Testing the origin of the UHECRs with neutrinos

Testing the origin of the UHECRs with neutrinos

Walter WinterDESY, Zeuthen, GermanyKavli Institute for Theoretical Physics (KITP), Santa Barbara, CA, USA

UHECR 2014,Springdale, UT, USAOct. 12-15, 2014

Walter Winter | UHECR 2014 | Oct. 13-15, 2014 |

Contents

> Introduction

>Can the observed neutrinos come from the same sources as the UHECRs?

>GRBs as test case for the UHECR-neutrino connection

> Summary


3

Cosmic messengers

Physics of astrophysical neutrino sources = physics of

cosmic ray sources


2014: 37 neutrinos in the TeV-PeV range

Science 342 (2013) 1242856; update by Gary Hill @ Neutrino 2014

Where do these come from?

Prompt atmospherics?Directional information: Clustering?

Isotropic/from Galactic plane/Galactic center?Why no events > few PeV?

Can these come from the sources of the ultra-highenergy cosmic rays?

Which source class? More than one? Flavor composition?

Requires more statistics


Connection with primary nuclei?

> In pp and pg interactions, the secondary pions take about 20% of the proton energy, the neutrinos about 5% (per flavor)

> PeV neutrinos must come from 20-500 PeV nuclei (depending on comp.)

>Observed cosmic ray composition non-trivial function of energy (at Earth!)

> Simple example: Neutrinos fromcosmic rayinteractions with hydrogenin the Milky Way[O(0.1-1) event]

>Connection with UHECR sources requires extrapolation over several orders of magnitude both in spectrum and composition

Gaisser, Stanev, Tilav, 2013

UHECRs

nprima-

ries

Joshi, Winter, Gupta, MNRAS, 2014


Fitting the observed neutrino spectrum

> Simplest possible model: Ap (or AA) interactions in sources;SFR evolution

> Possible fits to data:

WW, arXiv:1407.7536(PRD, accepted)

Protons a=2 B ~ 104 G(magnetic field effects on sec. pions, muons, kaons)

Nucleia=2, Emax=1010.1 GeVComposition at source

with b=0.4

Protonsa=2 Emax=107.5 GeV

Protons, a=2.5[Problem: Fermi diffuse g-ray bound Murase, Ahlers, Lacki, PRD 2013]


Connection to UHECRs?

WW, arXiv:1407.7536(PRD, accepted)

Protons a=2 B ~ 104 G

Nucleia=2, Emax=1010.1 GeVComposition at source

with b=0.4

Protonsa=2 Emax=107.5 GeV

Protons, a=2.5[Problem: Fermi diffuse g-ray bound Murase, Ahlers, Lacki, PRD 2013]

Yes, but: Energy input per decade very different in neutrino-relevant and UHECR energy ranges(Energetics seem to favor a~2, see e.g. B. Katz, E. Waxman, T. Thompson, and A. Loeb (2013), 1311.0287) will come up again later!

Yes, but: Synchrotron losses limit maximal proton energies as well. Need large Doppler factors (e. g. GRBs)

Yes, but: Need energy-dependent escape timescale leading to break/cutoff within source (diff. from ejection!)see e.g. Liu et al, PRD, 2004; arXiv:1310.1263

Yes, but: A(E) change somewhat too shallow to match observation; difference source-observation from propagation?


> Idea: Use timing and directional information to suppress atm. BGs

> Stacking limit exceeds observed neutrino flux (~10-8) by one order of magnitude; interesting to test specific modelsNature 484 (2012) 351

> Prediction (One zone model.based on fixed collision radius models) almost reached(some recent corrections!)

GRBs as a test case

(Source: NASA)

GRB gamma-ray observations(e.g. Fermi, Swift, etc)

(Sou

rce: IceCu

be)

Neutrino observations

(e.g. IceCube, …)Coincidence!

(Hümmer, Baerwald, Winter,

PRL 108 (2012) 231101;

method based on Guetta et al, 2004; Waxman, Bahcall

1997)


GRB - Internal shock model

(Source: SWIFT)

Prompt phaseCollision of

shells Shocks

Particle acc.

“Isotropic equivalent

energy“

~ 200-1000G

(Simulation by M. Bustamante)

Observable:Light curves

Engine(intermittent)


UHECR-neutrino connection: escape mechanisms?Baerwald, Bustamante, Winter, Astrophys. J. 768 (2013) 186

Optically thin(to neutron escape)

Optically thick(to neutron escape)

Direct proton escape(UHECR leakage)

nn

p

p

n

n

n

n

n

n

n

n

n

n

n

n

np

p

l‘ ~ c t‘pg l‘ ~ R‘L

n

n

p

p

p

p

p

One neutrino per cosmic ray

Protons magnetically confined

Neutron escape limited to edge of shells

Neutrino prod. relatively enhanced

pg interaction rate relatively low

Protons leaking from edges dominate


An example (before propagation)

For high enough acceleration efficiencies:R‘L can reach shell thickness at highest energies(if E‘p,max determined by t‘dyn)

Hard spectrum, aka “high pass filter“ (Globus et al, 2014)

Relative importance depends on optical thickness to pg interactions

(from: Baerwald, Bustamante, Winter, Astrophys. J. 768 (2013) 186)

Neutron spectrumharder than E-2

proton spectrum

(only adiabatic energy losses)


Combined source-propagation model: Ankle transition (ap=2, fit range 1010 ... 1012 GeV)

>Neutron-dominated cases can be constrained by neutrino emission

> Baryonic loading fe-1 (energy protons to photons) typically somewhat larger

than IceCube assume, to fit UHECR data (here Liso=1052 erg s-1, Eiso=3 1052 erg)

(Baerwald, Bustamante, Winter, Astropart. Phys. 62 (2015) 66; figures with TA data)

G=300

G=800


Combined source-propagation model: Dip transition (ap=2.5 with SFR evolution, fit range 109 ... 1012 GeV)

>Neutron-dominated cases even more extreme

>Required baryonic loading fe-1 extremely large; implication of unequal

energy output per decade (bolometric correction)


G=300

G=600 1050.5 erg/s


Parameter space constraints (ankle model, fit to TA data)

Example: Moderate acc. efficiency, escape by Bohm-like diffusion, SFR evolution of sources,ankle transition


Best-fit (shaded contours: TA UHECR fit)

Current IceCube limit

IceCube expectation (15yr)

log10 fe-1

(baryonic loading) obtained from fit

Optically thick pg

Direct escape

… but - maybe - assigning one parameter set to all shells is too simple?


The future: more dynamical collision models

> Set out a number of shells with a Lorentz factor distribution

> Shells collide, merge and cool by radiation of energy

> Light curve predictable (see below)

> Efficient energy dissipation (e. g. into gamma-rays) requires broad Lorentz factor distribution

(Bustamante, Baerwald, Murase, Winter, 2014; based on collision model Kobayashi, Piran, Sari, 1997; see Globus et al, 2014 for a similar approach)


Consequences for different messengers

>Collision radii reach from below photosphere to circumburst medium

>UHECR escape as neutrons (red) and directly (blue) at intermediate radii

> Energy output ~ no of collsions x energy per collision (counting important!)

> The burst looks different in different messengers!

(Bustamante, Baerwald, Murase, Winter, 2014)


Consequences for neutrino production

>Neutrino flux comes from a few collisions at photosphere

> Photospheric radius and photohadronic interactions both depend on particle densities (scale at same way)

> Super-photospheric (minimal?) prediction hardly depends on baryonic loading, G (different from earlier works!)

> Testable in high-energy extension of IceCube?

> Sub-photospheric contribution could be much larger. However: photons from below photosphere not observable

(Bustamante, Baerwald, Murase, Winter, 2014)

Eiso=1053 erg per GRB


Summary

>Neutrino observations open new window to cosmic ray source identification; data (discovery and constraints) become meaningful

>UHECR connection somewhat more challenging, as several orders of magnitude in energy between UHECRs and primaries leading to observed neutrino flux

>GRBs are an interesting test case, as The constraints are strongest on GRBs because of timing cuts

Well-motivated models for gamma-ray emission exist

IceCube data already test the parameter space

>Different messengers are produced in different regions of a GRB. Multi-messenger connections are more model-dependent than previously anticipated

>Heavy nuclei are anticipated to escape from larger radii than protons, as disintegration is to be avoided – but they can survive


BACKUP


Opticallythin

to neutrons

Neutrino production

from: Baerwald, Hümmer, Winter,

Astropart. Phys. 35 (2012) 508

Dashed arrows: kinetic equations include cooling and escape Q(E) [GeV-1 cm-3 s-1] per time frameN(E) [GeV-1 cm-3] steady spectrum

Input Object-dependent: B‘


> Energy losses in continuous limit:

b(E)=-E t-1loss

Q(E,t) [GeV-1 cm-3 s-1] injection per time frame (from sep. acc. zone)N(E,t) [GeV-1 cm-3] particle spectrum including spectral effects

NB: Need N(E) to compute particle interactions

> Simple case: No energy losses b=0:

> Special cases: tesc ~ R/c (leaky box)

tesc ~ E-a . Consequence: N(E) ~ Qinj(E) E-a, Escape: Qesc(E) = N(E)/tesc~ Qinj

(Neutrino spectrum from N(E) can have a break which is not present in escaping primaries Qesc(E))

Kinetic equations (steady state, one zone)

Injection EscapeEnergy losses


22

Peculiarity for neutrinos: Secondary cooling

Baerwald, Hümmer, Winter, Astropart. Phys. 35 (2012) 508; also: Kashti, Waxman, 2005; Lipari et al, 2007

Decay/cooling: charged m, p, K> Secondary spectra (m, p, K) loss-

steepend above critical energy

E‘c depends on particle physics only (m, t0), and B‘

Leads to characteristic flavor composition and shape

Decouples maximal neutrino and proton energies

E‘cE‘c E‘c

Pile-up effect Flavor ratio!

Spectralsplit

Example: GRB

Adiabatic

nm


From the source to the detector: UHECR transport

> Kinetic equation for co-moving number density:

> Energy losses UHECR must fromfrom our local environment (~ 1 Gpc at 1010 GeV, ~ 50 Mpc at 1011 GeV)

Photohadronics Hümmer, Rüger,

Spanier, Winter, 2010

Pair productionBlumenthal, 1970

Expansion ofUniverse

CR inj.z-dep!

(M. Bustamante)

[here b=-dE/dt=E t-1loss] GZK cutoff


Cosmogenic neutrinos

> Prediction depends on maximal proton energy, spectral index g, source evolution, composition

>Can test UHECR beyond the local environment

>Can test UHECR injection independent of CR production model constraints on UHECR escape

(courtesy M. Bustamante; see also Kotera, Allard, Olinto, JCAP 1010 (2010) 013)

Cosmogenic neutrinos

EeV

Protons


UHECR transition models

> Transition between Galactic (?) and extragalactic cosmic rays at different energies:

> Ankle model:

Injection index g ~ 2 possible ( Fermi shock acc.)

Transition at > 4 EeV

>Dip model:

Injection index g ~ 2.5-2.7 (how?)

Transition at ~ 1 EeV

Characteristic shape by pair production dip

Figure courtesy M. Bustamante; for a recent review, see Berezinsky, arXiv:1307.4043

Extra-galactic


26

More details: Gamma-ray observables?

>Redshift distribution>Can be integrated over.

Total number of bursts in the observable universe

Can be directly determined (counted)!

Order 1000 yr-1

(Kistler et al, Astrophys.J. 705 (2009) L104)

~ (1+z)a

Threshold correction

SFR


Consequence: Local GRB rate

> The local GRB rate can be written as

where fz is a cosmological correction factor:

(for 1000 observable GRBs per year and 30% of all bursts seen)

(Baerwald, Bustamante, Winter, arXiv:1401.1820)


Required baryonic loading (analytical)

>Required energy ejected in UHECR per burst:

> In terms of g-ray energy:

> Baryonic loading fe-1~50-100 for E-2 inj. spectrum (fbol ~ 0.2),

Eg,iso ~ 1053 erg, neutron model (fCR ~ 0.4)[IceCube standard assumption: fe

-1~10]

~1.5 to fit UHECR observations ~5-25

Energy in protons vs. electrons (IceCube def.)

How much energyin UHECR?

Fraction of energyin CR production?

Testing the origin of the UHECRs with neutrinos

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