Cosmic rays in early Star-Forming Galaxies and their effects on the Interstellar Medium

Ellis Owenellis.owen.12@ucl.ac.ukMullard Space Science Laboratory, University College London, United Kingdom

National Tsing Hua University, Taiwan (ROC)

Collaborators: Kinwah Wu (UCL-MSSL, UK)

Idunn Jacobsen (UCL-MSSL, UK)

Pooja Surajbali (MPIK, Heidelberg, Germany)

International Cosmic Ray Conference, Busan, Korea, July 2017

EQ J100054+023435 – Multiwavelength image with HST, Spitzer, Chandra, Keck, Galex, CFHT, Subaru, UKIRT, JCMT, VLA & IRAM. Credit NASA (2008)

Outline

• Early Star-Forming Galaxies• Propagation and Interaction of Cosmic Rays

– Direct– Indirect

• Energy Deposition and Cosmic Ray Heating• Remarks

2

Starburst Galaxies at High Redshift

• Starburst galaxies characterized by high star formation rates (SFR)

3

> 10 M� yr�1à many Supernovae à abundant cosmic rays

Starburst Galaxies at High Redshift

• Starburst galaxies characterized by high star formation rates (SFR)

• Why are high redshifts of interest?– Galaxies with very high SFRs seem to be abundant at high redshifts– Possible implications on cosmic reionization (Sazonov & Sunyaev 2015)

3

> 10 M� yr�1à many Supernovae à abundant cosmic rays

Starburst Galaxies at High Redshift

• Starburst galaxies characterized by high star formation rates (SFR)

• Why are high redshifts of interest?– Galaxies with very high SFRs seem to be abundant at high redshifts– Possible implications on cosmic reionization (Sazonov & Sunyaev 2015)

• Parametric model protogalaxy, very active to demonstrate concept• SFR = , environment defined by

3

> 10 M� yr�1

Density field Radiation field Magnetic field

à many Supernovae à abundant cosmic rays

1000 M� yr�1

Energy Transport by Cosmic Rays

• Cosmic rays may be influenced by magnetic fields– Low & Intermediate energies– Larmor radius

• Can hamper their propagation into intergalactic space– Containment vs. Diffusion

4

@n

@t= r · [D(E, r, t)rn] +Q(r, E)

• As a first estimate, assume Bohm diffusion ~1 scattering per gyro-radius

D =1

3c rL ' c rL

Cosmic Ray Diffusion & Containment

• Strong containment• Steady-state solution

with cosmic ray densities

• Around ~1012 times high than free-streaming case

5

10�2 10�1 100 101 102

r/kpc

10�38

10�35

10�32

10�29

10�26

10�23

10�20

10�17

dN

dE

dV/e

rg·c

m�

3 eV

�1

Free-streaming profile

Saturated magnetic field, steady-state profile

Cosmic Ray Interactions (Direct)

6

+ pion multiplicities at higher energies

p+ � ! p+ e+ + e� .

p+ � ! �+ !(p+ ⇡0 ! p+ 2�

n+ ⇡+ ! n+ µ+ + ⌫µ

n+ e+ + ⌫e + ⌫̄µ + ⌫µ

Photopion Interaction

Interactions with Radiation Fields (p𝛄)

Photopair Interaction

Interaction by particles scattering off ambient photons (starlight, CMB…)

p+ p !

8>>>>>><

>>>>>>:

p+�+ !

8><

>:

p+ p+ ⇡0

p+ p+ ⇡+

p+ n+ ⇡+

n+�++ !(n+ p+ ⇡+

n+ n+ 2(⇡+)

Cosmic Ray Interactions (Direct)

7

+ pion multiplicities at higher energies

Neutron and photon interactions produce pions

Pions decay to photons, muons, neutrinos, electrons, positrons,

antineutrinosn+ � ! ⇡’s ⇡ ! �, µ, e, ⌫ . . .

Interactions with Matter (pp)

p+ p !

8>>>>>><

>>>>>>:

p+�+ !

8><

>:

p+ p+ ⇡0

p+ p+ ⇡+

p+ n+ ⇡+

n+�++ !(n+ p+ ⇡+

n+ n+ 2(⇡+)

Cosmic Ray Interactions (Direct)

7

+ pion multiplicities at higher energies

Neutron and photon interactions produce pions

Pions decay to photons, muons, neutrinos, electrons, positrons,

antineutrinosn+ � ! ⇡’s ⇡ ! �, µ, e, ⌫ . . .

Interactions with Matter (pp)Dominates

Cosmic Ray Interactions (Indirect)

8

Electron Injection

Qe(�e) '⌥

6

400 me

mpQp(�p)

(Schober+ 2015, Lacki & Beck 2013)

Injection profile can be estimated from the CR source term

Cosmic Ray Interactions (Indirect)

8

Electron Injection

Qe(�e) '⌥

6

400 me

mpQp(�p)

(Schober+ 2015, Lacki & Beck 2013)

Sunyaev-Zel’dovich (SZ) Effect X-Ray Emission

Injection profile can be estimated from the CR source term

Inverse-Compton scattering off CMB CMB Photon

X-Ray Photon

Energetic electronsLSZ ⇡ 1048erg s�1

(upper limit)

Energy Deposition

• Absorption coefficient: energy absorbed at a point• Cross section depends on interaction (radiation/particles)

• In general, can account for attenuation from emission up to absorption point by RT

• Then heating is ~ energy absorbed at a point after attenuation

• Cross section: Klein-Nishina (X-rays)… Thomson limit with UV9

H(r) = F0 ↵(r) exp

✓�Z r

re

↵(r0) dr0◆

I⌫(r) = I⌫,0 exp

✓�Z r

r0

n(r0)�⌫dr0◆

↵(r) = n(r)�

Radiation

Energy Deposition

• Absorption coefficient: energy absorbed at a point• Cross section depends on interaction (radiation/particles)

10

↵(r) = n(r)�

Cosmic Rays

SZ X-rays

– Cross section is dominating ppinteraction

– Scale to account for the containment

– Emission profile from CR electron secondary injection

– Heating then as per conventional treatment (previous slide)

No B field

With B field

Energy Deposition

11

Stellar heating around 10-22 erg cm-3 s-1

X-ray heating

Cosmic Rays (saturated B field)

Note – Cosmic ray MC calculation using 1000 points

CR Heating: ISM

Summary & Remarks

• Cosmic rays are abundant in star forming galaxies– Of particular interest at high redshift

• Containment by magnetic field appears to be important global effect– Focuses CR heating into ISM above conventional stellar heating

• Accompanied by an X-ray heating effect due to SZ effect– Higher than direct CR heating outside the galaxy

• Impacts – Subsequent star formation (e.g. by heating star forming regions)– Thermal properties of surroundings– Pre-heating IGM during reionization

12

Backup: Cosmic Ray Sources/Hillas Criterion

• Cosmic rays: charged energetic particles (assume protons)

• Sources: supernova remnants (SNRs) can accelerate CRs up to 1017-18 eV

• Diffusive shock acceleration

13

Emax

qBR

Adapted from Jacobsen+2015

Backup: Star-Forming Galaxies at High-z

14

Magnetic Field

10�4 10�3 10�2 10�1 100 101

Age of Galaxy/Myr

10�31

10�27

10�23

10�19

10�15

10�11

10�7

10�3

Mag

neti

cF

ield

Stre

ngth

/G

• Two scale components:– Local scale, ~10-3 pc– Galactic ordered field ~1kpc

• SN driven

• Initial B field ~10-20 G permeates protogalaxy (Sigl+1997; Howard & Kulsrud 1997)

• Turbulent dynamo drives B field up to µG levels seen in current epoch (Schober+2013) Model follows J. Schober + 2013

SNe à Turbulence à B field

Backup: Cosmic Ray Interactions

15

1010 1012 1014 1016 1018 1020

Energy/eV

10�4

10�3

10�2

10�1

100

101

102

103

104

105

E↵e

ctiv

ePat

hLen

gth/

Mpc

1

2

3

4

5

6

7

8

Particle Path Lengths

CMB & cosmological

losses

Interactions with stellar radiation

fields

Interactions with density fields

Backup: Cosmic Ray Diffusion

• Fundamental diffusion solution (Gaussian)

• Principle of superposition

16

Z tmax

0n(t)dt

n1(t)

n2(t) = n1(t+ dt)

n3(t) = n2(t+ dt) = n1(t+ 2dt)

ni(t) = · · · = n1(t+ (i� 1)dt)

⌃i{Time (to deal with continuous injection)

n(r, t) =Q(rs)

(4⇡Dt)3/2exp

⇢� (r � rs)2

4Dt

�

Backup: Cosmic Ray Diffusion

• Fundamental diffusion solution (Gaussian)

• Principle of superposition

17

x

x

xx

x x

xx

xx

x x

xxx

x

x

xx

x

x

x

x

Space (to deal with source distribution –weighted by galaxy density profile)

n(r, t) =Q(rs)

(4⇡Dt)3/2exp

⇢� (r � rs)2

4Dt

�

Backup: Energy Deposition

18

��

RayEmission/ergcm

�3s

�1

Heating: Cross-Check with GALPROP

Stellar heating

X-ray heating

Cosmic Rays (saturated B field)

Cosmic Ray Heating Galprop Comparison

Cosmic Rays (initial)