Radio Pulsars – An Overview R. N. Manchester CSIRO Astronomy and Space Science Sydney Australia Summary Pulsar origins and discovery Basic pulsar timing Pulsar emission Interstellar dispersion, scattering and Faraday rotation A new pulsar distance model Pulsar timing and binary pulsars

What are pulsars? Pulsars are rotating neutron stars radiating beams of emission that we see as pulses as they sweep across the Earth Neutron stars are formed in supernova explosions by the collapse of the core of a massive star Millisecond pulsars are old pulsars that have been recycled by accretion of matter from a binary companion Almost all known pulsars lie within our Galaxy Neutron stars have a mass about 1.4 M Sun but a radius of only ~15 km Their surface gravity is enormous, ~10 12 g Earth, hence their break-up rotation speed is very high, ~1 kHz

A Pulsar Census Data from ATNF Pulsar Catalogue, V1.53 ( (Manchester et al. 2005) Currently 2524 known (published) pulsars 2354 rotation-powered Galactic disk pulsars 249 in binary systems 365 millisecond pulsars 142 in globular clusters 8 X-ray isolated neutron stars 21 magnetars (AXP/SGR) 28 extra-galactic pulsars MSPs “Normal”

Jocelyn Bell and Antony Hewish Bonn, August 1980 The Discovery of Pulsars Cambridge “4-acre” array

Searching for Pulsars Most pulsars have been found in searches at radio frequencies Pulsars have two main properties that are used to distinguish them from (most) other radio signals: periodicity and dispersion Multi-channel data sampled typically at ~100  s intervals Frequency Time Data summed in frequency with a range of dispersive delays and Fourier transformed Searched in modulation frequency with harmonic summing Candidates plotted and selected for confirmation Observed again at same position to confirm periodicity and dispersion

Discovery rate of pulsars since 1968 (Data from ATNF Pulsar Catalogue V1.53) Molonglo has twice discovered more than half the known pulsars – in 1968 and again in 1978 Parkes has discovered nearly two-thirds of the currently known pulsars More than 160 Fermi- related discoveries in past few years

Galactic Distribution of Pulsars (Data from ATNF Pulsar Catalogue V1.47)

Multibeam receiver - 13 beams at 1.4 GHz - very efficient for pulsar surveys Several independent surveys with different optimisations More than 1250 pulsars discovered with the multibeam system since 1997 Excellent database for studies of pulsar Galactic distribution and evolution The Parkes Multibeam Pulsar Surveys

Because of their large mass and small radius, NS spin rates - and hence pulsar periods – are extremely stable For example, in 2001, PSR J had a period of :  ms Although pulsar periods are very stable, they are not constant Pulsars are powered by their rotational kinetic energy They lose energy to relativistic winds and low-frequency electromagnetic radiation (the observed pulses are insignificant) Consequently, all pulsars slow down (in their reference frame) Typical slowdown rates are less than a microsecond per year For millisecond pulsars, slowdown rates are ~10 5 smaller Pulsars as Clocks

Start observation at a known time and average pulses to form a mean pulse profile Cross-correlate this with a standard template to give the arrival time at the telescope of a fiducial point on profile, usually the pulse peak – the pulse time-of-arrival (ToA) Measure a series of ToAs over days – weeks – months – years Transfer ToAs to an inertial frame – the solar-system barycentre Compare barycentric ToAs with predicted values from a model for the pulsar – the differences are called timing residuals. Fit the observed residuals with functions representing errors in the model parameters (pulsar position, period, binary period etc.). Remaining residuals may be noise – or may be science! Measurement of pulsar periods

Sources of Pulsar Timing “Noise”  Intrinsic noise Period fluctuations, glitches Pulse shape changes  Perturbations of the pulsar’s motion/spacetime Gravitational wave background Globular cluster accelerations Orbital perturbations – both Newtonian and relativistic  Propagation effects Wind from binary companion Variations in interstellar dispersion Scintillation effects  Perturbations of the Earth’s motion/spacetime Gravitational wave background Errors in the solar-system ephemeris  Clock errors Timescale errors Errors in time transfer  Receiver noise  Instrumental errors Radio-frequency interference and receiver non-linearities Digitisation artifacts or errors Calibration errors and signal processing artifacts and errors

The P – P Diagram. For most pulsars P ~ MSPs have P smaller by about 5 orders of magnitude Most MSPs are binary, but few normal pulsars are τ c = P/(2P) is an indicator of pulsar age Surface dipole magnetic field ~ (PP) 1/2.... P = Pulsar period P = dP/dt = slow-down rate. Great diversity in the pulsar population! (Data from ATNF Pulsar Catalogue V1.50) Galactic Disk pulsars

Pulsar Braking Braking torque ~ Ω 3 for magnetic-dipole braking; more generally torque ~ Ω n Braking index n can be measured from ν Elapsed time t since P = P 0 is a function of P 0, P, P and n. For P 0 << P, n = 3, have characteristic age τ c If true age known, can compute initial period P 0 B s, magnetic field at NS equator; at pole B = 2B s Measured braking indices are usually <3, e.g., for Crab pulsar, n = 2.55 Indicates that there is a braking contribution from a relativistic wind...

Pulsar Model (Bennet Link) Rotating neutron star Light cylinder R LC = c/  = 5 x 10 4 P(s) km Charge flow along open field lines Radio beam from magnetic pole (in most cases) High-energy emission from outer magnetosphere

For a typical pulsar, P = 1s and P = , B s ~ 10 8 T or G. Typical electric field at the stellar surface E ~  RB s /c ~ 10 9 V/cm Electrons reach ultra-relativistic energies in < 1 mm. Emit  -ray photons by curvature radiation. These have energy >> 1 MeV and hence decay into electron-positron pairs in strong B field. These in turn are accelerated to ultra-relativistic energies and in turn emit  -rays which pair-produce, leading to a cascade of e + /e - pairs. Relativistic pair-plasma flows out along ‘ open ’ field lines. Instabilities lead to generation of radiation beams at radio to  -ray energies. Pulsar Electrodynamics.

Frequency Dependence of Mean Pulse Profile Phillips & Wolsczcan (1992) Pulse width generally increases with decreasing frequency. Consistent with ‘ magnetic-pole ’ model for pulse emission. Lower frequencies are emitted at higher altitudes.

Pulsar Radio Spectra (Malofeev et al. 1994) Most pulsars have steep radio spectra Mean spectral index ~ -1.8 Low-frequency turnovers (~100 MHz) relatively common For many (most?) pulsars, spectrum steepens above about 1 GHz About 70 pulsars undetectable at radio freq; about 50:50 X-ray and γ-ray detections (Maron et al. 2000)

Magnetar Radio Emission Magnetars were discovered as X-ray/γ-ray burst sources with periodic emission modulation, P ~ few seconds About 20 now known; four have detected radio emission Timing showed huge P – implied B s ~ G Radio emission (intensity and pulse shape) extremely variable on short and long timescales Very flat radio spectrum: (Torne et al. 2015) PSR J (Camilo et al. 2006) PSR J J : single pulses at 42 GHz! -J detected at 225 GHz!.

Coherent Radio Emission Pulse timescale gives limit on source size ~ c  t Radio luminosity ~ S D 2 Δν Ω ~10 29 erg s -1 for 1 Jy, 1kpc, 10 9 Hz, 1 sr Brightness temperature: equivalent black-body temperature in Rayleigh- Jeans (low-frequency) limit Radio emission must be from coherent process!

Giant Pulses Intense narrow pulses with a flux density or pulse energy many times that of an average pulse Characterised by a power-law distribution of pulse energies First observed in the Crab pulsar – the pulsar was discovered through its giant pulses! Generally confined to narrow range of pulse phase, often associated with high-energy emission Arecibo 5.5 GHz obs. (Hankins et al. 2003) show pulse widths of a few ns – imply T b > K! Emission from plasma turbulence on 1-m scale? Crab Main Pulse Flux Density Distribution (Karuppusamy et al. 2010) 1.4 GHz

Emission beamed tangential to open field lines Radiation polarised with position angle determined by projected direction of magnetic field in (or near) emission region (Rotating Vector Model) Magnetic-Pole Model for Emission Beam

Mean pulse shapes and polarisation Lyne & Manchester (1988) P.A. Stokes I Linear Stokes V

Orthogonal-mode emission – PSR B P.A. Stinebring et al. (1984) V I %L

Mean pulse profile of PSR J Binary MSP P = 5.75 ms P b = 5.74 d Navarro et al. (1997) Stokes I Stokes V Linear I L V P.A. MSP pulse profiles are more complex – many components Complex PA variations, including orthogonal- mode transitions

Polarisation of PPTA Pulsars Polarisation profiles for 24 PPTA pulsars at 10cm (3100 MHz, PDFB4), 20cm (1400 MHz, PDFB3) and 50cm (730 MHz, CASPSR) All available data averaged, integration times 10h – 500h PSRCHIVE analysis calibration routines Phase-resolved spectral indices and RMs Spectral-index variations across pulse often related to components in profile Some deviations from λ 2 PA-variation Apparent RM variations across pulse most probably related to frequency- dependent differences in emission PA. Spectral Index RM PA L/I (Dai et al. 2015)

Drifting subpulses and periodic fluctuations PULSE LONGITUDE Drifting subpulses Taylor et al. (1975) Periodic fluctuations Backer (1973)

Pulsar Nulling (Wang et al. 2006) Parkes observations of 23 pulsars, mostly from PM survey Large null fractions (up to 96%) - mostly long-period pulsars Nulls often associated with mode changing

PSR B An extreme nuller (Kramer et al. 2006) Quasi-periodic nulls: on for 5-10 d, off for d Period derivative is ~35% smaller when in null state! Implies cessation of braking by current with G-J density Direct observation of current responsible for observed pulses

PSR J : A more extreme nuller (Camilo et al. 2012) Discovered at Parkes during 10cm observations of nearby magnetar P = 0.9s, B s ~3x10 14 G, τ c ~0.5 Myr Relatively strong, ~5 mJy at 2 GHz Turned off after ~ 1 yr of obs. Not detected for ~1.5 yr Turned on with smaller period than expected, but same P as before P smaller by ~ x5 in off state Implies dramatic change in pulsar magnetosphere that persists for times ~ 1 yr and is evidently bimodal Pulse is highly scattered at frequencies <~2 GHz; τ s (1GHz) ~ 2.3 s! Very large RM ~-3000 rad m -2, implies ISM B || ~ 7 μG.

Interstellar Scattering and Scintillation Small-scale irregularities in the ISM electron density deflect ray paths and distort the wavefront from the pulsar Longer deflected paths result in broadening of pulse profile Rays from different directions interfere resulting in modulation in space and frequency - diffractive ISS Motion of the pulsar moves the pattern across the Earth Larger-scale irregularities cause focussing/defocussing of wavefront - refractive ISS

(Lewandowski et al. 2015) Interstellar Scattering For expected (Kolmogorov) distribution of electron density fluctuations, scattering angle θ 0 ~ ν -2 Increases apparent source diameter Scattering delay ~ θ 0 2 ~ ν -4 – much greater at low frequencies Scattering Measure (SM ~ θ 0 d) parameterises the amount of scattering for a given path Wide range for a given DM – not very predictable Observed scattering mainly results from a few strongly scattering regions along path (Cordes & Lazio 2003)

Dynamic Spectra resulting from DISS (Bhat et al., 1999)

DISS Secondary Spectrum Take 2-D Fourier transform of dynamic spectra Secondary spectrum shows remarkable parabolic structures Main structure results from interference between core and outer rays (Stinebring 2006)

VLBI Imaging of DISS Secondary cross-spectrum for PSR B from 327 MHz VLBI between Green Bank and Arecibo Can image scattering screen on AU scales Screen is highly anisotropic – main scattering region ~ 16 AU x 0.5 AU Phase ambiguity for “1-ms” feature resolved using frequency dependence of delay/fringe rate (Brisken et al. 2010) AmpitudePhase

Extreme Scattering Events PPTA DR1 data set Correlated changes in DM, DISS timescale τ 0 and bandwidth ν 0 observed for PSRs J and J Highly turbulent zone ~5 AU in transverse size crossing line of sight. Mean density ~4 cm -3 if sheet extended in l-o-s direction or ~100 cm -3 if spherical Other examples of abrupt changes in DM observed, but no clear correlation with scintillation parameters Almost certainly related to “scintillation arcs” (Coles et al. 2015) ΔDM τ0τ0 ν0ν0

Faraday Rotation & Galactic Magnetic Field (Han et al. 2015) For pulsars can get mean l.o.s. field strength using RM & DM: 675 pulsar RMs 3800 ex-gal source RMs Can map Galactic field structure using RM/DM for different distance ranges

Galactic Magnetic Field from Pulsar RMs Disk field aligned with spiral structure Field in spiral arms generally counter- clockwise Field in inter-arm regions generally clockwise (Han et al. 2015)

Pulsar Distances Interstellar dispersion caused by free electrons along path to pulsar – delay Δt ~ ν -2 – easy to measure For pulsars with independent distances (e.g., parallax, SNR or globular cluster association) get = DM/d, along path. Typical values ~ 0.03 cm -3 Based on these, can build a model for the Galactic n e, e.g., NE2001 Can then use this model to determine distances to other pulsars (Cordes & Lazio 2002) NE2001