The Origin and Evolution of Cosmic Magnetism:Perspective from SKA

Luigina Feretti – IRA - Bologna MCCT-SKADS School, Medicina, 25–9-07

This topic is one of the 5 Key Science Projects of SKA, selected by the Science Working Group

Motivations:

1. Can address unanswered questions in fundamental (astro)physics

2. Is science which is unique to the radio band and to the SKA

3. Excites the broader community, & is of interest to funding agencies

… and from a phase-space perspective, will almost certainly yield new and unanticipated results!

Outline

- Importance of the study of cosmic magnetism

- Observation of large-scale magnetic fields

- Current ideas on the origin of cosmic magnetic fields

- Studies with SKA and SKA pathfinders

– cloud collapse / star formation

– stellar activity / stellar outflows

– ISM turbulence / gas motions – supernova remnants – stability of galactic disks– acceleration / propagation / confinement of cosmic rays– heating in galaxy clusters – AGNs / Jets

Cosmic Magnetism

Proplyd in Orion MHD turbulence

SN 1006 Merger in gal. cluster

Magnetism is one of theMagnetism is one of theFundamental forces in Fundamental forces in nature. It is crucial in :nature. It is crucial in :

Most bodies in the Universe are magnetized on all scales

Earth: 0.5 G Interplanetary Space: 50 G Sun: 10 G (poles) 1000 G (sunspots)Protostars: 1 mG White dwarfs: 106 G Neutron stars: 1012 G

Milky Way: 5 G (widespread) 1 mG (nucleus)Spiral galaxies: 10 G (average) 30 G (massive arms)Starburst galaxies: 50 G

Radio galaxies: G

Clusters of galaxies: 0.1-1 G

Intergalactic space: < 10-2 – 10-3 G

Large-scale fieldsChallenge to models

Magnetism and Radio Astronomy

Most of what we know about cosmic magnetism derives from radio observations

1 - Synchrotron emission

total intensity field strength

polarization orientation/degree of ordering

2 - Faraday rotation

1 - Synchrotron emission

Total intensity : measures the total field strength

Polarization: gives the orientation and the degree of ordering of field

By writing the synchrotron luminosity as the observed source brightness I0 at the frequency 0, and thesource depth d (to be inferred), applying the K-correction,assuming = 1 (same volume in particles and magnetic field), and expressing the parameters in commonly used units:

7/47/40

7/)412(7/40

7/412min )1()1(1023.1 dIzkxu

umin in erg/cm3

0 in MHz I0 in mJy/arcsec2

d in kpc

Constant computed for = 0.7, 1 = 10 MHz, 2 = 100 GHz

Usually k = 0 or k = 1 assumed for clusters

21

eq u7

24H

/

min

Equipartition magnetic field

BUT see Brunetti et al 1997, Beck and Krause 2005

The synchrotron radiation from a population of relativistic electrons in a uniform magnetic field is linearly polarized, withthe electric vector perpendicular to the magnetic field whichhas generated the synchrotron emission. In the optically thin case, for isotropic electron distribution,and electron power-law energy spectrum:

the degree of intrinsic linear polarization is

N(E)dE = N0E- dE

PolarizationPolarization

8075073

33PInt ..

The above value is reduced in the more realistic cases where

- the magnetic field is not uniform, since regions where the magnetic field has different orientations give radiation with different polarization angle orientations, which tend to average(or cancel) each other.

- there is Faraday rotation effect arising both from instrumental limitations (beamwidth – bandwidth) orwithin the source itself

(Sokoloff et al. 1998, 1999 : how fractional pol. is affected by magnetic field configurations)

Effelsberg 21cm (Reich et al 2003)

Synchrotron Emission from the Milky Way (Perseus - Auriga)

Polarized emissionl=166° l=150°

b=-4°

b=+4°

M51VLA +Effelsberg(Fletcher & Beck 2004)

Clusters of galaxies:

being the largest systems in the Universe, they represent an ideal laboratory to test theories for the origin of extragalactic magnetic fields

Reviews by Carilli & Taylor 2002,Govoni & Feretti 2004

COMA Cluster Beq 0.4 G

500 kpc

RA

DIO

: W

SR

T,

90 c

m (

Fere

tti et

al.

1998)

+Center

Cluster radio halos

Coma

A665

A2163

Cluster radio relics

A548b

0917+75

Abell 2256 I1.4 & B0

Clarke et al. (2004)

Projected magnetic field direction

Polarization degree: large scale order and generally follow the bright filaments

large regions (500 kpc) of fairly uniform magnetic field direction

Results

Filament of galaxies ZwCl 2341.1+0000

(Bagchi et al. 2002)

z 0.3 Size 4 Mpc

320 MHz VLA

Intergalactic Fields:

GRB 000131 at z = 4.5(Bloom et al 2001)

Radio galaxy at z = 5.2 (van Breugel et al 1999)

Upper limits of intergalactic fields from existing studies: BIGM < 10-9…-8 G (model dependent)

Intergalactic Fields (cont.)

2 - Rotation measure

gives an indirect measurement of the strength and structure of the field along the line of sight

Faraday Rotationrotation of the plane of polarization of linearly polarized emission as it passes through a magneto-ionic plasma

-- due to the different phase velocities of the orthogonal circular modes

2

Kronberg 2002

0

Rotation Measure

ne is the electron density in cm-3

L is the path length in kpc

B|| is the line of sight component of the field in G

Sources seen through a magnetized screen:

Infer B along the line of sight in the crossed medium by combining with info about ne from X-rays

Values derived for B are model dependent - analytical solution only for simplest models of the Faraday screen

Otherwise: - numerical techniques (Murgia, Govoni, 2004 - 2005)

- semianalytical approach (Ensslin, Vogt 2004-2005)

Numerical SimulationsPower spectrum analysis

(Ensslin and Vogt 2003

Murgia et al. 2004)

simulate a box with 3D multi-scale fields which have a radial decrease in field strength

resolution = 3 kpc, magnetic structures from 6 to 770 kpc

find n = 1 – 2 provide the best fit to the data: most of the magnetic field energy resides in the small scales

field strength using this approach are a factor ~ 2 lower than the analytical approach assuming smallest RM scale for coherence length

Murgia et al. (2004)

Milky Way

Pulsar RMs + spiral arm field (Han et al 2002)

RMs of 21 polarized sources (Han et al 1998)

M 31

All-sky RM map (Johnston-Hollitt et al 2002RED = POSITIVE RM, BLU = NEGATIVE RM

RM approximate range: -300, +300

Faraday mapping• extended, polarized radio sources can be mapped at several frequencies to produce RM maps

Image courtesy of NRAO/AUI

Cygnus A

cD in a poor cooling-core cluster

A2255 Govoni et al. 2006

Magnetic fields at the G level are ubiquitousin clusters : - coherence scales of 10-100 kpc - large degree of ordering - structure

ORIGIN ?

When and how were the first

magnetic fields generated ?

z 10

z 5

z 0.5

z 0.1

MAGNETIC FIELD

Primordial

Early stars

Protogalaxies

GalaxiesAGN

RECOMBINATION

Primordial Fields: (Olinto 1998, Grasso & Rubinstein 2001)

Created in the exotic ultra-dense stages of the Big Bang

physics poorly known, cannot exclude the creation of a magnetic field of the order 10-30 – 10-25 G

Remember present large scale fields : 10-6 G

Primordial fields would affect the cosmogonic process

anisotropic expansion

effects on nucleosynthesis (larger He abundance)

regulate structure formation

Post-recombination Fields:

1 – Early Stars (z 20)

2 – First AGN (z 5 ?)

3 – Protogalaxies and structure formation (z 5) (Kulsrud et al 1997, Kang et al. 1997)

Seed fields

Seed Fields

(Rees

20

04

)

Injection by galactic winds or active galaxies : Kronberg et al.1999, Völk & Atoyan 1999

Present-day fields of B ≥ 1 Present-day fields of B ≥ 1 μμG could have evolved G could have evolved from from B ~ 10B ~ 10-9-9–10–10-10-10 G G seed fields at z > 5 seed fields at z > 5

Large-scale fields represent a problem because the dynamo amplification time can be large so not many e-foldings at the present epoch

Amplification : dynamo actiondynamo action compressioncompression cluster mergerscluster mergers

Square Kilometer ArraySquare Kilometer Array

•Very powerful in the detection of total intensity and polarized emission and in RM measurements

• SKA: “instant” RMs and position angles:

= 1.4 GHz, = 400 MHz

- for t = 1 hour, 1 = 0.1 μJy

- for P = 1 μJy : RM 5 rad/m-2, 10o !

Adapted from

Gaensler et al. (2001) &

Hopkins et al. (2003)

• Five min observation with SKA at 1.4 GHz

• RMs down to P ~ 3 Jy (Stot ~ 0.1 mJy)

• Approx 500 RMs per deg2 (average separation ~2´-3´)

107 sources over the entire sky, spaced by 90” ( 20000 pulsars)

SKA Faraday Rotation SurveySKA Faraday Rotation Survey

Scientific breakthrough:

- magnetic field of the Galaxy

- magnetic field in nearby galaxies and clusters

- extended sources

Polarization from Fornax A (Fomalont et al 1989)

• Distant galaxies are too small Distant galaxies are too small to be probed by RM grid to be probed by RM grid

… … but can be probed by but can be probed by Faraday rotation and Faraday rotation and depolarization of depolarization of extendedextended background sourcesbackground sources

e.g. NGC 1310 against e.g. NGC 1310 against

Fornax A (Fomalont et al 1989)Fornax A (Fomalont et al 1989)

• Larger distances:Larger distances:

e.g. PKS 1229e.g. PKS 1229––021: absorber 021: absorber at at zz = 0.395 with B ~ 1= 0.395 with B ~ 1– – 4 4 μμGG (Kronberg et al 1992)(Kronberg et al 1992)

→ → powerful probe of powerful probe of evolution of galactic evolution of galactic magnetism as function magnetism as function of redshiftof redshift

Polarization SilhouettesPolarization Silhouettes

NGC 1310

Kronberg et al (1992)

• Large statistical samples can come from Large statistical samples can come from RMs and redshifts of quasarsRMs and redshifts of quasars

(e.g. Welter et al 1984; Oren & Wolfe (e.g. Welter et al 1984; Oren & Wolfe 1995)1995)

- trend of RM vs - trend of RM vs zz probes evolution probes evolution of of BB in Ly- in Ly-αα clouds clouds … … but Galactic contamination, but Galactic contamination, limited statisticslimited statistics

• Quasar RMs with SKA:Quasar RMs with SKA: - ~- ~101066 measurements measurements - identification & redshifts from - identification & redshifts from SDSS & successorsSDSS & successors - accurate foreground removal - accurate foreground removal using RM gridusing RM grid

Ly-Ly-αα Absorbers at Absorbers at zz ~ 1 – 3 ~ 1 – 3

→ magnetic field evolution in galaxies over cosmic time-scales

RRM ~ (1+z)-2

Residual RMs (Galaxy corrected) vs z of QSOs embedded in intervening clouds (Welter et al 1984) : marginal evidence of evolution !

Magnetic Fields in ProtogalaxiesMagnetic Fields in Protogalaxies

– thousands of “normal” spiral galaxies at z ~ 3 detectable with the SKA (1.4 GHz : size = 1 - 3” , flux ≥ 0.2 μJy )

– their radio flux strongly depends on field strength and on star formation rate (and may be polarized)

HDF galaxies with z > 4 (Driver et al 1998)

The Magnetized IGM: Cosmic WebThe Magnetized IGM: Cosmic Web

Existing limits (scale and model dependent): Existing limits (scale and model dependent):

|B|BIGMIGM| < 10| < 10-8-8-10-10-9-9 G G (e.g..Blasi et al 1999; Jedamzik et al 2000)(e.g..Blasi et al 1999; Jedamzik et al 2000)

RM pairs at separation needed to detect B = 1 nG at scale of 50 Mpc (Kolatt 1998)

z = 0.5

z = 1

z = 2

- Detection and polarimetry of very - Detection and polarimetry of very lowlowLevel synchrotron emissionLevel synchrotron emission

-RM measurements of extragalactic RM measurements of extragalactic sources are related to the amplitude sources are related to the amplitude and shape of the magnetic field and shape of the magnetic field power spectrum P(k) where k is the power spectrum P(k) where k is the wave number of the coherence scalewave number of the coherence scale

→ → SKA + z surveys can provide SKA + z surveys can provide magnetic power spectrummagnetic power spectrum of the of the UniverseUniverse

SKA SKA SpecificationsSpecifications for Polarimetry for Polarimetry

• Frequency: at least Frequency: at least 11–10 GHz–10 GHz, 0.3–20 GHz ideal, 0.3–20 GHz ideal• Large field of view: Large field of view: >1 deg>1 deg22 at a resolution of <1 at a resolution of <1"" • High sensitivity: High sensitivity: <0.1 mJy<0.1 mJy, confusion limited, confusion limited• Large bandwidth: Large bandwidth: >400 x 1>400 x 1 MHz MHz at 1.4 GHz at 1.4 GHz• Significant concentration ( Significant concentration ( > 50%> 50% ) of antennae in ) of antennae in

central corecentral core ( ~ ( ~ 5 km)5 km)

• High polarization purity ( High polarization purity ( ––40 dB40 dB at field center, at field center, ––30 dB30 dB at field edges) at field edges)

SKA pathfinders:

ATA (US) LOFAR (The Netherlands + Europe) LWA (US) KAT/MeerKAT (South Africa) MWA (Australia) MIRANDA (Australia + Canada) SKADS (Europe)

Low frequency

- Diffuse synchrotron emission of steep spectrum

- Polarized emission sources of low RM weak magnetic fields

2

= 10o

= 240 MHz, = 32 MHz

RM = 0.4 rad/m2

• Early primordial fields could have been generated by Early primordial fields could have been generated by battery effects, during inflation or phase transitions battery effects, during inflation or phase transitions

• A primordial intergalactic (IGM) field may have A primordial intergalactic (IGM) field may have regulated structure formation in the early Universeregulated structure formation in the early Universe

• ““Seed fields” at z > 5 may originate from primordial Seed fields” at z > 5 may originate from primordial fields or from post-recombination fieldsfields or from post-recombination fields

• Present-day large-scale fields of B ≥ 1 Present-day large-scale fields of B ≥ 1 μμGG could have could have evolved from evolved from BB00 ~ 10 ~ 10-9-9––1010-10-10 G G seed fields at z > 5 seed fields at z > 5

• Evolution from seed fields includes dynamo, Evolution from seed fields includes dynamo, compression, merger interactioncompression, merger interaction

Conclusions

THANK YOU

Biermann Battery effect

Electrostatic equilibrium

When gradients of electronthermodynamic quantities(e.g. density and temperature)are not parallel to thepressure gradient, theelectrostatic equilibrium is no longer possible. This leads to a current which generates A magnetic field restoringthe force balance.

Wid

row

20

02

First observed in the lab in1975 (Stamper & Ripin)

Zeeman effect

In a vacuum, the electronic energy levels of an atomare independent of the direction of its angular momentum.In the presence of magnetic fields, the atomic energy levels are split into a larger number of levels and the spectral lines are also split.

The Zeeman effect can be interpreted as due to the precessionof the orbital angular momentum vector in the magnetic field.The energy shift is proportional to the strength of the magneticfield.

Zeeman splitting in Hydrogen (1.4 GHz): 2.8 Hz G-1

Zeeman splitting in the H2O molecule (22 GHz): 10-3 Hz G-1

Lines are polarized, favouring their detection present detection only for strong magnetic fields (> mG) (sunspots + galactic objects)

Hydrogen

Bohr magneton