CENTAURUS A: THE EXTRAGALACTIC SOURCE OF COSMIC RAYS WITH ENERGIES ABOVE THE KNEE - IOPscience CENTAURUS A: THE EXTRAGALACTIC SOURCE OF COSMIC RAYS WITH ENERGIES ABOVE THE KNEE Peter L. Biermann1,3,4,5,6 and Vitor de Souza2 Published 2012 January 25 • © 2012. The American Astronomical Society. All rights reserved. 1 Max-Planck-Institute for Radioastronomy, Auf dem Hügel 69, D-53121 Bonn, Germany 2 Instituto de Física de São Carlos, Universidade de São Paulo, Brazil 3 Also at Physics Department, Karlsruhe Institute of Technology, Karlsruhe, Germany. 4 Department of Physics & Astronomy, University of Alabama, Tuscaloosa, AL, USA. 5 Department of Physics, University of Alabama at Huntsville, AL, USA. 6 Department of Physics & Astronomy, University of Bonn, Germany. Dates Sign up for new issue notifications Abstract The origin of cosmic rays at all energies is still uncertain.

In this paper, we present and explore an astrophysical scenario to produce cosmic rays with energy ranging from below 1015 to 3 × 1020 eV. We show here that just our Galaxy and the radio galaxy Cen A, each with their own galactic cosmic-ray particles but with those from the radio galaxy pushed up in energy by a relativistic shock in the jet emanating from the active black hole, are sufficient to describe the most recent data in the PeV to near ZeV energy range. Data are available over this entire energy range from the KASCADE, KASCADE-Grande, and Pierre Auger Observatory experiments. The energy spectrum calculated here correctly reproduces the measured spectrum beyond the knee and, contrary to widely held expectations, no other extragalactic source population is required to explain the data even at energies far below the general cutoff expected at 6 × 1019 eV, the Greisen-Zatsepin-Kuz'min turnoff due to interaction with the cosmological microwave background. We present several predictions for the source population, the cosmic-ray composition, and the propagation to Earth which can be tested in the near future. Export citation and abstract About Related Links 1.

INTRODUCTION Cosmic rays were originally discovered in 1912/1913 by Hess ( 1912 ) and Kohlhörster ( 1913 ) and still today we have no certainty where they come from. Their overall spectrum has been shown to be essentially a power law with a bend down near 1015 eV, called the knee, and a turn toward a new flatter component near ~3 × 1018 eV, called the ankle, with a final turnoff just around 1020 eV, as summarized in Gaisser & Stanev ( 2008 ). It is thought that the component below about 3 × 1018 eV is Galactic and the component above this energy is extragalactic on the basis that particles above such an energy would be hard to contain and isotropize in the magnetic fields in the interstellar medium (ISM) disk of the Galaxy. Different astrophysical scenarios have been proposed to explain the Galactic and extragalactic components of the cosmic radiation; see the overview in Stanev ( 2010a , 2010b ).

The basic paradigm for Galactic cosmic rays has been acceleration in the shock waves caused by supernova (SN) explosions (Baade & Zwicky 1934 ). The process of acceleration is diffusive shock acceleration (Fermi 1949 , 1954 ) and it is based on the compression experienced by particles that get reflected by magnetic irregularities from both sides of a shock in an ionized magnetic plasma (Drury 1983 ). SNe are exploding stars and they may explode either directly into the ISM or into the stellar wind of the predecessor star (Woosley et al. 2002 ). Lagage & Cesarsky ( 1983 ) showed that acceleration at the shocks caused explosions into the normal ISM that cannot reach even the energies at the knee. Heavy nuclei can be accelerated up to about 1018 eV in Galactic sources, such as SN explosions of massive stars which explode into their wind (Völk & Biermann 1988 ), OB-star super-bubbles, gamma-ray bursts (Dermer 2004 ) or micro-quasars, and active accreting black holes in stellar binary systems.

For higher energies (>1018 eV) extragalactic sources are the most accepted candidates. Nearby active radio galaxies with black hole activity were first proposed by Ginzburg & Syrovatskii ( 1963 ; see also, e.g., Lovelace 1976 ; Biermann & Strittmatter 1987 ) as possible sources. Gamma-ray bursts in other galaxies have also been suggested by Waxman ( 1995 ) and Vietri ( 1995 ). The radio galaxy Cen A is a prime example of a possible astrophysical source (Anchordoqui et al. 2011 ; Fargion & D'Armiento 2011 ). The interaction of high-energy particles with the microwave background limits the distance of the sources (Greisen 1966 ; Zatsepin & Kuz'min 1966 ; Allard et al. 2008 ; Stanev 2010b ).

For protons the energy at which the all-particle spectrum from many sources should turn off is estimated to be 6 × 1019 eV and for other chemical elements the energy turnoff is lower (Allard et al. 2008 ). As the interaction distance becomes very large for particles with energy between 3 × 1018 and 6 × 1019 eV, the calculations predict that a large number of extragalactic sources contribute to the measured flux in this energy range. This is a primary prediction of these calculations. Of special interest also is the transition from Galactic to extragalactic predominance which should happen in the energy range between 1016 and 1018 eV. Several experiments are presently taking data in this energy range: KASCADE-Grande, Telescope Array, IceTop, and the Pierre Auger Observatory. The transition is a very important feature because it is foreseen that breaks in the all-particle spectrum and in the composition can reveal the details of the particle production mechanisms, the source population, and propagation in the universe.

There are a number of recent attempts to explain the cosmic-ray spectrum in the transition range, e.g., by Hillas ( 2006 ) and by Berezinsky ( 2009 ). In this paper, we focus on the previously inaccessible continuous energy range above 10 PeV extending to the highest energies measured. Today it is possible to compare the predictions with high-precision data over the entire energy range. Therefore, it becomes important to have predictive power, i.e., to test quantitative hypotheses which were developed long before much of the new data were known. We revisit here an idea originally proposed in 1993 (Biermann 1993 ; Biermann & Cassinelli 1993 ; Biermann & Strom 1993 ; Stanev et al. 1993 ) and we show how our Galaxy and the radio galaxy Cen A can describe the energy spectrum from 10 PeV up to 3 × 1020 eV and describe the Galactic to extragalactic transition at the same time.

In the following sections, we first go through the tests the 1993 original model has undergone to date as regards spectra, transport, secondaries, and composition; second, we confirm the predictions of the original model with the newly available data beyond the knee energy, and finally we present the high-energy model which describes the transition between Galactic and extragalactic cosmic rays. 2. ORIGINAL MODEL AND ITS TESTS TO DATE In a series of papers started in 1993 (Biermann 1993 ; Biermann & Cassinelli 1993 ; Biermann & Strom 1993 ; Stanev et al. 1993 ; Biermann 1994 ) an astrophysics scenario was proposed which emphasized the topology of the magnetic fields in the winds of exploding massive stars (Parker 1958 ). In Stanev et al. ( 1993 ), a comprehensive spectrum was predicted for six element groups separately: H, He, CNO, Ne–S, Mn–Cl, and Fe. The key points of this original model are as follows.

(1) The shock acceleration happens in a region which is highly unstable and shows substructure, detectable in radio polarization observation of the shock region, which is also found in theoretical explorations (e.g., Bell & Lucek 2001 ; Caprioli et al. 2010 ; Bykov et al. 2011 ). Therefore, the particles go back and forth across the shock gaining momentum, while the scattering on both sides is dominated by the scale of these instabilities, which are assumed to be given by the limit allowed by the conservation laws of mass and momentum. (2) There are cosmic-ray particles which get accelerated by a shock in the ISM, produced by the explosion of a relatively modest high-mass star or, alternatively, by a low-mass SN Ia. This is most relevant for hydrogen and less so for helium and heavier nuclei. (3) Heavy cosmic-ray nuclei derive from very massive stars, which explode into stellar winds already depleted in hydrogen, and also in helium for the most massive stars. These explosions produce a two-part spectrum with a bend that is proposed to explain the knee.

In this scenario, the knee is due to the finite containment of particles in the magnetic field of the predecessor stellar wind, which runs as sin θ/r in polar coordinates (Parker 1958 ). Toward the pole region only lower energies are possible and the knee energy itself is given by the space available in the polar region. There is a polar cap component of cosmic rays associated with the polar radial field with a flatter spectrum. (4) Diffusive leakage from the cosmic-ray disk steepens all these spectra by 1/3 for the observer. (5) Very massive stars eject most of their zero-age mass before they explode and so form a very massive shell around their wind (Woosley et al. 2002 ). This wind shell is the site of most interaction for the heavy nuclei component of cosmic rays. For stellar masses above about 25 solar masses in zero-age main-sequence mass (Biermann 1994 ), the magnetic irregularity spectrum is excited by the cosmic-ray particles themselves. The spectral steepening due to the interactions is E−5/9 for the most massive star shells. The final spectrum is a composite of these components; see Figure 1 of Stanev et al.

( 1993 ). The spectra predicted by these arguments match the data such as shown by the recent Cosmic Ray Energetics And Mass (CREAM) results (Wiebel-Sooth et al. 1998 ; Biermann et al. 2009 ). This scenario has undergone detailed tests as regards propagation and interactions (Biermann 1994 ; Biermann et al. 2009 ) so as to describe both Galactic propagation and the spectra of the spallated isotopes as well as the resulting positron spectra, the flatter cosmic-ray positron and electron data, the Wilkinson Microwave Anisotropy Probe haze and the spectral behavior of its inverse Compton emission, and the 511 keV emission from the Galactic center region. New Transition Radiation Array for Cosmic Energetic Radiation (TRACER) results (Obermeier 2011 ) are also consistent in terms of (1) the low-energy source spectrum, (2) the energy dependence of interaction, (3) a finite residual path length at higher energy, and (4) a general upturn in the individual element spectra. The newest Pamela results (Adriani et al.

2011 ) are also consistent with the 1993 original model in which hydrogen was the only element to have a strong ISM–SN cosmic-ray component, and so has a steeper spectrum than helium. 2.1. A Test Beyond the Knee This original model was proposed to explain the particles observed above 109 eV per nuclear charge. Here we first test the original model with the KASCADE data. The most accurate measurement of the energy spectrum in the knee energy range has been done by the KASCADE experiment (KASCADE-Grande Collaboration 2010 ). Figure 1 shows for the first time the comparison of the original model to the measured data from KASCADE. KASCADE reconstructs the spectrum using two hadronic interaction programs (QGSJet and Sibyll) in the analysis procedure. In the figure we show the data and the original model, and also include the ratio of the difference between original model and data divided by the experimental error. For the ratio shown we use only one of these interaction codes; as an example we use QGSJet. The figure shows good agreement between data and the original model to within the errors of the data.

This confirms the original model in the last remaining energy range where it had not previously been tested for lack of good data. This is the first key result of this paper. Zoom In Standard image High-resolution image Export PowerPoint slide 2.2. Transport and Interaction Test One question which invariably comes up is how this model deals with propagation and spallation. For this line of reasoning it is important to note two aspects: (1) plasma simulations, the solar wind data (Matthaeus & Zhou 1989 ), the ISM data (Rickett 1977 ; Spangler & Gwinn 1990 ; Goldstein et al.

1995 ) as well as radio galaxy data (Biermann 1989 ) all are consistent with the interpretation that in an ionized magnetic plasma in near equipartition we have an approximate Kolmogorov spectrum (Kolmogorov 1941a , 1941b , 1941c ) running without break from very large scales down to dissipation scale; (2) very massive stars eject most of the initial zero-age main-sequence mass in their powerful winds, which then builds up a correspondingly massive shell at the outer boundary of the wind. The cosmic-ray-loaded SN shock encounters and interacts with this massive shell. All those cosmic-ray particles accelerated by an SN shock in the heavily enriched winds of Wolf–Rayet stars then excite magnetic irregularities, which can be described following Bell ( 1978a , 1978b ); these self-excited irregularities describe the path of the cosmic-ray particles through the massive shell. This gives (Biermann 1998 ; Biermann et al.

2001 ) a boron/carbon ratio energy dependence of E−5/9; this was found to be consistent with the data of Ptuskin et al. ( 1999 ), who determined E−0.54. Since the straight-line path gives a minimum path, this model also gives a finite path length of interaction at high energy, consistent with the new data as noted above. Other tests are as follows. The cosmic-ray electron spectrum has been determined to be E−3.23 ± 0.06 (Wiebel-Sooth & Biermann 1999 ) in the energy range up to a few TeV, a spectrum which is dominated by losses (Kardashev 1962 ). Therefore, the injection is with a spectrum flatter by unity. This has to be compared with the proton spectrum at a corresponding energy, which suggests E−2.66 ± 0.02 near TeV energies (Yoon et al.

2011 ); this spectrum has been steepened by the energy dependence of the diffusion, and so the difference to the inferred cosmic-ray electron spectrum gives this dependence as E−0.43 ± 0.06, consistent with E−1/3 as deduced from the Kolmogorov assumption. Another consistency check is the timescale inferred at the highest energies for leaking out of the kpc-thick cosmic-ray disk near the Sun (Biermann 1993 ; Biermann & Cassinelli 1993 ; Biermann & Strom 1993 ). Since we have about 107 yr at GeV energies for protons, we infer with the Kolmogorov assumption a timescale of 104.3 yr at 1017 eV—adopting the point of view that the highest energy cosmic-ray particles from the Galaxy are Fe (Stanev et al. 1993 ) at 1018.5 eV, matching in their scattering protons of 1017 eV; however, this is not yet finally settled.

This timescale is still significantly longer than the simple transit time across the thick disk of about 103.5 yr, so that the isotropy observed can be understood without already invoking the effect of the Galactic wind. Using three times the simple transit time as the minimum to give isotropy we can invert this line of reasoning and deduce a maximum energy dependence of the scattering of E−0.38 (Biermann 1993 ; Biermann & Cassinelli 1993 ; Biermann & Strom 1993 ), again under the assumption that the entire energy range is covered by the same power law. The Wolf–Rayet star model is also able to explain the positron spectra and positron production (Biermann et al. 2009 , 2010a , 2010b ) as noted above. 3. THE HIGH-ENERGY MODEL Based on the original model we propose here a high-energy model to explain the cosmic-ray data from 10 PeV to 300 EeV.

We analyze the possibility that the very same spectrum proposed in 1993, however shifted in energy, can explain all the data up to 300 EeV including the knee region, the highest energy range (>1018 eV), and at the same time the middle energy range (1016 eV < E < 1018 eV). The cosmic-ray particles as seen in our Galaxy were argued to provide the seed particles for further acceleration to ultrahigh energy by a relativistic shock emanating from an active black hole in the nearby radio galaxy Cen A (Gopal-Krishna et al. 2010 ). This idea is explored here beyond what has been proposed in Gopal-Krishna et al. ( 2010 ), which demonstrated that Cen A can provide a sufficient flux for the highest energy particles by working out the energetic particle flux traversed by a shock surface in the jet of the radio galaxy Cen A with the one-step further acceleration in a relativistic shock as proposed by Achterberg et al. ( 2001 ). They used the spectral shape of the original model (Stanev et al.

1993 ) to fit the Pierre Auger data; however the fit of the measured spectrum was not constrained by the low-energy spectrum as proposed in the original model (Stanev et al. 1993 ). The energy spectrum calculated here is simply a shift of the particle spectra proposed in the original model (Stanev et al. 1993 ) for low energies to the highest energies, preserving the relative abundances of the original model. This proposal is considerably stronger than the previous one presented in Gopal-Krishna et al. ( 2010 ). The energy shift corresponds to a factor of 2800 within the limits of a one-step acceleration by a relativistic shock as proposed by Achterberg et al. ( 2001 ). The original model (Stanev et al. 1993 ) has a number of parameters, which were set in 1992 (see Figure 6 in Stanev et al. 1993 ). None of these parameters had to be changed significantly in the analysis presented here. 4.

MATCH TO DATA FROM 10 PeV TO 300 EeV Finally, we can construct the energy spectrum of cosmic rays by adding the Galactic component to the extragalactic component as shown in Figure 1 . In this figure, we show differential flux × E3 versus energy per particle as predicted by this analysis compared to the data from KASCADE (KASCADE Collaboration 2009 ), KASCADE-Grande (KASCADE-Grande Collaboration 2010 ), and the Pierre Auger Observatory (Pierre Auger Collaboration 2010a ). In order to match the measurements from the two experiments, we have shifted, within the experimental uncertainties, the KASCADE and KASCADE-Grande flux down by 14% and the Auger flux up by 14%. We distinguish six groups: H, He, CNO, Ne–S, Cl–Mn, and Fe. Particles were subjected to losses in the intergalactic radiation field (Allard et al. 2008 ). The numbers above the lines correspond to error estimations (Model − Data)/(Experimental Error).

We note the good agreement between data and the model from below 1015 to 3 × 1020 eV. One extra assumption of the original model, namely the energy shift, allows a description of the energy spectrum above 1018 eV. Below the critical energy for interactions with the microwave background, it has been expected that we would observe a very large number of sources at large distances. However, no other source population is needed to describe the energy spectrum above 3 × 1016 eV. This is the second key new result of this paper. To summarize, a few results can be extracted from the proposal here. (1) There are no other sources necessary to provide extra flux in the energy range 3 × 1016 to 3 × 1020 eV, the second key result of this paper. A detailed analysis of radio galaxies (Caramete et al. 2011 ) shows that the next strongest radio galaxy to contribute is Virgo A, as already predicted by Ginzburg & Syrovatskii ( 1963 ).

We estimate the maximum possible extra flux from other sources within this distance limit at 0.1 in the log in Figure 1 , so at 25%. We note that the self-consistent MHD simulations for cosmological magnetic fields presented in Ryu et al. ( 2008 ; see also Cho & Ryu 2009 ; Das et al. 2008 ) were carried out for protons, and so for heavy nuclei the magnetic horizon in intergalactic space is small, less than 100 Mpc consistent with the measurements of the Pierre Auger Observatory (Pierre Auger Collaboration 2010c ) which suggest a heavier composition for energies above 1019 eV. However, at yet lower energies the sum of the more distant sources might exceed the flux predicted from single sources such as Cen A and Vir A; the magnetic scattering as predicted in Ryu et al. ( 2008 ; see also Cho & Ryu 2009 ; Das et al. 2008 ) seems to prevent this at all energies above the transition to Galactic cosmic rays.

(2) The dip near 3 × 1018 eV is explained by the switchover between the Galactic cosmic rays and the extragalactic cosmic rays (Rachen & Biermann 1993 ; Rachen et al. 1993 ). (3) The spectra of Galactic cosmic rays beyond the knee are adequately modeled by our approach, suggesting that the Wolf–Rayet star explosion model also matches the newest data beyond the knee. (4) There is no abrupt change in composition in the energy range from 3 × 1016 to 3 × 1018 eV. In order to describe the data (Pierre Auger Collaboration 2010a ) the high-energy model presented here requires minimal interaction along the path between the radio galaxy Cen A and us, and so indeed near-isotropic scattering in the magnetic wind of our Galaxy. 4.1. Isotropy? The most recent results from the Pierre Auger Observatory indicate an excess in the direction of Cen A (Pierre Auger Collaboration 2010b ) which might corroborate the analysis presented here.

The Pierre Auger Observatory measures an isotropic sky for energies below 6 × 1019 eV and a weakly anisotropic sky above this energy (Pierre Auger Collaboration 2010b ). The isotropy for energies below 6 × 1019 eV can be explained by a turbulent magnetic wind of our Galaxy (Everett et al. 2008 ). This wind is thought to be driven by cosmic rays and hot gas, and so is itself unstable, giving irregular magnetic fields, akin to radiation-driven winds of stars (Cassinelli 1994 ; Owocki 1990 ). The magnetic fields in the wind are strong enough to isotropize incoming heavy-element particles at very high energy and protons at lower energy. Scattering in a magnetic wind with B ~ 1/R as a function of radial distance R (Parker 1958 ) is strongly enhanced due to the extra factor derived from integrating the Lorentz force ln {Rmax/Rmin} ~ 5. The maximal energy for total bending can then be given by the magnetic field strength (Everett et al.

2008 ) at the base ~8 μG, the length scale at the base ~5 kpc, the logarithmic factor ~5, and so is given by 1020.2 Z eV, where Z is the charge of the cosmic-ray particle. Since at those energies the data measured by the Pierre Auger Observatory suggest that we actually have heavy nuclei, complete bending is assured. However, this would lead to a second problem in that we then might find a complete shielding for any particles coming from outside, so this magnetic wind must also have considerable irregularities; these irregularities in the wind need to be scale free (implying a saturated spectrum of irregularities or inverse cascade I(k)k ~ constant, where I(k) is the energy per wavenumber k per volume), so as to avoid a characteristic energy below which all particles are cut off; or, if such an energy exists, it must be low enough not to disturb the spectral sum. The key point is that Parker winds (Parker 1958 ) are very effective at bending orbits.

Obviously, the scattering might not be complete so that a small anisotropy is left, possibly explaining the Auger data clustering of events near the direction to the radio galaxy Cen A. 5. PREDICTIONS Astrophysical Origins of Ultrahigh Energy Cosmic Rays - D.F. Torres & L.A. Anchordoqui 1.4. Propagation of UHECRs In this section we briefly summarize the relevant interactions that CRs suffer on their trip to Earth. For a more detailed discussion the reader is refer to [ 1 , 70 , 71 , 72 , 73 ]. 1.4.1. The GZK-cutoff Ever since the discovery of the cosmic microwave background (CMB) standard physics implies there would be a cuttoff in the observed CR-spectrum.

In the mid-60's Greisen, Zatsepin, and Kuzmin (GZK) [ 74 , 75 ] pointed out that this photonic molasses makes the universe opaque to protons of sufficiently high energy, i.e., protons with energies beyond the photopion production threshold, (4) where mp (m ) denotes the proton (pion) mass and CMB ~ 10-3 eV is a typical CMB photon energy. After pion production, the proton (or perhaps, instead, a neutron) emerges with at least 50% of the incoming energy. This implies that the nucleon energy changes by an e-folding after a propagation distance ( y)-1 ~ 15 Mpc. Here, n 410 cm-3 is the number density of the CMB photons, p > 0.1 mb is the photopion production cross section, and y is the average energy fraction (in the laboratory system) lost by a nucleon per interaction. Energy losses due to pair production become relevant below ~ 1019 eV.

For heavy nuclei, the giant dipole resonance can be excited at similar total energies and hence, for example, iron nuclei do not survive fragmentation over comparable distances. Additionally, the survival probability for extremely high energy ( 1020 eV) -rays (propagating on magnetic fields >> 10-11 G) to a distance d, p( > d ) exp[- d / 6.6 Mpc], becomes less than 10-4 after traversing a distance of 50 Mpc. In recent years, several studies on the propagation of CRs (including both analytical analyses and numerical simulations) have been carried out [ 76 , 77 , 78 , 79 , 80 , 81 , 82 , 83 , 84 , 85 , 86 , 87 , 88 , 89 , 90 , 91 , 92 , 93 , 94 , 95 , 96 , 97 ]. A summary of the UHECR attenuation lengths for the above mentioned processes (as derived in these analyses) is given in Fig. 2 .

It is easily seen that our horizon shrinks dramatically for energies 1020 eV. Therefore, if UHECRs originate at cosmological distances, the net effect of their interactions would yield a pile-up of particles around 4 - 5 × 1019 eV with the spectrum droping sharply thereafter. As one can infer from Fig. 2 , the subtleties of the spectral shape depend on the nature of the primary species, yielding some ambiguity in the precise definition of the "GZK cutoff". In what follows we consider an event to superseed the cutoff if the lower energy limit at the 95% CL exceeds 7 × 1019 eV. This conforms closely to the strong criteria outlined in Ref. [ 98 ]. Figure 2. Attenuation length of 's, p's, and 56Fe's in various background radiations as a function of energy. The 3 lowest and left-most thin solid curves refer to -rays, showing the attenuation by infra-red, microwave, and radio backgrounds. The upper, right-most thick solid curves refer to propagation of protons in the CMB, showing separately the effect of pair production and photopion production.

The dashed-dotted line indicates the adiabatic fractional energy loss at the present cosmological epoch (see e.g. Appendix B of Ref. [ 1 ]). The dashed curve illustrates the attenuation of iron nuclei. 1.4.2. Propagation of CRs in a magnetized neighborhood of the Galaxy In addition to the interactions with the radiation fields permeating the universe, CRs suffer deflections on extragalactic and Galactic magnetic fields. Over the last few years, it has become evident that the observed near-isotropy of arrival directions can be easily explained if our Local Supercluster contains a large scale magnetic field which provides sufficient bending to the CR trajectories [ 99 , 100 ]. Intergalactic field strengths and coherent lengths are not well established, but it is plausible to assume that fields have coherent directions on scales 0.5 - 1 Mpc. The Larmor radius of a CR of charge Ze propagating in a magnetic field BnG B/10-9 G is given by (5) where E20 is the particle's energy in units of 1020 eV.

For rL >> the motion is not very different from quasilinear trajectory, with small deflections away from the straight line path given by (6) where LMpc is the propagation distance in units of Mpc. As the Larmor radius starts approaching the particles begin to diffuse. Diffusion has two distinctive regimes. Particles that are trapped inside magnetic subdomains (of size Mpc / Mpc) follow Kolmogorov diffusion. In such a case, the functional dependence of energy of the difussion coefficient is found to be [ 101 ] (7) With rising energy, rL , and there is a transition to Bohm diffusion. The diffusion coefficient in this regime is of order the Larmor radius times velocity (~ c) [ 102 ]. In this case the accumulated deflection angle from the direction of the source, can be estimated assuming that the particles make a random walk in the magnetic field [ 103 ] (8) Surprisingly little is actually known about the extragalactic magnetic field strength.

There are some measurements of diffuse radio emission from the bridge area between Coma and Abell superclusters that under assumptions of equipartition allows an estimate of 0.2 - 0.6 µG for the magnetic field in this region [ 104 ]. (2) Such a strong magnetic field (which is compatible with existing upper limits on Faraday rotation measurements [ 106 ]) could be possibly understood if the bridge region lies along a filament or sheet of large scale structures [ 107 ]. Faraday rotation measurements [ 108 , 106 ] have thus far served to set upper bounds of (10-9 - 10-8) G on extragalactic magnetic fields on various scales [ 106 , 109 ], as have the limits on distortion of the CMB [ 110 , 111 ]. The Faraday rotation measurements sample extragalactic field strengths of any origin out to quasar distances, while the CMB analyses set limits on primordial magnetic fields. Finally, there are some hints suggesting that the extragalactic field strength can be increased in the neighborhood of the Milky Way, BnG > 10 [ 112 ].

Now, using Eq. (5), one can easily see that because of the large uncertainty on the magnetic field strength, (nG) - (µG), all 3 different regimes discussed above are likely to describe UHECR propagation. If CRs propagate diffusively, the radius of the sphere for potential proton sources becomes significantly reduced. This is because one expects negligible contribution to the flux from times prior to the arrival time of the diffusion front, and so the average time delay in the low energy region, delay d2 / [4D(E)], must be smaller than the age of the source, or else the age of the universe (if no source within the GZK radius is active today, but such sources have been active in the past). Note that the diffuse propagation of UHE protons requires magnetic fields ~ 1µG. Therefore, for typical coherence lengths of extragalactic magnetic fields the time delay of CRs with E 1018.7 eV cannot exceed delay 14 Gyr, yielding a radius of d ~ 30 Mpc. In the case CR sources are active today, the radius for potential sources is even smaller d ~ 5 Mpc.

On the other hand, the sphere of potential nucleus-emitting-sources is severely constrained by the GZK cutoff: straightforward calculation, using the attenuation length given in Fig. 2 , shows that less than 1% of iron nuclei (or any surviving fragment of their spallations) can survive more than 3 × 1014 s with an energy 1020.5 eV. Therefore, the assumption that UHECRs are heavy nuclei implies ordered extragalactic magnetic fields BnG 15 - 20, or else nuclei would be trapped inside magnetic subdomains suffering catastrophic spallations. The large scale structure of the Galactic magnetic field carries substantial uncertainties as well, because the position of the solar system does not allow global measurements. The average field strength can be directly determined from pulsar observations of the rotation and dispersion measures average along the line of sight to the pulsar with a weight proportional to the local free electron density, <B||> 2 µG [ 113 , 114 , 115 , 116 ].

(We use the standard, though ambiguous notation, in which B refers to either the Galactic or extragalactic magnetic field, depending on the context.) Measurements of polarized synchrotron radiation as well as Faraday rotation of the radiation emitted from pulsars and extragalactic radio sources revealed that the global structure of the magnetic field in the disk of our Galaxy could be well described by spiral fields with 2 (axisymmetric, ASS) or (bisymmetric, BSS) symmetry [ 117 ]. In the direction perpendicular to the Galactic plane the fields are either symmetric (S) or antisymmetric (A). Discrimination between these models is complicated. Field reversals are certainly observed (in the Crux-Scutum arm at 5.5 kpc from the Galactic center, the Carina-Sagittarius arm at 6.5 kpc, the Perseus arm at 10 kpc, and possibly another beyond [ 118 ]). However, as discussed by Vallée [ 119 ], turbulent dynamo theory can explain field reversals at distances up to ~ 15 kpc within the ASS configuration.

More accurately, the field strength in the Galactic plane (z = 0) for the ASS model is generally described by [ 120 , 121 ] (9) (10) where is the azimuthal coordinate around the Galactic center (clockwise as seen from the north Galactic pole), is the galactocentric radial cylindrical coordinate, and (11) Here, 0 = 10.55 kpc stands for the galactocentric distance of the maximum of the field in our spiral arm, = 1 / tan p (with the pitch angle, p = - 10°), r0 = 8.5 kpc is the Sun's distance to the Galactic center, and 1 = 2 kpc.

The coordinates of the field are correspondingly, (12) The field strength above and below the Galactic plane (i.e., the dependence on z) has a contribution coming from the disk and another from the halo: (i) for A models (13) (ii) for S models, BS = BA( , , z) / tanh(z / z3); where z1 = 0.3 kpc, z2 = 4 kpc and z3 = 20 pc. With this in mind, the Galactic magnetic field produce significant bending to the CR orbits if E20 / Z = 0.03 [ 121 ]. Magnetic Fields Beyond The Milky Way Psychology Essay Magnetic Fields Beyond The Milky Way Psychology Essay Published: Last Edited: 23rd March, 2015 This essay has been submitted by a student. This is not an example of the work written by our professional essay writers. In this paper magnetic fields of extragalactic objects have been investigated. Their causes, structure and field strength have been discussed in detail. It was found that magnetic fields play an important role within the universe and occupy many galaxies and clusters.

The strengths of magnetic fields within galaxies range between 6-10 Î¼G and those within clusters are between 0.2-0.3 Î¼G. Future projects have also been discussed, such as the SKA which will improve understandings of magnetic fields and subsequently the Universe. 1.1 History The existence and function of magnetic fields beyond our galaxy has, until fairly recently, been an unknown aspect of astrophysics (Kronberg [1]). The first discoveries of magnetic fields as an influence within the Universe were made just after the Second World War (Ruzmaikin et al. [2]). Studies of the Sun around this time were the influential means by which magnetic fields were discovered; not just in our very own galaxy, the Milky Way, but further out into the cosmos and within other galaxies. By studying the Sun, a firm understanding of the processes of magneto plasma was gained by making direct measurements of its magnetic field strength and direction. Further investigations into this area of study led to the discovery of synchrotron radiation (Schwinger [3]); the starting point for astronomers to investigate further into the prospect of extragalactic magnetic fields.

Magnetic fields are measured in microgauss (Î¼G), which is a unit of magnetic flux density equal to one millionth of a gauss. 1 Gauss is equal to 10-4 kg C-1 s-1, so 1 microgauss will be 10-6 x 10-4 kg C-1 s-1. A few attributes of galactic magnetic fields are that they have a vast spatial scale, which is expressed in kiloparsecs (kpc), and a modest strength. 1 parsec in SI units is equal to 30.86 x 1012 km. 1.2 Structure of Paper Throughout this paper I will be discussing several aspects of extragalactic magnetic fields. I will begin by talking about where magnetic fields can be found within the universe and in which types of galaxies they are located. This will lead on to looking into how the fields were formed and what caused them to exist. A major part of understanding these magnetic fields is being able to detect them and therefore measure their strength, so I will be discussing ways in which to do this. The two main techniques are by observing synchrotron radiation and Faraday rotation.

The processes behind these techniques, their uses and any difficulties or problems using them will be explained in depth. When making investigations on extragalactic magnetic fields, it is helpful to have an idea of their structure. I will be looking into how the structures of the fields can be determined, what can be deduced from them and how they differ in different types of objects. This will then lead on to considering the magnetic field strengths themselves. What do the strengths convey to us and how do they differ in different extragalactic objects? Finally, I will discuss future prospects for investigations of magnetic fields beyond our own galaxy and theories about how they may further our understanding of this subject. Extragalactic Magnetic Fields 2.1 Location Magnetic fields can be found almost everywhere within the Universe. A few of the astronomical aspects that will be discussed in this report are: spiral galaxies, elliptical and irregular galaxies, radio galaxies, galaxy clusters and cosmological magnetic fields. The structure and strength of the magnetic fields within these different areas of interest vary greatly and will be discussed further in Â§Â§ 4-5. The most investigated area, and probably the most favoured, of extragalactic magnetic fields for astronomers are those within spiral galaxies.

Spiral galaxies have a prominent disc with a number of distinct spiral arms. The spiral arms are much brighter than the disc and are therefore easily distinguished because they contain young stars and star-forming regions. There are currently in excess of 100 known spiral galaxies which have estimated field strengths of roughly 10Î¼G and a small amount of which have had their magnetic field structures detailed (Widrow [4]). Techniques used to measure the magnetic field strength in spiral galaxies, particularly our own, are: synchrotron radiation, Faraday rotation, Zeeman splitting and optical rotation. Elliptical galaxies are ellipsoidal in shape and can contain between hundreds of millions and over a trillion stars. Although magnetic fields are everywhere in elliptical galaxies, they are hard to observe because of the scarcity of relativistic electrons, which are needed for the emission of synchrotron radiation. The strengths of the fields are very hard to estimate within these galaxies; however some have been determined using synchrotron radiation and Faraday rotation in polarised radio emission from objects behind the galaxies. Irregular galaxies are ones which do not fit into the spiral or elliptical criteria.

They do not have any definite structure and a particular feature is that they are rather rich in gas and dust. Chyzy et al [5] is a paper investigating the magnetic fields of NGC 4449, a dwarf irregular galaxy. The findings in this paper are particularly interesting as the fields are regular and comparable to those within many spiral galaxies. The term active galaxy refers to the objects with the highest luminosity in the radio wavelength in the night's sky. There are three main types of active galaxies: radio galaxies, radio-loud quasars and blazars. These galaxies have a very recognisable structure in that they are extremely large elliptical galaxies with two extended radio jets of plasma. These jets emerge at relativistic speeds and carry material from the galaxy out to large lobes millions of light years away from the galactic centre (Hardcastle [6]). As the material is travelling at relativistic speeds, synchrotron radiation is emitted by electrons. This is the process by which the magnetic field is determined from these particular galaxies. Radio emission and rotation measure data have been used to demonstrate the existence of magnetic fields within galaxy clusters. Several thousands of galaxies can be grouped together to form galaxy clusters.

In the spaces between the galaxies is a very hot and dilute gas which is called the inter-galactic medium. The plasma within this medium is at a temperature, T 107 - 108 K (Widrow [4]). As non-thermal radio emission is detected from clusters, the synchrotron radiation method can be used to calculate the magnetic field strength. The plasma in the inter-galactic medium emits x-rays and is understood as thermal bremsstrahlung radiation from the hot gas. The assumptions about the origins of extragalactic magnetic fields are subsequently best put to test whilst investigating cluster magnetic fields. Cosmological magnetic fields are a rather more hypothetical idea. It is thought that a magnetic field could exist at the largest possible scale in the Universe. The theory that the Universe itself is a magnet of gigantic proportions is something that is of great interest to scientists as there is no definite proof confirming their existence. There are many arguments for this theory. One of which essentially being that almost every other object within the Universe has a magnetic field - planets, stars, galaxies, clusters - why shouldn't the Universe itself have one? Evidence supporting the theory is Faraday rotation of emission from high red shift sources.

A full explanation of this can be found in the paper by Widrow [4]. 2.2 Sources of Magnetic Fields When scientist first discovered galactic magnetic fields, their origin was greatly speculated. There had to be something producing the magnetic fields and subsequently maintaining them. As the possibility of an astrophysical "battery" is far-fetched, a standard galactic dynamo was found to be a sufficient power source. It was Sir Joseph Larmor, at the 87th Meeting of the British Association for the Advancement of Science in 1919, who first proposed the theory of the dynamo being responsible for astronomical magnetic fields within the Sun and the Earth. Soon after, Steenbeck, Krause and Rädler (1966) developed the dynamo theory further as helical turbulence (ï„ƒ) was found to be an important factor within stars and planets. Thus the ï„ƒÏ‰ dynamo was developed as the means by which the sun sustained its magnetic field. After much investigation it was decided that this could be applied to larger scale galactic magnetic fields as well as those within the Sun.

The ï„ƒÏ‰ dynamo helps to maintain and amplify the magnetic field within a galaxy by using the collective action of the differential rotation (Ï‰) and the helical turbulence (ï„ƒ) to regenerate a new field. This is based on Faraday's electromagnetic induction law: an electromotive force is induced within a conductor moving in a magnetic field, which then induces an electric current, in turn producing a magnetic field which modifies the first. A full description of the magnetic dynamo model is given in Widrow's paper [4]. The ï„ƒÏ‰ dynamo is the means by which a magnetic field can be amplified and maintained but is not responsible for the origin of the field, especially small fields. The explanation of the origin of these fields is a major unresolved issue within the discipline of astronomy and is therefore highly sought after. Over the years there have been many proposals for their origin but nothing definite has been confirmed. Two new proposed mechanisms for the generation of seed magnetic fields have been made recently in Shukla [7]. Both of these models use a collision of a neutrino beam and an electron stream to create seed magnetic fields.

The first of these uses the basis that an electric current can flow if positive and negative charges are separated. The process works by having a closely packed neutrino or electromagnetic beam which randomly ejects electrons, leaving behind an area of positive charges. The charge separation generates the electric current which then easily gives rise to a magnetic field. As much of the Universe is made up of plasma containing positive ions, electrons and positrons, this theory of charge separation is reasonably viable. The second of these mechanisms is very similar to the first. It occurs by having electron beams within plasma which generates electron and positron currents, thus generating a magnetic field. Detecting Magnetic Fields 3.1 Synchrotron Radiation Synchrotron radiation is very important for helping to determine magnetic field strengths in stellar objects. It is emitted by relativistic electrons which are circulating in a storage ring. This means that the electrons change direction due to an external magnetic field and are therefore forced to follow a spiral trajectory around the magnetic field lines. By accelerating in this characteristic manner, the electrons emit electromagnetic energy: synchrotron radiation. Synchrotron radiation was originally discovered when electrons were first accelerated to ultra-relativistic energies in betatron experiments.

In 1940 Kerst performed the first successful betatron experiment and it wasn't until nearly ten years later that Langmuir recognised the synchrotron emission from the experiments. It was found that this particular type of radiation is non-thermal, which means that it is a continuum radiation of particles whose energy spectrum is not Maxwellian. A particular radiation is also referred to as non-thermal if it cannot be accounted for by the spectrum of thermal bremsstrahlung or black body radiation. Further information on this can be found in Longair [8]. But why is this synchrotron radiation and not cyclotron radiation? The answer to this is merely that it is because of the relativistic speeds of the electrons. Cyclotron radiation only occurs for electrons travelling at slow speeds. Essentially synchrotron radiation is the relativistic equivalent to cyclotron radiation. Once discovered within the astronomical field, it was found that synchrotron radiation is responsible for the radio emission from many astronomical objects, such as other galaxies, supernova remnants, particular nebulae such as the Crab Nebula and other extragalactic radio sources. It has also been found that it could be responsible for optical and x-ray emission from quasars.

So how can the magnetic field strength be determined from synchrotron radiation? The main points of how this can be achieved are as follows in the next paragraph; the fully descriptive method can be found in Widrow [4]. To begin we look at the emissivity of an electron within a magnetic field as a function of frequency () and energy (): (1) Here is the magnetic field component which is perpendicular to the line of sight, is the critical frequency and (x) is a cutoff function. Equation (1) means that the synchrotron emission at frequency is dominated by electrons with energy, . The electron energy distribution, , is an important factor when considering the total synchrotron emission from a source. It can be expressed thus: (2) This is a basic power law in which the exponent is called the spectral index. The value of the spectral index for spiral galaxies is typically between. The energy density in relativistic electrons is: . Therefore the synchrotron radiation is related to the energy density of relativistic electrons and the magnetic field strength can be calculated by either assuming equipartition or by minimising the total energy density. The equipartition field is when the magnetic and particle energies are equal.

In turn this means that the emitting plasma has the minimum total energy for producing the observed radiation. The calculation of can be determined by integrating over the frequency, : (3) Here is the total flux density, is the angular size of the source and is any frequency between the upper and lower frequency limits. Eventually, after much manipulation (found in Widrow [4]), the equipartition magnetic field strength can be estimated to be: (4) One of the constraints against using the synchrotron radiation method is that the beam of radiation from the source must be within a specific region with respect to the line of sight. The angle between the magnetic field and the line of sight must be approximately equal to the pitch angle in order for the beam to be detected. If the pitch angle is small then the radiated power from the source is small; hence why the perpendicular component of the magnetic field is used in equation (1). The mean pitch angle is more often used for radio sources as they tend to be rather turbulent and unidirectional. More information on this area and on synchrotron radiation in general can be found in Hughes [9]. 3.2 Faraday Rotation Polarised radio waves have two circularly polarised components - left and right.

When the polarised radio waves propagate through the interstellar medium, the angle at which the left and right components are propagating is rotated by a magnetic field. In order for this to happen, the magnetic field must be acting parallel to the direction of propagation. This is called Faraday rotation and is mainly observed in radiation from pulsars. Faraday rotation is mostly used to determine the magnetic field strength of clusters. This works by having a galaxy cluster in between the observer and a source producing radio emission. The magnetic field of the cluster alters the angle of the right and left circularly polarised components of the background source and thus allows the strength to be calculated. An account of how this was first successfully achieved can be found in Carilli and Taylor [10]. The left and right-circularly polarized components propagate with different phase velocities (Widrow [4]). The following equations can be found in Ruzmaikin et al [2], Widrow [4] and Carilli and Taylor [10]. The rotation of the plane of polarization is given by: (5) where Î” is the change in the rotation angle, is the initial polarisation angle, is the wavelength of the radiation and RM is the Faraday rotation measure.

The rotation measure is given by the integral over the distance between the source ( = and the observer ( = 0: (6) This can be simplified to: (7) where is the charge of an electron, is the electron mass, is the speed of light, is the electron density and is the magnetic field strength along the line of sight. From equation (7) it is clear that the electron density must be known in order to determine the magnetic field strength. This is simple enough to do when looking at pulsars within our own galaxy as the electron density can be determined from the dispersion measure which, in turn, is obtained by measuring the relative pulse delay against the frequency relation (Kronberg [1]). It is not possible to use this method for other galaxies; another similar technique must be used. By looking at the dispersion measure, obtained from light curves, of the galaxy as a whole, it can be compared with the rotation measure and thus allows the magnetic field to be determined. The electron density can also be determined using a deprojection technique from x-ray observations of hot gas in clusters. A full description of this can be found in Sun et al.

[11] and there is more detailed information on how to use light curves of extragalactic objects in Kronberg [1]. It is notable that the RM can be positive or negative; positive for a magnetic field directed towards the observer and negative when directed away. The RM is also a combined value of all the regions containing magnetic fields along the line of sight. These regions are the Galaxy, the source itself and the intergalactic medium. This therefore means that the RM can be expressed thus: (8) Further explanations of this can be found in Kronberg and Perry [12]. 3.3 Further Techniques In addition to synchrotron radiation and Faraday rotation there are a few other techniques used to determine the magnetic field strength of extragalactic objects. One of these is Zeeman splitting which is done by measuring the splitting of spectral lines of a magnetic field, usually the 21cm neutral hydrogen absorption line. Subsequently this was the first method ever used to observe cosmic magnetic fields (Widrow [4]). Although a rather unique and interesting technique, it does have one major downfall; observations allowing for its detection can take several thousands of hours. Another method used is x-ray inverse Compton emission.

This involves the scattering of a photon to high energies by means of a relativistic electron colliding with it. As this process involves relativistic speeds, the energy gain of the photon in the lab frame is approximately equal to the square of the Lorentz factor, . There is no energy change in the rest frame of the electron as this is Thomson Scattering. More information can be found on this method in Hardcastle [6]. Other methods include polarization of gas and dust and polarization of optical starlight. Information on these methods can be found in Ruzmaikin et al. [2], Widrow [4] and Longair [8]. Structure 4.1 Spiral Galaxies Figure 1 (Widrow [4])The structure of magnetic fields within other galaxies is determined by looking at the rotation measure and polarization of synchrotron emission. Fig. 1 shows simply how the magnetic field can be formed within a spiral galaxy. The top two pictures show the axisymmetric and bisymmetric field configurations near the equatorial plane. The two graphs below depict the rotation measure (RM), for observations made at the dotted circle, a chosen distance from the galactic nucleus, as a function of the azimuthal angle ().

The rotation measure can be plotted against the azimuthal angle at different distances from the centre of the galaxy and show how the magnetic field fluctuates with respect to distance. Although this method is very good, it does have its downfalls (Ruzmaikin et al. [13]). One of these being that it cannot determine a magnetic field structure which consists of superposition of many different modes. It also requires that observations of the rotation measure be taken for many different wavelengths. An alternative method is by modelling as a Fourier series (Widrow [4]). This is the polarization angle as a function of the azimuthal angle, which gives the series: (9) Thus, and then produce a diagram of the azimuthal structure of the magnetic field. Figure 2 (Kotarba et al. [14])Figure 2 shows an optical image of the spiral galaxy M51 (taken by Hubble) with the synchrotron intensity contour lines and the magnetic field vectors superimposed upon it. The synchrotron intensity contours are the white lines and the magnetic field vectors are the small yellow lines.

This image shows how the magnetic field structure appears to follow the structure of the galaxy itself, spiralling inwards from the outer reaches of the arms, into the very centre of the galaxy. It also displays how the gas and dust within the galaxy influences the magnetic field. Very much the same idea can be applied to other spiral galaxies such as M31. More information on this can be found in Beck's review of the magnetic field within the spiral galaxy M31 [15]. Gaseous halos around galaxies are also a good indicator of cosmic rays and magnetic fields within a galaxy. Studies of the galaxy NGC 253, an intermediate galaxy, show that the cosmic rays twist around the magnetic field lines and therefore produce a magnetic field component which is in the vertical direction from the disc of the galaxy up towards the halo. Heesen et al. [16] describes how the superwinds of the galaxy may also influence the structure of the magnetic fields. Figure 3 (Heesen et al. [16]) Figure 4 (Heesen et al. [16]) Figure 3 is the final comprised image from this paper depicting the field vectors around the centre of the galaxy.

It is visible here that the field vectors are following the structure of the two spiral arms of this galaxy, just as they are in M51. One more step has been taken with the investigations of NGC 253 and the structure of the poloidal magnetic field has also been determined. This is depicted in figure 4. It shows a noticeable X-shaped field structure over the centre of the x-ray emission image of the galaxy. This particular X-shaped halo field is mainly characteristic of edge-on views of some galaxies and is currently an area of study which is in deep investigation. The hypothesized theory as to the reason for this characteristic, given in Heesen et al. [16], is that it is due to the interactions of the halo gas with the superwinds from the galactic centre. 4.2 Radio Sources and Galaxy Clusters Laing [17] has written a paper in which a model for the magnetic field structure in extended radio sources has been designed. This model was found to correspond well with observations made of radio sources and with filaments of the Crab Nebula which have also been investigated. Figure 5 shows how this model has been created by taking a randomly orientated field within a certain volume and compressing it.

Figure 5 (Laing [17]) This therefore means that the field is directed still in a random way but only in one plane. The purpose of doing this is to show that synchrotron radiation emitted with a high degree of polarisation does not always mean that the magnetic field is uniform and thus can be limited to one plane containing the line of sight. It is also to show how the degree of polarised radiation from the magnetic field slab changes when the plane changes angle, Î², with respect to the line of sight. From investigations into the degree of polarisation it has been found that the degree is greatest at the edge of a source, which is when the plane of compression is parallel to the line of sight. A deeper insight into this area of investigation can be found in the paper by Laing [17]. Figure 6 (Laing[17])For the filaments of the Crab Nebula, the magnetic field and relativistic particles are moving away from the central pulsar. They collide with slower moving, dense filaments and create regions of high compressed field around them. It is thought that the uniform regions of this field are weak and thus a cylindrical shell model is created as shown in figure 6.

This very thin shell is wrapped around one side of the filament then allowing the field configuration model to be the same as that in figure 5 except with Î² changing across the filament (Laing [17]). As mentioned before in Â§ 3, Faraday rotation is the method by which magnetic fields have been discovered within galaxy clusters. It is thought that the structure of these galaxy clusters is influenced by the velocity field of the cluster formation (Dolag et al. [18]). The paper by Dolag et al. [18] describes how the evolution and final structure of the magnetic fields of clusters has been investigated. This was done by performing cosmological, magneto-hydrodynamic simulations of 100 different cluster models. By changing the specifications of the cluster models, i.e. mass, structure and strength of the primary field, many results illustrating the final strength and structure of the magnetic field can be found. In a similar paper by Donnert et al. [19] the possibility of galactic outflows being responsible for the magnetic fields within clusters is discussed.

By using the same method of performing the cosmological, magneto-hydrodynamic simulations, they have determined that magnetic galactic outflows and their evolution are entirely accountable for the magnetic fields observed within galaxy clusters today. However, although they are accountable for the fields themselves, they are not responsible for the final structure of the fields. This confirms that the structure of magnetic fields within galaxy clusters is indeed influenced by the velocity field caused by cluster formation. Magnetic Field Strengths As mentioned in Â§1, the magnetic fields are measured in microgauss (Î¼G) which is much smaller than the field strength in stars (roughly between 103 to 1013 G (Kunze [20])). The strengths of the fields within different objects in the Universe vary greatly, which can be seen in Table 1. Investigations into the field strengths of different extragalactic objects can help astronomers obtain a better understanding of evolution and origin of the fields. It can also help to improve knowledge about the early Universe. Table 1 depicts the estimated magnetic field strengths of a few different extragalactic objects. This allows comparisons to be made between the different objects.

Extragalactic object Average magnetic field strength/ Î¼G Spiral galaxies Table 1 (figures taken from Widrow [4]) *Here for Radio-halo clusters, L is the typical scale over which the field reverses direction. From table 1 it is clear that galaxy clusters have the smallest magnetic field and spiral and elliptical galaxies have the largest. Reasons for this are unclear, however as it was mentioned above that stars have field strengths of the scale 103-1013 G, and that these fields are getting weaker as the size of the object increases, it may be passable to say that the magnetic field strength of an object depends on its size. Future Prospects 6.1 Investigating in the future Looking into the future is something that is very important within the astronomical field. One of the proposed projects that will be undertaken on Earth is the Square Kilometre Array (SKA). There are five main investigations to be carried out using the SKA, one of which will be looking at the origin and evolution of cosmic magnetism.

The SKA will be able to outline the structure of magnetic fields in the innermost regions of galaxies in great detail and will also be able to detect polarized emission from distant galaxies (Gaensler [21]). Figure 7 shows an artist's impression of what the array will look like. It will occupy an area of 1 km2 and have fully steerable dishes and passive aperture tiles. There are two possible sites for this array and those are near Boolardy station in Australia and in the Karoo region of South Africa. The specifications of the instrument will be immense and the details of its flexibility can be found in Gaensler [21]. Figure 7 (Gaensler [21]) Details of other future projects similar to the SKA, including the EVLA (Extended Very Large Array), LOFAR (Low Frequency Array), MWA (Murchison Widefield Array) and the LWA (Long Wavelength Array) can be found in Beck [22]. 6.2 How will this further our understanding? There are three main questions that astronomers hope to answer using the SKA. These are: What is the strength and structure of magnetic fields in the Milky Way, in other galaxies and in galaxy clusters?

How have magnetic fields evolved in galaxies and clusters over time? How and when was the Universe magnetised? Details of how scientists intend to answer these questions can be found in Gaensler [21]. The SKA will improve the density of background rotation measures which in turn will allow for a new and improved three-dimensional Galactic magnetic field map. It will also allow a deeper understanding of magnetic fields within nearby galaxies and also help to produce any correlation between galaxy and magnetic field properties. Further information on the SKA can be found in Murphy [23]. All of the future methods mentioned in this section will help further understanding of magnetic fields within the Universe. Scientists hope to determine the causes of the magnetic fields and investigate any relation to star formation this may have. A deeper understanding of the evolution of the fields within galaxies and the ICM will also be determined. Concluding Thoughts In this paper many important and recently investigated aspects of magnetic fields beyond the Milky Way have been discussed. Magnetic fields are an interesting area of study within the astronomical field as until fairly recently not very much has been known about them. They occupy a great number of extragalactic objects and have varying strengths and structures.

A reason for the origin of these fields is something that is still being speculated today and will continue to be investigated in future projects. Although the source is as yet undecided, the ï„ƒÏ‰ dynamo is the means by which the magnetic fields are maintained. The two main processes by which the magnetic field strengths are calculated are synchrotron radiation and Faraday rotation. These have been used to determine the field strengths of many different extragalactic objects, such as spiral galaxies, radio sources and galaxy clusters. It was found that the average magnetic field strengths of the clusters are much weaker than those within galaxies themselves. Comparing these with the average magnetic field strength of stars, which in comparison to galaxies are small, compact objects, it could be concluded that as the object increases in size, the field is more spread out and therefore weaker. The structure of the fields within different galaxies has also been researched. Those within spiral galaxies appear to follow the spiral pattern of the galaxy itself, whereas in elliptical galaxies the fields are ubiquitous. This could mean that the shape of the galaxy is dictated by its magnetic field, especially within spiral galaxies. Extragalactic magnetic fields are one of the most interesting aspects of astronomy today.

Furthermore, findings from future projects, such as the SKA, may resolve issues such as the origin and evolution of the fields. These projects will also widen our knowledge of this area of study and give a better understanding of the Universe itself.

Question: Extragalactic refers to beyond the?
Answer: Milky Way