Basado en una presentación de Raffaella MorgantiASTRON

Radiogalaxias

Temas

1. ¿Qué son los AGNs y las radiogalaxias? - ¿Cómo encontrarlos? Una radiogalaxia prototipo - Mecanismos de emisión

2. Morfología de la emisión de radio: distintas morfologías, regiones nucleares, chorros altamente colimados – regiones calientes - lóbulos

¿Qué son los Active Galactic Nuclei?

Es difícil dar una definición única.

grandes cantidades de energía (hasta 104 veces más que en una galaxia normal) emitida desde una pequeña región (<1 pc3)

Un AGN puede tener una luminosidad que va de 1042 to 1044 erg/sec

Se cree que la gran energía liberada por los AGN se origina en un hoyo negro supermasivo (106 a 109 Msun en <<1pc)

Se han encontrado AGNs de menor luminosidad

¡Pero la presencia del hoyo negro supermasivo no es suficiente!

Algunas características

Comparison of the continuum emission from a Seyfert galaxy and a normal galaxy

Optical emission linesfor different AGNs

radio

optical X-ray UV

La luminosidad no es el único criterio:

emisión de continuo (comunmente azul) a lo largo de ~13 órdenes de magnitud en frecuencia

líneas de emisión

emisión (en alrededor de 10% de los AGNs)

Radio galaxies

Radio galaxies & radio-loud quasars: the most powerful radio sources

(Usually) extended (or very extended!) radio emission with common characteristics (core-jets-lobes)

Typically hosted by an elliptical (early-type) galaxy

Nevertheless, the radio contribute only to a minor fraction of the energy actually released by these AGNs.(ratio between radio and optical luminosity ~10-4)

Amazing discovery when they were identified with extragalactic, i.e. far away, objects

Unexpectedly high amount of energy involved!

They show most of the phenomena typical of AGNs(e.g. optical lines, X-ray emission etc.) very interesting objects in (almost) all wavebands

in addition they have spectacular radio morphologies

But they are quite rare!

Why are interesting?

How to find them?

Because of the variety of AGNs, there is also a variety of techniques to find them

(e.g. blue colours, strong emission lines etc.).

Here we focus on the way radio galaxies have been found: radio surveys

Radio surveys (some of them….)

3CR (Cambridge Telescope) 328 sources with > - 5o

flux above 9 Jy @ 178 MHz

4C 2Jy 178 MHz Cambridge (+5,6,7C)

PKS ~3Jy 408 MHz Parkes

      Molonglo

B2 0.25 408 MHz Bologna (+B3)

NRAO 0.8Jy 1.4-5GHz NRAO

PKS 0.7Jy 2.7 GHz Parkes

NVSS 2.5 mJy (45” res.) 1.4 GHz NRAO VLA Sky Survey

FIRST 1mJy (~5” res) 1.4 GHz Faint Images Radio Sky at Twenty centimeters

WENSS 300 MHz WSRT

(1 Jy= 10-26 W m-2 Hz-1)

85 mJy

Units that will be used for the radio data

Radio flux in “Jansky” 1 Jy = 10-26 W m-2 Hz-1

or 10-23 erg cm-2 sec -1 Hz-1

Radio power (usually estimated at a certain frequency e.g 1.4 or 5 GHz)

or integrated over a typical (radio) range of frequencies (107 to 1011 Hz)

)(4 12 HzWFDP

Radio power: source of 2 Jy flux (@ 1.4 GHz), z = 0.2 log P = 26.5 W/Hz source of 0.2 Jy flux, z = 0.2 log P = 25.5 W/Hz source of 10 Jy flux, z = 0.2 log P = 21.2 W/Hz

resolution /D 21 cm, D = 64 m 11 arcmin 21 cm, D= 3km 14 arcsec 21 cm, D= 3000 km 1 mas

Resolution important for the identification (radio surveys find not only radio galaxies!)

Difference in power limit for the different surveys

HIPASS beam

ATCA image, July 2001

NGC 6580 (S0)

IC 4933 (Sbc)

‘Confusion’ can be resolved by imaging at higher spatial resolution with large

interferometers (WSRT, VLA or ATCA)

Confusion

Optical identifications

NVSS

radio much larger than optical

resolution ~45 arcsec ~ 45 kpc(1 arcsec ~ 1 kpc at z = 0.04)

Radio galaxies are only found among the most powerful radio sources (together with radio-loud quasars).

radio emission from non-thermal synchrotron process

but (radio) AGNs can also be found at low radio powerhigh radio resolution is required to find a very compact core

(to distinguish non-thermal emission from thermal emission)

Going deeper and deeper

Green: WSRT finding chart at 1.4 GHz with an r.m.s. noise of 13 microJy/beam. Grey: NOAO optical R to a limiting depth of 26 magnitude.

VLBI nondetection at full sensitivity with an r.m.s. noise of 9 microJy/beam.

VLBI detections at full sensitivity with an r.m.s. noise of 9 microJy/beam.

(Morganti & Garrett, 2002, ASTRON Newsletter No. 17; Jannuzi & Dey, 1999, ASP Conference Series, 191, 111)

Deep Wide-Field VLBI Surveys

A prototypical radio galaxy

Any size: from pc to Mpc First order similar radio morphology (but differences depending on radio power, optical luminosity & orientation) Typical radio power 1023 to 1028 W/Hz

Lobes

Core

Jets

Hot-spots

How a radio galaxy works

Zoom-in ofthe central regions

to hot-spots and/or lobes

SupermassiveBlack Hole

accretion disk (UV, Xray)

torus (supposed to hide – for some orientation – the very central regions)

A prototypical radio galaxy

“cocoon” shocked jet gas

backflow

splash-point

bowshockundisturbed intergalactic gas

Observable Diagnostic Constituents Derived Properties 

Radio continuum Relativistic plasma   Thermal plasma

Energetic, Pressure, Jet propagation velocity, Internal magnetic field Ages, Faraday rotation, Magnetic fields

Radio absorption Lines (21cm)

Neutral gas Column density,kinematics

IR-mm continuum Dust Mass, Temperature

IR-mm emission lines (CO) Molecular gas Mass, density Temperature

UV/Optical/near IRContinuum 

Stars Scattered AGN light

Mass, Age,Star-formation ratePolarization properties

Optical emission lines: Ly , H ,[OIII]

Ionized gas (10^4 K) Mass, temperature,Ionized state kinematics

Ly absorption Neutral gas Column density Mass, covering factor

X-ray emission Non-thermal plasmaHot gas (10^7 K) 

Jet (and hot-spots) propertiesCluster properties

2

2

2

1

1

cv

cmE e

Electron energy distribution is a power law:

>>1

Relativistic electrons in a magnetic field

dEkEdEEN p)(

2

1

/

/

2/)1(2/)1( ........),(

ppBP

The radio spectrum is therefore a power law:

S 2/)1( p

Typical ~0.8 p~2.6

For one electron, max frequency

2 for slightly different covers the entire spectrum

Assuming the emission from each can be added up (optically thin case)

1. Energy loss2. Self-absorption in the relativistic electrons gas3. Absorption from ionized gas between us and the source (free-free absorption) torus!

Deviations from a constant spectral index

Theory Reality

Energy loss

The relativistic electrons can loose energy because of a number of process (adiabatic expansion of the source, synchrotron emission, invers-Compton etc.). the characteristics of the radio source and in particular the energy distribution N(E) (and therefore the spectrum of the emitted radiation) tend to modify with time.

Adiabatic expansion: strong decrease in luminosity but the spectrum is unchanged

Energy loss through radiation: characteristic electron half-life time (time for energy to half)

*2

8* 1064.1

tBE

After a time t* only the particle with E0<E* still survive while those with E0>E* have losttheir energy.

(Special case assuming

p=2) For the spectral index remains constant break

0

For break Single burstGHz~ 23 yrbreak tB

)5.0( 0 Continuous injection

These energy lost affect mainly the large scale structures (e.g. lobes).

Typical spectral index of the lobes = 0.7

Unless there is re-acceleration in some regions of the radio source!

)()(106.1)( 2/12/33* GHzGBMyrt break

Myrt

GHz

Myrt

GBGHz

break

break

50

1

18

108

*

*

Optically thick case: the internal absorption from the electrons needs to be considered the brightness temperature of the source is close to the kinetics temperature of the electrons.

The opacity is larger at lower frequency -> plasma opaque at low frequencies and transparent at high

Self-absorption in the relativistic electron gas

dBvS 2/12/5)(1

GHz)1()(~ 5/15/45/1 zSBpf mm Frequency corresponding to =1

Affects mainly the centralcompact region or very small radio sources

Higher “turnover” frequency smaller size of the emitting region.

Polarization

Characteristic of the synchrotron emission: the radiation is highly polarized.

For an uniform magnetic field, the polarization of an ensemble of electrons is linear, perpendicular to the magnetic field and the

fractional polarization is given by:

7333

(%)

pp

P 0.7- 0.8 for 2<p<4 never!

Typical polarization from few to ~20% Tangled magnetic field

Example of polarization

Polarization between 10 and 20%(some peaks at ~40% around the edge of the lobes)

Example of polarization in radio jets.

Energetics

Magnetic field strength (Bme) and minimum energy density (ume) Corresponding to equipartition of energy between the magnetic field And the relativistic particles in a synchrotron radio source

7/2

222.01.14 )1(1051.1

lS

zBme

Angular size in arcsec, flux in Jy and frequency in GHz l = path lengthMagnetic field in Gauss and minimum energy in erg/cm3

837 2

meme

Bu

Total energy (electrons and magnetic field) can be up to 1060 erg

Radio, optical, UV, X-ray ……

What is produced apart from the collimated radio jets:

UV radiation (likely coming from the accretion disk) that ionizes the gas optical emission lines

X-ray emission (also from the accretion disk)

The synchrotron spectrum can extend to the optical and X-ray wavelength. Life time of the electrons very short, needs re-acceleration

Gas around the AGN: HI, CO, etc. etc.

Centaurus A: example of emission inmany different wavebands