Star-forming Galaxies at High Redshift The feedback impact of energetic cosmic rays Ellis Owen Mullard Space Science Laboratory Dept. Space & Climate Physics University College London [email protected] Collaborators Kinwah Wu (MSSL), Xiangyu Jin (McGill), Pooja Surajbali (MPIK), Noriko Kataoka (Kyoto SU), Idunn Jacobsen (MSSL), Suetyi Chan (Nanjing), HST and ALMA image of MACS1149-JD1 (z=9.11) – NASA/ESA, Hashimoto+ 2018 Brian Yu (MSSL) IPMU, Tokyo – July 2019 Outline • Star-forming galaxies in the Universe • Cosmic rays in star-forming galaxies • Particle propagation • Energy deposition and heating by particles and radiation • Implications for star-formation and feedback • Summary 2 Star-forming Galaxies in the Universe Image of simulated Lyman-alpha emission around a high redshift group of protogalaxies – credit: Geach et al. Local Starbursts NGC 253 Arp 220 M 82 NASA/ESA 2008 ESO 2010 NASA/ESA 2006 1 1 1 SF 10 M yr− 220 M yr− 10 M yr− R ⇠ ⇠ ⇠ 1 1 1 0.1yr− 4yr− 0.1yr− RSN 4 See e.g. SN Rates: Fenech+2010, Lenc & Tingay 2006, Lonsdale+2006; SF Rate: Varenius+2016, Schreiber+2003, Bolatto+2013 High-redshift starbursts (z~6+) • Low mass, high SF rates $ – 10 M⨀ ,- – 10s − 100s M⨀ yr – SF efficiencies ~ tens of % MACS1149-JD1 EGSY8p7 • Some known to host Lyman-⍺ haloes – Multi-phase CGM… • Simulation work suggests possibility of filamentary GN-z11 EGS-zs8-1 inflows of gas (cf. works by Keres, Dekel, Birnboim…) MACS1149-JD1 (HST/ALMA) – NASA/ESA, Hashimoto+ 2018 EGSY8p7 (Hubble/Spitzer) – NASA, Labbe+ 2015 GN-z11 (HST) – NASA, Oesch+ 2015 EGSY-zs8-1(Hubble/Spitzer) – NASA/ESA, Oesch & Momcheva 2015 5 The high-redshift CGM environment Outflows ISM Cold inflows (operate mainly at high-redshift) Birnboim + Dekel 2003; Keres+2005; Dekel 2009 Figure based on Tumlinson+2017 6 Cosmic rays in star-forming galaxies Image credit: Crab Nebula, NASA, ESA 2005 Cosmic rays in the Milky Way Knee Knee Ankle Adapted from Gaisser 2007 8 Starbursts as cosmic ray factories 1015 Neutron stars • Hillas criterion 1010 GRBs Emax qBR 5 AGN /G 10 ! AGN jets • Cosmic rays sources White 100 dwarfs – Galactic (internal) in orange ISM hot spots IGM shocks 10-5 Supernova remnants 10-10 105 1010 1015 1020 1025 rsys/cm Fig. adapted from Owen 2019 (PhD thesis) See also Kotera & Olinto 2011; Hillas 1984 9 Cosmic ray acceleration • Shocks, e.g. SNR – First order Fermi acceleration • Each pass through the shock increases the energy • After n crossings, energy is E = E ⇠ n 0h i Downstream Upstream Magnetic shock 10 Cosmic ray acceleration • Some particles will escape after a n crossing – take P as probability of remain, N = N0P so number remaining after each crossing log P log ⇠ N E h i • Eliminate n and rearrange = N E 0 ✓ 0 ◆ Γ • Result is CRs accelerated to high E − energies, ~GeV and above, following a = E power-law ✓ 0 ◆ Γ 2.1 ⇡ 11 Cosmic ray interactions with radiation fields (p�) Interaction by particles scattering off ambient photons (starlight, CMB…) Photopion Interaction 0 + p + ⇡ p +2γ + pion multiplicities at p + γ ∆ ! higher energies ! ! n + ⇡+ n + µ+ + ⌫ ( ! µ + n + e + ⌫e +¯⌫µ + ⌫µ Photopair Interaction + p + γ p + e + e− . ! 12 Cosmic ray interactions with matter (pp) p + p + ⇡0 p + ∆+ p + p + ⇡+ 8 ! 8 p + p > <>p + n + ⇡+ + pion multiplicities ! > at higher energies > + < ++ >n + p + ⇡ n + ∆ : + ! (n + n + 2(⇡ ) > > > : Pions decay to photons, muons, Neutron and photon neutrinos, electrons, positrons, interactions produce pions antineutrinos n + γ ⇡’s ! ⇡ γ,µ,e,⌫ ... ! 13 Cosmic ray interactions Interactions with stellar radiation fields CMB & cosmological losses Interactions with typical ISM density fields Adapted from Owen+ 2018 (1808.07837) 14 Particle propagation Image credit: M25 Motorway, Carillon UK Transport The transport equation (hadrons) • The transport equation for protons (cooling/momentum diffusion assumed negligible) @n @ = [D(E,x) n]+ [b(E,r)n] [vn]+Q(E,x) S(E,x) @t r · r @E r· − 1015 Neutron stars 1010 GRBs 5 AGN /G 10 ! – Diffusion dominated, i.e. � = 0 AGN jets White 100 dwarfs ISM hot spots IGM shocks 10-5 Supernova remnants – Advection dominated, i.e. � �, � = 0 10-10 105 1010 1015 1020 1025 rsys/cm 16 The diffusion zone (‘stationary’ ISM) @n @ = [D(E,x) n]+ [b(E,r)n] [vn]+Q(E,x) S(E,x) @t r · r @E r· − M82 in H⍺ (WIYN) and optical (HST) Smith+ 2005 17 The diffusion zone (‘stationary’ ISM) @n @ = [D(E,x) n]+ [b(E,r)n] [vn]+Q(E,x) S(E,x) @t r · r @E r· − • Approximate ISM as a sphere • Absorption depends on – density of CRs – density of ISM gas – interaction cross section (dominated by pp process) • CR injection as a BC (for now) – restate problem as individual linearly independent events (t’ since inj. event) t0 2 n0 r0 n = @n 3/21exp@ 2 cdtσˆ@pn⇡ nISM exp [4⇡D(E,r )t =] r−D0(E,r) cnσˆp⇡ nISM−4D(E,r0)t0 @0 t 0 r2 @r ( Z @r − ) ⇢ 18 The diffusion zone (‘stationary’ ISM) • Solution for steady-state from continuous injection from single source (time integral) • Numerically convolve with a source ensemble (MC distribution) – Follows ISM gas profile x x x x x x x x x x x x x x x x x x x x x x x 19 The diffusion zone (‘stationary’ ISM) • Solution for steady-state from continuous injection from single source (time integral) • Numerically convolve with a source ensemble (MC distribution) – Follows ISM gas profile 21 10− 1 22 − 10− eV 3 23 10− − cm 24 · 10− erg 25 10− / V N d d · 26 E 10− E d 27 10− 1 0 1 2 10− 10 10 10 r/kpc 20 Secondary electrons • Injection by the pp attenuation process 0 + pp ⇡ , ⇡ , ⇡− ! ⇡− e− +¯⌫ + ⌫ +¯⌫ ! e µ µ ⇡+ e+ + ⌫ +¯⌫ + ⌫ ! e µ µ • Transport equation (electrons) @n @ e = [D(E,x) n ]+ [b(E,r)n ] [vn ]+Q (E,x) S (E,x) @t r · r e @E e r· e e − e – Negligible advection – No absorption – Diffusion coefficient same as for protons (depends on charge) 21 Cosmic ray distribution in ‘stationary’ ISM • Electrons calculated for each proton injection diffusion profile (MC distribution) • Then convolved across proton source distribution (i.e. the SN events) 21 10− 1 22 − 10− eV 3 23 10− − cm 24 · 10− erg 25 10− / V N d d · 26 E 10− E d 27 10− 1 0 1 2 10− 10 10 10 r/kpc 22 The advection zone (outflows) @n @ = [D(E,x) n]+ [b(E,r)n] [vn]+Q(E,x) S(E,x) @t r · r @E r· − M82 in H⍺ (WIYN) and optical (HST) Smith+ 2005 23 The advection zone (outflows) Hydrodynamical outflow model 300 250 1 200 − s · Outflow velocity, v 150 Circular velocity, Vc km v/ 100 50 0 2 4 6 8 10 h/kpc Owen+ 2019a; Samui+ 2010 24 The advection zone (outflows) Particle propagation @n • Steady state solutions to transport equation, i.e. =0 @t 1 102 − sr 2 1 10− − s 2 10 6 − − cm 10 10− 1 − 14 10− GeV / 18 10− 50 kpc ⌦ d Φ d E 22 d 10− 100 101 102 103 104 105 106 Owen+ 2019a EEp//GeVGeV 25 Energy deposition and heating Image credit: National Bunsen Burner Day (March 31st), McGill University 2016 Heating mechanisms & global impacts • Direct Coulomb (DC) “quenching” Also operates in magnetized CGM via CR streaming instability + Alfvén wave dissipation (similar rate) • Indirect Inverse Compton X-rays (IX) “strangulation” Also advected CRs may also slow inflows, or heat filaments in CGM by DC, Alfvén dissipation… Figure based on Tumlinson+ 2017 27 Direct Coulomb heating Zone comparisons Diffusion Zone B Zone A 26 10− CRs (Thermal) X-rays Stellar radiation 1 − 29 Free-streaming CRs 10− sr 1 Zone B Zone B − s 3 10 32 Outflow schematic − − cm · 35 10− erg / ⌦ d 38 10− t E d d V d 41 Advection Zone A 10− CRs (extended injection) Extended injection region 44 10− 1 0 1 10− 10 10 h/kpc Adapted from Owen+ 2019a (1901.01411) 28 Indirect heating ISM reference case 26 10− 1 − s 29 3 10− − z = 12; t = 370 Myr cm 32 10− · 35 erg 10− / t d E V d 38 d 10− Now z = 0; t = 13.7 Gyr 41 10− 1 0 1 2 10− 10 10 10 r/kpc CGM minimum shown here: value in inflow filaments would be higher (proportional to density) Owen+ 2019b (1905.00338) 29 CGM structure and implications Strangulation vs. quenching • Scaling relation to deal with filaments Protogalaxy • Effective heating power ~1-2 orders of H II region Interstellar medium magnitude lower than ISM value A B Circum-galactic medium Cold inflowing filaments 26 10− 1 − s 29 H II region 3 10− − z = 12; t = 370 Myr cm 32 10− · 35 erg 10− / t Owen+ 2019b (1905.00338) d E V d 38 d 10− Now z = 0; t = 13.7 Gyr 41 10− 1 0 1 2 10− 10 10 10 r/kpc 30 Timescales • Estimate by considering condition for them Time for ISM to exceed Tvir to no longer be gravitationally bound – due to DC heating upper-limit ⌧Q = ⌧mag + ⌧DC ⌧S = ⌧mag + ⌧IX Time for filament region to exceed Tvir Magnetic containment time; due to IX heating required for CR effects to develop 1 ⌧mag SFR− Application to real galaxies… / 31 Feedback and star-formation Image credit: HST image of N90 Star forming region in SMC, NASA/ESA 2007 Inferred behavior of MACS1149-JD1 • Spectroscopic z=9.11 (t = 550 Myr) +0.8 1 SF 4.2 1.1 M yr− R ⇡ − • Two populations of stars – One from observed SF activity – Other from activity ~100 Myr earlier • Earlier burst of Star-formation at z=15.4; t=260 Myr (Hashimoto+2018) • Quenched fairly quickly – distinct inferred age of older stellar HST and ALMA image of MACS1149-JD1 (z=9.11) – NASA/ESA, Hashimoto+ 2018 population – SED, size of Balmer break Can CRs account for the rapid 100 Myr quenching after initial burst? 33 Can CRs do the job? • Hashimoto+ 2018 star-formation burst models (intense, medium, slow).