Cosmic ray backgrounds for dark matter searches

(1)

Subir Sarkar

Rudolf Peierls Centre for Theoretical Physics, Oxford

&

Niels Bohr Institute, Copenhagen [with Philipp Mertsch, KIPAC Stanford]

Cosmic ray backgrounds for dark matter searches

AMS Days @ CERN: The future of cosmic ray physics , Geneva, 15-‐17 April 2015

(2)

If O(10%) of the shock K.E. of ~10⁵¹ erg can be converted into hadronic cosmic rays, then the observed ~3 SN/century can maintain the energy density of ~0.3 eV/cm³ in galactic cosmic rays

Supernova remnants are believed to be Pevatrons … responsible for the acceleration of Galactic cosmic rays up to the ‘knee’ at ~10

³

TeV

γ-‐ray spectrum consistent with accelera3on of hadrons

W44, Fermi W44, Herschel

(3)

Adriani et al, Nature 458:607,2009

PAMELA measured a rising positron fraction:

Anomaly ⟹ excess above astrophysical background

Source of anomaly:

• 

Dark matter?

⁽⁵⁰⁰⁺^models)

• 

Nearby pulsars?

• 

Nearby supernova remnants?

The PAMELA anomaly

(Gast & Schael, ICRC 09)

(4)

The Fermi ^anomaly

… and con`irmed that the positron fraction is rising – using the Earth s magnetic

`ield to do charge separation

Abdo et al, PRL 108:011103,2012

Fermi -‐ LAT also saw excess e

^±

over the expectation from the cosmic ray propagation

(GALPROP) model

Abdo et al , PRL 102:181101,2009

(5)

First result from AMS-02

Aguilar et al, Phys.Rev.Lett.110:141102,2013

“As the most precise measurement of the cosmic ray positron flux to date, these results show clearly the power and capabilities of the AMS detector,” said AMS spokesperson, Samuel Ting. “Over the coming months, AMS will be able to tell us conclusively whether these positrons are a signal for dark matter, or whether they have some other origin.”

(6)

Rate

(e.g. few hundred GeV neutralino LSP or Kaluza-‐Klein state)

⇤

⇤⇤

⇤

⇤⇤

⇧⇧⇧

⇧

⇧⇧ ⇧

⇧ ⇧

⌥

⌥⌥⌥

⌥⌥ ⌥

⌥

⇤⌥

⇤

1 10 10² 10³ 10⁴

1 10

0.3 3 30

Positron energy in GeV

Positronfraction

background?

PAMELA 08

⇤⇤⇤

⇤⇤

⇤

⇧

⇧ ⇧ ⇧

⇧

⌥

⌥⌥⌥ ⌥ ⌥ ⌥

⌥

⌅

⌅⌅

⌅⌅⌅⌅⌅⌅⌅⌅

⌅⌅

⇥

⇥⇥

⇥

⇥⇥⇥⇥⇥⇥⇥⇥⇥ ⇥

⇥⇥

⇥

⌃

⌃⌃⌃ ⌃⌃⌃ ⌃ ⌃

⌃

10 10² 10³ 10⁴

10^⇤3 10^⇤2 10^⇤1

Energy in GeV E3 e⇤ ⇥e⇥ ⇥GeV2 ⇤cm2 sec

HESS08 ATIC08

PPB⇤BETS08 EC

background?

⇤⇤⇤⇤⇤⇤⇤⇤⇤⇤

⇤

⇤⇤⇤⇤⇤

⇤

1 10 10² 10³ 10⁴

10^⇤5 10^⇤4 10^⇤3 10^⇤2

p kinetic energy in GeV

p⇤p

background? PAMELA 08

DM with M ⌅ 4 TeV that decays into ⇧^⇥⇧^⇤

Indeed dark matter has been widely invoked as the source of the excess e

⁺

.

DM annihilation DM decay

Rate

(lifetime ~10⁹x age of universe e.g.

dim-‐6 operator suppressed by M_GUT for a TeV mass techni-‐baryon)

Nardi, Sannino & Strumia, JCAP 0901:043,2009 Bergström, Bringmann & Edjsö, PR D78:127850,2008

(7)

But DM annihilation needs a huge boost factor to match flux

è Such a large annihilation #-‐section would imply a negligible relic abundance unless an inverse velocity dependence is invoked e.g. Somerfeld enhancement (this requires hypothetical light gauge bosons to provide a new long range force)

Cirelli, Kadastik, Raidal & Strumia, Nucl.Phys.B813:1,2009

Arkani-‐Hamed et al, PR D79:015014,2009

(8)

⇤⇤

⇤

⇤⇤

⇤

⇤⇤

⇧⇧⇧

⇧

⇧⇧ ⇧

⇧ ⇧

⌃

⌃⌃⌃

⌃⌃ ⌃

⌃

⇤⌃

⇤

1 10 10² 10³ 10⁴

1 10

0.3 3 30

Positronfraction

background?

PAMELA 08

⇤ ⇤⇤⇤⇤⇤⇤⇤⇤ ⇤⇤⇤⇤⇤

⇤⇤

⇤

⇧

⇧⇧⇧⇧ ⇧ ⇧ ⇧

⇧

⌃ ⌃

⌃⌃

⌃

⌃⌃

⌃

⌥

⌥⌥

⌥⌥⌥⌥⌥⌥⌥⌥

⌥⌥

⌅

⌅⌅

⌅

⌅⌅⌅⌅⌅⌅⌅⌅⌅ ⌅

⌅⌅

⌅

⇥

⇥⇥⇥⇥⇥⇥⇥ ⇥

⇥

10 10² 10³ 10⁴

10^⇤3 10^⇤2 10^⇤1

ATIC

PPB⇤BETS08 EC

background?

⇤⇤⇤⇤⇤⇤⇤⇤⇤⇤

⇤⇤⇤⇤⇤⇤

⇤

1 10 10² 10³ 10⁴

10^⇤5 10^⇤4 10^⇤3 10^⇤2

p⇤p

DM with M ⌅ 150 GeV that annihilates into W^⇥W^⇤

But the observed antiproton flux is consistent with the background expectation (from standard cosmic ray propagation in the Galaxy)

Cirelli et al, Nucl.Phys.B813:1,2009

⇤⇤

⇤

⇤⇤

⇤

⇤⇤

⇧⇧⇧

⇧

⇧⇧ ⇧

⇧ ⇧

⌃

⌃⌃⌃

⌃⌃ ⌃

⌃

⇤⌃

⇤

1 10 10² 10³ 10⁴

1 10

0.3 3 30

Positronfraction

background?

PAMELA 08

⇤ ⇤⇤⇤⇤⇤⇤⇤⇤ ⇤⇤⇤⇤⇤

⇤⇤

⇤

⇧

⇧⇧⇧⇧ ⇧ ⇧ ⇧

⇧

⌃ ⌃

⌃⌃

⌃

⌃⌃

⌃

⌥

⌥⌥

⌥⌥⌥⌥⌥⌥⌥⌥

⌥⌥

⌅

⌅⌅

⌅

⌅⌅⌅⌅⌅⌅⌅⌅⌅ ⌅

⌅⌅

⌅

⇥

⇥⇥⇥⇥⇥⇥⇥ ⇥

⇥

10 10² 10³ 10⁴

10^⇤3 10^⇤2 10^⇤1

ATIC

PPB⇤BETS08 EC

background?

⇤⇤⇤⇤⇤⇤⇤⇤⇤⇤

⇤⇤⇤⇤⇤⇤

⇤

1 10 10² 10³ 10⁴

10^⇤5 10^⇤4 10^⇤3 10^⇤2

p⇤p

background?

PAMELA 08

DM with M ⌅ 150 GeV that annihilates into W^⇥W^⇤

⇤⇤

⇤

⇤⇤

⇤

⇤⇤

⇧⇧⇧

⇧

⇧⇧ ⇧

⇧ ⇧

⌃

⌃⌃⌃

⌃⌃ ⌃

⌃

⇤⌃

⇤

1 10 10² 10³ 10⁴

1 10

0.3 3 30

Positronfraction

background?

PAMELA 08

⇤ ⇤⇤⇤⇤⇤⇤⇤⇤ ⇤⇤⇤⇤⇤

⇤⇤

⇤

⇧

⇧⇧⇧⇧ ⇧ ⇧ ⇧

⇧

⌃ ⌃

⌃⌃

⌃

⌃⌃

⌃

⌥

⌥⌥

⌥⌥⌥⌥⌥⌥⌥⌥

⌥⌥

⌅

⌅⌅

⌅

⌅⌅⌅⌅⌅⌅⌅⌅⌅ ⌅

⌅⌅

⌅

⇥⇥⇥⇥⇥⇥⇥⇥ ⇥

⇥

10 10² 10³ 10⁴

10^⇤3 10^⇤2 10^⇤1

ATIC

PPB⇤BETS08 EC

background?

⇤⇤⇤⇤⇤⇤⇤⇤⇤⇤

⇤⇤⇤⇤⇤⇤

⇤

1 10 10² 10³ 10⁴

10^⇤5 10^⇤4 10^⇤3 10^⇤2

p⇤p

background?

PAMELA 08

DM with M ⌅ 1 TeV that annihilates into ⇧^⇥⇧^⇤

⇤

⇤⇤

⇤

⇤⇤

⇧⇧⇧

⇧

⇧⇧ ⇧

⇧ ⇧

⌃

⌃⌃⌃

⌃⌃ ⌃

⌃

⇤⌃

⇤

1 10 10² 10³ 10⁴

1 10

0.3 3 30

Positronfraction

background?

PAMELA 08

⇤ ⇤⇤⇤⇤⇤⇤⇤⇤ ⇤⇤⇤⇤⇤

⇤⇤

⇤ ⇤

⇤

⇧

⇧⇧⇧⇧ ⇧ ⇧ ⇧

⇧

⌃ ⌃

⌃⌃

⌃

⌃ ⌃

⌃

⌥

⌥⌥

⌥⌥⌥⌥⌥⌥⌥⌥

⌥⌥

⌅

⌅⌅

⌅

⌅⌅⌅⌅⌅⌅⌅⌅⌅ ⌅

⌅⌅

⌅

⇥⇥⇥⇥⇥⇥⇥⇥ ⇥

⇥

10 10² 10³ 10⁴

10^⇤3 10^⇤2 10^⇤1

ATIC

PPB⇤BETS08 EC

background?

⇤⇤⇤⇤⇤⇤⇤⇤⇤⇤

⇤⇤⇤⇤⇤⇤

⇤

1 10 10² 10³ 10⁴

10^⇤5 10^⇤4 10^⇤3 10^⇤2

p⇤p

DM with M ⌅ 1 TeV that annihilates into ⇧^⇥⇧^⇤

Can cit with DM decay or annihilation only if dark matter particles are leptophilic

(Seems rather contrived nevertheless many such models were proposed)

This is a serious constraint on dark matter models of the PAMELA/AMS-‐02

positron anomaly

(9)

DM annihilation/decay energy release would increase the ionisation fraction of the intergalactic medium and broaden the ‘last scattering surface’ of the CMB

Monthly Notices of the Royal Astronomical Society 301, 210-214 (1998)

Cosmic microwave background anisotropy in the decaying neutrino cosmology J. A. Adams, S. Sarkar and D. W. Sciama

This would result in damping of the ‘acous3c’ peaks in the power spectrum of

CMB ﬂuctua3ons – as we noted originally for a model of decaying dark ma@er

(10)

The results are easily generalised to any source of ionising photons (E >13.6 eV) e.g. generated in the annihilation of dark matter particles (and resulting cascade)

Now that the CMB TT power spectrum is known to O(%) accuracy, Planck data sets a strong limit on this, disfavouring DM interpreta3ons of the PAMELA/AMS-‐02 anomaly

Ade et al, arXiv:

(11)

The inclusive jet differential cross section has been measured for jet transverse energies, E

_T

, from 15 to 440 GeV, in the pseudorapidity region 0.1≤|η|≤0.7. The results are based on 19.5 pb

^-1

of data collected by the CDF Collaboration at the Fermilab Tevatron collider. The data are

compared with QCD predictions for various sets of parton distribution functions. The cross section for jets with E

_T

> 200 GeV is significantly higher than current predictions based on O(α

_s³

)

perturbative QCD calculations. Various possible explanations for the high-E

_T

excess are discussed.

Abe et al, PRL 77:438,1996

This is not the first time an anomalous excess over background has been seen …

… it turned out to be a mis-‐estimation of

the QCD background – not new physics!

(12)

What particle physicists have learnt through experience ( UA1 monojets, NuTeV anomaly, CDF high E

_T

excess, … )

Yesterday s discovery is today s calibration

Richard Feynman

… and tomorrow s background!

Val Telegdi

… is also now a major issue for astroparticle physics viz.

just how well do we know the astrophysical background for

signals of apparently new particle physics?

(13)

The background is the production of secondary e ^± during propagation of nuclear cosmic rays in Galaxy

Acceleration of protons

interstellar medium

~90% H, ~10% He

…

¯

p

(14)

Diffusion of galactic cosmic rays

Transport equation:

energy losses

diffusion injection

Boundary conditions:

Green s function: describes clux from a discrete, burst-‐like source

… integrate over spatial distribution and time-‐variation of injection

GALPROP (Moskalenko & Strong, ApJ 493:694,1998, 509:212,1998) solves time-‐dependent transport equation … yields ~the same answer for equilibrium cluxes as the leaky box model in which cosmic rays have small energy-‐dependent probability of escape from Galaxy

⇒ exponential distribution of path lengths between cosmic ray sources and Earth Expectation: secondary/primary ratio E

^-δ

, where the diffusion co-‐efcicient D E

^δ

… `it to nuclear ratios (e.g. B/C) gives: δ ~ 0.4-‐0.7

(15)

All measured ratios consistent with τ

_esc

~ E

^-‐δ

, δ ~ 0.3-‐0.7

NB: Kolmogorov spectrum for interstellar magnetic cield turbulence would imply δ = 1/3, while Kraichnan spectrum yields δ = 1/2

Secondary-to-primary ratios (using GALPROP code)

Trotta et al, ApJ 729:106,2011

The Astrophysical Journal, 729:106 (16pp), 2011 March 10 Trotta et al.

D₀ (10²⁸ cm² s⁻¹

δ

68%, 95% contours

4 6 8 10 12

0.2 0.25 0.3 0.35 0.4

D₀ (10²⁸ cm² s⁻¹ v Alf (km/s)

4 6 8 10 12

30 35 40 45 50

δ v Alf (km/s)

0.2 0.3 0.4

30 35 40 45 50

D₀ (10²⁸ cm² s⁻¹ z h (kpc)

4 6 8 10 12

0 2 4 6 8 10 12

v_Alf (km/s) z h (kpc)

30 40 50

0 2 4 6 8 10 12

D₀ (10²⁸ cm² s⁻¹

ν 1

4 6 8 10 12

1.7 1.8 1.9 2 2.1

δ

ν 1

0.2 0.3 0.4

1.7 1.8 1.9 2 2.1

v_Alf (km/s)

ν 1

30 40 50

1.7 1.8 1.9 2 2.1

z_h (kpc)

ν 1

0 5 10

1.7 1.8 1.9 2 2.1

δ

ν 2

0.2 0.3 0.4

2.2 2.3 2.4 2.5 2.6

v_Alf (km/s)

ν 2

30 40 50

2.2 2.3 2.4 2.5 2.6

ν₁

ν 2

1.7 1.8 1.9 2 2.1

2.2 2.3 2.4 2.5 2.6 (

(

Figure 3. Two-dimensional marginalized posterior probability distributions for some parameter combinations. The yellow and blue regions enclose 68% and 95%

probability, respectively. The encircled red cross is the best fit, the filled green dot the posterior mean.

(A color version of this figure is available in the online journal.)

5.3. Quality of Best-fit Model

We now assess the quality of our best-fit model. Define the χ² as

χ² ≡

!5 j=1

N_j

!

i=1

"ΦX(Ei,Θ, φ) − ˆΦ^ijX

#2

σ_ij²/τ_j , (18) i.e., we compute the χ² using the rescaled error bars for the data points (note that the χ² ̸= −2 log P (D|Θ, φ, τ ), i.e., the χ² is not minus twice the log-likelihood because of the pre- factor containing τ appearing in Equation (8)). There are N = 76 total data points and M = 16 fitted parameters, including both the modulation and the error rescaling parameters. Therefore, the number of degrees of freedom (dof) is 60, and for the best-fit model we find χ² = 69.3, which leads to a reduced chi-squared χ²/dof = 68/60 = 1.15. This is not surprising, since by construction the error bar rescaling parameters, τ , are adjusted dynamically during the global fit to achieve this. A more detailed breakdown of the contribution to the total χ² by data set is given in Table 3.

The predictions for the fitted CR spectra of the best-fit model parameters are shown in Figures 4–6, including an error band delimiting the 68% and 95% probability regions. The species shown are B/C and ¹⁰Be/⁹Be ratios, and the spectra of carbon and oxygen. In each plot, we show the spectrum modulated with the potential corresponding to our best-fit parameters from our global fits for each of the data sets employed. We also show the data sets, each with error bars enlarged by the best-fit value of our scaling parameters, τ , as given in Table 2. The yellow/

blue band delimits regions of 68% and 95% probability, and is modulated according to the potential given in the each panel.

We emphasize that the power of our statistical technique is such that we can, for the first time, provide not only a best- fit model but also an error band with a well-defined statistical meaning.

In order to better visualize the comparison of our best-fit model to the fitted data, we plot in the bottom part of each panel the best-fit residuals, i.e., the difference between data and best- fit model, divided by the experimental error bar (enlarged by the

10

0 0.1 0.2 0.3 0.4

Ratio

0 2

10² 10³ 10⁴ 10

(16)

However e

^±

lose energy readily during propagation, so only nearby sources dominate at such high energies

… the usual background calculation is then irrelevant

Delhaye et al, A&A 501:821,2009

Might there be primary sources of e⁺ (with a hard spectrum) in our Galactic neighbourhood?

⌧ ' 5 ⇥ 10

⁵

yr

✓ 1TeV E

◆

100 MeV 10 GeV

1 TeV 100 GeV

1 GeV

Kobayashi et al, ApJ 601:340,2004

(17)

A nearby cosmic ray accelerator?.

Rise in e

⁺

fraction could be due to secondaries produced during acceleration … which are then accelerated along with the primaries

(Blasi, PRL 103:051104,2009)

.

.. generic feature of a stochastic acceleration process, if

τ

_1➛2

<

τ

acc

(Cowsik 1979, Eichler 1979)

This component has a harder spectrum so naturally

`its the PAMELA / AMS-‐02 positron excess!

Ahlers, Mertsch & Sarkar,PRD80:123017,2009

(18)

Flux:

Conservation equation:

Steady state:

Diffusive (1 ^st -order Fermi) shock acceleration

density change acceleration convection injection

log f

log p

i.e. = 4 for strong shock (u

₁

/u

₂

= 4)

Courtesey: Luke Drury

(19)

Acceleration determined by compression ratio:

Solve transport equation:

Diffusive (1 ^st -order Fermi) shock acceleration

u f

x = D ²f

x² + 1 3

du

dxp f p

f

^x^{⇥ ⇤}

⇤ f

^inj

(p), lim

x⇥⇤

f ⇥ ⌅

Solution for:

where,

f

⁰

(p) = ⇤

p 0

dp

^⇥

p

^⇥

p

^⇥

p

⇥

f

_inj

(p

^⇥

) + Cp

As long as is softer than at high energies:

f(x, p) p

f

_inj

(p) p

(20)

● 

Secondaries have same spectrum as primaries (Feynman):

● 

Only particles with are accelerated

● 

Bohm diffusion:

⇒ Fraction of accelerated secondaries is

i.e. steady state spectrum is:

DSA with secondary production

p₂ > p₁

… thus yielding a r ising positron fraction

log n

log p

(21)

Diffusion near shock front

Ø 

Diffusion coef`icient not known

a priori in neighbourhood of shock

Ø 

‘ Bohm diffusion’ sets a limit:

Ø 

Actual rate may be parametrised by:

Ø 

Can try and determine diffusion rate from simulations (dif`icult!)

Ø 

So we `ix K

_B

by `itting to the Fermi e

^±

excess … can then predict e

⁺

/ ( e

⁺

+

e

^-‐

)

for PAMELA/AMS-‐02, and other

secondary-‐to-‐primary ratios (e.g B/C)

D = D

^Bohm

/K

_B

, K

_B

= B

²

/ B

²

(22)

where

is downstream volume

The downstream spectrum, integrated over SNR lifetime …

Mertsch, Sarkar, Phys.Rev.D90, 061301(R),2014

(23)

It is not just the few (optically) observed SNRs which contribute to observed cosmic rays … there must be many other hidden SNRs

(if there are ~3 SN/century and cosmic rays diffuse in Galaxy for ~10

⁷

yr)

!" ^! ! !"

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)*+,-./0 12/

,*3045

602705

84/9:

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H

!" ^! ! !"

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!"⁽

)*+,-./0 12/

,*3045

!"# $%&%'() %! *+"#"!& ,)#

10 GeV 100 GeV 1 TeV

Known Simulated

Ahlers, Mertsch & Sarkar,PRD80:123017,2009

(24)

Electrons in the Galaxy

The behavior of cosmic ray electrons in our Galaxy is very energy dependent because of radiative losses.

Some movies made of the electron density in our Galaxy based on the Galprop simulation code have been made by Simon Swordy. Each of these starts with an empty Galaxy and runs for about 10 million years (although ~20 seconds in real time!) with standard SNR rate and distribution used in GALPROP. One shows

electrons at 1 GeV, where radiative losses are minimal, and the other with electrons at 100TeV, where radiative losses remove electrons very quickly. The end frames of each simulation are shown below, together with links to the Galactic electron density movies. The purpose of these is to show how measured

electron cluxes are subject to the time and positional distribution of local SNR at high energies (>~1TeV) (Please cite Simon Swordy as author of these movies)

http://astroparticle.uchicago.edu/resources.html

(25)

×

!"#

!"$

!#

$

#

"$

"#

!"# !"$ !# $ # "$ "#

%&'()

! " #! #" $! $"

!

"

#!

#"

$!

Statistical distribution of SNRs in our neighbourhood

•  Draw source positions from this distribution

•  Inject e

^-‐

^&

e

⁺

normalized to observables

•  Propagate to Earth accounting for synchrotron and inverse-‐Compton scattering energy losses

•  Confront total e

^-‐

+e

⁺

clux at Earth with data

•  The best Hit to data is closest to real distribution

(Case & Bhattacharya, ApJ 504:761,1998)

Ahlers, Mertsch, Sarkar,PRD80:123017,2009

(26)

Ahlers, Mertsch & Sarkar, PRD80:123017,2009

The propagated primary e

^-‐

spectrum is much too steep to match the Fermi LAT data ... but the accelerated total secondary ( e

⁺

+

e

^-‐

)

component has a harder

spectrum so does `it the bump !

Fitting the e ⁺ + e ^- flux

(27)

The ‘ postdicted’ positron fraction

1 10 10² 10³ 10⁴

10 ² 10 ¹ 1

Energy GeV⇥

Positronfraction

PAMELA

Standard Solar modulation

Charge-‐sign dependent Solar modulation

Ahlers, Mertsch, Sarkar, PRD80:123017,2009

(28)

Our proposal is different from:

“Inhomogeneity in the SNR distribution

as the origin of the PAMELA anomaly”

Shaviv, Nakar & Piran, PRL 103:111302,2009

Idea: Electrons from nearby SNRs cool above ~ 20 GeV (through synchrotron and

inverse-‐Compton losses) before reaching us … but protons do not cool, so secondary positron

production is less affected

⇒ enhancement of e

⁺

/e

^-‐

But with usual propagation parameters

(D

₀

~ 10

²⁸

cm

²

s

^-‐1

, δ ~ 0.6, τ

_esc

~ 10

¹⁶

s)

the break energy is 2 TeV, not 20 GeV …

also nearby invisible SNRs (e.g. Geminga)

will cill in the dips in the energy spectrum

(29)

Nuclear secondary-to-primary Ratios

Dark matter ^✗ Pulsars ^✗ Acceleration of

secondaries

^(TBD)

Since nuclei are accelerated in the same sources, the ratio of secondaries (e.g.

Li, Be, B) to primaries (C, N, O) must also rise with energy beyond ~100 GeV

If this is con`irmed, both dark matter and

pulsar origin models would be ruled out!

Mertsch, Sarkar, PRL 103:081104,2009

spallation during propagation only spallation during acceleration

Energy per nucleon, GeV

10 10² 10³

B/C ratio

0 0.05 0.1 0.15 0.2 0.25 0.3

0.35 ATIC, experiment

HEAO-3, experiment [1]

Osborn & Ptuskin, leaky box model [4]

HEAO-3 model, leaky box model [1]

?

(30)

We cit the AMS-‐02 p, He cluxes to cix the spectral indices and normalisation, and the e

^-‐

clux (in accordance with radio data)

10⁰ 10¹ 10² 10³ 10⁴ kinetic Energy E [GeV/n]

10² 10³ 10⁴

E2.7 Jp,E2.7 JHe[(GeV/n)1.7 m2 s1 sr1 ]

Rmax = 10³GV Rmax = 3 ⇥ 10³GV R_max = 10⁴GV

p (AMS-02) He (AMS-02)

10⁰ 10¹ 10² 10³ 10⁴

kinetic Energy E [GeV]

10⁰ 10¹ 10²

E3 J e,E3 J e+[GeV2 m2 s1 sr1 ]

R_max = 10³ GV Rmax = 3 ⇥ 10³GV Rmax = 10⁴ GV

e (AMS-02) e⁺ (AMS-02)

(31)

10 ¹ 10⁰ 10¹ 10² 10³ 10⁴ kinetic Energy E [GeV]

10 ¹

positronfraction

R_max = 10³ GV Rmax = 3 ⇥ 10³ GV R_max = 10⁴ GV

(AMS-02)

10 ¹ 10⁰ 10¹ 10² 10³ 10⁴ 10⁵ kinetic Energy E [GeV/n]

10 ¹

B/C

R_max = 10³GV Rmax = 3⇥ 10³GV R_max = 10⁴GV

AMS-02 CREAM TRACER

We can then predict secondary to primary ratios – the only free parameter is the maximum energy of the cosmic accelerator (taken to

be 1, 3, 10 TeV for illustration)

10 ¹ 10⁰ 10¹ 10² 10³ 10⁴

10 ⁶ 10 ⁵ 10 ⁴ 10 ³

pbar/p

Rmax = 10³GV Rmax = 3⇥ 10³GV Rmax = 10⁴GV (PAMELA)

Measurements of B/C and

by AMS-‐02 at higher energies

then calibrate/test our model

(32)

10 ¹ 10⁰ 10¹ 10² 10³ 10⁴ 10⁵ kinetic Energy E [GeV/n]

10 ¹

B/C

Rmax = 10³ GV Rmax = 3⇥ 10³ GV Rmax = 10⁴ GV

AMS-02 CREAM TRACER

10 ¹ 10⁰ 10¹ 10² 10³ 10⁴

10 ⁶ 10 ⁵ 10 ⁴ 10 ³

pbar/p

Rmax = 10³GV Rmax = 3⇥ 10³GV R_max = 10⁴GV (PAMELA)

… and here are our predictions confronted with the latest AMS-‐02 data Looks like the cutoff is closer to ~500 GV?

(33)

We have been trying (late last

night!) to get better cits to the new data

but it is not easy … perhaps

our model is too simple and

some further recinements are

necessary.

This is justicied

now that we

have precision

data from AMS!

(34)

Cosmic ray backgrounds for dark matter searches

Subir Sarkar

Rudolf Peierls Centre for Theoretical Physics, Oxford

&

Niels Bohr Institute, Copenhagen [with Philipp Mertsch, KIPAC Stanford]

Cosmic ray backgrounds for dark matter searches

AMS Days @ CERN: The future of cosmic ray physics , Geneva, 15-­‐17 April 2015

Supernova remnants are believed to be Pevatrons … responsible for the acceleration of Galactic cosmic rays up to the ‘knee’ at ~10

TeV

PAMELA measured a rising positron fraction:

Anomaly ⟹ excess above astrophysical background

Source of anomaly:

Dark matter?

Nearby pulsars?

Nearby supernova remnants?

The PAMELA anomaly

The Fermi anomaly

… and con`irmed that the positron fraction is rising – using the Earth s magnetic

`ield to do charge separation

Fermi -­‐ LAT also saw excess e

over the expectation from the cosmic ray propagation

(GALPROP) model

Abdo et al , PRL 102:181101,2009

First result from AMS-02

Rate

Indeed dark matter has been widely invoked as the source of the excess e

.

DM annihilation DM decay

Rate

But DM annihilation needs a huge boost factor to match flux

è Such a large annihilation #-­‐section would imply a negligible relic abundance unless an inverse velocity dependence is invoked e.g. Somerfeld enhancement (this requires hypothetical light gauge bosons to provide a new long range force)

But the observed antiproton flux is consistent with the background expectation (from standard cosmic ray propagation in the Galaxy)

Can cit with DM decay or annihilation only if dark matter particles are leptophilic

(Seems rather contrived nevertheless many such models were proposed)

This is a serious constraint on dark matter models of the PAMELA/AMS-­‐02

positron anomaly

DM annihilation/decay energy release would increase the ionisation fraction of the intergalactic medium and broaden the ‘last scattering surface’ of the CMB

Cosmic microwave background anisotropy in the decaying neutrino cosmology J. A. Adams, S. Sarkar and D. W. Sciama

This would result in damping of the ‘acous3c’ peaks in the power spectrum of

CMB ﬂuctua3ons – as we noted originally for a model of decaying dark ma@er

The results are easily generalised to any source of ionising photons (E >13.6 eV) e.g. generated in the annihilation of dark matter particles (and resulting cascade)

Now that the CMB TT power spectrum is known to O(%) accuracy, Planck data sets a strong limit on this, disfavouring DM interpreta3ons of the PAMELA/AMS-­‐02 anomaly

The inclusive jet differential cross section has been measured for jet transverse energies, E

, from 15 to 440 GeV, in the pseudorapidity region 0.1≤|η|≤0.7. The results are based on 19.5 pb

of data collected by the CDF Collaboration at the Fermilab Tevatron collider. The data are

compared with QCD predictions for various sets of parton distribution functions. The cross section for jets with E

> 200 GeV is significantly higher than current predictions based on O(α

)

perturbative QCD calculations. Various possible explanations for the high-E

excess are discussed.

This is not the first time an anomalous excess over background has been seen …

… it turned out to be a mis-­‐estimation of

the QCD background – not new physics!

What particle physicists have learnt through experience ( UA1 monojets, NuTeV anomaly, CDF high E

excess, … )

Yesterday s discovery is today s calibration

Richard Feynman

… and tomorrow s background!

Val Telegdi

… is also now a major issue for astroparticle physics viz.

just how well do we know the astrophysical background for

signals of apparently new particle physics?

The background is the production of secondary e ± during propagation of nuclear cosmic rays in Galaxy

Acceleration of protons

interstellar medium

~90% H, ~10% He

…

…

…

¯

p

Diffusion of galactic cosmic rays

Transport equation:

energy losses

diffusion injection

Boundary conditions:

Green s function: describes clux from a discrete, burst-­‐like source

… integrate over spatial distribution and time-­‐variation of injection

⇒ exponential distribution of path lengths between cosmic ray sources and Earth Expectation: secondary/primary ratio E

, where the diffusion co-­‐efcicient D E

AMS Days @ CERN: The future of cosmic ray physics , Geneva, 15-‐17 April 2015

The Fermi ^anomaly

Fermi -‐ LAT also saw excess e

è Such a large annihilation #-‐section would imply a negligible relic abundance unless an inverse velocity dependence is invoked e.g. Somerfeld enhancement (this requires hypothetical light gauge bosons to provide a new long range force)

This is a serious constraint on dark matter models of the PAMELA/AMS-‐02

Now that the CMB TT power spectrum is known to O(%) accuracy, Planck data sets a strong limit on this, disfavouring DM interpreta3ons of the PAMELA/AMS-‐02 anomaly

… it turned out to be a mis-‐estimation of

The background is the production of secondary e ^± during propagation of nuclear cosmic rays in Galaxy

Green s function: describes clux from a discrete, burst-‐like source

… integrate over spatial distribution and time-‐variation of injection

, where the diffusion co-‐efcicient D E

… `it to nuclear ratios (e.g. B/C) gives: δ ~ 0.4-‐0.7

, δ ~ 0.3-‐0.7

`its the PAMELA / AMS-‐02 positron excess!

Diffusive (1 ^st -order Fermi) shock acceleration

Diffusive (1 ^st -order Fermi) shock acceleration