agn grad lectures 2007[1]

AGN, accretion and Jets

Overview

• AGN classification and unification theories

• Determining black hole masses

• Jets: emission mechanisms and environmental impact

• Accretion

Journal club reading

2nd lecture (Feb. 26th): Faranoff and Riley, 1974, MNRAS, 167, 31

3rd lecture (March 4th): Arav et al., 1998, MNRAS, 297, 990

Further reading

• AGN Active Galactic Nuclei, J. Krolik, Princeton

• Continuum Radiation Processes High Energy Astrophysics, M. Longair, CUP (2 vols); Radiation Processes in Astrophysics, Rybicki & Lightman, Wiley.

• Optical line radiation Astrophysics of Gaseous Nebulae and Active Galactic Nuclei, D. Osterbrock, University Science Books.

• Jets Beams and Jets in Astrophysics, ed. P. Hughes, CUP, Physics of Extragalactic Radio Sources, D De Young, Chicago U.P.

• Accretion Accretion Power in Astrophysics, Frank, King & Raine, CUP.

Active Galactic Nuclei

What is an AGN?1. Active galactic nucleus: A galaxy for which the bolometric

luminosity cannot be explained by the stars contained therein.

2. Paradigm: extra luminosity is explained by the emission from gas falling into a super massive (> 106 M0) black hole, i.e. accretion. Note most, if not all, galaxies contain super massive black holes but the accretion rate is too low for them to be called AGN (for instance the Milky Way).

3. Observationally, an AGN is a galaxy with some of the following phenomena associated with its nucleus:

(a) Radio: very small angular size and high surface brightness.

(b) Bright extended radio emission.

(c) Broad-band continuum: bright from radio to X-rays or γ-rays.

(d) Strong ionized narrow (broad) emission lines (optical to X-rays).

(e) Luminosity variability on minutes to months timescale.

(f) High linear polarization.

The galaxy tree

Roy, 2002 PhD thesis

How do we find AGN?• Optical: from colour (narrow bandpass imaging) and

emission lines (spectroscopy).• Infra-red emission, but difficult to distinguish from

starburst galaxies.• X-rays luminosity• γ-rays luminosity• Radio luminosity and extended emission.

So, detectable in all wavelength bands!

All of these techniques have their own selection effects, thereby detecting different AGN. The selection is redshift/orientation/environment-dependent as well as dependent on intrinsic

differences.

Classification of AGN

Four key parameters:

1. Luminosity from the accretion process: thermal continuum of the accretion disk, and emission lines in the IR to X-ray resulting from photo-ionization from the accretion process (and potentially the inner part of the jet).

2. Existence and power of the jet: produces non-thermal emission, particularly radio but sometimes also optical to γ-ray.

3. The existence of lobes and hotspots in sources where the jet stay collimated, versus tails and lack of hotspots is sources where jet de-collimates along its trajectory.

3. Orientation (obscuration and beaming).

The relative amount of energy radiated by jet varies by orders of magnitude: radio-loud radio-quiet. As does the luminosity from the accretion process.

AGN ZOO

name point-like

broad-band

broad lines

narrow lines

radio Varia-ble

Polar-ized

radio loud quasar

yes yes yes yes yes some some

Radio quiet quasar

yes yes yes yes weak weak weak

Broad line RG FRII

yes yes yes yes yes weak weak

Narrow line FR I+II

no no no yes yes no no

Optical violent var.

yes yes weak weak yes yes yes

BL Lac objects

yes yes no no yes yes yes

Seyfert 1 yes yes yes yes weak some weak

Seyfert 2 no yes no yes weak no some

LINERs no no no yes no no no

Seyfert 1’s have detected host galaxies of similar L to the nucleus, radio quiet quasars have a nucleus that is much brighter than the host galaxy.

Properties measured in the optical band, except radio. Adapted from Krolik

optical spectra: one method of classification

Radio quiet

Radio loud

LINERs and normal galaxies are also radio quiet.

Seyfert Galaxies• Bright stellar nucleus, host galaxy easily detected.• Strong optical emission lines used for classification.• Usually (not always)In spiral host galaxies.• Bright in X-rays.•First described by Seyfert in 1943.•Type 1: broad and narrow lines in optical.•Type 2: narrow lines only in optical spectra.

Circinus spiral (Seyfert 2)

Seyfert galaxies: soft X-ray spectra

NGC 5548, Kaastra et al. 2002

Seyfert 1

NGC 1068, Kinkhabwala et al. 2002

Seyfert 2

Absorption lines are blueshifted, indicating an outflow.

Seyfert galaxy spectral classification• Seyfert 1: In optical and UV: broad (up to 10000 km/s

FWHM) permitted lines (H, H, CIV, etc.); narrow forbidden lines ([O II], [O III], [Ne V], [O VII]); in X-ray strong continuum (black body: disk? and a power-law: Comptonized disk emission?) with narrow absorption lines (O V – O VIII, N VII, C VI, Ne IX– Ne X).

• Seyfert 2 Narrow permitted (up to 1000 km/s FWHM) and forbidden lines, very weak emission in UV. Little continuum emission in soft X-rays. Fe Kα line and hard X-ray continuum similar to Seyfert 1’s.

• From the width of the optical lines there is a continuum between Seyfert 1 and Seyfert 2: e.g. Seyfert 1.5.

Generalized to type 1 (unobscured) and type 2 (obscured) AGN.

Unification of Seyferts, type 1 and 2 AGN

• A cartoon model of interior of an AGN (Urry & Padovani 1995), very approximate!

• Seyfert 1: see along the jet axis (note in general there are no jets in Seyferts).

• Seyfert 2: see through the obscuring torus (unlikely in the form of a doughnut).

• Explains the lack of non-polarized broad lines and soft X-ray continuum in Seyfert 2’s.

Obscured Type 2 AGN

Antonucci et al. 1994

Ionization cone and obscuring torus

NGC 5252, Morse et al. 1998 VLTI MIDI, NGC 1068, Jaffe et al. 2004

9 10 11 12 13

Wavelength (μm)

Spatial resolution is 26 mas, size is 2.1x3.4 pc

800 K

320 K

Evidence for the obscuring ‘torus’

• Ionization cones: Narrow line region (NLR) emission at large distances is roughly collimated.

• Polarization: Broad lines seen in polarized light in Type 2 object. Scattering by dust (or free electrons).

• Ionizing continuum observed from Type 2 objects is insufficient to photoionize the NLR.

• Soft X-ray absorption in Type 2 objects.

• The opening angle of the Obscuring torus is determined by the ratio between type 1 and 2 Seyferts/AGN.

Physical parameters of the broad line region• Line widths >> (kT/mp)1/2 10 kms-1 for T = 104 K. Bulk

motions therefore required. Sometimes double peak line profiles rotation of the emitting gas.

• Density [OIII] lines absent. Critical density for collisional de-excitation of the upper level producing these lines is 1014 m-3, so this is a lower limit for BLR density. CIII] is present, implying density < 1016 m-3. Detailed modelling gives a few x 1015 m-3. (Note there are other mechanisms to de-excite, possibly invalidating density diagnostic).

• Size from reverberation mapping: ~0.01 pc.• Structure in the form of an outflowing wind, from the

smoothness of the line profiles and the size of the BLR.• Mass typically 2 solar masses.

Physical parameters: narrow line region

• Density from [SII] doublet 108 - 1010 m-3

• Temperature from [OIII] line ratios 10000 - 25000K.

• Size Characteristic scale from imaging is 50 pc - 20 kpc. Distinction sometimes made between NLR and E(xtended)NLR. Possibility of smooth transition to BLR at inner boundary.

• Filling factor f 10-6 -10-5.

• Velocity field Rotation usually a good first approximation, with velocities up to 2000 kms-1.

• Mass typically 107 solar masses - vastly more than BLR.

NLR morphology and kinematics

• Morphology tends to be axisymmetric rather than spherically symmetric.

• Ionization cones seen in (e.g.) [OIII]. Wedge-shaped structures with opening angles 30 - 100 deg.

(extended) NLR of NGC 4151, Seyfert 1.5 observed with HST.

LINERS• LINER: Low Ionization Nuclear Emission line Region.• Optical: low ionization lines: strong O I, already weak [O III] low

temperature gas, sometimes broadened lines.• X-rays: strong continuum with same lines in emission, but broadened, as

Seyfert galaxies plus Fe XXV and Fe XXVI. • X-ray emission lines come from an extended region: emission from SNR’s

and X-ray binaries.• Have a rather low luminosity compared to other AGN. Are believed to be

between Seyferts and inactive galaxies.

M 81, Page 2003

Quasars: the nucleus outshines the host galaxy

Quasi-stellar (radio) sources: bright nucleus; strong emission lines.

The radio-loud versus radio-quiet• Dichotomy originally found was a selection effect of the

surveys. A continuous radio luminosity function is observed in volume limited surveys.

• Radio-quiet radio-silent: Radio-quiet objects sometimes have pc-scale jets, and have similarly bright nuclei in radio.

• Radio-quiet galaxies seem to miss the bright lobes observed in radio-loud galaxies.

• Radio-loudness is probably a short (less than a few 108 years) phase in the life of a (radio-quiet) galaxy.

Faranoff and Riley I (FR I) galaxy

Tail

TailJets

• 3C 31: VLA 1.4GHz image; bends in jet and tail are common.• Both jets are visible and are brighter, but less collimated than in FR II galaxies.• Tails are generally less bright than the lobes in FR II galaxies.

Faranoff and Riley II (FR II) galaxy

• The jet (weaker than in FR I’s) is mostly one-sided (beaming) and straight.• The lobes are back-flowing gas and are brighter than the jet. • There are hotspots.

Ingredients for FR I and II• Central component: core or nucleus. The parsec-scale jet

base? (RQQ without jets have similar core luminosities).

• Jet: often the luminosity varies with distance from the core, appears in bright knots.

• Hotspot or termination point (FR II), major disruption of the jet, multiple hotspots in are often observed.

• Lobe (FR II): back-flowing plasma from the hotspot, generally at higher frequencies is not observed all the way back to the core.

• Tail (FR I): emission from the non-collimated jet.

Faranoff and Riley II (FR II) galaxy

FRII hotspot characteristics• Jet terminates in a strong shock (the hot-spot).

Compression, field amplification and in some cases Fermi acceleration give enhanced emission.

• Flow around the hot-spot is complex and 3-dimensional. Post-shock flow speeds may still be relativistic.

• Particles escape from the hot-spot and in general flow back towards the nucleus, forming the lobe. The external medium is pushed further out, creating what is sometimes called a cavity, i.e. the lobe.

• Advance speed of hotspot < 0.1c (I.e. << jet speed ) but supersonic -> bow shock in IGM.

Evolution of FRII sources

• Youngest FRII sources observed so far (compact symmetric objects or CSO’s) are 10 - 30 pc long and have measured advance speeds of 0.2c, so are inferred to be a few hundred years old.

• Models suggest that the advance speed falls and the radio luminosity decreases with time.

• Typical ages inferred for FRII sources are 107 - 108

years from advance speed of the hotspots and their distance from the core.

• General, the jet is an intermittent phase in the live of an AGN (double-double radio galaxies).

The effects of motion and buoyancy in a galaxy cluster

Looking in the jets: BL Lac objects and Optical Violently Variable quasars

• Mrk 421: BL Lac

• Spectral energy distribution (SED) plot.

• No or weak lines, as we don’t see the nucleus, just the jet.

• Blazars: both classes.

Further unification models of AGN

Obscuration (torus) and beaming (for jets coming towards us) both operate.

• Existence of broad/narrow lines: obscuration (see Seyferts).

• BL Lac’s are thought to be FR I’s seen through the jet (similar luminosity properties after taking beaming effects into account).

• OVV quasars and flat spectrum radio quasars are thought to be FR II’s seen through the jet (similar as above).

• Steep spectrum radio quasars are FR II’s seen at ~38º from LOS.

• FR II’s seem to have higher beaming factors (one sided-jets) than FRI’s, but have similar viewing angles.

• The relation between FR I and FR II is poorly understood.

Exceptions and contradictions

None of the classifications and unifications are completely water-tight:

In some AGN, Hβ (used to classify type I versus type II) is narrow but O II is not: problem in narrow-line versus broad-line classification. Due to mass of the BH?

Some Seyfert 1’s seem to become Seyfert 2-like in X-rays for days to months. Slight difference in opening angle obscuration? Or does the accretion stop for that period?

FR I’s are less luminous in radio, but the only extra galactic sources detected in ultra high (TeV) γ-ray emission.

Black hole mass determination• Individual stellar velocities Milky Way (3 x 106 solar

masses within 0.01 pc).• Water masers in NGC 4258: 3.6 x 107 solar masses

within 0.1 pc (VLBA).• Resolved gas kinematics e.g. M87: 3.2 x 109 solar

masses within 18 pc (HST).• Stellar velocity dispersion• Reverberation mapping

Black hole - bulge luminosity correlation:

MBH / Mbulge 10-3

Orbital motion around the Galactic centre

Water masers in NGC 4258

Argon et al. 2007

Water maser: results for NGC 4258

• Water masers are point tracers of mass. They emit at 1.35 cm and can be observed with VLBI (angular resolution 200as; spectral resolution 0.2 kms-1).

• Masers in nearly edge-on disk show Keplerian rotation in a warped disk, so v = (GM/r)1/2.

• M = (3.9±0.1) x 107 solar masses, for D=7.2Mpc, and inclination of 82º.

• Small distance of the masers to the centre central mass can only be a black hole.

• Several other nearby AGN with measured masers, not all masers seem to lie only in the disk.

Gas kinematics in M87• Same principle as masers, but poorer spatial resolution.• Spatially-resolved HST spectroscopy of gas disk, emitting the Hα line.• Assume Keplerian rotation; ignore warping, non-circular rotation.• Central mass 3.2 x 109 solar masses.

Stellar kinematics and reverberation mapping

• Stellar kinematics: Need to assume a model for the stellar orbits and build a stellar orbit library, gives systematic uncertainty. Use observed velocity as well as velocity dispersion from absorption lines, possible for less well resolved galaxies.

• Reverberation mapping: Using light-travel time lag between continuum and line flux and the line width, again assuming Keplerian orbits. Light-travel time lag is geometry (unknown) dependent, leading to a ~30% uncertainty. But, not dependent on spatial resolution, so available to much larger z.

Jets

Physical processes in jets

• Radiation mechanisms

• Radiative transfer

• Relativistic flow

• Particle acceleration

Jets are energetically dominated by non-thermal plasma containing relativistic electrons, magnetic field (and perhaps protons). The particle energy

distributions are approximate power laws.

Synchrotron radiation summary

• Relativistic e-/e+ (Lorentz factor ≫) spiralling in a magnetic field at the relativistic gyro frequency g = eB/2mec. I.e. relativistic version of cyclotron radiation.

• Unlike the cyclotron radiation, the synchrotron process produces a continuous spectrum without lines.

• Energy loss rate for a single particle = 2Tc2sin2, where is the pitch angle and T is the Thomson cross-section.

• Critical frequency c = (3/2)gsin. Spectrum falls off at about this frequency.

• Maximum emission from a single e- is at 0.29c, with c= 1.22 x1010 2 B (frequency in Hz; B in T).

• The broad-band spectrum is determined by the energy spectrum of the radiating particles.

Properties of synchrotron radiation

• Energy distribution e- : N(dn0-p dover a wide

energy range (from observation and shock theory). Spectral index , defined by S() = S0, is given by p = 2 + 1. Typically, 0.5 < < 3 for optically-thin emission.

• Polarization: up to (3+3)/(3+5) 0.7 (normally much less is observed) for a uniform field and isotropic pitch angles, E-vector perpendicular to projected magnetic field.

• Synchrotron cooling:higher energy e- lose energy faster that lower energy e-. This leads to downward curvature of the spectrum at high frequencies - details depend on pitch-angle scattering and particle history.

Synchrotron emission

Emission cones at 2 points of an accelerated particle’s trajectory (Rybicki & Lightman)

e- path in uniform B field, and polarization E vector.

Theoretical spectrum for optically thick/thin synchrotron radiation

Optically thick

Optically thin

Assuming a power-law electron distribution.

Minimum energy density from synchrotron radiation

• Synchrotron radiation requires particles and fields, both of which have energy: how much?

Energy spectrum

Total emitted power perelectron

Total emitted power perunit frequency

Approximate emissionfrequency

Energy in particles

Emission per unit volume and frequencyinterval

Integrate over energydistribution (ignore upper limit for k > 2)

Combine, addingfudge factor K fornon-radiating particles

.. And field

Observed spectralindex = (k-1)/2;Lower frequency limit

Add the two and differentiate w.r.t. B -> minimum where

Minimum total energy per unitvolume

Jets• Where to find jets Young stars; binary stars;

pulsars; GRB and supernovae Ib/c,II; AGN.• Parsec-scale jets Initial collimation and

propagation; variability; shocks and superluminal motion: radio VLBI/VLBA observations. Only AGN.

• High-energy emission: X-rays to TeV emission.• Large-scale jets Interaction with the IGM, only

AGN.• Jet formation why have some AGN jets and others

not?

What are jets and why are they important?

• Fast, (highly-)collimated outflows from accreting (compact) objects.

• Occur in some accreting: young stars, neutron stars and black hole binaries (‘micro-quasars’), AGN and long duration gamma-ray bursts, Supernova Ib/c, II.

• Outflow speeds can be highly relativistic (>0.99c), but for young stars is only few 100 km/s.

• Radiation from radio - TeV gamma-rays for BL Lac’s.• Enormous energy output: luminosity and kinetic.• Major influence on galaxies and clusters through

energy and entropy input of AGN jets.

Jets in stars and pulsars

HH 212, H2 line emission (VLT). Note that the star is invisible due to the thick disk. Codella et al. 2007

Chandra image of the Crab pulsar. NASA/CXC/SAO

Jet resulting from the 1987A supernova

Type II supernova in the LMC, progenitor was a blue supergiant.

Ring of material is from a previous pre-supernova outburst.

The 2 bright spots is the indication for a jet in this source. Jet speed is quite substantial.

SINS website, picture from 1998.

pc-scale jet in NGC 1068 (Seyfert 2)

NGC 1068 is a radio quiet quasar. The jet is less than 70 pc long, and probably a young jet.

Jets in RQQ are not very common.

Sub-pc to kpc scale jet in M 87

Difference scales of jet emission in the M 87 (FR I) or Virgo A.

Note that collimation of the jet must happen within 0.8 pc of the core.

Note the ‘knotty’ character of the jet on smaller scales.

Image courtesy of NRAO/AUI

Collimation and initial propagation• Jets are initially relativistic in all radio-loud (and

quite possibly radio-quiet) AGN, and are perpendicular to the accretion disk.

• They originate very close to the central black hole: best limits 50 - 100 RSCHWARZSCHILD in nearby radio galaxies M87 and Cen A (VLBI observations).

• Parsec-scale jets are mostly one-sided: mainly the effects of Doppler boosting, but free-free absorption also detected in a few cases.

• The radio nucleus is generally assumed to be the optically-thick photosphere at the base of the jet, and appears to be stationary (superluminal expansion is measured from the nucleus).

Radio spectral properties • Observation: power-law spectrum with index >0.5 in

frequency (>2 in energy).• Acceleration is required on all scales because of the

very short synchrotron and inverse Compton loss time-scales.

• Shocks increase emissivity by adiabatic compression of particles and fields, and accelerate particles to ultrarelativistic speeds from a mildly relativistic seed population. Generally applicable for acceleration.

• MHD winds will collimate and accelerate a jet if B-field footpoints are fixed in the accretion disk. Less well studied (Blanford-Znajek mechanism).

Fermi acceleration at strong shocks

1st order Fermi acceleration: energy gain is of order vp/c. vp ≫ vshock, particle gyroradius is larger than shock thickness: collisionless shock. Because of turbulence behind the shock and irregularities ahead of it, particles scatter and there is isotropic velocity distribution.

1st order Fermi acceleration at shocks• Strong shock in region already occupied by fast

particles.• Particles are overtaken by the shock (energy gain

v/c).• Then scattered by irregularities behind and ahead of

the shock, so some of them cross the shock again: energy gain is repeated.

• Non-relativistic strong shock

• N(E)dE E-2dE (= 0.5)• Relativistic shock: α < 0.5

• Observed: α ~ 0.7

Geometry for superluminal motion

Superluminal motion

• Maximum apparent speed for given real velocity c occurs when = cos (differentiate with respect to ) for an approaching component.

• Thus maximum app = with = (1- 2)-1/2.

• Therefore, high superluminal speeds require fast jets at a small angle to the line of sight.

• Observed range of superluminal speeds in extragalactic sources up to >20c (hence > 20)

• In superluminal extragalactic sources we see components moving away from a stationary nucleus on one side only: the approaching side.

Superluminal motion in a microquasar

Superluminal motion in a powerful quasar

• Motion is not always instraight lines.• Components have differentspeeds; some appear to be stationary. • Components are not necessarily blobs of plasma, might be the shock front.• This is a pattern speed.• Apparent speeds up to 15c (hence minimum Lorentz factor) in this OVV quasar.

Doppler beaming and anisotropy• Jet emission is assumed to be intrinsically

symmetrical and bipolar, separated by 180º.• The approaching jet appears brighter because of

Doppler beaming.• A jet pointed towards us appears much brighter than

it does in the rest frame of the emitting plasma.• Consequently, a bright, compact source selected

from a flux-limited sample is more likely to be directed at our line of sight, and show superluminal motion: the hypothesis of isotropy is incorrect.

• But, some YSO jets are also 1-sided, and this is cannot be due to beaming (v ~ few 100 km/s).

Sample statistics: correlation of core prominence and jet sidedness ratio

What are jets made of?

We see only the radiating electrons, but the plasma must be neutral. Hence 3 possibilities; not exclusive:

Relativistic e+e-: Problems with annihilation, Compton drag and gamma-ray opacity on very small scales, but favoured by circular polarization arguments.

Relativistic e-, thermal p+: Avoids Compton drag because protons carry momentum.

Relativistic p+e-: Energetically expensive. No evidence in favour.

Also magnetic field, Poynting flux (energy flux in a propagating EM wave), thermal plasma.

X-ray emission on large scales

• X-ray emission is detected from jets, hot-spots and lobes in a few cases.

• FRI jets: bright in X-ray, often knots, synchrotron(?).

• FRII jets: synchrotron or beamed inverse Compton scattering of CMB photons. Much rarer than FR I X-ray jets.

• Hotspots: some SSC (equipartition); beamed synchrotron?

• Lobes: various estimates of B = (0.2 - 1) Beq; scattering of CMB or accretion disk (IR) photons (note the different photon energy requires different Lorentz factors: i.e. electron energies).

Note the thin jet, which is due to synchrotron radiation, and the very bright core and hotspot. Note the

lack of thermal and lobe emission. FR II’s are normally not in clusters.

X-ray emission from jet and hotspot in Pictor A (FR II)

Wilson et al. 2000

X-ray emission from Cygnus A (FR II)

Note the very bright core, the 3 larger hotspots and emission from thermal gas, as well as weak emission from a possible relic counterjet, probably inverse-Compton radiation of CMB.

0.2-10 keV image, green circle is indication of size of optical detected galaxy.

FR II jet and ICM/IGM interaction schematic

Kino, Kawakatu 2005

Optical X-ray RadioX-ray is brightest near the core, where radio is weakest, and weakest

far from the core where radio is brightest. Probably emission in all bands is from synchrotron radiation, but IC is also possible.

The 3C 273 jet in optical, X-ray and radio

Marshall et al. 2001

VHE emission from PKS 2155-304 (BL Lac)

Fast variability (~250 s) in the very high energy (VHE), i.e. >200 GeV, measured by HESS (High Energy Stereoscopic System), during a flare. The average flux is >10 more than standard. Emission mechanism and origin not understood: IC or synchrotron.

Aharonian et al. 2007

Environmental impact: 3C 84 (FRI)

Central elliptical galaxy in Perseus cluster. Does cluster gas thwart the jet? X-ray (mostly hot cluster gas) false colour and radio contours.

Fabian et al. 2000

Cygnus A (FR II) cluster in X-ray

The black ellipse indicated X-ray and radio emission from Cygnus A. Smooth hot cluster emission surrounds it.

Jet power

Does the jet power come from the accretion disk or the black hole?

Outflow from disk surface powered by gravitational energy release.

Electromagnetic energy extraction from the black hole if the event horizon is threaded by a magnetic field.

How are jets collimated and formed?

Magnetic collimation - a natural mechanism, as toroidal field is expected from a spinning system, and this provides a confining force. However, not too well understood.

Furthermore, in some FR I galaxies we see the jet being re-collimated at kpc-scale from the nucleus. Probably some re-collimation mechanism operates in FR II galaxies along the whole jet, several 100’s kpc long.

Analytical and 2/3D MHD simulations of jets have had some success, but:

Fundamental questions about the power source remain, there are some problems with stability and the initial and boundary conditions are ad hoc.

Accretion

Overview

• Accretion onto black holes and energy extraction.

• Accretion disks

• Accretion in binary stars

Schwarzschild black holes• Non-rotating black holes are described by the

Schwarzschild metric

rs = 2GM/c2 = 3 km (M/MSUN)

is the Schwarzschild radius.

• Event horizon For r < rs, there is no photon trajectory which allows escape.

• Last stable orbit occurs at r = 3 rs.

• Efficiency of energy release from accretion onto a Schwarzschild black hole is related to the binding energy of the last stable orbit. The maximum efficiency is 1 - 81/2 = 0.057.

Kerr black holes

• Kerr metric describes all rotating black holes. They are characterised by the mass M and angular momentum J = aMc (0 a 1) only.

• Dragging of inertial frames If a 0 then there are no stationary observers: every physically realisable reference frame must rotate.

• Last stable orbit More complicated forms. Radii are different for prograde and retrograde orbits (minimum GM/c2 for a = 1).

• Efficiency of energy extraction is higher than for non-rotating holes because the last stable orbit is closer in. Maximum value = 1 - 31/2 = 0.42.

Electromagnetic energy extraction

• Basic idea Large-scale magnetic fields anchored in the disk extract rotational energy.

• Disk re-supplied by fresh infalling material.

• Blandford-Znajek mechanism Field lines are also anchored on the black hole, allowing its rotational energy to be tapped.

• Power / J ~ 1038 (a/m)2 (B/T)2 (M/108 Msun)2

• Process is poorly understood.

Photon propagation near black holes

• Special relativistic Doppler boosting

• Gravitational redshift If dt is the proper time interval seen by a distant observer and dt’ is that seen by an observer close to the black hole, then

• dt’ = (1-rs/r)1/2 dt

• As r -> rs, events which take a finite amount of time as measured near the black hole appear to take divergently long times when observed at large distances (and radiation is redshifted).

• Curvature in photon trajectories

-> emission line skewed to higher energies.

Predicted Fe Kα relativistic line profile

Fe Kα 6.4 keV line is a reflection line of neutral material, supposedly in the accretion disk.

A relativistic line profile was observed with ASCA for the Seyfert galaxy MCG-6-30-15, but is not confirmed by either Chandra or XMM-Newton.

Accretion disks

• Angular momentum is difficult to lose for infalling material. Orbit of minimum energy at constant angular momentum is circular - hence disk.

• Viscosity causes loss of angular momentum, so disk material gradually sinks towards the central object, dissipating energy which can potentially be radiated away.

• What is the viscosity?

• Turbulence

• Magnetic fields

Accretion disk and torus of NGC 4261 (LINER)

Jaffe & Ford

Thin disks• Standard model, well established for accreting

binary stars. Geometrically thin, optically thick disk (alias Shakura-Sunyayev; -disk).

- prescription: csH, where is the kinematic viscosity, cs is the sound speed, H is the disk scale height and 1is assumed to be constant.

• Temperature: If the emission is black-body and comes from close to the last stable orbit, then

• T 106L39-1/4(L/LEdd)1/2 K

• where L39 is the luminosity in units of 1039 W. Hence characteristic temperatures in UV for AGN; X-ray for binary stars.

What is the viscosity mechanism?

Reynolds number, where V is the flow speed,L is a typical length scale and = / is thekinematic viscosity. Low R => viscous flow;high R => turbulence.

From kinetic theory ( = mean free path)=> R 1012.

Therefore, kinetic viscosity is irrelevant, and flow is turbulent.Turbulence and magnetic fields provide an effective viscosity, but are difficult to calculate. Hence the α prescription of Shakura & Sunyaev.

Energy loss rate and spectrum

• Heat dissipation and disk luminosity (assuming Keplerian disk/potential):

• and

• With rm the radius where the stress on the disk is zero, and for a non-relativistic, steady-state disk.

• Energy loss rate is independent of viscosity (which is why we have been able to make progress despite lack of knowledge of the viscosity prescription).

• Spectrum: L is proportional to:• 2 for kT/h ≫ frequency (Rayleigh-Jeans).• 1/3 for frequencies corresponding to temperatures of

material in the disk.• exp(-h/kT) at frequencies above kTin/h.

r

r

rMGM

Q m143

3m

total rMGM

L

21

Radiatively inefficient accretion

Observed accretion with L << LEdd

One explanation is that the accretion rate is low.

Another explanation is an accretion flow in which the radiation rate is very low. This will happen if the disk is:

optically thin and geometrically thick, such that cooling time is much longer than infall time. Therefore the gas falls into the black hole before it has time to radiate. T i ≫ Te.

Radiatively inefficient accretion.

An example: Advection-dominated accretion flow (ADAF), but other models are also possible (ADIOS, CDAF). Some have potential problems with disk stability and boundary conditions.

AGN formation and fuelling

• Formation is fast as quasars exist at z 6.

• Heavy-element abundances Quasar abundances are close to solar. Fe present => time for SNI to have happened in large numbers.

• Merging processes: provide fuel by disrupting gas in the galaxy, and gas capture from the merging galaxy.

• Bar-driven inflow, spiral galaxies show bars often on different scales, which allows gas to fall to the nucleus.

X-ray accretion disk spectrum

• The X-ray continuum spectrum of the accretion disk has at least 2 components:

1) a hard X-ray power-law component with photon index between 1.5 and 2. This is also called the corona and is assumed to be inverse-Compton emission from gas just above the accretion disk. An alternative is that it is magnetically energized gas.

2) a soft excess: an excess of emission in the soft X-rays, well fit by a black body component and potentially direct emission from the accretion disk, although then the temperature of the accretion disk is higher than theory predicts.

Accretion onto compact object in binaries

• Compact object can be a white dwarf (Cataclysmic variables), a neutron star or a stellar mass black hole.

• Primary star (sometimes also called secondary) can be a low mass star (low mass X-ray binary) or a high mass star, O or B star (high mass X-ray binary).

• Different formation and accretion mechanism, and evolution.

• HMXB have a harder X-ray spectrum than LMXBs, and are more likely to show eclipses.

• Accretion onto neutron star can show type I bursts: thermo-nuclear flashes, common in LMXBs; or X-ray pulses from pulsars mostly in HMXBs.

Evidence for black holes

• Analysis of binary star orbits from primary star spectral lines (dependent on inclination, but the mass function gives a lower limit).

• Mass of unseen companion > 3 solar masses (sometimes by large factors).

• Hence cannot be neutron stars supported by neutron degeneracy pressure.

• Lack of type I bursts or X-ray pulses is not conclusive evidence.

Accretion mechanisms in binary systems

• Roche-lobe overflow occurs in a binary system containing a compact object (white dwarf, neutron star or black hole) and a primary star which is on the giant branch.

• Primary expands (due to stellar evolution) so that its surface reaches the inner Lagrange point (saddle point in the gravitational potential between the stars).

• Material can then flow from the giant to the compact companion. The Roche lobe is the equipotential surface which meets the inner Lagrange point.

• Most likely accretion mechanism for LMXBs.

Accretion mechanisms - stellar winds

• If the primary star is smaller than its Roche lobe, but is losing mass rapidly via a stellar wind, then some fraction of the wind can be captured by the compact companion.

• Typical mass-loss rates are between 10-7 and 10-5 solar masses per year for stars between 15 and 60 solar masses.

• These systems are high-mass X-ray binaries, and have X-ray luminosities of 1029 - 1031 W.

Accretion in Cataclysmic variables

• The WD and companion are in a very close orbit.

• The secondary star has not expanded to fill its Roche lobe, and is normally a red dwarf, instead of a giant in the standard Roche lobe overflow picture.

• Due to the close orbit the WD distorts the secondary (donor) star, this distortion will allow gas to flow across the inner Lagrangian point towards the WD.

Effects of magnetic fields• Compact stars (neutron stars or white dwarfs; not black holes)

often have a strong surface magnetic field. This can have a major effect on accretion.

• Wind accretion onto a compact star. Assume field is dipolar, hence energy density

umag ~ (B2/20)(R/r)6 (R is secondary radius). • Numbers for a 1.4 solar mass neutron star accreting at the

Eddington rate:

L = 1.8 x 1031 W, accretion efficiency = 0.1, B = 108 T and R = 10 km.

• Hence the immediate vicinity of the neutron star is magnetically dominated and gas must flow close to the poles of the dipole field in an accretion column.

• In extreme cases, no accretion disk forms.

Magnetic neutron stars

For neutron star with strong magnetic field, disk disrupted in inner parts. This is where most radiation is produced.

Compact object spinning => X-ray pulsar; spun up by disk.

Material is channeled along field lines and falls onto star at magnetic poles

Observational tests of disk accretion• Eclipse mapping Use eclipse by companion star to study

spatial and velocity structure of disk (primarily accretion onto white dwarfs).

• Doppler tomography Observe velocity structure of spectral line; use change of direction caused by orbital motion to reconstruct emission distribution and velocity field.

• Integrated disk spectra

• Lyman edges For face-on disks, expect a discontinuity in the spectrum at the wavelength of Lyman alpha because of the abrupt change in opacity.

• Quasi-periodic oscillations

Doppler Tomography of IP Pegasus

IP Pegasus is a cataclysmic variable. In the accretion disk there are spiral shocks due to the gravitational field of the donor star. The donor star is encircled, and the line indicates the gas infall trajectory.

Steeghs 2003

Quasi-periodic oscillations (QPO) observed in a low-mass X-ray binary

First observed in CV’s with frequencies of a few Hertz.

kHz QPO’s are present in LMXB’s. QPO’s are only observed in the X-ray flux.

Assumed that higher frequency means the gas emitting is closer to the compact object.

Therefore, kHz QPO’s should tell us something about the inner accretion disk. But the emission process and why it is not periodic are poorly understood.

Sco-X1, van de Klis 2000

High-mass X-ray binaries

• Young population, short-lived OB primaries, mostly in spiral arms.

• Some have X-ray pulsars as compact object.

• Spin periods 66 ms - 1000 s; orbital periods > 1 day.

• Spin-up and spin-down are both possible: as there is mass loss from the system as well as mass transfer.

• Roche-lobe overflow, and wind accretion.

• Magnetic field confirmed from cyclotron absorption.

• First stellar mass BH detected is in Cygnus X-1, is in a HMXB.

Low-mass X-ray binaries

• Brightest X-ray sources in the Galaxy.

• Neutron star or black hole as compact object.

• Few contain pulsars (either low magnetic field or magnetic and spin axes are aligned).

• All Roche-lobe overflow.

• Eclipses and dips => orbital period: minutes to years.

• Bursts with typical duration 10 - 30 s (thermonuclear runaway) => not a black hole.

• Example: Sco X-1, a neutron star binary, was the 1st extra-solar X-ray source detected.

Achievements (what we think we know)• Accretion onto a massive black hole is the power source

for AGN.• Synchrotron radiation is the main non-thermal radiation

mechanism for jets and associated emission.• Inverse Compton scattering by non-thermal electrons is

partly responsible for high-energy radiation from jets.• Jet flow is relativistic in AGN.• Obscuring tori exist.• Photoionization is the main source of ionization for the

gas emitting emission, even at kpc distance from the nucleus.

Ignorance• Why do AGN occur at all? What triggers an AGN? How long does it

last?

• How are they fuelled?

• Radiative versus non-radiative accretion.

• Jets: how are they formed? How are particles accelerated? What are they made of? Why do they sometimes occur, sometimes not?

• What is the difference between FR I and FR II galaxies?

• Why are there obscuring tori? What is the torus made off, and what is its geometry?

• How is hot plasma generated near accretion disks to produce AGN hard X-ray emission (coronae)? How are the winds observed in UV/X-ray accelerated?

• What is the origin of the cool gas that form the BLR and NLR?