Transcript Document

Radio Galaxies
Part 1
Raffaella Morganti
ASTRON
The plan
1.
What are AGNs and radio galaxies - How to find them
A prototype radio galaxy - Emission mechanisms
2.
Morphology of the radio emission:
various morphologies, nuclear regions,
highly collimated jets - hot-spots - radio lobes
3. Gas in radio galaxies.
Neutral hydrogen: looking at some data and what can we derive
4.
A look at the host galaxy
Ionized gas: from pc to tens of kpc
5. Origin and evolution of radio galaxies
Unified Schemes for AGNs: trying to make sense of the different types
What are the Active Galactic Nuclei
Difficult to have an unique definition.
huge amount of energy (up to 104 times more
than a normal galaxy)
emitted from a tiny region (<1 pc3)
An AGN can have a luminosity ranging from 1042 to 1044 erg/sec
wide energy range covered wide range of different objects!
now even lower power AGN are found
The high energy released by an AGN is believed to originate from a
supermassive black-hole (106 to 109 Msun in <<1pc)
but the presence of a supermassive BH is not enough!
Some characteristics
Optical emission lines
for different AGNs
Luminosity is not the only criteria:
 continuum emission (often quite blue)
across ~13 order of magnitude in frequency
 emission lines
 radio emission (in about 10% of AGNs)
optical UV
radio
X-ray
Comparison of the
continuum emission
from a Seyfert galaxy
and a normal galaxy
Observables to classify an object as AGN
Very small
angular size
Wavelength/resolution dependent
Radio VLBI, HST cores
High luminosity
Broad-band
continuum
Composite spectrum for AGNs 
Strong emission
lines
Variability
Polarization
Radio emission
Powerful way to detect AGNs,
BUT only a minority are strong radio
sources and the radio accounts for at
most 1% of the total energy output
There are many different types of AGNs,
depending which characteristic is dominant
 radio loud vs radio quiet
 strong optical lines  narrow vs broad
 core dominance
 weird & extremely variable objects
Different types of AGNs: a summary
Optical Emission Line Properties
Radio quiet
Type2
Type 1
narrow line
broad line
Seyfert 2
Seyfert 1
?
Quasars
Broad
absorption
line QSO
BL Lac
BL Lac
Radio loud
low power
Narrow-line
Blazar, OVV
and many
other weird
objects
radio galaxies
high power
broad-line RG
lobe/core
dominated QSR
Decreasing angle to line of sight
Type 0
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
 Amazing discovery when they were identified with
extragalactic, i.e. far away, objects
Unexpectedly high amount of energy involved!
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)
Why are interesting?
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!
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
PKS
2Jy
~3Jy
(1 Jy= 10-26 W m-2 Hz-1)
178 MHz
Cambridge
(+5,6,7C)
408 MHz
Parkes
Molonglo
B2
NRAO
PKS
NVSS
FIRST
WENSS
0.25
408 MHz
Bologna (+B3)
0.8Jy
1.4-5GHz
NRAO
0.7Jy
2.7 GHz
Parkes
2.5 mJy (45” res.)
1.4 GHz
NRAO VLA Sky
Survey
1mJy (~5” res)
1.4 GHz
Faint Images Radio Sky
at Twenty centimeters
85 mJy
300 MHz
WSRT
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)
P  4 D2 F (W Hz 1 )
or integrated over a typical (radio) range of frequencies (107 to 1011 Hz)
 Resolution important for the identification
(radio surveys find not only radio galaxies!)
resolution /D
21 cm, D = 64 m
11 arcmin
21 cm, D= 3km
21 cm, D= 3000 km
14 arcsec
1 mas
 Difference in power limit for the different surveys
Radio power: source of 2 Jy flux (@ 1.4 GHz),
source of 0.2 Jy flux, z = 0.2
source of 10 Jy flux, z = 0.2
z = 0.2
log P = 26.5 W/Hz
log P = 25.5 W/Hz
log P = 21.2 W/Hz
Confusion
‘Confusion’ can be resolved by imaging at higher spatial
resolution with large interferometers
(WSRT, VLA or ATCA)
ATCA image,
July 2001
NGC 6580
(S0)
IC 4933
(Sbc)
HIPASS
beam
Optical identifications
radio much larger than optical
NVSS
resolution ~45 arcsec ~ 45 kpc
(1 arcsec ~ 1 kpc at z = 0.04)
Going deeper and deeper
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 power
high radio resolution is required to find a very compact core
(to distinguish non-thermal emission from thermal emission)
Deep Wide-Field VLBI Surveys
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 detections at full
sensitivity with an r.m.s. noise
of 9 microJy/beam.
VLBI nondetection 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)
A prototypical radio galaxy
Lobes
Hot-spots
Core
Jets
 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
How a radio galaxy works
to hot-spots and/or lobes
torus (supposed to hide – for some orientation –
the very central regions)
Zoom-in of
the central regions
Supermassive
Black Hole
accretion disk
(UV, Xray)
A prototypical radio galaxy
bowshock
splash-point
backflow
Observable Diagnostic
Constituents
Derived Properties
Radio continuum
Relativistic plasma
Energetic, Pressure, Jet propagation
velocity, Internal magnetic field
Ages, Faraday rotation, Magnetic
fields
Thermal plasma
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 IR
Continuum
Stars
Mass, Age,
Star-formation rate
Polarization properties
Scattered AGN light
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 plasma
Hot gas (10^7 K)
Jet (and hot-spots) properties
Cluster properties
Relativistic electrons in a magnetic field
E   me c 2
 
1
v2
1 2
c
>>1
 For one electron, max frequency   2
for slightly different  covers the entire spectrum
 Electron energy distribution is a power law:
N (E )dE  kE pdE
 Assuming the emission from each can be added up (optically thin case)
P (, )  B ( p 1) / 2  ( p 1) / 2
 / 2
........

 
/
1

The radio spectrum is therefore a power law:
S   
Typical ~0.8
   (p  1) /2
p~2.6
Deviations from a constant spectral index
1. Energy loss
2. Self-absorption in the relativistic electrons gas
3. Absorption from ionized gas between us
and the source (free-free absorption)  torus!
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)
E *
1.64 108
B 2t *
(Special case assuming p=2)
After a time t* only the
particle with E0<E* still
survive while those with
E0>E* have lost
their energy.
For  break the spectral index remains constant   0
break ~ B 3 tyr2 GHz Single burst
For  break
  ( 0  0.5)
Continuous injection
 These energy lost affect
mainly the large scale
structures (e.g. lobes).
 Typical spectral index of the
lobes   = 0.7
1 / 2
t* (Myr) 1.6103 B 3/ 2 (G )break
(GHz )
break  8GHz B  10 G
t*  18 Myr
break  1GHz
t*  50 Myr
 Unless there is re-acceleration in some regions of the radio source!
Self-absorption in the relativistic electron gas
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
  1
S ( ) v 5 / 2 B 1 / 2d
Frequency corresponding to =1
m ~f ( p ) B 1 / 5 Sm  4 / 5 (1  z )1 / 5 GHz
Affects mainly the central
compact 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:
3p  3
P (%) 
3p  7
Typical polarization from few to
0.7- 0.8 for 2<p<4
never!
~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.
Faraday rotation
Travel through a plasma+magnetic field (that can be internal or external
to the source) changes the polarization angle
  2.6  1017 2  ne B dl



Rotation measure (RM)
Ne = electron density of the plasma
dl = depth
B = component of the magnetic field
parallel to line of sight
RM can be derived via observations
at different wavelengths
 If the medium is in front of the radio
source: no change in the fractional polarization
 If the medium is mix in the radio source:
depolarization
dependence on
wavelength (if due to Faraday rotation)
thermal electrons with density ~ 10-5 cm-3
Depolarization happens also if the magnetic field is tangled on the scale of the
beam of the observations
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
Bme
S 
 1.51104 (1  z )1.1  0.22  2 
 l 
2/ 7
ume
2
7 Bme

3 8
Angular size in arcsec, flux in Jy and frequency in GHz
l = path length
Magnetic field in Gauss and minimum energy in erg/cm3
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 in
many different wavebands