Chapter 9: Cosmic Rays

Chapter 9 Cosmic Rays

9.1 Composition and energy distribution Cosmic rays can be broadly deﬁned as the massive particles, photons (γ rays, X-rays, ultraviolet and infrared radiation, ...), neutrinos, and exotics (WIMPS, axions,...) striking the earth. The primary cosmic rays are those entering the upper atmosphere, the cosmic rays of the interstellar medium. Secondary cosmic rays are those produced by the interactions of the primary rays in the atmosphere or in the earth. Also products of cosmic ray interactions in the interstellar medium (e.g., spallation products from cosmic ray - cosmic ray collisions) are also labeled as secondary cosmic rays. Cosmic rays can be of either galactic (including solar) or extragalactic origin.

If we conﬁne ourselves to the particle constituents (protons, nuclei, leptons), their motion in the galaxy has been roughly randomized by the galactic magnetic ﬁeld. (We will mention some exceptions to this below.) Thus they provide very little information about the direction of the source. The peak of the distribution in energy is in the range of 100 MeV - 1 GeV. The intensity of cosmic rays of energy 1 Gev/ nucleon or greater is about 1/cm2sec sr. An approximate formula is

α E 2 IN (E) 1.8 nucleons/cm sec sr (1) ∼ GeV where E is the energy per nucleon (rest and kinetc) and α -2.7. The energy density cor- responding to this is thus about 1 ev/cm3. This can be compared∼ to the energy density of stellar light of 0.3 eV/cm3.

The principal components of the (primary) cosmic rays are shown Figure 1. This abundance distribution is approximately independent of energy, at least over the dominant energy range of 10 MeV/nucleon through several GeV/nucleon. By mass about 79% of nucleons in cosmic rays are free protons, and about 80% of the remaining nucleons are bound in helium. The composition has been measured by instruments mounted on balloons, satellites, and spacecraft. Figure 2 shows the chemical distribution of the elements in our solar system diﬀers from that of the cosmic rays in some remarkable ways. The most dramatic diﬀerence is an enormous enrichment in the cosmic ray abundances for the elements Li/Be/B. Note also that there is enrichment is even Z elements relative to odd Z, when normalized to solar system abundances. Finally the cosmic rays are enriched in the heaviest elements relative to H and He.

As can be seen from Figure 2, many elements heavier than the iron group have been mea- 5 sured with typical abundances of 10− relative to iron. Much of this information was gained from satellite and spacecraft measurements over the last decade. Some of the conclusions: 1) Abundances of even Z elements with 30 < Z < 60 are in reasonable agreement with solar system abundances. ∼ ∼

1 20. Cosmic rays 3

0.1 He 1 − 10−2 H C

10−3

Fe −4 sr s MeV/nucleon) 10 2

10−5

10−6 Differential ﬂux (m

10−7

10−8

10−9 10102 103 104 105 106 107 Kinetic energy (MeV/nucleon)

Figure 1: Major components of the primary cosmic rays (from Simpson).

Figure 20.1: Major components of2 the primary cosmic radiation (from Ref. 1).

Most measurements are made at ground level or near the top of the atmosphere, but there are also measurements of muons and electrons from airplanes and balloons. + Fig. 20.3 includes recent measurements of negative muons [3,13,14,15]. Since µ (µ−)are produced in association with νµ(νµ), the measurement of muons near the maximum of the intensity curve for the parent pions serves to calibrate the atmospheric νµ beam [16].

June 14, 2000 10:39 2.1. COMPOSITION 11

FigureFigure 2.1: 2: Relative The abundances abundance of cosmicof elements rays inare cosmic compared rays with and solarin the abundances. solar system.

3 Z Element Relative Abundance 1 H 485 2 He 26 3-5 Li-B 0.40 6-8 C-O 2.20 9-10 F-Ne 0.30 11-12 Na-Mg 0.22 13-14 Al-Si 0.19 15-16 P-S 0.03 17-18 Cl-Ar 0.01 19-10 K-Ca 0.02 21-25 Sc-Mn 0.05 26-28 Fe-Ni 0.12

Table 1: Relative abundances of cosmic ray nuclei at 10.6 GeV/nucleon normalized to oxy- 6 2 gen (=1). The oxygen ﬂux at kinetic energy 10.6 GeV/nucleon is 3.26 10− /cm sec sr 1 × (GeV/nucleon)− . From Boezio et al. and Engelmann et al., as quoted by Gaisser and Stanev.

2) In the region 62 < Z < 80, which includes the platinum-lead region, abundances are enhanced relative to solar∼ by∼ about a factor of two. This suggests an enhancement in r-process elements, which dominate this mass region. This is consistent with r-process scenarios we have discussed (e.g., if the r-process site is core-collapse supernovae, then one would expect enrichment in r-process nuclei as supernovae are also believed to be the primary acceleration mechanism for lower energy cosmic rays). However, one also needs to consider whether heavier nuclei with larger charges might be confined in the galaxy for a longer time, another effect that could differentiate heavy from light nuclei.

Galactic cosmic rays are fully ionized: the acceleration mechanisms fully strip the ions. Cos- mic rays also have an antimatter component, as measured in the Space Shuttle Discovery AMS (Alpha Magnetic Spectrometer) experiment. The AMS detected about 200 antipro- tons above 1 GeV, generally attributed to nuclear collisions of CR particles with interstellar matter.

The energy distribution of cosmic rays from about 1010eV to about 1015eV is smooth, with a power-law distribution particles/cm2 sec MeV/n Es (2) ∝ with s between -1.6 and -1.7. However there is a break or “knee” in the curve at about 1015 eV. The slope sharpens above this knee (see Figure 3), falling with s ranging from -2.0 to -2.2, eventually steeping to an exponent of above -2.7. The knee is generally is attributed to the fact that supernova acceleration of cosmic rays is limited to about this energy. This would argue that the cosmic rays above this energy either have a diﬀerent origin, or were

4 cr-knee.gif (GIF Image, 468x333 pixels) http://imagine.gsfc.nasa.gov/docs/features/topics/snr_group/images...

Figure 3: Energy spectrum of cosmic rays, showing a kink or ”knee” near 1015 eV.

further accelerated after production. While this break is commonly attributed to the inabil- ity of supernova shocks to accelerate particles to energies beyond the knee, the sharpness of the knee has troubled many of the experts: in is difficult to find natural models that produce such a defined break.

There is an additional, somewhat less distinct feature at an energy of about 1019 eV, termed the ankle. Characteristics of this high-energy feature will be discussed later.

9.2 Propagation and origin The most commonly used toy model for galactic cosmic rays is called the ”leaky box” model. It assumes that the cosmic rays are confined within the galactic disk, where the mass density is high, but with some gradual leaking out of the disk. The confining force is the galactic 7 magnetic field, which is on the order of 10− Gauss. A relativistic particle moving in a magnetic field executes a helical path. Using the relation between the momentum p, field B, and magnetic radius R, 4 p (MeV/c) = 3 10− BR(gauss cm) (3) ⊥ × one sees that a 1014 eV proton would have a radius of 3 1018 cm, or 1pc, which is much × 5

1 of 1 03/10/2004 10:17 AM less than the distance to the Crab nebulae, a potential accelerator for cosmic rays relatively near earth. Thus cosmic rays from the knee and below will have no memory of their origin when they reach earth. A 1018 eV proton corresponds to 10 kpc, about the galactic radius. Clearly any cosmic ray much above this energy could be presumed to be extragalactic, unless it can be associated with a local source, which would be possible as the direction of such cosmic rays would point back to their origin. For example, a cosmic ray of energy 1020 eV would be minimally perturbed by the galactic magnetic ﬁeld and thus would ”point back” to the extragalactic source from which the cosmic ray originated.

The leaky box models cosmic ray production in the galaxy, cosmic ray trapping by magnetic ﬁelds and eventual escape, and cosmic ray interactions in the intragalactic medium. This model does a good job in explaining the energy-dependence of the life of cosmic rays (more on this later). But others have argued for other models, including closed models where cosmic rays are fully conﬁned, then explaining isotope lifetimes (see below) through devices such as a combination of a few nearby and many distant cosmic ray sources.

The conventional explanation for the most dramatic isotopic anomaly in the cosmic rays, the enrichment in Li/Be/B by about six orders of magnitude, is that these isotopes are produced in the interstellar medium when accelerated protons collide with C, N, and O. We mentioned this process earlier. The enrichment of odd-A nuclei (these also tend to be relatively rare in their solar distribution since stellar processes tend to favor production of more stable even A nuclei) is also often attributed to spallation reactions off more abundant even-A nuclei. These associations immediately lead to some interesting physics conclusions because, from the known density of cosmic rays (at least in the earth’s vicinity) and from known spallation cross sections, one can estimate the amount of material through which a typical cosmic ray propagates. Although the estimates are model dependent - and probably not sufficiently interesting to go through in detail - the resulting values for the effective thickness are typically 4 - 6 g/cm2. Now the mass density within intragalactic space is 3 24 3 about 1 proton/cm , or about 1.7 10− g/cm . Thus taking a velocity of c, we can crudely estimate the cosmic ray lifetime ·

24 3 10 2 1.7 10− g/cm (3 10 cm/sec) t = (4 6)g/cm (4) × × · × − So this gives t 3 106y (5) ∼ · This calculation assumes an average galactic mass density that is not known by direct measurement. Thus it is nice that a more direct estimate of the galactic cosmic ray lifetime is provided by cosmic ray radioactive isotopes. The right chronometer is one that has a lifetime in the ballpark of the estimate above. 10Be, with a lifetime of 1.51 106 y, is thus quite suitable. It is a cosmic ray spallation product: this guarantees that it× is born as a cosmic ray. Its abundance can be normalized to those of the other, stable Li/Be/B isotopes: the spallation cross sections are known. Thus the absence of 10Be in the cosmic ray spectrum would indicate that the typical cosmic ray lifetime is much larger than 1.51 106 y. The · 6 survival probability should also depend on the 10Be energy, due to time dilation eﬀects. One observes a reduction in 10Be to about (0.2-0.3) of its expected instantaneous production, relative to other Li/Be/B isotopes. From this one concludes

t (2 3) 107years ∼ − · This suggests that the mass density estimate used above (in our ﬁrst calculation) may have been too high by a factor of 5-10.

In modeling the origin of cosmic rays, the first conclusion, given their richness in metals, is that must come from highly evolved stars such as those that undergo supernovae. We have already noted the abundance of r-process nuclei, which could be taken as a “smoking gun” of supernova dominance, for those who accept that supernovae are the r-process site. However this is clearly not the full picture. Studies of the isotopic composition as the knee is approached shows that the composition changes: the spectrum of protons becomes notice- ably steeper in energy, while the iron group elements do not show such a dramatic change. This is qualitatively consistent with the notion that the galactic ”accelerator” producing cosmic rays tops out at some maximum confining field strength. Because the acceleration likely scales as ZB, where Z is the charge, the very highest energies should be dominated by the largest Zs, the nuclei.

Above the knee - at energies above 1016 eV the galactic magnetic ﬁeld is too weak to ap- preciably trap particles. Thus it is probable that at these high energies the character of the cosmic rays changes from primarily galactic to primarily extragalactic: the trapping that enhances the abundances of lower-energy galactic cosmic rays would not enhance very high energy ones.

The cosmic rays appear to be approximately isotropic – once one gets above about 50 GeV to escape local magnetic eﬀects. One exception is a small anisotropy measured by AGASA, excess events at about 1018 eV pointing back to the galactic center - which generated the speculation that these are neutrons, which because of time dilation can reach the solar system once they reach 1018 eV. If this speculation were true, it would argue for a central galactic accelerator that (most likely) is capable of accelerating ions up to 1018 eV per nucleon. Re- cent results from Pierre Auger ﬁnd no evidence for the AGASA anomaly.

9.3 Energetics and origin of galactic cosmic rays We now have the basic information needed to calculate the energetics of galactic cosmic rays. If we take the galactic radius as 10 kpc, we have a volume of 1068 cm3. We noted that the energy density of cosmic rays is about 1 eV/cm3. Thus the cosmic ray energy content of the galaxy is about 1068 eV. We have argued above that the lifetime of cosmic rays might be about 107 years. Thus the energy production in cosmic rays must be about

1068eV/107y 3 1054eV/sec 5 1042ergs/sec. (6) ∼ × ∼ ×

7 We mentioned in passing above that there are cosmic ray models that, unlike the leaky box, confine cosmic rays longer – and thus need to contrive localize sources to account for Be, etc. Thus our energetics calculation could be modified, replacing the leaky box model with one where the cosmic rays were effectively confined over times up to the age of the galaxy. This would extend the dwell times by up to three orders of magnitude, and thus reduce the energetics requirement proportionally. So that in principle would lower the galactic energetic requirement to about 1040 ergs/sec. The result is a very generous range for the necessary energetics, requiring sources capable of generating between 1040 and 5 1042 ergs/sec. · One can quickly show that stellar winds – e.g., the flares our sun produces – are insufficient energetically to produce energy at this rate, even if integrated over the 1011 potential sources in the galaxy – the integrated energy in stellar flares falls short of the lower bound above by at least one and perhaps two orders of magnitude. But an interesting source is supernovae – they eject large quantities of material, that material includes both protons and nuclei, and the shock wave and associated fields are an acceleration mechanism.

We have seen that the explosion energy of a supernova is about 1051 ergs, and the estimated rate of Type II supernovae is 1/30 years. Additional contributions would come from rarer classes of core-collapse supernovae, SNIb SNIc. Thus the energy production rate is about 1042 ergs/sec – close to that estimated above. There are many arguments based on the chemical patterns of cosmic rays that support core-collapse supernovae as a principal mechanism for generating and accelerating cosmic rays. Novae are powered by the accretion onto a white dwarf. The matter is sucked from a large companion star, perhaps a red giant or main sequence star, under conditions that allow periodic thermonuclear runaway reactions on the white dwarf surface. Novae are another plausible contributor. Novae outbursts within our 4 galaxy occur with a frequency of about 100/year with an energy output about 10− that of the ejecta from a core-collapse supernova. Thus the integrated nova output is similar to that of galactic supernovae.

9.4 Solar-system, terrestrial, and atmospheric environmental effects Cosmic rays measured on earth reflect not only the galactic inventory of cosmic rays, but also the effect of the local environment of our solar system, the earth, and the atmosphere. The sun’s heliosphere – the region of space altered by the solar wind – extends to about 200 AU. The magnetic structure of the heliosphere shields the region against energetic charged particles. This is apparent experimentally from the correlations between induced cosmic ray activity measured on earth and the 11-year solar cycle, as shown in Fig. 4. This correlation is reflected in changes in the chemistry of the atmosphere, such as 14C production, an important radio-isotope for dating. 14C is a product of cosmic ray interactions with O in the atmosphere, and the production is anticorrelated with the shielding, and thus with the intensity of solar activity. Interesting, 14C production also appears to correlate with various long-term climate anomalies of the past 1000 years, such as the Maunder Minimum (14C high) and the 12th-Century Maximum (low), leading to some interesting speculations about

8 Figure 4: The top curve is the cosmic ray flux from the neutron monitor in Climax, Colorado (1953 – 1996). The middle curve is the annual mean variation in the cosmic ray flux as measured by ionization chambers (1937 – 1994). The bottom curve is the relative sunspot number. While there is a clear solar cycle modulation of the cosmic ray flux, the amplitudes are not well correlated. From Svensnak.

the solar role in climate change.

There are also more localized effects due to the earth’s magnetic field, a dipolar field that extends from the magnetic poles and thus is roughly parallel to the earth surface at the equator (Fig. 5). If we envision a cosmic ray proton come into the earth at the equation and use 1000 km as a rough trapping radius, then our previous formula

4 p (MeV/c) = 3 10− BR(gauss cm) ⊥ × for B=0.3 Gauss would give a p E 10 GeV as the energy of the trapped proton. Thus this defines an energy below which∼ charged∼ cosmic rays would not penetrate to the surface, but instead be sharply deflected. This is the reason that so many cosmic ray balloon experiments are done at the poles, where the magnetic fields are weaker away from the earth and more perpendicular to the surface: one can then more easily sample cosmic rays character- istic of the environment away from the immediate terrestrial neighborhood.

The trapping of lower-energy cosmic rays by the earth’s magnetic ﬁeld in also part of the

9 Figure 5: Rough representation of the earth’s magnetic ﬁeld, which has an average value near the surface of about 0.3 Gauss.

third effect of the local environment, interactions of the cosmic rays with the atmosphere. The trapping leads to a larger time in the vicinity of upper atmosphere and thus enhances nuclear interactions than can produce all sorts of cosmic ray secondaries. Such secondaries are also produced by higher energy cosmic rays that are not perturbed significantly by the magnetic field, but do interact in the upper atmosphere. The secondaries include pions, muons, new nuclei created by spallation, electrons and positrons, neutrinos, and the high- energy showers produced by very energetic cosmic ray interactions.

The density proﬁle of the atmosphere is approximately exponential with a scale height of about 7.6 km, r/7600m 3 ρ(r) = 1.205 e− kg/m . (7) × Thus the total amount of matter that a cosmic ray encounters on approaching the earth grows exponentially with decreasing altitude.

9.5 Muons and neutrinos at the earth’s surface and below. The muons and neutrinos, the most penetrating components of the cosmic ray secondaries, result from strong interaction mechanisms that produce pions and kaons through decay chains such as π+ µ+ + ν → µ µ+ e+ + ν +ν ¯ (8) → e µ as we have discussed before. Figure 6 illustrates some important features of these and other cosmic ray secondaries. The ﬁgure was generated by cascade code calculations that begin

10 4 20. Cosmic rays ] 2 500 electrons + positrons GeV 300 –1

sr 200 –1 s

–2 100 [m 70 dE / 50 dN 3 E 10 100 1000 E [GeV]

Figure 20.2: Diﬀerential spectrum of electrons plus positrons multiplied by E3 (data summary from Ref. 7). The dashed line shows the proton spectrum multiplied by 0.01.

Altitude (km) 15 10 5 3 2 1 0 10000

1000

] _ –1 νµ + νµ

sr + − 100 µ + µ –1 s –2 10 p + n

1 e+ + e−

Vertical ﬂux [m π+ + π− 0.1

0.01 0 200 400 600 800 1000 Atmospheric depth [g cm–2]

Figure 6: The solid lines show vertical fluxes of cosmic rays in the atmosphere with energy above 1 GeV, as estimated from parameterized initial nucleon primary fluxes. The points are from measurements of µ−. From Gaisser and Stanev. Figure 20.3: Vertical fluxes of cosmic rays in the atmosphere with E>1GeV estimated from the nucleon flux of Eq. (20.2). The points show measurements of negative muons with Eµ > 1 GeV [3,13,14,15].11

Because muons typically lose almost two GeV in passing through the atmosphere, the

June 14, 2000 10:39 6 20. Cosmic rays 0.2 ] 1.7 )

c 0.1

(GeV/ 0.05 –1 sr –1 s 0.02 –2

[cm 0.01 µ

0.005 dN/dp 2.7 µ p 0.002 1 10 100 1000 pµ (GeV/c)

Figure 7: The spectrum of muons at the earth’s surface at θ = 0 degrees (solid points) and

75 degrees (open diamonds). From Gaisser and Stanev.

N Figure 20.4: Spectrum of muons at θ =0◦ ( [18], [20], [21], [22], X,+ [23]), with the primary cosmic ray spectrum at the top of the atmosphere, then propagate that

and θ =75 [24]). spectrum and◦ the associated secondaries through the atmosphere, taking into account the various reactions that produce new particles and degrade their energies. In contrast to the Figurevery large 20.4 flux/energy shows the losses muon encountered energy for spectrum other secondaries, at sea levelthe muons for two and neutrinosangles. At large angles lowgenerally energy penetrate muons decay the atmosphere before well. reaching The atmosphere the surface (and and earth) high are transparent energy pions to decay before theyneutrinos, interact, while thus muons the typically average lose 2muon GeV in energy ionization increases. energy while An transiting approximate the atmo- extrapolation sphere. formula valid when muon decay is negligible (Eµ > 100/ cos θ GeV) and the curvature of theConsequently Earth can bemuons neglected dominate ( theθ< charged-particle70◦)is spectrum at the earth’s surface. They are the most easily measured component of the cosmic rays and also the primary background in a lot of neutrino experiments – the reason such neutrino experiments must go deep underground. The effects of energy loss0.14 areE illustrated2.7 in Figure 7, where the vertical flux of muons at the earth’sdN surfaceµ is compared withµ− the much harder spectrum that results for 2 muons impinging thedE earthµ at≈ 75cm degreesssrGeV from vertical. The latter must penetrate much more matter, which then hardens the spectrum. 1 0.054 + , (20.5) × 1.1Eµcos θ 1.1Eµcos θ  1+ 12 1+   115 GeV 850 GeV  where the two terms give the contribution of pions and charged kaons. Eq. (20.5) neglects a small contribution from charm and heavier flavors which is negligible except at very high energy [27]. + The muon charge ratio reflects the excess of π over π− in the forward fragmentation region of proton initiated interactions together with the fact that there are more protons than neutrons in the primary spectrum. The charge ratio is between 1.1 and 1.4 from 1 GeV to 100 GeV [18,23]. Below 1 GeV there is a systematic dependence on location due to geomagnetic effects. [23]

June 14, 2000 10:39 + Muons lose energy by ionization, bremsstrahlung, production of e e− pairs, photonuclear reactions like Compton production, etc. Their range R in standard rock (Z 11, ρ 2.65 g/cm3) is given in the table. R is given in meters-of-water-equivalent (mwe),∼ or 102 g/cm∼ 2. The energy loss is given in terms of α, dE µ = α. (9) − dx The table shows that the energy loss per g/cm2 increases sharply between 100 GeV and 1 TeV, a result of direct bremsstrahlung, pair production, and nuclear reactions taking over from ionization losses, which dominate at lower energy. The former grow linearly with the muon energy. Consistent with the table, a rough rule of thumb is that the net ﬂux of muons underground decreases by an order of magnitude for every 1500 mwe, or about 500m of standard rock.

Some of you may know that there has been great interest in establishing underground laboratories in the U.S. for sensitive dark matter, neutrino, and nucleon decay experiments. While Europe has facilities like Gran Sasso and Frejus, Russia has Baksan, and Japan has Kamioka, the deepest current site in the US is the Soudan Laboratory, operated by the University of Minnesota, at a depth of 710 m. Access underground at Soudan is challenging because of the small lift that provides vertical access to that site.

One wonderful development for North America, however, is SNOLab, an expansion of the SNO experimental area in the Sudbury Mine, including both additional underground rooms and a signiﬁcant surface support facility. This is the world’s deepest underground laboratory, two kilometers underground and providing 6010 mwe in overburden (more than two kilometers of rock overburden). Furthermore, as the nickel deposits that interest the mine’s owner have so far continued to be increasingly rich at depth, it quite conceivable that the mine could some days reach depths of three kilometers or more. As discussed below, this would approach the ultimate depth for an underground laboratory, the point at which atmospheric neutrinos become an irreducible background.

Some next-generation experiments could require 5000 mwe or more of overburden. Examples include certain next-generation dark matter, solar neutrino, and double beta decay eﬀorts. Overburden is often important in experiments focused on measuring low-energy events: penetrating muons themselves can often be vetoed by ”turning oﬀ” the detector after a muon is detected. (Kamiokande, for example, has about a 10% deadtime due to this technique for background suppression.) But muons, in addition to producing prompt secondaries like knock-out neutrons, can also activate nuclei in the detector that, some time later, decay to produce a signal. In cases where there is a long delay before the induced activity is detected, conventional vetoing techniques can lead to an unacceptable deadtime. In such cases there is no alternative but great depth to remove the initiating muon events.

13 6 2 Eµ (GeV) R (1000 mwe) α (10− MeV cm /g)

10 50 2.19 100 410 2.74 1000 2450 6.60 10000 6090 46.43

Table 2: From Gaisser and Stanev: the calculated ranges R and energy losses (per gm/cm2) α are both sharply dependent on the muon energy, and lead to a systematic hardening of the muon spectrum as one goes underground.

In other cases, backgrounds may be controllable at moderate depth, making great depth less important than other factors, such as the economies that come from ease of access: the ability to ”drive-in” large equipment and eﬃciently operate large experiments may be the priority. Both Gran Sasso and Kamioka, the world’s premier underground laboratories, are drive-in facilities at moderate depth.

While Europe has established several underground laboratories (Gran Sasso, Canfranc, Fre- jus,...) and while Japan has steadily developed Kamioka, the US has lacked a comparable facility. Experiments were done at Homestake before the mine’s closing ten years ago, and the Soudan Mine (Minnesota) and WIPP (New Mexico) both currently host experiments. But a site oﬀering good and frequently experimental access and depth comparable to our better than Gran Sasso has been lacking. The NSF headed a ten-year eﬀort o create a Deep Underground Science and Engineering Laboratory in the US, settling early on Homestake as the preferred site, but very recently (and oddly) decided the operation of such a facility was not in its mandate. A DOE-appointed committee is now trying to sort things out, though the cost of developing Homestake is a major obstacle, given how tight science budgets are expected to be in 2012 and beyond, and given DOE obligations to existing as well as new, previously approved projects.

The relevance to the present discussion is the importance of being able to calculate the muon flux at depth, folding the energy and angular spectrum at the surface with the site topog- raphy. An example shown here were the investigations of University of Washington physics grad student Kregg Philpott to determine the effective depth of a tunnel in the Cascades, which several years ago was a DUSEL candidate. The resulting contours of the muon flux under the Steven Pass’s Cowboy and Big Chief Mountains is shown in Fig. 8, for a horizontal slice at the Pioneer tunnel’s elevation. This tunnel proves to be at just the right place for convenient access to the point of greatest overburden, under Cowboy Mt. The resulting flux, 7 2 7 2 2.07 10− /cm sec, is about 90% that at Kamioka (2.28 10− /cm sec). . × × While we have discussed going deep underground to find an environment sufficiently clean to do neutrino physics, there are some absolute limits to what can be achieved. Atmospheric

14 Figure 8: The cosmic ray muon ﬂux under Cowboy Mt., Stevens Pass, at the 743 m elevation of the Pioneer Tunnel. Calculation done by Kregg Philpott.

15 neutrinos provide a underground source of high-energy muons and electrons that does not attenuate: these are the muons that the energetic atmospheric neutrinos produce as secondaries. As Figure 9 shows, at a depth of about 10 km mwe (3.5 km of rock) these become a constant background that is not reducible by going further underground. It also follows that upward going muons (or electrons, though they have a shorter range) in deep detectors thus can be used as a signal of neutrino reactions in the rock below. Figure 10 is a reminder of our earlier discussion describing how these neutrino ﬂuxes were exploited to probe neutrino oscillations, through the muons and electrons they produce in underground detectors.

Given that we have discussed so many aspects of the terrestrial neutrino flux, including the changes induced by neutrino oscillations, it is amusing to display the local neutrino spectrum in its entirety in Figure 11. The atmospheric neutrinos, despite their importance due to their high energies and larger cross sections in matter, do not appear on the graph because their fluxes are so low. Two sources we have not discussed in detail are the geoneutrinos, produced from terrestrial radioactivity and perhaps recently seen in KamLAND, and the thermal flux ( keV) of solar neutrinos of all flavors, generated in our sun from neutral current processes, analogous∼ to those we discussed (at higher energies) in the context of supernovae cooling.

9.6 The highest energy cosmic rays, the GZK cutoff, and Pierre Auger Some of the most curious observations in cosmic ray physics have to do with the highest energy cosmic rays. A number instruments, such as the Fly’s Eye and AGASA detectors, and most recently the Pierre Auger observatory, are designed with sensitivity to ultrahigh- energy cosmic rays. For example, the Fly’s Eye uses air fluorescence to detect UHE cosmic rays. An extensive air shower is generated when a primary cosmic ray interacts with the atmosphere. This is imaged using the fluorescence light produced by excitation of the nitrogen molecules by the secondaries in the extensive air shower. The shape of the shower allows the experimenters to reconstruct the energy of the primary. It also may provide some information on composition. The Fly’s Eye can provide stereo information on the developing shower because of detector arrays located 3.4 km apart.

The rare high energy events show some structure, particularly around 1019eV, where the spectrum flattens from a slope of about -3.0 to one of about -2.6. Some years ago there were claims that events were been seen above the so-called Greisen-Zatsepin-Kuzmin cutoff of about 1020 eV. At one time there were tens of such events extending up to about 1021 eV )(see Figs. 12 and 13). The GZK cutoff is explained below, the new results from Pierre Auger are described.

The cosmic medium is ﬁlled by background radiation of relic photons, left over from the back bang and noninteracting since the time electrons and nuclei recombined to form atoms. 3 Their typical energy is about 10− eV. We consider the propagation of a high energy proton through this medium.

16 20. Cosmic rays 9

10−6 ) 1

− −8 s 10 1 − sr 2 −

10−10

10−12 Vertical intensity (cm

10−14 1 10 100 Depth (km water equivalent)

Figure 9: Vertical muon intensity as a function of depth (in km.w.e.) The shaded region at depth represents neutrino-induced muons above 2 GeV. The upper line is for horizontal- induced muons, the lower one for vertically upward muons. From Gaisser and Stanev.

5 2

Figure 20.5: Vertical muon intensity vs depth (1 km.w.e. = 10 gcm− of standard

rock). The experimental data are from: : the compilations of Crouch [32], :

f v

Baksan [33], : LVD [34], : MACRO [35], : Frejus [36]. The shaded area at large depths represents neutrino induced muons of energy above 2 GeV. The upper 17 line is for horizontal neutrino-induced muons, the lower one for vertically upward muons.

Contained and semi-contained events reflect neutrinos in the sub-GeV to multi-GeV region where the product of increasing cross section and decreasing flux is maximum. In the GeV region the neutrino flux and its angular distribution depend on the geomagnetic location of the detector and, to a lesser extent, on the phase of the solar cycle. Naively, we expect νµ/νe = 2 from counting neutrinos of the two flavors coming from the chain of pion and muon decay. This ratio is only slightly modified by the details of the decay kinematics, but the fraction of electron neutrinos gradually decreases above a GeV as parent muons begin to reach the ground before decaying. Experimental measurements have to account for the ratio ofν/ν ¯ , which have cross sections different by a factor of 3 in this energy range. In addition, detectors generally have different efficiencies for detecting muon neutrinos and electron neutrinos which need to be accounted for in comparing

June 14, 2000 10:39 rmblw hog h at.Terslsaefi eywl by well very fit are neu results for The events of muon-like earth. tion the the through in below, depletion obser from sharp and a lines) but (blue events, predicted the between result agreement neutrino atmospheric lent SuperKamiokande The 1: Figure htpntaignurnspoueairdcbeflxo lcrn n un tdepth. at muons fact and electrons the of exploited flux which irreducible – a results produce neutrino neutrinos atmospheric penetrating SuperKamiokande that The 10: Figure ν µ → Number of Events ν 100 200 300 400 100 120 140 τ 20 40 60 80 0 0 siltoswt aia iig(e lines). (red mixing maximal with oscillations 1-.0 .1 1 05 05 1 0.5 0 -0.5 -1 1 0.5 0 -0.5 -1 1-.0 .1 1 05 05 1 0.5 0 -0.5 -1 1 0.5 0 -0.5 -1 ut-e -iemulti-GeV multi-GeV e-like u-e -iesub-GeV sub-GeV e-like cosine ofzenithangle 4 18 100 150 200 250 300 100 200 300 400 500 50 0 0 µ -like +PC e electron-like ved hwn excel- showing s µ -like rnscoming trinos h assump- the Figure 11: The fluence of neutrinos at earth, ranging from cosmic background neutrinos to those produced by high-energy cosmic ray interactions in the atmosphere. From Haxton and FIG.Lin. 4. Natural neutrino sources. The terrestrialν ¯e flux and continuous flux of extragalactic supernova neutrinos of all flavors are from Krauss et al. [12]. The solar (fusion) νe flux is the standard solar result of Bahcall et al. [11]. The thermal solar neutrinos are for a single flavor. spectrum contains a great deal of information on the temperature distribution within the sun. These remarks are made because the most likely opportunity for measuring the thermal neutrino spectrum is a process that depends on flux density, not on total flux, and which samples that flux at a precise energy, the19 resonant reaction

ν¯ + e− +(A, Z) (A, Z 1). (23) e → − This reaction has been discussed previously in connection with terrestrialν ¯e sources [12]. Cross sections can be large in high Z atoms, where the electron overlap with the nucleus is favorable. Because nuclear level widths are very narrow, this process samples theν ¯e flux density at a discrete energy. The are several possible candidate transitions with energies between 2 and 20 keV. (One that has been studied in connection with neutrino mass measurements is the decay of long-lived 163Ho to 163Dy, which has a positive q-value of less than 3 keV: either a neutrino mass orν ¯e inducement of electron capture alters the atomic orbits that participate in the capture.) The heavy-flavor neutrino flux also contains interesting information: if the existence of this flux were established, it would immediately impose kinematic mass limits of 1 keV on ∼ the νµ and ντ . Unfortunately there is no obvious possibility for measuring these species. The problem could well prove as difficult as in the case of the cosmic microwave neutrinos, where existing experimental bounds exceed the expected flux by about 15 orders of magnitude [14].

11 14 20. Cosmic rays 10 ] 1.7 GeV –1

sr 1 –1 s –2 [cm

0.1 dN/dE 2.7 E

1011 1012 1013 1014 1015 1016 1017 1018 1019 1020 1021 E [eV/nucleus] 16 20. Cosmic rays

1026 ]

H O Figure 20.9: The all-particle spectrum: [52], N [54], [55], [56], [57],

eV 4 + [58], X–1 [59], [60]. References for the high energy portion of the spectrum are 25 given in Fig. sr 10 20.10. –1 s of the muons–2 (ρµ) is larger and depends on the transverse momenta of the muons at

production as [cm well as multiple scattering.

There are10 large24 ﬂuctuations in development from shower to shower, even for showers of the same energy and primary mass—especially for small showers, which are usually well dN/dE 3

past maximumE development when observed at the ground. Thus the shower size Ne and primary energy E0 are only related in an average sense, and even this relation depends on depth in the1023 atmosphere. One estimate of the relation is [52] 1018 1019 1020 1021 E [eV]6 6 0.9 E0 3.9 10 GeV (Ne/10 ) (20.13) ∼ × 14 17 2 for verticalFigure showers 12: From with about 10 ten

increasesFigureand the 20.10: 10 shower21 eV.Expanded An maximum expanded viewspectrum (on of average) the is highest also moves shown. energy down From portion Gaisser into theof and the atmosphere Stanev. cosmic-ray and the

H relationspectrum: between o [71]Ne (stereo),and E0 changes. [71] (monocular) At the maximum[72], of[73]. shower development, there are approximately 2/3 particles per GeV of primary energy. Detailed simulations and cross-calibrations20 between diﬀerent types of detectors are necessary to establish the primary energy spectrum from air-shower experiments [52,53]. Figure 20.9 shows the “all-particle” spectrum. In establishing this spectrum, eﬀorts have been made to minimize the dependence of the analysis on the primary composition. In the energy range above 1017 eV, the Fly’s Eye technique [71] is particularly useful because it can establish the primary energy in a model-independent way by observing most of the longitudinal development of each shower, from which E0 is obtained by integrating the energy deposition in the atmosphere.

June 14, 2000 10:39

June 14, 2000 10:39 Figure 13: A comparison of the HIRES and AGASA high-energy cosmic ray events. The latter data set tends to show more events at or above 1020 eV.

21 Let pµ = (, ~p) be the photon four-momentum. Then ~p = . Let Pµ = (ω, P~ ) be the proton four-momentum. We evaluate in the center-of-mass | |

µ µ 2 2 (P + p )(Pµ + pµ) = (ω + ) = CM (10) which we recognize as the square of the center-of-mass energy. Note that if CM exceeds

mπ + MN (11) then clearly the reaction γ + N π + N (12) → can occur, degrading the nucleon energy. But the center-of-mass energy is a Lorentz invariant quantity, so it can be evaluated in the laboratory frame

2 = M 2 + 2ω 2P~ ~p (13) CM N − · As the cosmic background photons are moving in all directions, we are free to maximize the RHS by taking cos θ 1. Noting that the incident nucleon is highly relativitistic ∼ − (max)2 M 2 + 4ω (14) CM ∼ N max Thus the requirement for photoproduction is CM > mπ + MN ∼ ⇒ m2 + 2M m M m ω > π N π N π ∼ 4 ∼ 2 2.7K 3 1020 eV (15) ∼ · Eγ !

This results in a mean free path for protons of about 108 light years for protons at 1020 eV or higher energies, a distance substantially smaller than the horizon. Thus if the∼ origin of such cosmic rays is all of extragalactic space, there should be a very sharp cutoff in the cosmic ray flux at about this energy. This is called the Greisen-Zatsepin-Kuzmin cutoff, and is shown in the figure. Yet the very high energy events from AGASA show no evidence for any such cutoff.

The AGASA events created a lot of controversy. Fairly mundane solutions were offered, such that these high energy events are exclusively nuclear (CNO elements and then iron), which can achieve higher energies at lower velocities. It has been suggested that the events are secondaries from the collisions of still higher energy neutrinos. However this requires new physics in that the Standard Model predicts a declining neutrino cross section above the Z mass that would not be sufficient to produce the needed event rate, it is thought. It was pointed out that ultrahigh energy cosmic rays both above and below the GZK cutoff could be of relatively local origin, thus evading the cutoff – but also raising new questions of what confines the cosmic rays if they are extragalactic but somehow produced primarily in a local

22 region about us. Perhaps we are somehow near some remarkable local source or sources: our position in the cosmos is special.

In fact the expect GZK cutoff involves a lot of physics. Pion production may be the key issue for protons, but in nuclei additional processes have to be considered. At a center-of-mass energy considerably below pion production theshold, nuclei can absorb (in their rest frame) a photon of energy 20 MeV, resulting in photodistintegration. This leads to a GZK cutoff for iron nuclei of ∼5 1020eV. ∼ · Very high energy cosmic ray interactions with atmospheric nuclei lead to an extensive air shower of secondary particles that, by the time the shower reaches the earth’s surface, may number in the billions and cover an area as large as 25 square kilometers. The event rate of very high energy collisions, above 1019 eV, is about 1 per 100 square kilometers per year. Because of this daunting number, the recent construction of the Pierre Auger Observatory was a major step forward for the field. This detector is sited in western Argentina and its surface sites are spread over an area of about 4000 square kilometers. The detector is a hybrid, with two detection schemes. One scheme uses 1600 water-tank detectors that form a grid with spacings of about 1.5 km. Energetic particles scattering in the detector produce Cerenkov light that is recorded by phototubes mounted on the tanks. So a shower covering up to 25 square kilometers might be recording simultaneously in 5-10 tanks. The energy of the incident particles can be estimated from the amount of light produce in the tanks, and then this energy can be extrapolated over the entire detection area (that is, one corrects for the small fraction of the surface that is covered by tanks) to determine the total shower energy. The second scheme uses the upper atmosphere as the detection medium. The charged particles in the shower interact with atmospheric nitrogen, leading to atomic excitation and subsequent emission of ultraviolet light. The resulting N fluorescence is recorded in Pierre Auger’s optical detectors – a grid of focusing mirrors to collect light and cameras that can then view the light at distances up to 15 km. The cameras can follow the path of high energy cosmic ray showers as the shower progresses through the atmosphere. Since all of the shower is viewed, a more accurate estimate of total shower energy can be made – provided atmospheric conditions allow the shower to be seen.

The Pierre Auger Observatory began operations in 2004. It has gathered signiﬁcant information on events with energies in the vicinity of the GZK cutoﬀ. Shower development depends on the nature of the collision, with the depth to which the shower develops sensitive to whether the incident cosmic ray was a proton or a nucleus like Fe. Variations in observed showers are compared to simulations – a tricky business because the most energetic cosmic ray collisions correspond to center-of-mass energies two orders of magnitude beyond those directly probed at the LHC/FermiLab and RHIC. Because one has the shower energy and a test of composition, one may be able to tell whether cosmic ray composition is changing with energy. Cross checks are possible. For example, shower-to-shower variations are predicted to decrease as the typical mass of the incident nucleus increases.

23 ] ] 2 QGSJET01 2 80 QGSJETII 850 Sibyll2.1 proton 70 proton

EPOSv1.99 ) [g/cm > [g/cm 60 max

max 800 50

750 RMS(X 40 30 700 iron 20 iron 650 10

18 19 0 10 10 1018 1019 E [eV] E [eV]

FIGURE 3. Xmax and RMS(Xmax) compared with air shower simulations [22] using different hadronic interaction models [23]. Figure 14:h Pierrei Auger data for shower depth (in g/cm2, left panel) and in the RMS variation in shower depth (right panel) both show trends that suggest the primary cosmic rayrange. spectrum For this purpose is trending the effective away upper from and primarily lower field protonsFD stations at energies within the around quoted uncertainties. 1018 eV to nuclei at of view limits are calculated for each event and Xmax is The measured values are displayed in Fig. 3. A fit of h i energiesmeasured near as a functionthe GZK of these cutoff. limits. From An example M. Unger of andXmax thedata PA with Collaboration, a constant elongation arXiv:1103.5857. rate does not de- h i 2 the Xmax dependence on the lower FOV limit is shown scribe our data (χ /Ndf=34.9/11), but using two slopes in Fig.h 2 fori data and simulated events. As can be seen, yields a satisfactory fit (χ2/Ndf=9.7/9) with an elonga- +35 2 18.24 0.05 the Xmax is asymptotically unbiased for events with a tion rate of (106 ) g/cm /decade below 10 eV h i 21 ± Thesufficiently Pierre deep Auger FOV data limit, of but Fig. it is systematically 14 show a too trendand from (24 protons3)g/cm−2/decade to nuclei above as this one energy. approaches If the prop- theshallow GZK when cutoff. the FOV This boundary is reminiscent starts cutting of the into similarthe erties trend of± hadronic that was interactions seen near do not the change “knee” significantly that istails usually of the attributedXmax distribution. to theObviously, cutoff the for unbiased galactic re- supernovaover less than cosmic two orders ray of acceleration, magnitude in primary with energy SN- gion depends on the distribution itself and can thus not (< factor 10 in center of mass energy), this change of associatedbe determined acceleration by simulations. mechanisms Instead, welike fit the shock data waves∆ able+ to35 accelerate2 nuclei to somewhat D10 =(82 21) g/cm /decade would imply a change in higherwith the energies mean of a than one-sided protons. truncated Thus normal nuclei distribu- tend tothe dominateenergy dependence− at the ofvery the composition end of the around energy the an- spectrum.tion (shown Thus as solid something lines in Fig. 2)similar and reject might all events be happeningkle, supporting in the the extragalactic hypothesis of a transition accelerators from galac- re- that have a FOV limit for which the measured Xmax de- tic to extragalactic cosmic rays in this region. sponsible for the highest2 energy cosmich rays.i parts by more than 5g/cm from its asymptotic value. The shower-to-shower fluctuations, RMS(Xmax), are After all cuts, 3754 events are selected for the Xmax obtained by subtracting the detector resolution in quadra- Whatanalysis. about The X super-GZKmax resolution events? as a function Pierre of energy Auger for findsture a from suppression the width of the the observed event rateXmax settingdistributions in abovethese events3 1019 is estimatedeV, consistent using a detailed with the simulation GZK ofcutoff.resulting It does in a not correction find evidence of 6g/cm for2. As super-GZK can be seen in the FD and· the atmosphere. The resolution, defined by the right panel of Fig. 3, we≤ observe a decrease in the eventsthe full – standard suggesting deviation, that is atperhaps the 20g/cm earlier2 level first-generation above experiments may not have fully un-2 fluctuations with energy from19 about 55 to 26g/cm as derstooda few EeV. their The energy difference calibrations between the (see reconstructed Fig. 15). At energies above 5.5 10 eV a correlation the energy increases. Assuming· again that the hadronic withXmax thevalues distribution in events that hadof nearby a sufficiently extragalactic high energy objectsinteraction isfound, properties including do not change an much excess within in the the ob- directionto be detected of Centaurus independently A, by two the or nearest more FD radio-loudstations served active energy galaxy. range, these Centaurus decreasing A fluctuations is about are 15 an is used to cross-check these findings and as it was shown million light years from earth, and has a prominentindependent relativistic signature jet that of an is increasing thought average to extract mass of in [13], the simulations reproduce the data well. the primary particles. energy from the vicinity of a supermassive black holeFor theat theinterpretation galaxy’s of center. the absolute values of Xmax h i and RMS(Xmax) a comparison to air shower simulations 9.7 High energy neutrinos,RESULTS GZK Neutrinos, and Beyondis needed.As GZK can beNeutrinos seen in Fig. 3, there are considerable The field of very high energy neutrino astronomydifferences is in its between infancy. the results The of Amanda calculations neutrino using dif- The measured Xmax and RMS(Xmax) are measured in ferent hadronic interaction models. These differences are h i telescopeenergy bins was of ∆ alg prototypeE = 0.1 below south 10 EeV Pole and detector,∆lgE = andnot now necessarily a much exhaustive, larger sincesecond-generation the hadronic interactio de-n tector,0.2 above IceCube, that energy. is operating. The last bin Neutrinos starts at 1019 penetrating.4 eV, models the do ice not cancover be the detected full range of via possible the muons extrapo- integrating up to the highest energy event (E = (59 lations of low energy accelerator data. If, however, taken they produce. The muons produce Cherenkov± light in the ice, which can be detected by 8) EeV). The systematic uncertainty of the FD energy at face value, the comparison of the data and simulations photomultipliersscale is 22% [24]. placeUncertainties on detector of the calibration, “strings” at- placedleads in to the same deep conclusions ice from as the above, surface. namely a Deep gradual icemospheric - ice a conditions,mile below reconstruction the surface and - event is exceptional selection increase pure and of the transparent, average mass of with cosmic light rays attentua- with energy tiongive lengths rise to a of systematic 100 meters uncertainty or more. of 13g/cm Thus2 afor relativelyup to 59 spare EeV. array of photomultiplier strings, 2 ≤ Xmax and 6g/cm for the RMS. The results were It is illustrative to compare the data with predictions h i ≤ found to be independentof zenith angle, time periods and for a simple two-component proton/iron model using 24 Latest results from the Auger Observatory Esteban Roulet

log (E/eV) 10 18 18.5 19 19.5 20 20.5 ] 2

eV σ sys(E)=22% -1 38

sr 10 -1 yr -2 J(E) [km 3

E HiRes 1037 Auger power laws power laws + smooth function

1018 1019 1020 Energy [eV]

Figure 1: Spectrum measured above 1 EeV (solid dots) and power law fits with breaks (dotted line). For Figure 15: The solid dots are the Pierre Auger measurements at the GZK cutoff. Also comparison also the HiRes measurements are displayed (open dots). shown are data from a second observatory, HiRes. From Roulet and the PA Collaboration, arXiv:1101.1825. The physical origin of the ankle is still uncertain, being the main candidate scenarios to explain this feature those relating it to the transition from a dying galactic component to a harder extragalactic component becoming dominant, or alternatively the so-called dip-scenario [3], in which cosmic rays are assumed to be extragalactic protons down to energies below 1 EeV and the concave shape observed arises from the effect of energy losses by pair creation with cosmic microwave background (CMB) photons. To properly fit the observed spectral shape this last scenario requires soft spectra at the sources (α at the source being typically αg 2.4-2.7) and/or strong evolution of ≃ the sources with redshift, what makes the distant sources intrinsically brighter (or more abundant) so that a larger fraction of the observed protons come from25 far away and are hence more affected by interactions with CMB photons. Also an upward shift of the energy scale of Auger by 40% ∼ would be required in this scenario to fit the location of the dip in the spectrum. Anisotropy measurements will help to distinguish among the two scenarios, because the galactic/extragalactic transition may lead to measurable dipole type patterns in the arrival direction distribution resulting from the diffusive escape of the galactic cosmic rays, and already significant constraints have been obtained by Auger at EeV energies [4]. Also composition measurements are important because galactic cosmic rays at EeV energies are expected to be dominated by heavy nuclei, since their confinement by galactic magnetic fields is a rigidity dependent effect. Enhance- ments of the Auger observatory to improve the sensitivity down to energies of 1017 eV, such as the HEAT fluorescence detectors or the AMIGA infill and muon detectors, will help to shed light on these issues in the near future. The second feature mentioned, i.e. the suppression observed at the highest energies, is similar to the expectations from the so called GZK effect associated to the attenuation of extragalactic protons by photo-pion production off CMB photons (this suppression was predicted by Greisen and Zatsepin and Kuz’min just after the discovery of the CMB). Also the photodisintegration of nuclei as heavy as iron would lead to similar features, while lighter nuclei would instead show a suppression down to a lower energy threshold, in approximate proportion to their masses. Note that a change in the injection spectrum at the sources may also contribute in part to the observed

3 placed in the ice by using hot water to drill very deep holes in the ice, can view vast masses of ice eﬃciently. IceCube is the world’s largest high energy neutrino detector, viewing one cubic kilometer of deep South Pole ice at depths of 1.4-2.4 kilometers via detectors placed on 86 strings. LBNL played an important role in the design of the detectors and the associated data acquisition system. IceCube construction was completed at the end of 2010. While many muons penetrate the ice, almost all of these come from above the detector, the secondaries from cosmic ray interactions. But because only neutrinos can penetrate the earth, upward going muons must be the daughters of upward going neutrinos that interact in or near the detector. Thus by studying upward going muons, the experimenters can determine the rate of neutrinos penetrating the earth from below.

The IceCube mass is sufficient to detect neutrinos up to about 1016 eV. Beyond this point, the event rate is too low: the km3 volume is insufficient. One of IceCube’s important scien- tific justifications is using neutrinos to located possible specific point sources – so that we can connect very energetic events with candidate astrophysical accelerators. Since neutrinos are unaffected by the CMB or by magnetic field, they point back to their sources.

But neutrino astronomy with still higher energy neutrinos is clearly important. One of the important consequences of the Pierre Auger observation of the GZK cutoff is the associ- + + ated prediction that very high energy neutrinos, produced for example by π µ νµ + → → νµ +e +ν ¯µ +νe. In addition, it is quite plausible that the astrophysical accelerators producing cosmic rays in the vicinity of 1019 1020 eV are in fact capable of accelerating protons or nuclei to still higher energies – though− we would not see such energetic primary cosmic rays because they cannot penetrate the CMB. Nevertheless, in the vicinity of the accelerator, these cosmic rays would be converted to neutrinos, which would reach earth. The flux of GZK neutrinos and the flux of those neutrinos associated with cosmic rays accelerated beyond the GZK cutoff can be estimated, if one assumes a spectrum of primary cosmic rays α 23 (say a proton flux declining as E− ); a maximum proton energy (10 eV was chosen for one of the calculations used in defining the band in Fig. 16), and a cosmological model (effectively the z at which the typical neutrino was produced, so that redshift effects can be accounted for). In this way, reasonable assumptions yield curves within the band of Fig. 16. The bulk of these neutrinos are beyond the reach of IceCube: the low flux of the highest energy neutrinos requires detection volumes of 100 km3 or more, not 1 km3.

Yet these neutrinos are important: they may be our only tool for investigating the universe at asymptotically high energies and large distances. How can they be detected? First, plausible estimates of ﬂuxes quickly lead to the conclusion that the detector must be huge, on the order of 100 km3 or more. Clearly such a detector cannot be a conventional one. Second, one needs some kind of neutrino interaction that can be detected in such a large volume, at reasonable cost.

There is an attractive possibility, based on the Askaryan eﬀect, proposed by Gurgen Askaryan

26 arXiv:0904.1309v1 [astro-ph.HE] 8 Apr 2009 rpitsbitdt Elsevier to submitted Preprint bevto fchrn ai msinuiga c target ice Direct an fluxes. using neutrino emission UHE radio coherent on cur- of limits has observation [4] best experiment the balloon the set ANITA rently in the and transmission ice, the radio polar of many of work established characteristics has experimental fundamental [3] significant work collaboration RICE theoretical then, the on by Since based described Askaryan. [2], originally Zhelezhykh by was and medium dense Gusev by a in emission dio neu- individual UHECR the the of indicate shed trinos. direction could the only via also not themselves but would sources issues, flux these neutrino on Mea- GZK light the factors. UHECR of unknown distribution, surement currently source other the and of flux composition, details the the of characteristics flux on the “guaranteed” depend a although to neutrinos, lead UHE background interactions microwave of Such cosmic 1). the Fig. with (see UHECRs interac- the Greisen-Zatsepin- of from the tion resulting [1], of suppression presence (GZK) the Kuzmin established has eV PACS: which environmen words: station radio IceRay Key the testbed of a monitoring of year-round status provide the detecting will on of detecto neutrino report goal IceCube also to the the We designed with with is coincidence ice, in array the Operating The in Pole. interactions South neutrino the at detector neutrino iheeg omcry(HC)setu bv 10 above spectrum (UHECR) ray cosmic energy high Introduction 1. fo results simulation and considerations design discuss We Abstract .Allison P. naryt eetUEnurnsvaterchrn ra- coherent their via neutrinos UHE detect to array An otne rgesi h eemnto fteultra- the of determination the in progress Continued ∗ mi address: Email Presenter. .Karle A. 46L,9.5V,9.0S,84.40.-x 98.70.Sa, 95.55.Vj, 14.60Lm, cRy nIeuecnee Radio- IceCube-centered An IceRay: a .Beatty J. , etiodtcin sayneet ai frequency radio effect, Askaryan detection, neutrino e .Kelley J. , [email protected] b .Chen P. , h ∗ d et fPyisadAtooy nv fAaaa Tuscaloos Alabama, of Univ. Astronomy, and Physics of Dept. ,e g et fPyis&Atooy nv olg odn odnW London London, College Univ. Astronomy, & Physics of Dept. et fPyisadAtooy nv fDlwr,Nwr,D Newark, Delaware, of Univ. Astronomy, and Physics of Dept. .Landsman H. , a et fPyisadAtooy nv fHwi,Mna I96 HI Manoa, Hawaii, of Univ. Astronomy, and Physics of Dept. b f et fPyis nv fMrln,CleePr,M 20742, MD Park, College Maryland, of Univ. Physics, of Dept. et fPyis hoSaeUiest,Clmu,O 43210 OH Columbus, University, State Ohio Physics, of Dept. e et fPyis nv fWsosn aio,W 30,USA 53703, WI Madison, Wisconsin, of Univ. Physics, of Dept. c c et fPyis ainlTia nv,Tie 0,Taiwan 106, Taipei Univ., Taiwan National Physics, of Dept. .Connolly A. , J Kelley) (J. .Seckel D. e .Learned J. , d .DuVernois M. , g .Varner G. , a .Miki C. , 17 ol lo opeeclrmtyo usto h events. the of subset a of calorimetry complete allow would r uha iia at n c:svrlo h e xeiet eedn tSA.Teeare There detectors. SLAC. radio very at surface a done instrumenting with were on ice based experiments Pole most key South materials, idea, the of several this volume of around in large several developed observed produce being been ice: can ideas has and detector – the several effect 10cm salt, to Askaryan of silica, compared The order as large wave.) the is such charge wavelength on the (The thickness of coherently. a thickness waves have matter, radio may in or light which microwave of – Cherenkov speed charges local of the front than moving faster move charges The c/n electrons. energy high duce produce sec- matter into of in annihilate shower reactions electrons neutrino a radiation Cherenkov energy coherent produce high of electron/ Indeed, can cone an a frequencies. matter emit microwave that through or and radio traveling charge at net particle a carry energetic that very particles ondary A various are 1962. points in data The resulting Graph CMB. neutrinos Auger. Pierre the al. energy from et data with high including Allison rays of spectrum, from ray cosmic spectrum cosmic hadronic predicted primary the of the of measurements is interaction band the hashed from The 16: Figure tteSuhPole. South the at t cRy rpsdlresaeutahg nry(UHE) energy ultra-high large-scale proposed a IceRay, r h omgncnurn u ihraoal vn rates. event reasonable with flux neutrino cosmogenic the ny Z etiomdl r rmPohre&Jhsn[0 a [10] Johnson & Protheroe from are Kalashev models neutrino GZK only. h aus 5,HvrhPr 6,teFysEe[] GS [8 AGASA repre points [7], Data Eye f Fly’s collaborations. data [9], the flux including ferential Auger c [6], and 2007, predicted Park [1], Haverah early and HiRes ray [5], of Yakutsk as cosmic the energy spectrum ultra-high neutrino World mogenic 1: Figure eekvGKNurn Detector Neutrino GZK Cerenkov ˇ a eettechrn saynrdoeiso rmUHE from emission radio Askaryan coherent the detect where .Williams D. , noprtsbt NT n cCb ehooyand technology IceCube and ANITA both incorporates a a .Gorham P. , .Morse R. , n tal. et γ steidxo ercin tln aeegh ai rmcoae the – microwaves or radio – wavelengths long At refraction. of index the is hwrcnann ehp 0 oeeetosta oirn:pstosand positrons positrons: than electrons more 20% perhaps containing shower dI ( [11]. E a h .Nichol R. , ) a /dE .Halzen F. , ,A 58,USA 35487, AL a, 91,USA 19716, E utpidby multiplied , γ 1 B,UK 6BT, C1E ,but s, 2,USA 822, USA , USA d γ .Rott C. , e a opo cte ffeetosi h atrt pro- to matter the in electrons off scatter Compton can s .Hanson K. , E 2 ro asaestatistical are bars Error . b .Ruckman L. , 27 e .Hoffman K. , pi ,2009 8, April etdif- sent a , rom f os- nd , ], 9.8 Gamma rays: general issues, GRO The properties of the earth’s atmosphere divides gamma ray astronomy into two halves. From the ultraviolet to gamma rays of energy 20 GeV the atmosphere is opaque. Thus observations in this energy range must be done∼ with instruments mounted in satellites or carried by balloons. The ease of detecting the radiation generally goes up with energy, but the strengths of typical astrophysical sources go down. The net result is that one can do a lot over this range: this motivated the design of the four instruments on board the Gamma Ray Observatory. For gamma rays above 20 GeV, interactions in the atmosphere produce showers that can be observed either by the Cerenkov light produced by the secondaries or, at high altitude, by direct detection of the secondaries. Of course, observations can also be (and are) made by space-bound detectors, too.

Many of the concerns of this field are driven by instrumental issues. Some of the central issues are clear: Producing detectors with greater sensitivity. The dynamic range of existing detectors - the• gap between the brightest nearby sources such as the Crab and the faintest detectable sources - is typically about a factor of 100. Thus the situation is equivalent to being able to see no stars fainter than the 5th magnitude. Improvements in sensitivity can thus greatly extend the horizon of our observations, while also allowing much more detailed spectral studies of known bright sources. Greater sensitivity can be achieved with great collection areas and by reducing ambient backgrounds. The push towards greater size is clearly an expensive challenge given the necessity for observing outside the atmosphere. Enhancing spectral resolution. This includes both enhanced spectral resolution and enhanced• spatial resolution, the latter required to better identify sources. Enhanced temporal and spatial coverage and resolution. • The technology in the field has advanced greatly over he past 20 years. The second of NASA’s Great Observatory projects (following the Hubble Space Telescope) was the Comp- ton Gamma Ray Observatory. Launched in 1991 from the shuttle Atlantis, the GRO operated for over nine years. Its instruments included BATSE, OSSE, COMPTEL, and EGRET. BATSE - the Burst and Transient Source Experiment - was an all-sky monitor with sharp timing capability: its NaI detectors were designed to detect gamma rays with energies between 20 and about 600 keV. This proved decisive in demonstrating that gamma ray bursts are distributed approximately isotropically. This detector also led to the discovery of new pulsars, x-ray novae, and the identification of one soft gamma ray repeater. OSSE - the Oriented Scintillation Spectrometer Experiment - employed NaI and CsI detectors to look for gammas between 50 keV and 10 MeV, and could point to sources. COMPTEL - the Imagining Compton Telescope – detected gamma between 0.75 and 30 MeV, with an angular resolution of 1 degree. EGRET - the Energetic Gamma Ray Experiment - was able to correlate high energy gammas (20 MeV-30 GeV) with the lower energy bursts detected by BATSE. It also measured high energy gammas from active galaxies.

28 9.9 Nuclear Gamma Rays We have touched on this theme before, but here I’d like to gather together several examples of what is becoming possible. One of these examples is 26Al, which 720,000 y lifetime for decay to 26Mg. The decay of the 5+ ground state populates the ﬁrst two excited states of 26Mg, which are 2+ states with energies of 2.938 and 1.809 MeV. The latter state is populated 97% of the time, so the primary signature is a 1.809 MeV γ. The 2.938 MeV state decays to the 1.809 MeV level, so a small number of 1.129 MeV γs are also produced.

The primary site for producing 26Al is thought to be Type II and IIb supernovae. Addi- tional aluminum may come from the winds of very massive stars. COMPTEL has produced a galactic map of the 1.809 MeV γs. That map diﬀers from higher energy ( 100 MeV) maps in that there is marked clumpiness to the production, including intense sources∼ associated with Cygnus, Vela, etc. This map is interpreted as an indicator of recent supernova activity. The conclusions drawn from such a map are clearly model dependent because one obtains an angular distribution but no spatial depth information. But under the assumption that the entire galaxy is contributing in the expected way, one deduces a recent supernova rate of 3.4 2.8/century. The model dependence also includes the uncertainty in the aluminum production± per event. The attribution of the total ﬂux to the Al injection rate of massive stars yields an upper bound on the recent star formation rate of 5 4 M per year. This

calculation, of course, requires not only a model of the Al production± per supernova, but also a model of the range of stellar masses that undergo core collapse and a model for the distribution of stars with mass (populations roughly decline exponentially with an exponent of about -2.35).

In a similar way, 44Ti decay proceeds with a 60 year half life to the ground state of 44Sc, 44 ∼ + which in turn electron captures to Ca. The order of the states in Sc is 2 (gs), 1− (68 keV), and 0− (146 keV). The decay feeds the second excited state 98% of the time, which then decays through the 1− state to the ground state, producing γs of 78 and 68 keV. The subsequent decay to 44Ca has a 4 hour lifetime and produces a 1.157 MeV γ, as the 2+ ﬁrst excited state of 44Ca is populated 99% of the time. The 100 keV line for the the source Cas A is within the detection abilities of OSSE, while the∼ 1.157 MeV line can be seen by COMPTEL.

Various types of supernovae are thought to produce 44Ti, including both types I and II. As in the case of the 26Al line, the galaxy is eﬀectively transparent to the produced γ ray, so the detection provides a measure of the very recent supernova rate free from worries about obscuration. With COMPTEL one had an instrument sensitive to nrearby supernovae occurring within the past 1000 years, or 15 half lives. Given a supernova rate of some several per century, it is clear that the distribution should be from quite localized sources, representing recent events. Typical productions of 44Ti from supernovae are, according to modelers, on 4 the order of 10− M per event.

29 Figure 17: COMPEL’s map of the galactic plane in the 1.809 MeV gamma ray emitted by 26Al. The distribution is not uniform. It is thought to represent over 100,000 events, supernovae or novae, producing something on the order of 1-3 solar masses of 26Al.

The youngest known galactic supernova remnant is Cas A, noted optically by John Flam- steed in 1680. COMPTEL reported the observation of 44Ti γs from Cas A in 1994. The results of a survey extending over a six-year period beginning in 1991 found only one additional source, identified with a young supernova remnant not observed either optically or in the radio. The source is located in the direction of the Vela constellation. Constraints on the doppler broadening of the 1.16 MeV line limits the velocity of the ejecta to no more than 19000 km/sec. The COMPTEL results are confirmed by ROSAT x-ray data of the vela region, which found a shell-type supernova remnant at the same location. The shell temperature is on the order of several keV, which also indicates a young remnant. Finally, COMPTEL previously saw 26Al in this region, attributing this to the known Vela supernova remnant. However the centroid of that distribution has been argued to better fit the new supernova remnant. Modelers are currently engaged in arguments about the nature of the progenitor, using both the ejection velocity bound and the 44Ti yield to bound models. One possibility is a core-collapse supernova of a massive star that had previously lost its hydrogen envelope. It has been claimed that a supernova at this distance (100-300 pc, derived from the Ti γ flux and age estimates of 600-1100 years based on the expansion velocity) could have been as bright as the moon, leaving an interesting question as to why it was not observed.

Finally, there have been recent papers suggesting that future generations of gamma ray detectors might see γs in the energy range of 100-700 keV associated with elements speciﬁc to the r-process. Establishing a correlation between such γs and known supernova remnants could thus establish the r-process site and, potential, constrain the total r-process production per site. One of the candidates, 126Sb, has a 144,000 year lifetime, easily long enough to allow a galaxy survey. It produces lines at 415, 666, and 695 keV. The proposed detector ATHENA possibly could detect 126Sb lines from Vela.

30 9.10 Astrophysics of Intermediate Energy Gammas EGRET conducted the ﬁrst all-sky search for gamma sources above 100 MeV, succeeding in identifying 271 high-energy gamma ray sources, most of which are associated with active galaxies and characterized by nonthermal spectra. 170 these were unknown prior to the survey. Such ”blazars” are highly variable and are bright radio sources. The radio structure consists of knotty jets moving outward at high velocities. The luminosity in gamma rays can exceed that from other wavelengths by up to two orders of magnitude. The variability of the emission can be fast, less than a week. The density of high energy gammas at the source are suﬃcient that photon-photon pair production would keep them trapped, unless the gammas are highly beamed, as in a relativistic jet. Thus the hope is that the gamma ray spectrum can yield information on the nature of the jet and of the acceleration processes occurring there. It is thought that the gamma rays may originate from a region of the jet closer to the central engine, than in the case of the radio emission.

A few blazars have been observed producing very high energy gamma rays, up to TeV scales, which then can be readily observed from earth. One, Markarian 421, was measured in the TeV range by the Whipple Observatory, which detects Cerenkov radiation from air showers. It was not seen at GeV energies by EGRET. The high energy flare had a rise time of about two days. One day after the flare began this source was seen in the X-ray. Both of these signals differ from typical blazars in their higher energies: most blazars have lower energy gamma and low-energy radiation that does not extend into the x-ray. The interpretation is that blazars accelerate electrons to high energies, which then radiate soft synchrotron radiation and hard gammas by inverse Compton scattering. Markarian 421 is thus exceptional in the energies to which it accelerates electrons, accounting for the higher energy of its x- ray/gamma emission.

Markarian 421, which was first observed in May, 1994, is not unique. The second closest blazar of this type, Markarian 501 (z=0.034), was seen in 1997. It flared to become the brightest TeV source in the sky, outshining the Crab Nebulae by an order of magnitude. The periods of flaring lasted a few days. The elevated activity spanned a period from about March through June. The high energy spectrum (observed by Whipple) was flat from below a TeV to at least 10 TeV. The x-ray cutoff in Markarian 421 is about 1 keV; in Markarian 501 it is above 100 keV. As in 421, it is assumed that this spectrum comes from relativistic acceleration of electrons along a jet closely aligned with our line of sight.

What does this tell us about the electron energies? First, the plasma in the jet is moving towards us, boosting the energy of the emitted gammas, relative to the jet rest frame. The Doppler factor is (1 v2/c2) D = − (16) (1 +qz)(1 v cos θ/c) − where θ is the observation angle relative to the jet axis. This is the standard relativistic

31 Doppler shift corrected by the redshift. The observed energy of the gamma rays is

E Dγm c2 (17) max ∼ e 2 where γmec is the maximum energy of the electrons in the jet rest frame. The Doppler factor also appears in the calculation of the photon density in the blob, and thus of the blobs opacity to high energy γs. The argument goes as follows: if ∆tobs is the fastest observed TeV gamma ray ﬂare variability, then the radius of the blob emitting the photons must be less than cD∆tobs. Thus a larger D means a lower photon density. From this one concludes D > 30.∼ Since the highest energy gammas from Markarian 501 is about 20 TeV, one derives for∼ the maximum energy of electrons in the jet frame

γmc2 0.65TeV (18) ∼ Since the maximum electron velocity is now known, one can deduce the magnetic field in the jet required to produce synchrontron radiation with a maximum energy of 200 keV. That yields E DBγ2 B 0.7G (19) max ∝ ⇒ ∼ where B is the field in the jet rest frame. Other aspects of the blazar dynamics - such as the physics responsible for the short timescale of the flares - is less clear.

Another exciting result involving high energy γs are the gamma ray burst observations of EGRET. While the vast majority of the bursts are seen by BATSE, on the order of 10% of the events produce spectra that extend into EGRET’s range of above 30 MeV, typically producing about ﬁve counts. A great deal more about such events is now known because of the Fermi Telescope.

These high energy tails place a lot of constraints on gamma ray burst models. Since high energy gammas were also seen early in the burst, there must be high-energy particle acceleration simultaneous to the keV emission. The lack of attenuation of high energy γs from γ γ pair production oﬀ lower energy γs place a particularly strong constraint on the source and− on models involving beaming.

Finally, a third interesting source of high energy γs are both isolated and binary pulsars. Gamma ray emission can be very strong, representing 10% or more of the spin-down energy of some pulsars. There is relatively little consensus on the mechanism or even the precise site of the gamma ray production (neutron star surface? accretion disk? etc?)

Centaurus X-3 is a well-studied high-mass accreting X-ray binary. EGRET measured a smoothly declining spectrum that extended from 100 MeV to its detection limit. Results have been reported from the Durham Mark 6 gamma ray telescope in the vicinity of a TeV: the ﬂux is consistent with a linear extrapolation of the EGRET ﬂux. Centaurus X-3 contains a 4.8s pulsar in a 2.1 day orbit about an O-type supergiant V779 Centaurus. The pulsar

32 period has been shortening since its discovery almost 30 years ago, which is attributed to spin-up from matter accreting on the neutron star from the more rapidly rotating inner edge of its accretion disk. The EGRET GeV burst observations were pulsed in agreement with the X-ray period. Initial TeV gamma ray observations also indicated pulsation near the pulsar period and localized in But later measurements at high energy indicated unpulsed emission that one would then associated with radiation over an extended volume encompassing the orbit.

9.11 Fermi Telescope A new high-energy gamma ray telescope – the Fermi Telescope, or the telescope-formerly- known-as-GLAST – was launched into space in June of 2008. Its purpose is to probe intermediate energy gamma ray sources, such as active galactic nuclei, supernova remnants, and pulsars, in order to catalog these sources and determine the nature of these accelerators. The telescopes can detect gammas with energies between 20 MeV and 300 GeV – so it can be viewed as somewhat of a successor to EGRET. The telescope looks for transient events as well as exotic phenomena, such as dark matter annihilation into high energy γs. One instrument, the Gamma Ray Burst Monitor, is designed to view the entire sky not occulted by earth, to look for and identify the location of gamma ray bursts. It is sensitive to gammas between 8 keV and 40 MeV. The second instrument, the Large Area Telescope, has an aperture that can view 20% of the sky at one time. It can detect gammas from 20 MeV to 300 GeV. In its ﬁrst year, Fermi extended the EGRET catalogue of 271 point sources to 1451 sources. Among Fermi discoveries to date are the observation of a gamma-ray-emitting pulsar but emits in no other band; the observation of the most energetic gamma-ray burst yet recorded; and the observation of very energetic cosmic rays from the vicinity of supernova remnants. It has also placed important constraints on the diﬀuse gamma ray background, an isotropic gamma ray component of unknown origin, and on gamma-ray decay channels of dark matter candidates.

9.12 Gamma ray bursts Gamma ray burst have been mentioned several times. Gamma ray bursts are short-lived burst lasting from a few milliseconds to several minutes. They are detected roughly once per day, from all directions in the sky. They represent tremendous energy, several hundred times brighter than a typical supernova (were they radiated in 4π). They were originally detected in the 1960s by military satellites monitoring nuclear testing. They are not local, but originate at cosmological distances. In the past few years it has become clear they are associated (at least some of the time) with certain supernovae.

Recently fast-response telescopes - like NASA’s High-Energy Transient Explorer or the Japanese Automated Response Telescope - have been able to view the gamma-ray burst area within a couple of minutes of the burst. The ”afterglows” have been observed for on the order of an hour to a day after the burst – representing a lot of additional energy.

33 Figure 18: The Fermi Telescope, prior to launch.

34 Figure 19: The Fermi Telescope’s map of gamma rays and gamma-ray point sources in the plane of the galaxy. Compare to Figure 17.

General ideas exist that gamma ray bursts are associated with ﬁreball phenomena where tremendous energy and high velocities are produced in a region relatively free of baryons – to keep the gamma rays from thermalizing. The motion leads to a beaming of the burst, reducing the energy requirements but increasing the frequency of events (many of which are not beamed at earth and thus are not observed).

A specific idea is the hypernova. A massive star’s core collapses into a black hole. The black hole’s spin or magnetic fields may act like a slingshot, flinging material outward as a beamed blast wave or jet. The gamma rays are created when the blast wave collides with stellar material still inside the star. The gamma rays burst out of the star just in front of the blast wave. Behind the gamma rays, the blast wave pushes the stellar material outward.

The blast wave then sweeps through space, colliding with gas and dust, producing additional radiation – progressing from less energetic gammas, to x-rays, to visible light, and to radio. This forms the afterglow, lasting perhaps up to days.

35 Figure 20: A long-duration gamma ray burst event.