Looking for Dark Matter in Cosmic Rays

Andrea VittinoUniversita’ di Torino and IPhT, CEA/Saclay

Under the supervision of Nicolao Fornengo (Unito) and Marco Cirelli (CEA)

Seminario del secondo annoTorino, February 14 2014

A dark universeThere are overwhelming evidences for the existence of dark matter:

If DM is made of particles that interact among themselves and with standard

model particles we may hope to detect it

This talk is about indirect detection

• We are looking for the most visible products of DM annihilation: antiparticles

• We are looking for them in the cosmic rays flux

• In particular, we are interested in:

‣ Antiprotons (Fornengo, Maccione, AV, “Constraints on particle dark matter from

cosmic-ray antiprotons”, arXiv: 1312..3579[hep-ph])

‣ Anti-nuclei (Fornengo, Maccione, AV, “Dark Matter searches with cosmic antideuterons: status and perspectives” JCAP 09 (2013) 031, Cirelli, Fornengo, Taoso, AV “Anti-helium from Dark

Matter annihilations”, arXiv: 1401.4017)

‣ Positrons (Di Mauro, Donato, Fornengo. Lineros, AV, “Interpretation of AMS-02 electrons and positrons data ”, arXiv:1402.0321)

Antimatter as a tool for indirect detection

• We are looking for the most visible products of DM annihilation: antiparticles

• We are looking for them in the cosmic rays flux

• In particular, we are interested in:

‣ Antiprotons (Fornengo, Maccione, AV, “Constraints on particle dark matter from

cosmic-ray antiprotons”, arXiv: 1312..3579[hep-ph])

‣ Anti-nuclei (Fornengo, Maccione, AV, “Dark Matter searches with cosmic antideuterons: status and perspectives” JCAP 09 (2013) 031, Cirelli, Fornengo, Taoso, AV “Anti-helium from Dark

Matter annihilations”, arXiv: 1401.4017)

‣ Positrons (Di Mauro, Donato, Fornengo. Lineros, AV, “Interpretation of AMS-02 electrons and positrons data ”, arXiv:1402.0321)

Antimatter as a tool for indirect detection

• Some years ago, antideuterons have been proposed as an almost background free channel for DM indirect detection [F. Donato, N. Fornengo, P.Salati, Phys.Rev. D62 (2000) 043003]

• A reaction that generates many antideuterons also generates a lot of antiprotons. We can derive very strong bounds on DM annihilation cross section from PAMELA measurements on the antiprotons flux

Antideuterons are a good discovery channel for DM!

Antiprotons are an excellent tool to constrain DM properties

Antiprotons and antideuterons

• Some years ago, antideuterons have been proposed as an almost background free channel for DM indirect detection [F. Donato, N. Fornengo, P.Salati, Phys.Rev. D62 (2000) 043003]

• A reaction that generates many antideuterons also generates a lot of antiprotons. We can derive very strong bounds on DM annihilation cross section from PAMELA measurements on the antiprotons flux

Antideuterons are a good discovery channel for DM!

Antiprotons are an excellent tool to constrain DM properties

Antiprotons and antideuterons

How do we calculate antiprotons and antideuterons fluxes?

1 - Production (DM annihilation or decay)

2 - Propagation in the galaxy

3 - Solar modulation

From the Source to the Earth

1 - Production (DM annihilation or decay)

2 - Propagation in the galaxy

3 - Solar modulation

From the Source to the Earth

We work in a “model

indipendent” way (i.e. pure annihilation

channels)

Hadronization is modeled with a MC event generator (Pythia)

The spectra can be taken directly from Pythia, while for the spectra we

have to merge two antinucleons

DM model Hadronization Coalescence

p

d

np

p

d

�

�

q,W,Z

What can we say about coalescence?

Production

A simple idea: antinucleons coalesce if they are close enough (in the phase space)

is the probability that the antinucleons are formed:

dNd

dT/

Zd3~kpd

3~kn Fpn(ps,~kp,~kn)C(~� = ~kp � ~kn)

F(pn)(ps,~kp,~kn) =

dN(pn)

d3~kpd3~kn

F(pn)

The function is the probability that the antinucleons merge:

We take (radius of the antideuteron)

C

C(~�) = ✓(�2 � p20) ✓(�r2 � r20)is a free

parameter. Which isits value?

p0

The Coalescence puzzle

r0 ⇡ 2 fm

i.e. we don’t consider antinucleons coming from the decay of long-living particles

We tune to reproduce ALEPH data:

0.1 0.12 0.14 0.16 0.18 0.2 0.222x10<6

5x10<6

10<5

p0 [GeV]

Rd

[ALEPH collaboration, Phys. Lett. B 369 (2006) 192]

production rate in e+e- collisions at the Z resonance

p0 = (195 ± 22) MeV

Basically, a is formed if(|�(~k)| < 195 MeV

|�(~r)| < 2 fm

p0

d

d

The Coalescence puzzle

p0 = (195 ± 22) MeV

1 - Production (DM annihilation or decay)

2 - Propagation in the galaxy

3 - Solar modulation

From the Source to the Earth

Two-zone diffusion model

K(r, z, E) = �K0

✓R

1 GV

◆�

~Vc = sign(z)Vc

is the DM halo density profile

R

L

DM halo

To propagate both the and the we have to solve a transport equation:dp

1

2< �v >

dNp,d

dT

✓⇢(r, z)

mDM

◆2

⇢(r, z)

diffusivezone

Galactic propagation

�r[K(r, z, E)rnp,d

(r, z, E)] + Vc

(z)@

@znp,d

(r, z, E) + 2h�(z)�annp,d

np,d

(r, z, E)+

2h�(z)@E

(�KEE

(E)@E

np,d

(r, z, E) + btot

(E)np,d

(r, z, E)) = qp,d

(r, z, E)

Source term

DiffusionConvection

Annihilation

Reacceleration Energy losses

For DM annihilation:

If we neglect both reacceleration and energy losses we can factorize the flux:

� K0 (kpc2/Myr) L (kpc) Vc (km/s)MIN 0.85 0.0016 1 13.5MED 0.70 0.0112 4 12MAX 0.46 0.0765 15 5

0.01 0.1 1 10 100 1000 10000104

105

106

107

108

109

1010

T [GeV]

R [y

r]

EinastoNFWcored isothermalMAX

MED

MIN

All the astrophysics is confined here!

K0, Vc and δ constrained by B/C data

p

The two-zone diffusion model is defined by these parameters:

�d.p(E) =�d,p

4⇡nd,p(r = r�, z = 0, E) =

�d,p

4⇡

✓⇢�

mDM

◆2

Rd,p(E)1

2< �v >

dNd,p

dE

F.Donato, N.Fornengo, D.Maurin, P.Salati and R.Taillet Phys.Rev.D 69 (2004) 063501

Energy losses and reacceleration(Work in progress!)

We write the solution as a Bessel series:

A

0

BBBBBB@

...nj�1i

nji

nj+1i...

1

CCCCCCA=

0

BBBBBB@

...n0,j�1i

n0,ji

n0,j+1i...

1

CCCCCCA

ni

+2h

Bi

d

dE

"btot

(E)ni

�KEE

(E)dnp,d

i

dE

#= n0

i

np,d(r, z, E) =X

i

ni(z, E)J0

✓⇣ir

R

◆

To include reacceleration and energy losses one has to solve:

Which is equivalent to:

Solution can be found by (numerically) invertingA

1 - Production (DM annihilation or decay)

2 - Propagation in the galaxy

3 - Solar modulation

From the Source to the Earth

The Sun’s magnetic field (SMF) has the form of a large rotating spiral

An heliospheric current sheet (HCS), whose shape varies with time according to solar activity,

separates field lines directed towards or away from the Sun

How can we model the motion of a charged particle inside the SMF?

Generally, this is done by using the force field approximation:

�TOA(TTOA) =2mTTOA + TTOA

2mTIS + TIS�IS(TIS) TTOA = TIS � '

Solar modulation

• The tilt angle α : it describes the spatial extent of the HCS. It is proportional to the intensity of the solar activity (α ∈ [20◦,60◦])

• The mean free path λ of the CR particle along the magnetic field direction

We exploit the code HELIOPROP to solve numerically the transport equation and explore the solar parameters space

@f

@t= �(~Vsw + ~vd) ·rf +r · (K ·rf) +

P

3(r · ~Vsw)

@f

@P

We vary 2 parameters:

The propagation in the heliosphere is described by the following equation:

Convection DriftsDiffusion

(random walk)Adiabatic losses

L. Maccione, Phys.Rev.Lett. 110, 081101(2013)

[E. N. Parker, P&SS 13, 9 (1965)]

Charge dependent solar modulation

Solid lines Force field approximation (pbar, dbar)

Dots CD solar modulation (pbar, dbar)

Above 10 GeV solar mod is negligible

In our sample, energy losses vary significantly from particle to particle (they depend on the path):

Charge dependent solar modulation

The modulated flux at ETOA is given by a weighted

average of LIS fluxes in the corresponding range of ELIS

●

●

●●

● ●●

● ●● ●

●●

●

●

●

●

●

●

●

●

●

●

1 10 10010−6

10−5

10−4

10−3

10−2

10−1

T [GeV]

φ p [(

m2 s

sr G

eV)−

1 ]

MED flux

● PAMELABESSbackground

Antiprotons bounds

The measured flux appears to be very well fitted by a pure

astrophysical background

Very little room left for dark matter!

To compute the 3σ bounds on <σv> we associate to the

background flux a theoretical uncertainty of the 40%

The antiprotons flux has been measured recently by PAMELA and BESS-Polar

1 10 100 100010−28

10−27

10−26

10−25

10−24PAMELA bounds − EINASTO MED − annihilating DM

mDM [GeV]

<σan

nv>

[cm

3 s−1 ]

uussbbccttW+W−

ZZ

Antiprotons bounds

MED

1 10 100 100010−29

10−28

10−27

10−26

10−25PAMELA bounds − EINASTO MAX − annihilating DM

mDM [GeV]

<σan

nv>

[cm

3 s−1 ]

uussbbccttW+W−

ZZ

1 10 100 100010−27

10−26

10−25

10−24

10−23PAMELA bounds − EINASTO MIN − annihilating DM

mDM [GeV]

<σan

nv>

[cm

3 s−1 ]

uussbbccttW+W−

ZZ

Antiprotons bounds

MIN

MAX

The uncertainty related to the propagation model is the largest one

Quite large effect, in particular for high DM masses

In our calculations we have used an Einasto profile

Here we consider also an NFW and a cored isothermal profiles

The role of the DM profile

10 100 100010−27

10−26

10−25

10−24bb channel − annihilating DM

mDM [GeV]

<σan

nv>

[cm

3 s−1 ]

EinastoNFWCored isothermal

For the solar modulation we have used parameters compatible with the

PAMELA data-taking period(negative polarity,

α = 20° λ=0.15 AU)

We vary the value of the mean free path in the SMF direction

Not so large effect!(It reaches at most the 30-50%)

The role of solar modulation-I

10 100 1000

10−26

10−25

10−24bb channel − annihilating DM

mDM [GeV]

<σan

nv>

[cm

3 s−1 ]

λ = 0.15 AUλ = 0.20 AUλ = 0.25 AU

The role of solar modulation-II

1 10 100 10000.01

0.1

1uu channel − decaying DM

mDM [GeV]

Rbo

unds

λ = 0.20 AUλ = 0.25 AU

100 200 500 10000.01

0.1

1W+W− channel − annihilating DM

mDM [GeV]

Rbo

unds

λ = 0.20 AUλ = 0.25 AU

So what can we say?

★ Bounds coming from antiprotons measurements are really strong: for hadronic channels they are comparable or even stronger than the ones coming from gamma ray measurements.

★ The main source of uncertainty is represented by the galactic propagation and the choice of the DM profile. The solar modulation brings a smaller but unavoidable uncertainty.

★ The question now is: given these strong bounds, will we be able to see an antideuteron flux in the near future?

DM in the antideuterons channelAMS-02 GAPS

It’s operative since 2011(here we will refer to a 3 year

data-taking period)

It’s a proposed experiment which will begin its science flights in 2017/2018(in the LDB+ setup a data-taking period

of 210 days)

0.1 1 1010−8

10−7

10−6

10−5

10−4uu channel − mDM = 10 GeV

T [GeV/n]

φd [

(m2s

sr G

eV

/n)−

1]

GAPSLDB+ AMS

MED fluxes

Force FieldCD_60_0.60_1CD_60_0.15_0.5background

0.1 1 1010−8

10−7

10−6

10−5

10−4 bb channel − mDM = 20 GeV

T [GeV/n]

φd [

(m2s

sr G

eV

/n)−

1]

GAPSLDB+ AMS

MED fluxes

Force FieldCD_60_0.60_1CD_60_0.15_0.5background

0.1 1 1010−8

10−7

10−6

10−5

10−4 W+W− channel − mDM = 100 GeV

T [GeV/n]

φd [

(m2s

sr G

eV

/n)−

1]

GAPSLDB+ AMS

MED fluxes

Force FieldCD_60_0.60_1CD_60_0.15_0.5background

With the maximal cross sections allowed by antiprotons constraints:

We can have a flux on the reach of both experiments!

Background [F.Donato, N.Fornengo and D.Maurin, Phys. Rev. D 78 (2008)]

uu - mDM = 10 GeV bb - mDM = 20 GeV WW - mDM = 100 GeV

Fluxes at Earth

<σv> = 0.1 x <σv>th <σv> = 0.3 x <σv>th <σv> = 2 x <σv>th

Reachability = curve in the (mDM,<σv>) plane that corresponds to a detection (with a 3σ C.L.)

10 100 100010−28

10−27

10−26

10−25

10−24

bb channel 3σ curves + PAMELA bounds

mDM [GeV]

<σv>

[cm

3 s−1 ]

MIN−MED−MAX

tilt angle = 60°mean free path = 0.60 AUspectral index = 1

10 100 100010−29

10−28

10−27

10−26

10−25

10−24

uu channel 3σ curves + PAMELA bounds

mDM [GeV]

<σv>

[cm

3 s−1 ]

MIN−MED−MAX

tilt angle = 60°mean free path = 0.60 AUspectral index = 1

Experimental Reachability

With antideuterons we can explore vast regions of the DM parameter space

Number of expected eventsFor a WIMP annihilating in the bb channel

10 100 10000.1

1

10

100bb channel

mDM [GeV]

NG

APS

Force FieldCD_60_0.15_0.5CD_60_0.20_1CD_60_0.60_1

10 x <σv>th <σv>th

0.1x<σv>th

Solid lines: configurations compatible with PAMELA bounds on antiprotons

From 3 to 5 events depending on solar modulation

Dot-dashed lines: configurations not compatible

with PAMELA bounds on antiprotons

Number of expected events

10 100 10000.1

1

10

100uu channel

mDM [GeV]

NG

APS

Force FieldCD_60_0.15_0.5CD_60_0.20_1CD_60_0.60_1

100 200 500 10000.01

0.1

1

10

100W+W− channel

mDM [GeV]

NG

APS

Force FieldCD_60_0.15_0.5CD_60_0.20_1CD_60_0.60_1

uu channel WW channel

0.1x<σv>th 0.1x<σv>th

<σv>th <σv>th 10x<σv>th

10x<σv>th

So what can we say?

★ Bounds coming from antiprotons measurements are really strong: for hadronic channels they are comparable or even stronger than the ones coming from gamma ray measurements.

★ The main source of uncertainty is represented by the galactic propagation and the choice of the DM profile. The solar modulation brings a smaller but unavoidable uncertainty.

★ The question now is: given these strong bounds, will we be able to see an antideuteron flux in the near future?

Hopefully yes!

10 100 100010−28

10−27

10−26

10−25

10−24EINASTO MED − bb channel − annihilating DM

mDM [GeV]

<σan

nv>

[cm

3 s−1 ]

PAMELAAMS

But...life is never easy!In the close future, a better knowing of the antiprotons background and a higher experimental precision will lead to much stronger bounds on <σv> coming from

antiprotons data

1 10 100 100010−29

10−28

10−27

10−26

10−25

10−24EINASTO MED − uu channel − annihilating DM

mDM [GeV]

<σan

nv>

[cm

3 s−1 ]

PAMELAAMS

uncertainty on the backgroundgoes from 40% (upper curves)

to 5% (lower curves)

AMS-02 bounds are calculated by using mock data

(for T > 1 GeV)

A possible way out: the coalescence mechanism

0.1 1 10

10−7

10−6

bb channel − mDM = 20 GeV

T [GeV/n]

φ d [(

m2 s

sr G

eV/n

)−1 ]

10 100 10000.01

0.1

1

10

100

bb channel

mDM [GeV]

NG

APS

p0 = 195 MeVp0 = 217 MeVp0 = 239 MeVp0 = 261 MeV

We have to remember that our results strongly depends on the value of the coalescence momentum which we don’t know very well

We can have enhancements of the expected signals!

N.Fornengo,L.Maccione, A.Vittino, JCAP 09 (2013) 31

What about antihelium?

What as said for the antideuterons is obviously true also for antihelium. Some remarks:

• We expect both the signal and the background to be much fainter (being a 3-body coalescence process) this is why we consider as a possible DM signal antiHe-3 and not antiHe-4

•We don’t have any experimental data on anti-Helium production to rely on

�¯Ad3N

¯A

d3k¯A=

✓4

3p3coal

◆A�1

�p

d3Np

d3kp

����kp=kA/A

!A Roughly a 10-3 suppression for each additional nucleon

The coalescence momentum is now a completely free parameter

0.1 1 10 100

DM DM → bb mDM = 40 GeV

T [GeV/n]

φ [(m

2 s sr

(GeV

/n))−

1 ]

AMS−02 reach

BESS excluded

PAMELA excluded

AMS−01 excluded

pcoal = 195 MeV

pcoal = 300 MeVpcoal = 300 MeVpcoal = 300 MeVpcoal = 300 MeVpcoal = 300 MeVpcoal = 300 MeVpcoal = 300 MeVpcoal = 300 MeVpcoal = 300 MeV

bkg

10−16

10−14

10−12

10−10

10−8

10−6

10−4

10−2

<σv> = 3x10−26 cm3s−1

EIN profile

0.1 1 10 100

DM DM → W+W− mDM = 1000 GeV

T [GeV/n]φ

[(m2 s

sr (G

eV/n

))−1 ]

AMS−02 reach

BESS excluded

PAMELA excluded

AMS−01 excluded

pcoal = 195 MeV

pcoal = 300 MeV

pcoal = 300 MeV

pcoal = 300 MeV

pcoal = 300 MeV

pcoal = 300 MeV

pcoal = 300 MeV

pcoal = 300 MeV

pcoal = 300 MeV

pcoal = 300 MeVpcoal = 600 MeV

pcoal = 600 MeV

pcoal = 600 MeV

pcoal = 600 MeV

pcoal = 600 MeV

pcoal = 600 MeV

pcoal = 600 MeV

pcoal = 600 MeV

pcoal = 600 MeV

bkg

10−16

10−14

10−12

10−10

10−8

10−6

10−4

10−2

<σv> = 3x10−25 cm3s−1

EIN profile

Prospects for detection are rather weak, unless the coalescence momentum is really large (~600 MeV)

What about antihelium?

A different (but complementary) channel : positrons

What’s interesting in positrons? Some years ago we had a scent of dark matter (the PAMELA excess, recently confirmed by AMS-02)

Hektor et al. Phys.Lett. B728 (2014)

An interpretation in terms of Dark Matter is possible and exciting (even

if not so easy for classical DM models)

However, it seems that also boring astrophysics can do the job (or at least can be seen as an irreducible background for DM interpretations)

Two words on astrophysical primary sources

A positron flux can be produced by both Pulsar Wind Nebulae and Supernova Remnants

Q(E) = Q0

✓E

E0

◆��

exp

✓� E

Ec

◆

★ The spectral index is estimated from radio observations: it’s expected to be [2 ± 0.3] for SNRs and [1.65±0.35] for PWNe

★ is the cut-off energy. It’s expected (from gamma-ray observations) to lay in the TeV range but, at present, it’s not so well known

★ is the normalization; it’s related to the fraction of the total emitted energy that goes into e± pairs (the efficiency for PWNe is expected to be ~few %, while for SNRs should be [10-4,10-5])

Ec

�

Q0

Our results We performed a detailed modeling of the astrophysical contribution from the known population of SNRs and PWNe to the four observables measured by

AMS-02

10-2

10-1

100

100 101 102 103

e+/(

e++

e- )

E [GeV]

e+/(e++e-)

3! bandII

PWNTOT

AMS-02PAMELA

FERMIAMS-01

HEATCAPRICE

10-3

10-2

10-1

100 101 102 103 104

E3 !

[G

eV

2 c

m-2

s-2

sr-1

]

E [GeV]

e-

3" bandII

PWNSNR local

SNR d > 3 kpcTOT

AMS-02PAMELA

FERMIAMS-01

HEATCAPRICE

10-4

10-3

10-2

100 101 102 103 104

E3 !

[G

eV

2 c

m-2

s-2

sr-1

]

E [GeV]

e+

3" bandII

PWNTOT

AMS-02PAMELA

FERMIHEAT

CAPRICE

10-3

10-2

10-1

100 101 102 103 104

E3 !

[G

eV

2 c

m-2

s-2

sr-1

]

E [GeV]

e+ + e-

3" bandII

PWNSNR local

SNR d > 3 kpcTOT

AMS-02FERMI

ATICHEAT

CAPRICEBETSHESS

Positrons fraction

Positrons flux Positrons + electrons flux

Electrons flux

Our results

10-2

10-1

100

100 101 102 103

e+/(

e++

e- )

E [GeV]

e+/(e++e-)

3! bandII

PWNTOT

AMS-02PAMELA

FERMIAMS-01

HEATCAPRICE

10-3

10-2

10-1

100 101 102 103 104

E3 !

[G

eV

2 c

m-2

s-2

sr-1

]

E [GeV]

e-

3" bandII

PWNSNR local

SNR d > 3 kpcTOT

AMS-02PAMELA

FERMIAMS-01

HEATCAPRICE

10-4

10-3

10-2

100 101 102 103 104

E3 !

[G

eV

2 c

m-2

s-2

sr-1

]

E [GeV]

e+

3" bandII

PWNTOT

AMS-02PAMELA

FERMIHEAT

CAPRICE

10-3

10-2

10-1

100 101 102 103 104

E3 !

[G

eV

2 c

m-2

s-2

sr-1

]

E [GeV]

e+ + e-

3" bandII

PWNSNR local

SNR d > 3 kpcTOT

AMS-02FERMI

ATICHEAT

CAPRICEBETSHESS

Positrons fraction

Positrons flux Positrons + electrons flux

Electrons flux

𝜸PWN = 1.9 ± 0.03

𝜸SNR = 2.382 ± 0.004

Q0,SNR = (2.748 ± 0.027)1050 GeV-1

𝜼PWN = 0.0320 ± 0.0016

We performed a detailed modeling of the astrophysical contribution from the known population of SNRs and PWNe to the four observables measured by

AMS-02

Our results

★ We found out that astrophysical sources (i.e. SNRs and PWNe), with reasonable parameters, can reproduce quite well the observables measured by the AMS-02 experiment

★ Thus, one in principle has no need to include Dark Matter in the analysis; however, since the astrophysical signal is not really well known some space is still there

★ Some further measurements (higher energy fluxies, anisotropy of the signal) can help to solve the positrons puzzle

Main Conclusions ★ Different channels carry different messages and implications for dark

matter searches

★ Antiprotons represent today the most powerful channel to constrain dark matter particles that annihilate (or decay) into hadrons. Soon-to-come AMS-02 data will likely close some windows in BSM physics

★ Despite antiprotons bounds, anti-deuterons are still a possible and very promising smoking-gun for the prensent and the near future. Antihelium is currently out of reach but can in principle be a promising signature for the future (with the caveat that we need more informations on the production mechanisms).

★ Positrons are still a puzzle. An excess is clearly there but at least a fraction of it is likely due to standard astrophysical sources. To draw meaningful conclusions about dark matter is not an easy task

★ Analysis of the constraints on DM properties that arise from the AMS-02 measurements (with M. Di Mauro, F. Donato, N.Fornengo and R. Lineros)

★ Determination of the flux of neutrinos generated by cosmic-rays interactions in the solar corona (with M. Cirelli, N. Fornengo and M. Piibeleht)

★ Study of the properties of the point sources in the gamma-ray sky with a pixel count analysis (with A.Cuoco, N. Fornengo, M. Regis and H.S. Zechlin)

★ Constraints on two-Higgs-doublet models from LHC and DM detection experiments (with M.S. Boucenna and N. Fornengo)

Ongoing Projects

• 4 papers: ‣ Interpretation of AMS-02 electrons and positrons data, M. Di Mauro, F. Donato,

N. Fornengo, R.Lineros, A. Vittino, [arXiv:1402.0321]

‣ Anti-helium from Dark Matter annihilations, M. Cirelli, N.Fornengo, M.Taoso, A.Vittino, [arXiv:1401.4017]

‣ Constraints on particle dark matter from cosmic rays antiprotons N.Fornengo, L.Maccione, A.Vittino [arXiv:1312.3579]

‣ Dark matter searches with cosmic antideuterons: status and perspectives N.Fornengo,L.Maccione, A.Vittino, JCAP 09 (2013) 31, [arXiv:1306.4171]

• 2 proceedings: ‣ Perspectives of dark matter searches with antideuterons AV, N.Fornengo,

L.Maccione [arXiv:1308.4848] To appear in NIM A

‣ The role of antiprotons and antideuterons in DM indirect detection M. AV,

N.Fornengo, L.Maccione to appear in World Scientific

My PhD - I

• 5 talks in international Conferences and workshops: ‣ APS meeting, Paris December 16-18 2013

‣ New perspectives in DM, Lyon October 22-26 2013

‣ 14th ICATPP Conference, Como September 23-27 2013

‣ 13th RICAP Conference, Roma May 22-24 2013

‣ IDAPP two-days meeting, Ferrara October 29-31 2012

• 1 invited seminar: ‣ IFIC (Valencia), May 14 2013

• 3 summer schools attended:

‣ ISAPP 2013, Djuronaset (Stockholm) July 29 - August 6 2013

‣ UNILHC, IFIC (Valencia) August 26 - September 5 2012

‣ ISAPP 2012, APC (Paris), July 2-13 2012

My PhD - II