obtained by assuming a power-law spectral model with a photon number index of 4.9 (see Appendix B, § 2.3). BATSE continuously monitors five pulsars that accrete from high-mass evolved supergiants: Cen X-3, OAO 1657—415, Vela X-1, 4U 1538—52, and GX 301—2. The long-term spin frequency evolution of these pulsars has revealed several surprises that challenge the standard model of such systems, as we discuss in § 5. For example, Cen X-3 (the only Roche lobe-filling, high-mass supergiant system
Fic. 9—4U 1626-67 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have not been orbitally corrected since the orbit is unknown, were determined at 5 day intervals by searching the Fourier power spectrum of the 20-50 keV DISCLA data for the strongest signal in a small range around 7.7 s (see Appendix B, § 1.1). The pulsed fluxes were obtained by assuming a power-law spectral model with a photon number index of 4.9 (see Appendix B, § 2.3). BATSE continuously monitors five pulsars that accrete from high-mass evolved supergiants: Cen X-3, OAO 1657—415, Vela X-1, 4U 1538—52, and GX 301—2. The long-term spin frequency evolution of these pulsars has revealed several surprises that challenge the standard model of such systems, as we discuss in § 5. For example, Cen X-3 (the only Roche lobe-filling, high-mass supergiant system
Fic. 8—Her X-1 frequency and pulsed flux measurements from BATSE. A mean frequency has been determined for each Main High state for which adequate scheduled folded-on-board data were available. The frequencies, which have been orbitally corrected using parameters from Wilson et al. (1994a), were obtained using linear fits to pulse phases of 20-70 keV data (see Appendix B, § 1.3). Pulsed fluxes are the pulse phase- averaged flux obtained during one orbital period (averaged from eclipse egress to eclipse ingress) during the brightest portion of each available Main High state. An exponential model, with the e-folding energy allowed to vary, was used to obtain the fluxes (see Appendix B, § 2.3).
Fic. 8—Her X-1 frequency and pulsed flux measurements from BATSE. A mean frequency has been determined for each Main High state for which adequate scheduled folded-on-board data were available. The frequencies, which have been orbitally corrected using parameters from Wilson et al. (1994a), were obtained using linear fits to pulse phases of 20-70 keV data (see Appendix B, § 1.3). Pulsed fluxes are the pulse phase- averaged flux obtained during one orbital period (averaged from eclipse egress to eclipse ingress) during the brightest portion of each available Main High state. An exponential model, with the e-folding energy allowed to vary, was used to obtain the fluxes (see Appendix B, § 2.3).
Fic. 11—GRO J1744—28 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have been orbitally corrected using parameters from Finger et al. (1996a), were determined at 2 day intervals by fitting pulse phases derived from pulsar folded-on-board data in the 20-40 keV band (see Appendix B, § 1.3). The pulsed fluxes were obtained by assuming a spectra of the form F(E) = AE~* exp (—E/kT) with 4 = 2.0 and kT = 15 keV, as determined by OSSE measurements (Strickman et al. 1996) (see Appendix B, § 2.3) Detailed discussions of the BATSE observations of Cen X-3 have appeared elsewhere (Finger et al. 1992; Finger,
Fic. 11—GRO J1744—28 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have been orbitally corrected using parameters from Finger et al. (1996a), were determined at 2 day intervals by fitting pulse phases derived from pulsar folded-on-board data in the 20-40 keV band (see Appendix B, § 1.3). The pulsed fluxes were obtained by assuming a spectra of the form F(E) = AE~* exp (—E/kT) with 4 = 2.0 and kT = 15 keV, as determined by OSSE measurements (Strickman et al. 1996) (see Appendix B, § 2.3) Detailed discussions of the BATSE observations of Cen X-3 have appeared elsewhere (Finger et al. 1992; Finger,
Fic. 10—GX 1+4 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have not been orbitally corrected since the orbit is unknown, were determined at 5 day intervals by searching the Fourier power spectrum of the 20-50 keV DISCLA data for the strongest signal in the pulse period range 110s < P,,i. S 130s (see Appen- dix B, § 1.1). The pulsed fluxes were determined at 5 day intervals by assuming a power-law spectral model with a photon number index of 2.5 (see Appendix B, § 2.3).
Fic. 10—GX 1+4 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have not been orbitally corrected since the orbit is unknown, were determined at 5 day intervals by searching the Fourier power spectrum of the 20-50 keV DISCLA data for the strongest signal in the pulse period range 110s < P,,i. S 130s (see Appen- dix B, § 1.1). The pulsed fluxes were determined at 5 day intervals by assuming a power-law spectral model with a photon number index of 2.5 (see Appendix B, § 2.3).
Fic. 12—Cen X-3 frequency and pulsed flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Finger et al. (1993), were determined at 2.1 day intervals by epoch-folding the 20-50 keV DISCLA data (see Appendix B, § 1.2). The pulsed fluxes were determined at 2.1 day intervals by assuming an exponential spectrum with an e-folding energy of 12 keV (see Appendix B, § 2.3).
Fic. 12—Cen X-3 frequency and pulsed flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Finger et al. (1993), were determined at 2.1 day intervals by epoch-folding the 20-50 keV DISCLA data (see Appendix B, § 1.2). The pulsed fluxes were determined at 2.1 day intervals by assuming an exponential spectrum with an e-folding energy of 12 keV (see Appendix B, § 2.3).
Fic. 13—OAO 1657—415 frequency and pulsed-flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitall) corrected using parameters from Chakrabarty et al. (1993) and a refined orbital period of P,,,, = 10.44809(30) days (see text below for details), were measured at 1 day intervals from the power spectra of the 15-55 keV CONT data (see Appendix B, § 1.1). The pulsed fluxes were measured at 1 day intervals by assuming an exponential spectrum with a e-folding energy of 20 keV (see Appendix B, § 2.3). OAO 1657 —415.—HEAO 1 discovered 38.22 s pulsations from OAO 1657—415 in 1978 (White & Pravdo 1979). See Figure 13. BATSE observations revealed a 10.4 day binary orbit with a 1.7 day eclipse by the stellar companion (Chakrabarty et al. 1993), making it the seventh eclipsing
Fic. 13—OAO 1657—415 frequency and pulsed-flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitall) corrected using parameters from Chakrabarty et al. (1993) and a refined orbital period of P,,,, = 10.44809(30) days (see text below for details), were measured at 1 day intervals from the power spectra of the 15-55 keV CONT data (see Appendix B, § 1.1). The pulsed fluxes were measured at 1 day intervals by assuming an exponential spectrum with a e-folding energy of 20 keV (see Appendix B, § 2.3). OAO 1657 —415.—HEAO 1 discovered 38.22 s pulsations from OAO 1657—415 in 1978 (White & Pravdo 1979). See Figure 13. BATSE observations revealed a 10.4 day binary orbit with a 1.7 day eclipse by the stellar companion (Chakrabarty et al. 1993), making it the seventh eclipsing
Fic. 14—Vela X-1 spin frequency and pulsed flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Finger et al. (1997), were determined by fitting pulse phase measurements (see Appendix B, § 1.3). Each point uses the uneclipsed data of the 8.96 day binary orbit. The pulsed fluxes were deter- mined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3). Vela X-1 is the brightest persistent accretion-powered pulsar in the 20-50 keV energy band. Individual pulses are often visible in the raw data (Fig. 1). BATSE observations showed Vela X-1 alternating between spin-up and spin- down, with no long-term trend in spin frequency, consistent with pre-BATSE observations of a random walk in spin
Fic. 14—Vela X-1 spin frequency and pulsed flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Finger et al. (1997), were determined by fitting pulse phase measurements (see Appendix B, § 1.3). Each point uses the uneclipsed data of the 8.96 day binary orbit. The pulsed fluxes were deter- mined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3). Vela X-1 is the brightest persistent accretion-powered pulsar in the 20-50 keV energy band. Individual pulses are often visible in the raw data (Fig. 1). BATSE observations showed Vela X-1 alternating between spin-up and spin- down, with no long-term trend in spin frequency, consistent with pre-BATSE observations of a random walk in spin
Detailed discussion of BATSE observations have appeared elsewhere (Koh et al. 1997). BATSE observed two
Detailed discussion of BATSE observations have appeared elsewhere (Koh et al. 1997). BATSE observed two
Fic. 18—4U 0115+634 frequency and flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using the orbital parameters discussed in the text, were determined at 1 day intervals by epoch-folding the 20-50 keV DISCLA data at a range of trial frequencies (see Appendix B, § 1.2). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 15 keV (see Appendix B, § 2.3). The gaps in Fig. 18 are extended intervals when the source was undetectable with BATSE.
Fic. 18—4U 0115+634 frequency and flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using the orbital parameters discussed in the text, were determined at 1 day intervals by epoch-folding the 20-50 keV DISCLA data at a range of trial frequencies (see Appendix B, § 1.2). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 15 keV (see Appendix B, § 2.3). The gaps in Fig. 18 are extended intervals when the source was undetectable with BATSE.
Fic. 17.—Outburst times for all transient accreting pulsars observed by BATSE
Fic. 17.—Outburst times for all transient accreting pulsars observed by BATSE
Fic. 19.—Measurements of the 4U 0115+ 63 periastron epoch. The plot shows the periastron epochs minus the linear ephemeris MJD 49,279.2677 + 24.317037 x n, where n is an integral number of orbits. This ephemeris is discussed in the text. The epoch measurements have been determined from Uhuru Kelley et al. 1981b), SAS-3 (Rappaport et al. 1978), Ariel 6 (Rickets et al. 1981), Ginga (Tamura et al. 1992), and this work. Excluding the Ginga measurement, the observations are consistent with a constant orbital period.
Fic. 19.—Measurements of the 4U 0115+ 63 periastron epoch. The plot shows the periastron epochs minus the linear ephemeris MJD 49,279.2677 + 24.317037 x n, where n is an integral number of orbits. This ephemeris is discussed in the text. The epoch measurements have been determined from Uhuru Kelley et al. 1981b), SAS-3 (Rappaport et al. 1978), Ariel 6 (Rickets et al. 1981), Ginga (Tamura et al. 1992), and this work. Excluding the Ginga measurement, the observations are consistent with a constant orbital period.
Fic. 21—GS 0834—430 frequency and flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Wilson et al. (1997), were determined at 1 day intervals from the power spectra of the 20-70 keV CONT data (see Appen- dix B, § 1.1). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 14 keV (see Appendix B, § 2.3). 2S 1417—624—SAS-3 discovered 17.6 s pulsations from 2S 1417—624 in 1978 (Kelley et al. 1981a), and the compan- ion was later identified to be a 17th mag OB star (Grindlay, Petro, & McClintock 1984). See Figure 22. Detailed dis- cussion of BATSE observations have appeared elsewhere (Finger, Wilson, & Chakrabarty 1996b). BATSE observed a large outburst of 2S 1417—624 from 1994 August 29 to December 11 (MJD 49,593—49,697), followed by a sequence of five smaller outbursts of diminishing amplitudes occurring approximately every 40 days (Finger et al. 1996b),
Fic. 21—GS 0834—430 frequency and flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Wilson et al. (1997), were determined at 1 day intervals from the power spectra of the 20-70 keV CONT data (see Appen- dix B, § 1.1). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 14 keV (see Appendix B, § 2.3). 2S 1417—624—SAS-3 discovered 17.6 s pulsations from 2S 1417—624 in 1978 (Kelley et al. 1981a), and the compan- ion was later identified to be a 17th mag OB star (Grindlay, Petro, & McClintock 1984). See Figure 22. Detailed dis- cussion of BATSE observations have appeared elsewhere (Finger, Wilson, & Chakrabarty 1996b). BATSE observed a large outburst of 2S 1417—624 from 1994 August 29 to December 11 (MJD 49,593—49,697), followed by a sequence of five smaller outbursts of diminishing amplitudes occurring approximately every 40 days (Finger et al. 1996b),
Fic. 20—GRO J1750—27 spin frequency and pulsed flux measure- ments from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Scott et al. (1997), were deter- mined at 1 day intervals from fits of phase measurements of the 20-50 keV DISCLA data (see Appendix B, § 1.3). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3).
Fic. 20—GRO J1750—27 spin frequency and pulsed flux measure- ments from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Scott et al. (1997), were deter- mined at 1 day intervals from fits of phase measurements of the 20-50 keV DISCLA data (see Appendix B, § 1.3). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3).
Fic. 23.—GRO J1948 +32 frequency and pulsed flux measurements from BATSE. The pulse frequencies, which have not been orbitally cor- rected as the orbital parameters are unknown, were determined at 1 day intervals from the power spectra of the 20-70 keV CONT data (see Appen- dix B, § 1.1). The pulsed fluxes were obtained by assuming an exponential spectrum with an e-folding energy of 15 keV (see Appendix B, § 2.3).
Fic. 23.—GRO J1948 +32 frequency and pulsed flux measurements from BATSE. The pulse frequencies, which have not been orbitally cor- rected as the orbital parameters are unknown, were determined at 1 day intervals from the power spectra of the 20-70 keV CONT data (see Appen- dix B, § 1.1). The pulsed fluxes were obtained by assuming an exponential spectrum with an e-folding energy of 15 keV (see Appendix B, § 2.3).
Fic. 22.—2S 1417—624 frequency and pulsed flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Finger et al. (1996b), were determined at 1 day intervals by epoch-folding the 20-50 keV DISCLA data (see Appendix B, §1.2). The pulsed fluxes were determined at 1 day intervals by assuming a Comptonized spectrum model of the form F(E) = AE* exp (— E/kT), with A= 1.6 and kT = 11.9 keV (Finger et al. 1996b) (see Appendix B, § 2.3). The orbit was determined assuming a correlation between pulsed flux and accretion torque, which could potentially introduce modulations in the apparent rate of spin-up during the sequence of outbursts following the main outburst.
Fic. 22.—2S 1417—624 frequency and pulsed flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Finger et al. (1996b), were determined at 1 day intervals by epoch-folding the 20-50 keV DISCLA data (see Appendix B, §1.2). The pulsed fluxes were determined at 1 day intervals by assuming a Comptonized spectrum model of the form F(E) = AE* exp (— E/kT), with A= 1.6 and kT = 11.9 keV (Finger et al. 1996b) (see Appendix B, § 2.3). The orbit was determined assuming a correlation between pulsed flux and accretion torque, which could potentially introduce modulations in the apparent rate of spin-up during the sequence of outbursts following the main outburst.
Fic. 24—EXO 2030+375 frequency and pulsed flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Stollberg et al. (1994), were determined at 1 day intervals from the power spectra of the 20-70 keV CONT data (see Appendix B, § 1.1). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3). The source was quiescent for 2.5 yr before being detected by BATSE again in 1996 April and May (Stollberg et al. 1996). These two outbursts occurred ~5 days prior to
Fic. 24—EXO 2030+375 frequency and pulsed flux measurements from BATSE. The intrinsic spin frequencies, which have been orbitally corrected using parameters from Stollberg et al. (1994), were determined at 1 day intervals from the power spectra of the 20-70 keV CONT data (see Appendix B, § 1.1). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3). The source was quiescent for 2.5 yr before being detected by BATSE again in 1996 April and May (Stollberg et al. 1996). These two outbursts occurred ~5 days prior to
Fic. 25.—GRO J1008—57 frequency and pulsed flux measurements from BATSE. The pulse frequencies, which have not been orbitally cor- rected as the orbital parameters are unknown, were determined at 1 day intervals from the power spectra of the 20-70 keV CONT data (see Appen- dix B, § 1.1). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3). Detailed discussions of BATSE observations of A0535+26 have appeared elsewhere (Finger et al. 1994a, 1996c). BATSE has observed six outbursts spaced roughly at the orbital period, the fourth of which is a “ giant” out- burst that occurred from 1994 January 28 to March 20
Fic. 25.—GRO J1008—57 frequency and pulsed flux measurements from BATSE. The pulse frequencies, which have not been orbitally cor- rected as the orbital parameters are unknown, were determined at 1 day intervals from the power spectra of the 20-70 keV CONT data (see Appen- dix B, § 1.1). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3). Detailed discussions of BATSE observations of A0535+26 have appeared elsewhere (Finger et al. 1994a, 1996c). BATSE has observed six outbursts spaced roughly at the orbital period, the fourth of which is a “ giant” out- burst that occurred from 1994 January 28 to March 20
Fic. 26.—A0535+ 262 frequency and pulsed flux measurements from BATSE. The frequencies, which have been orbitally corrected using orbital parameters from Finger et al. (1994a), were determined at 3 day intervals from fits to phase measurements made using the 20-50 keV DISCLA data (see Appendix B, § 1.3). The pulsed fluxes were measured at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3).
Fic. 26.—A0535+ 262 frequency and pulsed flux measurements from BATSE. The frequencies, which have been orbitally corrected using orbital parameters from Finger et al. (1994a), were determined at 3 day intervals from fits to phase measurements made using the 20-50 keV DISCLA data (see Appendix B, § 1.3). The pulsed fluxes were measured at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3).
Fic. 28.—GRO J2058+42 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have not been orbitally cor- rected as the orbital parameters are unknown, were determined at 4 day intervals by epoch-folding the 20-50 keV DISCLA data at a range of trial frequencies (see Appendix B, § 1.2). The pulsed fluxes were determined at 4 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3). and has an equatorial rotational velocity of v sin i = 290 km s_ 1 (Hammerschlag-Hensberge et al. 1980; Bianchi & Ber- nacca 1980). An orbital period of 186.5 days was inferred from the recurrence times of outbursts, which typically last x10 days (Watson, Warwick, & Ricketts 1981; Priedhorsky & Terrell 1983). Pulse frequency variations over multiple EXOSAT observations imply an eccentricity of e 2 0.6 (Cook & Warwick 1987). To date, BATSE has observed seven outbursts, of which three are shown. The separation between the BATSE outbursts is in good agree- ment with the 186.5 day period.
Fic. 28.—GRO J2058+42 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have not been orbitally cor- rected as the orbital parameters are unknown, were determined at 4 day intervals by epoch-folding the 20-50 keV DISCLA data at a range of trial frequencies (see Appendix B, § 1.2). The pulsed fluxes were determined at 4 day intervals by assuming an exponential spectrum with an e-folding energy of 20 keV (see Appendix B, § 2.3). and has an equatorial rotational velocity of v sin i = 290 km s_ 1 (Hammerschlag-Hensberge et al. 1980; Bianchi & Ber- nacca 1980). An orbital period of 186.5 days was inferred from the recurrence times of outbursts, which typically last x10 days (Watson, Warwick, & Ricketts 1981; Priedhorsky & Terrell 1983). Pulse frequency variations over multiple EXOSAT observations imply an eccentricity of e 2 0.6 (Cook & Warwick 1987). To date, BATSE has observed seven outbursts, of which three are shown. The separation between the BATSE outbursts is in good agree- ment with the 186.5 day period.
Fic. 27.—Pulse profile as a function of pulsed flux and energy for A0535+26. Pulse profiles of A0535+262 during the giant outburst in 1994 February—March, obtained by epoch-folding CONT data. Profiles in three energy bands are given for four time intervals. The mean luminosity L in ergs s~ 4 is given for each time interval. The time intervals are February 15.1-17.6 (L = 9.1 x 10°”), February 25.0—-March 1.5 (L = 4.5 x 10°”), March 5.0-8.6 (L = 1.6 x 10°”), and March 13.1-15.6 (L = 4.0 x 10°°). Luminosities were calculated from 20-100 keV fluxes based on occultation measurements by assuming a distance of 2 kpc and assuming the 20-100 keV band contains 45% of the bolometric flux.
Fic. 27.—Pulse profile as a function of pulsed flux and energy for A0535+26. Pulse profiles of A0535+262 during the giant outburst in 1994 February—March, obtained by epoch-folding CONT data. Profiles in three energy bands are given for four time intervals. The mean luminosity L in ergs s~ 4 is given for each time interval. The time intervals are February 15.1-17.6 (L = 9.1 x 10°”), February 25.0—-March 1.5 (L = 4.5 x 10°”), March 5.0-8.6 (L = 1.6 x 10°”), and March 13.1-15.6 (L = 4.0 x 10°°). Luminosities were calculated from 20-100 keV fluxes based on occultation measurements by assuming a distance of 2 kpc and assuming the 20-100 keV band contains 45% of the bolometric flux.
Fic. 29.—4U 1145-619 frequency and pulsed flux measurements from BATSE. The pulse frequencies, which have not been orbitally corrected as not all the orbital parameters are known, were determined at 1 day inter- vals from the power spectra of the 20-70 keV CONT data (see Appendix B, § 1.1). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 15 keV (see Appendix B, § 2.3). A1118—616.— Ariel 5 discovered 406.5 s pulsations from A1118—616 in 1974 (Ives, Sanford, & Bell-Burnell 1975), and the optical companion was later identified to be the Be star He 3-640/Wray 793 (Chevalier & Ilovaisky 1975; Heinze 1976; Wray 1976). See Figure 30. Since the initial discovery, no outbursts were observed until BATSE detected one from 1991 December 30 to 1992 January 10 (MJD 48,621—48,633). This outburst reached a 20-70 keV pulsed intensity of ~14 counts s~' on 1992 January 3, followed by approximately 50 days of erratic flaring behav- ior with a maximum on 1992 February 1 (MJD 48,654) (Coe
Fic. 29.—4U 1145-619 frequency and pulsed flux measurements from BATSE. The pulse frequencies, which have not been orbitally corrected as not all the orbital parameters are known, were determined at 1 day inter- vals from the power spectra of the 20-70 keV CONT data (see Appendix B, § 1.1). The pulsed fluxes were determined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 15 keV (see Appendix B, § 2.3). A1118—616.— Ariel 5 discovered 406.5 s pulsations from A1118—616 in 1974 (Ives, Sanford, & Bell-Burnell 1975), and the optical companion was later identified to be the Be star He 3-640/Wray 793 (Chevalier & Ilovaisky 1975; Heinze 1976; Wray 1976). See Figure 30. Since the initial discovery, no outbursts were observed until BATSE detected one from 1991 December 30 to 1992 January 10 (MJD 48,621—48,633). This outburst reached a 20-70 keV pulsed intensity of ~14 counts s~' on 1992 January 3, followed by approximately 50 days of erratic flaring behav- ior with a maximum on 1992 February 1 (MJD 48,654) (Coe
Fic. 30.—A118—616 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have not been orbitally corrected as the orbital parameters are unknown, were determined by Coe et al. (1994b) at 4 day intervals by epoch-folding the 20-50 keV DISCLA data at a range of trial frequencies (see Appendix B, § 1.2). The pulsed fluxes were deter- mined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 15 keV (see Appendix B, § 2.3).
Fic. 30.—A118—616 frequency and pulsed flux measurements from BATSE. The spin frequencies, which have not been orbitally corrected as the orbital parameters are unknown, were determined by Coe et al. (1994b) at 4 day intervals by epoch-folding the 20-50 keV DISCLA data at a range of trial frequencies (see Appendix B, § 1.2). The pulsed fluxes were deter- mined at 1 day intervals by assuming an exponential spectrum with an e-folding energy of 15 keV (see Appendix B, § 2.3).
Fic. 31.—Frequency histories of Cen X-3 . Upper panel: The long-term frequency history of Cen X-3. Lower panel: High-resolution BATSE mea- surement of the intrinsic spin frequencies of Cen X-3. Interestingly, this switching behavior is very common. At least four of the eight persistent pulsars observed by BATSE show torque reversals between steady spin-up and steady spin-down. The 7.6 s pulsar 4U 1626—67 underwent a reversal to smooth spin-down at aratev ~ —7 x 10°13 Hz
Fic. 31.—Frequency histories of Cen X-3 . Upper panel: The long-term frequency history of Cen X-3. Lower panel: High-resolution BATSE mea- surement of the intrinsic spin frequencies of Cen X-3. Interestingly, this switching behavior is very common. At least four of the eight persistent pulsars observed by BATSE show torque reversals between steady spin-up and steady spin-down. The 7.6 s pulsar 4U 1626—67 underwent a reversal to smooth spin-down at aratev ~ —7 x 10°13 Hz
Fic. 33.—Observed relationships between flux and pulsar spin-up rate v. The upper panel shows the spin-up rate of A0535 +262 during the 1994 giant outburst vs. the measured total 20-100 keV flux, determined from Earth occultation measurements. The bottom panel shows the spin-up rate of GRO J1744—28 during the 1995 December—1996 March outburst vs. the 20-50 keV rms pulsed flux. For both sources, the square symbols are from the outburst rise, and the diamond symbols are from the outburst decline. The dotted curves are power laws with the expected index of 6/7, while the dashed curves are best-fit power laws. The best-fit index for A0535 + 262 is 0.951(26), and that for GRO J1744 — 28 is 0.957(26). All three systems suggest a larger y than the naive predic- tion of y = §. In particular, the system with y > 1, indicates an increase in r,, with M, if it is proportional to the mea- sured flux. However, one must remember that BATSE does not measure bolometric flux but only the 20-50 keV pulsed flux. In addition, the large changes in beaming fraction implied by the changing pulse profiles need to be modeled and accounted for. The observed flux may thus be related
Fic. 33.—Observed relationships between flux and pulsar spin-up rate v. The upper panel shows the spin-up rate of A0535 +262 during the 1994 giant outburst vs. the measured total 20-100 keV flux, determined from Earth occultation measurements. The bottom panel shows the spin-up rate of GRO J1744—28 during the 1995 December—1996 March outburst vs. the 20-50 keV rms pulsed flux. For both sources, the square symbols are from the outburst rise, and the diamond symbols are from the outburst decline. The dotted curves are power laws with the expected index of 6/7, while the dashed curves are best-fit power laws. The best-fit index for A0535 + 262 is 0.951(26), and that for GRO J1744 — 28 is 0.957(26). All three systems suggest a larger y than the naive predic- tion of y = §. In particular, the system with y > 1, indicates an increase in r,, with M, if it is proportional to the mea- sured flux. However, one must remember that BATSE does not measure bolometric flux but only the 20-50 keV pulsed flux. In addition, the large changes in beaming fraction implied by the changing pulse profiles need to be modeled and accounted for. The observed flux may thus be related
Root MEAN SQUARE SPIN DERIVATIVES FOR PERSISTENTLY ACCRETING NEUTRON STARS
Root MEAN SQUARE SPIN DERIVATIVES FOR PERSISTENTLY ACCRETING NEUTRON STARS
Fic. 34.—Power density spectra of torque fluctuations in persistent pulsars, computed using methods described in Deeter et al. (1987). Open squares indicate power from apparent torque variations introduced by counting noise, which have been subtracted from all measured values. The integrated rms torque variations for these observations are provided in Table 5.
Fic. 34.—Power density spectra of torque fluctuations in persistent pulsars, computed using methods described in Deeter et al. (1987). Open squares indicate power from apparent torque variations introduced by counting noise, which have been subtracted from all measured values. The integrated rms torque variations for these observations are provided in Table 5.
Fic. 35.—Galactic location of BATSE-detected giant outbursts from Be transients. The circles denote the location within the galaxy of those Be transients that were detected by BATSE during giant outbursts.
Fic. 35.—Galactic location of BATSE-detected giant outbursts from Be transients. The circles denote the location within the galaxy of those Be transients that were detected by BATSE during giant outbursts.
Fic. 36—Cumulative b distribution (left-hand panel) and cumulative ¢ distribution (right-hand panel) of the seven Be-transient systems detected by BATSE during giant outbursts: A0535+262, A1118—61, GRO J1750—27, GRO J2058+42, 2S 1417-624, 4U 0115+ 634, and 4U 1145-619. The histogram and solid line respectively represent the actual data and the model prediction. Since our modeling has +/ and +b symmetry, we use the absolute values of ¢ and b in this analysis. A sampling distance of 40z), with z) = 100 pc gives the best fit to the b distribution and yields an acceptable Kolmogorov-Smirnov statistic. The relatively poor correspondence between the / distribution and the model is due to the clustering of transients at 60° <<? <90°.
Fic. 36—Cumulative b distribution (left-hand panel) and cumulative ¢ distribution (right-hand panel) of the seven Be-transient systems detected by BATSE during giant outbursts: A0535+262, A1118—61, GRO J1750—27, GRO J2058+42, 2S 1417-624, 4U 0115+ 634, and 4U 1145-619. The histogram and solid line respectively represent the actual data and the model prediction. Since our modeling has +/ and +b symmetry, we use the absolute values of ¢ and b in this analysis. A sampling distance of 40z), with z) = 100 pc gives the best fit to the b distribution and yields an acceptable Kolmogorov-Smirnov statistic. The relatively poor correspondence between the / distribution and the model is due to the clustering of transients at 60° <<? <90°.
Fic. 37—Cumulative distribution of the mass functions of the Be- transient pulsars. The measurements are given by the stairstep curve. The remaining curves give the expected distribution for random orientation and a constant companion mass of 4 M , (dotted line), 6 M (dashed line), and 10 M 4 (dot-dashed line). The widest resulting orbits (and those of longest period)
Fic. 37—Cumulative distribution of the mass functions of the Be- transient pulsars. The measurements are given by the stairstep curve. The remaining curves give the expected distribution for random orientation and a constant companion mass of 4 M , (dotted line), 6 M (dashed line), and 10 M 4 (dot-dashed line). The widest resulting orbits (and those of longest period)
Fic. 38.—Eccentricities of Be/X-ray pulsars systems plotted against the orbital period. The dotted lines give the possible range of the orbital period for each system prior to the supernova that formed the neutron star, assuming an initially circular orbit and an asymmetric supernova explo- sion. The kick velocities are likely to be larger than the orbital velocities in the presupernova systems (+100 kms“). Lyne & Lorimer (1994) found from a study of radio pulsar proper motions a mean kick velocity of 450 +90 km s 1. The majority of systems are therefore disrupted. Kalogera (1996) gives analytic expressions for the distribution of orbital parameters of the undisrupted systems, assuming an ini- tially circular orbit and a Gaussian kick velocity distribu- tion. The form of the eccentricity distribution in the limit of large kick velocity relative to orbital velocity is found to be independent of all other parameters. In Figure 39, the cumulative distribution of the observed eccentricities is
Fic. 38.—Eccentricities of Be/X-ray pulsars systems plotted against the orbital period. The dotted lines give the possible range of the orbital period for each system prior to the supernova that formed the neutron star, assuming an initially circular orbit and an asymmetric supernova explo- sion. The kick velocities are likely to be larger than the orbital velocities in the presupernova systems (+100 kms“). Lyne & Lorimer (1994) found from a study of radio pulsar proper motions a mean kick velocity of 450 +90 km s 1. The majority of systems are therefore disrupted. Kalogera (1996) gives analytic expressions for the distribution of orbital parameters of the undisrupted systems, assuming an ini- tially circular orbit and a Gaussian kick velocity distribu- tion. The form of the eccentricity distribution in the limit of large kick velocity relative to orbital velocity is found to be independent of all other parameters. In Figure 39, the cumulative distribution of the observed eccentricities is
Fic. 39.—Cumulative distribution of Be/X-ray pulsar binary eccentric- ities. The observed distribution is given by the solid line. The predicted distribution, normalized for a total of eight sources with e < 0.5, is given by the dashed line (see text). The dotted lines given the eccentricity range allowed by the upper limit for the eccentricity of 2S 1553—54 and the errors on the eccentricity of GS 0834 — 430.
Fic. 39.—Cumulative distribution of Be/X-ray pulsar binary eccentric- ities. The observed distribution is given by the solid line. The predicted distribution, normalized for a total of eight sources with e < 0.5, is given by the dashed line (see text). The dotted lines given the eccentricity range allowed by the upper limit for the eccentricity of 2S 1553—54 and the errors on the eccentricity of GS 0834 — 430.
Note—These channel boundaries are approximate and are averaged over the eight detectors. Each detector has slightly dif- ferent edges, as shown in Fig. 40. The CONT edges are program- mable; the displayed values are typical, and are computed using the calibration discussed in Preece et al. 1997. We have made use of CONT data through channel 7 in our analyses. ENERGY CHANNELS IN BATSE DISCLA AND CONT DaTA TABLE 6
Note—These channel boundaries are approximate and are averaged over the eight detectors. Each detector has slightly dif- ferent edges, as shown in Fig. 40. The CONT edges are program- mable; the displayed values are typical, and are computed using the calibration discussed in Preece et al. 1997. We have made use of CONT data through channel 7 in our analyses. ENERGY CHANNELS IN BATSE DISCLA AND CONT DaTA TABLE 6
Fic. 40.—The energies of the Large Area Detector (LAD) channel boundaries for the CONT data. We display the lower and upper boundaries of CONT channels 1—6 for all detectors. The gain is stabilized on board using the 511 keV background feature. These bin energies apply to two different intervals, MJD 48,406.11-49400.69 and MJD 49,419.69 — 50062.82, which together comprise most of the CGRO mission.
Fic. 40.—The energies of the Large Area Detector (LAD) channel boundaries for the CONT data. We display the lower and upper boundaries of CONT channels 1—6 for all detectors. The gain is stabilized on board using the 511 keV background feature. These bin energies apply to two different intervals, MJD 48,406.11-49400.69 and MJD 49,419.69 — 50062.82, which together comprise most of the CGRO mission.
Fic. 41.—Typical power spectrum of the BATSE LAD background in the 25-33 keV range. The data shown are for CONT Channel 1 (25-33 keV) from LAD 2 on MJD 48,851—48,852 when the mean count rate was 423.7 counts s~‘. The top curve is the unprocessed raw data, while the bottom curve is after background subtraction.
Fic. 41.—Typical power spectrum of the BATSE LAD background in the 25-33 keV range. The data shown are for CONT Channel 1 (25-33 keV) from LAD 2 on MJD 48,851—48,852 when the mean count rate was 423.7 counts s~‘. The top curve is the unprocessed raw data, while the bottom curve is after background subtraction.
Fic. 42.—The 20-75 keV BATSE LAD angular response to an incident photon power-law spectrum dN/dE oc E~*, for « = 2 and « = 5. For comparison, cos 6 response (dotted line) and cos” 6 response (dashed line) are also shown. A consequence of the variation of angular response with energy is that the response to a pulsar depends upon its energy spectrum. Typical accreting pulsar spectra can be represented as dN/dE « E~*, where 2 < a < 5 in the 20-100 keV range. Figure 42 shows the angular response to a 20-75 keV photon power-law spectrum for a = 2 and « = 5. For comparison, cos 8
Fic. 42.—The 20-75 keV BATSE LAD angular response to an incident photon power-law spectrum dN/dE oc E~*, for « = 2 and « = 5. For comparison, cos 6 response (dotted line) and cos” 6 response (dashed line) are also shown. A consequence of the variation of angular response with energy is that the response to a pulsar depends upon its energy spectrum. Typical accreting pulsar spectra can be represented as dN/dE « E~*, where 2 < a < 5 in the 20-100 keV range. Figure 42 shows the angular response to a 20-75 keV photon power-law spectrum for a = 2 and « = 5. For comparison, cos 8
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