Planck 2015 results
XIX. Constraints on primordial magnetic fields
^{1} APC, AstroParticule et Cosmologie, Université Paris Diderot, CNRS/IN2P3, CEA/lrfu, Observatoire de Paris, Sorbonne Paris Cité, 10 rue Alice Domon et Léonie Duquet, 75205 Paris Cedex 13, France
^{2} Aalto University Metsähovi Radio Observatory and Dept of Radio Science and Engineering, PO Box 13000, 00076 AALTO, Finland
^{3} African Institute for Mathematical Sciences, 6–8 Melrose Road, Muizenberg 7945, Cape Town, South Africa
^{4} Agenzia Spaziale Italiana Science Data Center, via del Politecnico snc, 00133 Roma, Italy
^{5} Aix Marseille Université, CNRS, LAM (Laboratoire d’Astrophysique de Marseille) UMR 7326, 13388 Marseille, France
^{6} Astrophysics Group, Cavendish Laboratory, University of Cambridge, J J Thomson Avenue, Cambridge CB3 0HE, UK
^{7} Astrophysics & Cosmology Research Unit, School of Mathematics, Statistics & Computer Science, University of KwaZuluNatal, Westville Campus, Private Bag X54001, Durban 4000, South Africa
^{8} CGEE, SCS Qd 9, Lote C, Torre C, 4o andar, Ed. Parque Cidade Corporate, CEP 70308200, Brasília, DF, Brazil
^{9} CITA, University of Toronto, 60 St. George St., Toronto, ON M5S 3H8, Canada
^{10} CNRS, IRAP, 9 Av. colonel Roche, BP 44346, 31028 Toulouse Cedex 4, France
^{11} CRANN, Trinity College, Dublin 2, Ireland
^{12} California Institute of Technology, Pasadena, California, CA 91125, USA
^{13} Centre for Theoretical Cosmology, DAMTP, University of Cambridge, Wilberforce Road, Cambridge CB3 0WA, UK
^{14} Centro de Estudios de Física del Cosmos de Aragón (CEFCA), Plaza San Juan, 1, planta 2, 44001 Teruel, Spain
^{15} Computational Cosmology Center, Lawrence Berkeley National Laboratory, Berkeley, California, CA 94720, USA
^{16} Consejo Superior de Investigaciones Científicas (CSIC), 28006 Madrid, Spain
^{17} DSM/Irfu/SPP, CEASaclay, 91191 GifsurYvette Cedex, France
^{18} DTU Space, National Space Institute, Technical University of Denmark, Elektrovej 327, 2800 Kgs. Lyngby, Denmark
^{19} Département de Physique Théorique, Université de Genève, 24, Quai E. Ansermet, 1211 Genève 4, Switzerland
^{20} Departamento de Astrofísica, Universidad de La Laguna (ULL), 38206 La Laguna, Tenerife, Spain
^{21} Departamento de Física, Universidad de Oviedo, Avda. Calvo Sotelo s/n, 33007 Oviedo, Spain
^{22} Department of Astronomy and Astrophysics, University of Toronto, 50 Saint George Street, Toronto, Ontario, Canada
^{23} Department of Astrophysics/IMAPP, Radboud University Nijmegen, PO Box 9010, 6500 GL Nijmegen, The Netherlands
^{24} Department of Physics & Astronomy, University of British Columbia, 6224 Agricultural Road, Vancouver, British Columbia, Canada
^{25} Department of Physics and Astronomy, Dana and David Dornsife College of Letter, Arts and Sciences, University of Southern California, Los Angeles, CA 90089, USA
^{26} Department of Physics and Astronomy, Johns Hopkins University, Bloomberg Center 435, 3400 N. Charles St., Baltimore, MD 21218, USA
^{27} Department of Physics and Astronomy, University College London, London WC1E 6BT, UK
^{28} Department of Physics, Florida State University, Keen Physics Building, 77 Chieftan Way, Tallahassee, Florida, USA
^{29} Department of Physics, Gustaf Hällströmin katu 2a, University of Helsinki, 00014 Helsinki, Finland
^{30} Department of Physics, Princeton University, Princeton, New Jersey, NJ 08544, USA
^{31} Department of Physics, University of California, Santa Barbara, California, CA 93106, USA
^{32} Department of Physics, University of Illinois at UrbanaChampaign, 1110 West Green Street, Urbana, Illinois, USA
^{33} Dipartimento di Fisica e Astronomia G. Galilei, Università degli Studi di Padova, via Marzolo 8, 35131 Padova, Italy
^{34} Dipartimento di Fisica e Astronomia, ALMA MATER STUDIORUM, Università degli Studi di Bologna, Viale Berti Pichat 6/2, 40127 Bologna, Italy
^{35} Dipartimento di Fisica e Scienze della Terra, Università di Ferrara, via Saragat 1, 44122 Ferrara, Italy
^{36} Dipartimento di Fisica, Università La Sapienza, P. le A. Moro 2, 00185 Roma, Italy
^{37} Dipartimento di Fisica, Università degli Studi di Milano, via Celoria 16, 20133 Milano, Italy
^{38} Dipartimento di Fisica, Università degli Studi di Trieste, via A. Valerio 2, 34127 Trieste, Italy
^{39} Dipartimento di Matematica, Università di Roma Tor Vergata, via della Ricerca Scientifica 1, 00133 Roma, Italy
^{40} Discovery Center, Niels Bohr Institute, Blegdamsvej 17, 2100 Copenhagen, Denmark
^{41} Discovery Center, Niels Bohr Institute, Copenhagen University, Blegdamsvej 17, 2100 Copenhagen, Denmark
^{42} European Space Agency, ESAC, Planck Science Office, Camino bajo del Castillo, s/n, Urbanización Villafranca del Castillo, 28692 Villanueva de la Cañada, Madrid, Spain
^{43} European Space Agency, ESTEC, Keplerlaan 1, 2201 AZ Noordwijk, The Netherlands
^{44} Gran Sasso Science Institute, INFN, viale F. Crispi 7, 67100 L’ Aquila, Italy
^{45} HGSFP and University of Heidelberg, Theoretical Physics Department, Philosophenweg 16, 69120 Heidelberg, Germany
^{46} Helsinki Institute of Physics, Gustaf Hällströmin katu 2, University of Helsinki, 00100 Helsinki, Finland
^{47} INAF–Osservatorio Astrofisico di Catania, via S. Sofia 78, 95123 Catania, Italy
^{48} INAF–Osservatorio Astronomico di Padova, Vicolo dell’Osservatorio 5, 35122 Padova, Italy
^{49} INAF–Osservatorio Astronomico di Roma, via di Frascati 33, 00040 Monte Porzio Catone, Italy
^{50} INAF–Osservatorio Astronomico di Trieste, via G.B. Tiepolo 11, 40127 Trieste, Italy
^{51} INAF/IASF Bologna, via Gobetti 101, 40129 Bologna, Italy
^{52} INAF/IASF Milano, via E. Bassini 15, 20133 Milano, Italy
^{53} INFN, Sezione di Bologna, via Irnerio 46, 40126 Bologna, Italy
^{54} INFN, Sezione di Roma 1, Università di Roma Sapienza, Piazzale Aldo Moro 2, 00185 Roma, Italy
^{55} INFN, Sezione di Roma 2, Università di Roma Tor Vergata, via della Ricerca Scientifica 1, 00185 Roma, Italy
^{56} INFN/National Institute for Nuclear Physics, via Valerio 2, 34127 Trieste, Italy
^{57} IPAG: Institut de Planétologie et d’Astrophysique de Grenoble, Université Grenoble Alpes, IPAG, CNRS, IPAG, 38000 Grenoble, France
^{58} IUCAA, Post Bag 4, Ganeshkhind, Pune University Campus, Pune 411 007, India
^{59} ImperialCollege London, Astrophysics group, Blackett Laboratory, Prince Consort Road, London, SW7 2AZ, UK
^{60} Infrared Processing and Analysis Center, California Institute of Technology, Pasadena, CA 91125, USA
^{61} Institut Néel, CNRS, Université Joseph Fourier Grenoble I, 25 rue des Martyrs, Grenoble, France
^{62} Institut Universitaire de France, 103, bd SaintMichel, 75005 Paris, France
^{63} Institut d’Astrophysique Spatiale, CNRS, Univ. ParisSud, Université ParisSaclay, Bât. 121, 91405 Orsay Cedex, France
^{64} Institut d’Astrophysique de Paris, CNRS (UMR 7095), 98 bis boulevard Arago, 75014 Paris, France
^{65} Institut für Theoretische Teilchenphysik und Kosmologie, RWTH Aachen University, 52056 Aachen, Germany
^{66} Institute for Space Sciences, 077125 BucharestMagurale, Romania
^{67} Institute of Astronomy, University of Cambridge, Madingley Road, Cambridge CB3 0HA, UK
^{68} Institute of Theoretical Astrophysics, University of Oslo, Blindern, 0371 Oslo, Norway
^{69} Instituto de Astrofísica de Canarias, C/vía Láctea s/n, La Laguna, 38205 Tenerife, Spain
^{70} Instituto de Física de Cantabria (CSICUniversidad de Cantabria), Avda. de los Castros s/n, 39005 Santander, Spain
^{71} Istituto Nazionale di Fisica Nucleare, Sezione di Padova, via Marzolo 8, 35131 Padova, Italy
^{72} Jet Propulsion Laboratory, California Institute of Technology, 4800 Oak Grove Drive, Pasadena, California, CA 31109, USA
^{73} Jodrell Bank Centre for Astrophysics, Alan Turing Building, School of Physics and Astronomy, The University of Manchester, Oxford Road, Manchester, M13 9PL, UK
^{74} Kavli Institute for Cosmological Physics, University of Chicago, Chicago, IL 60637, USA
^{75} Kavli Institute for Cosmology Cambridge, Madingley Road, Cambridge, CB3 0HA, UK
^{76} Kazan Federal University, 18 Kremlyovskaya St., 420008 Kazan, Russia
^{77} LAL, Université ParisSud, CNRS/IN2P3, 91898 Orsay, France
^{78} LERMA, CNRS, Observatoire de Paris, 61 Avenue de l’Observatoire, 75014 Paris, France
^{79} Laboratoire AIM, IRFU/Service d’Astrophysique – CEA/DSM – CNRS – Université Paris Diderot, Bât. 709, CEASaclay, 91191 GifsurYvette Cedex, France
^{80} Laboratoire Traitement et Communication de l’Information, CNRS (UMR 5141) and Télécom ParisTech, 46 rue Barrault, 75634 Paris Cedex 13, France
^{81} Laboratoire de Physique Subatomique et Cosmologie, Université GrenobleAlpes, CNRS/IN2P3, 53 rue des Martyrs, 38026 Grenoble Cedex, France
^{82} Laboratoire de Physique Théorique, Université ParisSud 11 & CNRS, Bâtiment 210, 91405 Orsay, France
^{83} Lawrence Berkeley National Laboratory, Berkeley, California, CA 94720, USA
^{84} Lebedev Physical Institute of the Russian Academy of Sciences, Astro Space Centre, 84/32 Profsoyuznaya st., Moscow, GSP7, 117997 Moscow, Russia
^{85} Leung Center for Cosmology and Particle Astrophysics, National Taiwan University, 10617 Taipei, Taiwan
^{86} MaxPlanckInstitut für Astrophysik, KarlSchwarzschildStr. 1, 85741 Garching, Germany
^{87} McGill Physics, Ernest Rutherford Physics Building, McGill University, 3600 rue University, Montréal, QC, H3A 2T8, Canada
^{88} National University of Ireland, Department of Experimental Physics, Maynooth, Co. Kildare, Ireland
^{89} Nicolaus Copernicus Astronomical Center, Bartycka 18, 00716 Warsaw, Poland
^{90} Niels Bohr Institute, Blegdamsvej 17, 2100 Copenhagen, Denmark
^{91} Niels Bohr Institute, Copenhagen University, Blegdamsvej 17, 2100 Copenhagen, Denmark
^{92} Nordita (Nordic Institute for Theoretical Physics), Roslagstullsbacken 23, 106 91 Stockholm, Sweden
^{93} Optical Science Laboratory, University College London, Gower Street, WC1E 6 BT London, UK
^{94} SISSA, Astrophysics Sector, via Bonomea 265, 34136 Trieste, Italy
^{95} SMARTEST Research Centre, Università degli Studi eCampus, via Isimbardi 10, 22060 Novedrate (CO), Italy
^{96} School of Physics and Astronomy, Cardiff University, Queens Buildings, The Parade, Cardiff, CF24 3AA, UK
^{97} School of Physics and Astronomy, University of Nottingham, Nottingham NG7 2RD, UK
^{98} Sorbonne UniversitéUPMC, UMR 7095, Institut d’Astrophysique de Paris, 98 bis boulevard Arago, 75014 Paris, France
^{99} Space Research Institute (IKI), Russian Academy of Sciences, Profsoyuznaya Str, 84/32, Moscow 117997, Russia
^{100} Space Sciences Laboratory, University of California, Berkeley, California, CA 94720, USA
^{101} Special Astrophysical Observatory, Russian Academy of Sciences, Nizhnij Arkhyz, Zelenchukskiy region, 369167 KarachaiCherkessian Republic, Russia
^{102} Stanford University, Dept of Physics, Varian Physics Bldg, 382 via Pueblo Mall, Stanford, California, CA 943054060, USA
^{103} SubDepartment of Astrophysics, University of Oxford, Keble Road, Oxford OX1 3RH, UK
^{104} The Oskar Klein Centre for Cosmoparticle Physics, Department of Physics, Stockholm University, AlbaNova, 106 91 Stockholm, Sweden
^{105} Theory Division, PHTH, CERN, 1211, Geneva 23, Switzerland
^{106} UPMC Univ Paris 06, UMR7095, 98 bis boulevard Arago, 75014 Paris, France
^{107} Université de Toulouse, UPSOMP, IRAP, 31028 Toulouse Cedex 4, France
^{108} University Observatory, Ludwig Maximilian University of Munich, Scheinerstrasse 1, 81679 Munich, Germany
^{109} University of Granada, Departamento de Física Teórica y del Cosmos, Facultad de Ciencias, 18071 Granada, Spain
^{110} University of Granada, Instituto Carlos I de Física Teórica y Computacional, 18071 Granada, Spain
^{111} Warsaw University Observatory, Aleje Ujazdowskie 4, 00478 Warszawa, Poland
Received: 5 February 2015
Accepted: 19 February 2016
We compute and investigate four types of imprint of a stochastic background of primordial magnetic fields (PMFs) on the cosmic microwave background (CMB) anisotropies: the impact of PMFs on the CMB temperature and polarization spectra, which is related to their contribution to cosmological perturbations; the effect on CMB polarization induced by Faraday rotation; the impact of PMFs on the ionization history; magneticallyinduced nonGaussianities and related nonzero bispectra; and the magneticallyinduced breaking of statistical isotropy. We present constraints on the amplitude of PMFs that are derived from different Planck data products, depending on the specific effect that is being analysed. Overall, Planck data constrain the amplitude of PMFs to less than a few nanoGauss, with different bounds that depend on the considered model. In particular, individual limits coming from the analysis of the CMB angular power spectra, using the Planck likelihood, are B_{1 Mpc} < 4.4 nG (where B_{1 Mpc} is the comoving field amplitude at a scale of 1 Mpc) at 95% confidence level, assuming zero helicity. By considering the Planck likelihood, based only on parityeven angular power spectra, we obtain B_{1 Mpc} < 5.6 nG for a maximally helical field. For nearly scaleinvariant PMFs we obtain B_{1 Mpc} < 2.0 nG and B_{1 Mpc} < 0.9 nG if the impact of PMFs on the ionization history of the Universe is included in the analysis. From the analysis of magneticallyinduced nonGaussianity, we obtain three different values, corresponding to three applied methods, all below 5 nG. The constraint from the magneticallyinduced passivetensor bispectrum is B_{1 Mpc} < 2.8 nG. A search for preferred directions in the magneticallyinduced passive bispectrum yields B_{1 Mpc} < 4.5 nG, whereas the compensatedscalar bispectrum gives B_{1 Mpc} < 3 nG. The analysis of the Faraday rotation of CMB polarization by PMFs uses the Planck power spectra in EE and BB at 70 GHz and gives B_{1 Mpc} < 1380 nG. In our final analysis, we consider the harmonicspace correlations produced by Alfvén waves, finding no significant evidence for the presence of these waves. Together, these results comprise a comprehensive set of constraints on possible PMFs with Planck data.
Key words: magnetic fields / cosmic background radiation / early Universe
© ESO, 2016
1. Introduction
1.1. Cosmic magnetism
Magnetic fields are one of the fundamental and ubiquitous components of our Universe. They are a common feature of many astrophysical objects, starting from the smallest up to the largest observed scales (for reviews see Ryu et al. 2012; and Widrow et al. 2012). In particular, largescale magnetic fields are observed in almost every galaxy, starting from the Milky Way, with possible hints of their presence in highredshift galaxies too (Beck 2000; Bernet et al. 2008; Wolfe et al. 2008), suggesting an early origin for the galactic fields.
Largescale magnetic fields are also probed in galaxy clusters, both through the measurement of the Faraday rotation effect on the light of background galaxies and through radio emission from the halos and relics of the clusters (Govoni & Feretti 2004; Ferrari et al. 2008; Feretti et al. 2012). These largescale magnetic fields have measured amplitudes that range from a few to several microgauss. Recent Faraday rotation measurements from lowdensity intercluster regions also suggest the presence of largescale magnetic fields in cosmic structure filaments (Neronov et al. 2013a). A recent addition to the largescale magnetic field constraints comes from the interpretation of FermiLAT data (Neronov & Vovk 2010). Highenergy γrays (in the TeV band) emitted by blazars generate electronpositron pairs when interacting with the optical and infrared background light. These pairs reemit at lower energies (GeV), but FermiLAT data do not show this reemitted flux in the GeV band. One possible explanation for this observation is a deflection of the electronpositron pairs owing to the presence of a diffuse magnetic field (see Neronov & Vovk 2010; Taylor et al. 2011; Tavecchio et al. 2011; Vovk et al. 2012; and Neronov et al. 2013b for details on this explanation; and see Broderick et al. 2012 for alternative scenarios). Constraints on the GeV emission provide lower limits on the amplitude of intergalactic fields of the order of 10^{18}–10^{15} G (Tavecchio et al. 2010; Taylor et al. 2011; Dermer et al. 2011; Vovk et al. 2012), if this scenario is correct.
The origin of largescale magnetic fields is strongly debated. Several mechanisms have been proposed and one popular hypothesis is that the observed largescale fields are remnants of fields that existed from the earliest times, i.e., primordial fields. During structure formation the adiabatic compression and turbulent shock flows would naturally lead to an amplification of initial seeds (which may act in addition to astrophysical mechanisms of largescale magnetic field generation, like AGN ejection and galactic dynamos; for reviews see Widrow 2002; and Giovannini 2004b). A Kolmogorovlike magnetic power spectrum has been observed in the central region of the Hydra cluster (Kuchar & Enßlin 2011), supporting the idea that the observed extragalactic magnetic fields are largely shaped and amplified by hydrodynamical processes. Primordial magnetic fields (PMFs) can naturally provide the initial seeds to be amplified into the observed largescale fields. Several earlyUniverse scenarios predict the generation of cosmological magnetic fields, either during inflation (Ratra 1992), with a suitable breaking of conformal invariance of electromagnetism (Turner & Widrow 1988), during phase transitions motivated by particle physics (Vachaspati 1991; Grasso & Riotto 1998), or via other physical processes (Durrer & Caprini 2003; Ichiki et al. 2006). The importance of PMF studies not only lies in the possibility of PMFs being the progenitors of the observed cosmic magnetic fields, but also in them providing a new potential observational window to the early Universe (for reviews see Kahniashvili 2005; Giovannini 2008; Kunze 2013; and Durrer & Neronov 2013).
1.2. Imprints of primordial magnetism
PMFs leave imprints on several cosmological observables and can be constrained with different cosmological data sets. In particular, interesting constraints come from their influence on Big Bang Nucleosynthesis (BBN), which provides upper limits of the order of 0.1 μG (Grasso & Rubinstein 1995; Kahniashvili et al. 2010), and from their impact on largescale structure formation (see, e.g., Shaw & Lewis 2012 and Fedeli & Moscardini 2012).
Another limit on PMFs based on BBN is derived by Caprini & Durrer (2002). Gravitational waves can be produced by PMFs, in particular before neutrino free streaming. The upper limit on the amount of gravitational waves allowed at nucleosynthesis to not spoil the BBN predictions poses a constraint on the amplitude of PMFs, which is especially strong for causal magnetogenesis mechanisms.
It has been suggested that PMFs could have an influence on the formation of the filamentary largescale structure, in the pioneering work of Wasserman (1978) and by Kim et al. (1996). In particular, Battaner et al. (1997) conclude that magnetic fields with comoving strengths lower than 1 nG have negligible effect and that for strengths higher than 10 nG the influence is too high to be compatible with observations.
Other constraints on PMFs come from their impact on the thermal spectrum of the cosmic microwave background (CMB) radiation. PMFs can induce spectral distortions of both early and late type via injection of dissipated magnetic energy into the plasma through damping processes (Jedamzik et al. 2000; Sethi & Subramanian 2005; Kunze & Komatsu 2014). PMFs also imply photon emission and absorption through the cyclotron process, which, depending on the amplitude of PMFs, could in principle play a role in the generation and evolution of spectral distortions and in particular in the thermalization process (for a general overview see Burigana et al. 1991; Hu & Silk 1993; and Chluba & Sunyaev 2012). On the other hand, the cyclotron process is only relevant at very long wavelengths. Thus, for realistic shapes of distorted spectra and PMFs with amplitudes compatible with current constraints, the cyclotron contribution is found to be much less important than radiative Compton and bremsstrahlung contributions in reestablishing a blackbody spectrum (Burigana & Zizzo 2006) and also extremely small in generating a polarized signal (Zizzo & Burigana 2005). Thus, the current constraints derived from COBEFIRAS present good prospects for possible future observations (Jedamzik et al. 2000; Kunze & Komatsu 2014; Chluba et al. 2015) while polarization anisotropies directly induced by PMFs are not significantly affected by the cyclotron process associated with PMFs.
Many of the stronger and more robust constraints come from the impact of PMFs on the CMB anisotropies. CMB data are a crucial source of information for investigating and constraining PMF characteristics, understanding their origin, and exploring the possibility of them being the seeds that generated the observed largescale magnetic fields.
The first release of Planck^{1} data in 2013 has led to some of the most stringent constraints on PMFs (see Planck Collaboration XVI 2014). The scope of this paper is to provide the “Planck constraints on PMFs” through combined analyses of the temperature and polarization data. The results of this paper are derived from Planck products, which are based on the work done in Planck Collaboration I (2016), Planck Collaboration II (2016), Planck Collaboration III (2016), Planck Collaboration IV (2016), Planck Collaboration V (2016), Planck Collaboration VI (2016), Planck Collaboration VII (2016), and Planck Collaboration VIII (2016). Before going into the description of our analysis, we briefly discuss the most important PMF models.
The simplest PMF model is that of a homogeneous field. This model cannot be included in a homogeneously and isotropically expanding cosmological model. It needs to be analysed in the context of an anisotropic cosmological model with associated isotropybreaking predictions (Kahniashvili et al. 2008), and has already been strongly constrained by COBE data (Barrow et al. 1997). More recently, Adamek et al. (2011) have reconsidered the impact of a homogeneous largescale magnetic field on the CMB anisotropies, with the addition of the contribution from freestreaming particles like neutrinos. The presence of the anisotropic neutrino stress induces a compensation of the magnetic field effect, allowing for a larger homogeneousfield amplitude than concluded from previous analyses.
The most widely used model of PMFs is a stochastic background modelled as a fully inhomogeneous component of the cosmological plasma (we neglect PMF energy density and anisotropic stress contributions at the homogeneous level), with the energy momentum tensor components (quadratic in the fields) on the same footing as cosmological perturbations. Within this model, PMFs leave several signatures in the CMB temperature and polarization anisotropy patterns, also inducing nonGaussianities.
In this paper, we predict and analyse four different types of PMF signatures: the impact on the CMB power spectra in temperature and polarization; the impact on polarization power spectra induced by Faraday rotation; the impact on CMB nonGaussianities and the related nonzero magneticallyinduced bispectra; and finally the impact on the statistics of CMB anisotropies and in particular the breaking of statistical isotropy, given by induced correlations in harmonic space.
The energy momentum tensor of PMFs sources all types of cosmological perturbations, i.e., scalar, vector, and tensor perturbations. Magneticallyinduced pertubations have some crucial differences with respect to the primary perturbations. Firstly, PMFs generate vector perturbations that are sourced by the Lorentz force and, unlike the primary ones, are not decaying. Secondly, magneticallyinduced perturbations are not suppressed by Silk damping (Hu & White 1997; Subramanian & Barrow 1998a). Their impact on the CMB temperature power spectrum is dominant on small angular scales, where the primary CMB is suppressed, and therefore they can be strongly constrained by highresolution CMB data. Of additional interest are helical PMFs, which may be generated during inflation through mechanisms like pseudoscalar coupling and may be crucial for connecting these PMFs to the magnetic fields observed on large scales. A helical PMF generates parityviolating correlations, such as TB and EB (Caprini et al. 2004; Kahniashvili & Ratra 2005; Kahniashvili et al. 2014), which can be used to constrain PMFs with CMB polarization data. After recombination, PMFs are dissipated via two additional effects that take place in the magnetized, not fully ionized, plasma: ambipolar diffusion and magnetohydrodynamical (MHD) turbulence. The dissipation injects magnetic energy into the plasma, heating it and thereby modifying the optical depth of recombination, with an impact on the primary CMB power spectra (Sethi & Subramanian 2005; Kunze & Komatsu 2014, 2015; Chluba et al. 2015).
PMFs have another effect on the primary polarization anisotropies. They induce Faraday rotation, which rotates Emodes into Bmodes, thus generating a new Bmode signal, and vice versa. This signal grows with decreasing observational frequency and therefore is a good target for Planck’s lowfrequency channels.
PMFs modelled as a stochastic background have a nonGaussian contribution to CMB anisotropies, even if the magnetic fields themselves are Gaussian distributed, since the components of the energy momentum tensor are quadratic in the fields and therefore approximately follow χ^{2} statistics (Brown & Crittenden 2005). In particular, PMFs generate nonzero higherorder statistical moments. The thirdorder moment, the CMB bispectrum, can be used as a probe to derive constraints on PMFs that are complementary to the previously mentioned ones. Planck polarization data thus provide a new way of probing PMFs, namely through the magneticallyinduced polarization bispectrum.
The presence of PMFs induces and sustains the propagation of Alfvén waves. These waves have an impact on the statistics of the CMB anisotropies and in particular induce specific correlations between harmonic modes (Kahniashvili et al. 2008). It is possible to use this effect to constrain the amplitude of Alfvén waves and thus indirectly constrain the PMF amplitude (Durrer et al. 1998; Kim & Naselsky 2009; Planck Collaboration XXIV 2014).
In the MHD limit, assuming that the fields are only modified by cosmic expansion, the magnetic field strength decreases as a^{2}, where a is the cosmological scale factor. Throughout, we will use a “comoving” magnetic field, defined as B = a^{2}B^{(phys)}, where B^{(phys)} is the physical strength of the magnetic field.
1.3. Structure of the paper
The paper is structured as follows. In Sect. 2 we describe the analysis of the impact of helical and nonhelical PMFs on CMB power spectra in temperature and polarization and we derive the constraints on the PMF amplitude and spectral index that come from Planck data. In Sect. 3 we present three different analyses of the magneticallyinduced bispectrum, specifically two analyses of the magneticallyinduced passive bispectrum and an analytical treatment of the magneticallyinduced scalar bispectrum. In all cases we derive the constraints on the amplitude of PMFs with a scaleinvariant spectrum using Planck nonGaussianity measurements. In Sect. 4 we present our analysis of the Faraday rotation signal induced by PMFs and we derive constraints from Planck low frequency polarization data. In Sect. 5 we present the analysis of the impact of Alfvén waves on statistical correlations in harmonic space and the associated constraints on Alfvén waves derived from Planck data. We summarize our conclusions in Sect. 6
2. Impact of primordial magnetic fields on the CMB power spectra
PMFs affect cosmological perturbations and may leave significant imprints on the CMB power spectra in temperature and polarization. Accurate prediction of these signatures allows us to derive constraints on PMF characteristics from CMB anisotropy data from Planck using the Planck likelihood. In this section we derive the predictions for the magneticallyinduced power spectra in temperature and polarization, considering helical and nonhelical PMFs, and present the resulting constraints on PMFs.
2.1. Magnetic modes
When considering a stochastic background of PMFs, we can neglect the contribution of energy density and anisotropic stress at the homogeneous level. The magnetic energy momentum tensor can be seen as describing perturbations carrying energy density and anisotropic stress and inducing a Lorentz force on the charged particles of the plasma. PMFs source all types of perturbations; scalar, vector, and tensor. In the past years, several different analyses of magneticallyinduced perturbations have been performed. Some examples from the wide literature of the field concern magneticallyinduced scalar perturbations (Giovannini 2004a; Kahniashvili & Ratra 2007; Yamazaki et al. 2007, 2008; Finelli et al. 2008; Giovannini & Kunze 2008c,a,b; Bonvin & Caprini 2010; Bonvin 2010; Kunze 2011), while other treatments also include magneticallyinduced vector and tensor perturbations (Subramanian & Barrow 1998b, 2002; Durrer et al. 2000; Kahniashvili et al. 2001; Mack et al. 2002; Caprini & Durrer 2002; Subramanian et al. 2003; Lewis 2004; Caprini 2006; Paoletti et al. 2009; Shaw & Lewis 2010). We can identify three different classes of initial conditions for magneticallyinduced perturbations; compensated (Giovannini 2004a; Finelli et al. 2008), passive (Lewis 2004; Shaw & Lewis 2010), and inflationary (Bonvin et al. 2013).
In this paper we focus on the two magneticallyinduced modes that are present for all types of PMFs produced prior to decoupling, independent of their generation mechanism, i.e., the compensated and passive modes. We do not consider specific inflationary initial conditions (Bonvin et al. 2013) to maintain the generality of the PMFs we constrain. For the same reason we neither consider a possible crosscorrelation between the magneticallyinduced and the adiabatic mode motivated by inflation (Jain & Sloth 2012).
2.1.1. Compensated modes
The compensated modes are the regular magneticallyinduced modes. These are the regular (finite at τ → 0) solutions of the perturbed EinsteinBoltzmann equations, including the magnetic contributions after neutrino decoupling. These modes are called “compensated” because the magnetic contributions to the metric perturbations in the initial conditions are compensated by fluid modes to leading order. The initial conditions are the solutions of the EinsteinBoltzmann equation system for large wavelengths at early times, with the perturbed quantities expanded in power series of kτ (where k is the perturbation wavenumber and τ is the conformal time). When performing this calculation for the magneticallyinduced modes, the growing regular mode requires the source terms in the equations for the metric perturbations to vanish at the lowest order. This can only be realized by a compensation between the magnetic terms and the perturbed quantities of the fluid.
2.1.2. Passive modes
The second class, the passive modes, is generated by the presence of a PMF before neutrino decoupling. Without neutrinos freestreaming, there is no counterpart in the fluid to balance the anisotropic stress of the PMF. This generates a logarithmically growing mode (in conformal time, which diverges for early times). After neutrino decoupling, the anisotropic neutrino stress compensates the anisotropic stress, which is due to the PMF, leading back to the compensated case described before. But an imprint of this logarithmically growing mode survives neutrino decoupling in form of a constant offset on the amplitude of the inflationary nonmagnetic mode (the primary cosmological pertubations of the standard model without PMFs). This amplitude offset is due to the continuity condition for the matching of the initial conditions before and after neutrino decoupling. Passive modes have a logarithmic dependence on the ratio between the neutrino decoupling time and the generation time of the PMF, i.e., their amplitudes grow as h(k) ∝ ln(τ_{ν}/τ_{B}) (where τ_{ν} is the neutrino decoupling time and τ_{B} is the PMF generation time). The passive modes, unlike the compensated ones, evolve following the standard nonmagnetic equations and only influence scalar and tensor perturbations.
2.2. Impact of nonhelical PMFs on the CMB angular power spectra
Our analysis is based on previous treatments of magneticallyinduced compensated and passive scalar, vector, and tensor modes presented by Lewis (2004), Finelli et al. (2008), Paoletti et al. (2009), and Shaw & Lewis (2010).
At linear order, PMFs evolve like a stiff source and we can therefore discard the backreaction of gravity onto the stochastic background of PMFs. Prior to the recombination epoch the electric conductivity of the primordial plasma is very large. We therefore consider the limit of infinite conductivity, in which the induced electric field is zero. In this limit, the temporal evolution of the PMF reduces to B^{(phys)}(x,τ) = B(x) /a(τ)^{2}, where B(x) is the comoving field^{2}.
We model a nonhelical stochastic PMF with a powerlaw power spectrum, where the twopoint correlation function is described by^{3}(1)where P_{B}(k) = A_{B}k^{nB} and denotes a cartesian component of a normalized wave vector. In this model, the PMF is characterized by two quantities, the amplitude of the power spectrum, A_{B}, and the spectral index n_{B}. The latter is one of the major discriminating factors between generation mechanisms, since different mechanisms generate fields with different spectral indices (for example, causal mechanisms generate fields with n_{B} ≥ 2; Durrer & Caprini 2003). While magneticallyinduced compensated perturbations do not suffer from Silk damping, PMFs are nevertheless suppressed on small scales by radiation viscosity (Jedamzik et al. 1998; Kahniashvili et al. 2001). To account for this damping we introduce a sharp cutoff in the PMF power spectrum at the damping scale k_{D}.
For the amplitude we use the convention to smooth over a comoving scale of λ = 1 Mpc, (2)For the damping scale we use (Subramanian & Barrow 1998a; Mack et al. 2002)^{4}(3)where h is the reduced Hubble constant, H_{0} = 100 h km s^{1} Mpc^{1}, and Ω_{b} is the baryon density parameter^{5}. Magneticallyinduced scalar, vector, and tensor perturbations are sourced by the energy momentum tensor components due to PMFs, together with the Lorentz force contribution. The energy momentum tensor of the PMFs is where the components are all quadratic in the magnetic field. The power spectra of the perturbations are therefore fourthorder in the magnetic field and given by convolutions of the magnetic power spectrum. The twopoint correlation function of the spatial part of the energy momentum tensor is^{6}We can then obtain scalar, vector, and tensor correlation functions,
where the Π^{(X)} are the scalar, vector, and tensor components of the energy momentum tensor, , and we sum over repeated indices. Such convolutions can be written in terms of spectra as where ℳ_{ijtl} = P_{it}P_{jl} + P_{il}P_{jt}−P_{ij}P_{tl}. With this convention, the relevant components of the energy momentum tensor become where , , , and Ω denotes the volume with p<k_{D}.
Fig. 1 Magneticallyinduced CMB TT (top left), TE (top right), EE (bottom left), and BB (bottom right) power spectra. The solid lines represent primary CMB anisotropies, the dotted lines represent magneticallyinduced compensated scalar modes (except for the BB panel, where it represents the lensing contributions and the solid line represents primary tensor modes with a tensortoscalar ratio of r = 0.1), the dashed lines represent vector modes, whereas the dotdashed lines represent magneticallyinduced compensated tensor modes. We consider PMFs with B_{1 Mpc} = 4.5 nG and n_{B} = −1. 

Open with DEXTER 
The conservation equations for the fields give a relation between the scalar projection of the anisotropic stress, the energy density, and the Lorentz force, ^{7} (which reduces to a simple relation between Lorentz force and anisotropic stress for vector modes, ). This relation simplifies the treatment, reducing the number of correlators to be computed by a factor of 2. We use the analytic solutions to the convolutions derived by Finelli et al. (2008) and Paoletti et al. (2009) for fixed spectral indices. For general n_{B}, we use the fits to the generic analytic solutions provided by Paoletti & Finelli (2011). These fits simplify the computation owing to the presence of hypergeometric functions in the analytic solutions. The infrared behaviour of the spectra, which is relevant for CMB anisotropies, depends on the spectral index. In particular, the spectra describe white noise for indices greater than , whereas they are infrareddominated, as k^{2nB + 3}, for smaller indices. We use the initial conditions derived by Lewis (2004), Paoletti et al. (2009), Paoletti & Finelli (2011), and Shaw & Lewis (2010) for scalar and compensated tensor modes, vector modes, and passive modes, respectively. We use an extended version of the CAMB code (Lewis & Challinor 2011) that includes all magnetic contributions to calculate predictions for the CMB power spectra in temperature and polarization.
2.2.1. Compensated modes
In Fig. 1, we show the predictions for magneticallyinduced compensated modes. This shows that the dominant compensated contributions to the angular power spectra are given by the scalar and vector modes. In particular, because magneticallyinduced perturbations are not suppressed by Silk damping, a significant contribution of magneticallyinduced modes arises on small angular scales, where the primary CMB fluctuations are suppressed. As we will show, the impact of PMFs on the CMB power spectrum at high multipoles is particularly relevant for the highprecision Planck data, allowing us to derive strong constraints on the PMF amplitude. Since magneticallyinduced perturbations are solely sourced by energy momentum tensor components that are due to PMFs, the shape of the magneticallyinduced spectra strongly depends on the PMF spectral index. In Fig. 2, we show this dependence for the temperature power spectrum for scalar and vector perturbations. We note the qualitatively different dependence for two regimes. For n_{B}> −3 / 2 the angular power spectrum remains flat (i.e., ℓ(ℓ + 1)C_{ℓ} ∝ ℓ^{2}) with a rescaling of the amplitude owing to the amplitude of the Fourier spectra, whereas for n_{B}< −3 / 2 the shape varies according to the infrared domination of the energy momentum tensor of the PMFs.
Fig. 2 Dependence of the magneticallyinduced CMB power spectrum on the spectral index. For all plotted cases, the amplitude is B_{1 Mpc} = 4.5 nG. The black lines show primary CMB anisotropies; for the other colours we refer to the legend. Left: scalar contributions, right: vector contributions. 

Open with DEXTER 
2.2.2. Passive modes
In addition to compensated initial conditions, we also consider passive tensor modes. The magneticallyinduced passive modes are not completely determined by the amplitude and spectral index of the PMF, but also depend on the ratio τ_{ν}/τ_{B}. For tensors we specifically have h(k) ∝ Π^{(T)}(k)ln(τ_{ν}/τ_{B}), where h(k) is the tensor metric perturbation. This ratio may vary between 10^{17} and 10^{6} for fields originating at the grand unification energy scale (GUT) and at later phase transitions (Shaw & Lewis 2010). We consider the passive tensor modes, since, for red spectra, they give the dominant contribution on large angular scales, where the compensated modes are subdominant. In Fig. 3, we compare the magneticallyinduced passive tensor modes for a nearly scaleinvariant spectrum, n_{B} = −2.9, for PMFs generated at the GUT scale (for which τ_{ν}/τ_{B} = 10^{17}) with the corresponding dominant compensated modes. In Fig. 4, we show the dependence of the CMB power spectrum due to passive tensor modes on the spectral index, for GUT scale PMFs, and its dependence on the time ratio, showing the two extreme values of the possible range. We note that the compensated vector modes dominate at small angular scales, but the passive tensor modes give a contribution at low and intermediate multipoles for red spectra. For bluer spectra, the passive spectrum becomes steeper and therefore subdominant with respect to primary CMB fluctuations on large angular scales and with respect to vector modes on small angular scales. A similar behaviour can be observed in the scalar passive mode, which has the same origin as the tensor one but in the scalar sector. Shaw & Lewis (2010) have shown that, just like the tensor mode, the scalar passive mode becomes relevant on large angular scales for nearly scaleinvariant power spectra. We will show below that the dominant passive contribution to the constraints on the PMF amplitude is given by tensor modes.
Fig. 3 Magneticallyinduced CMB TT (top left), TE (top right), EE (bottom left), and BB (bottom right) power spectra due to passive tensor modes, compared with the ones that are due to compensated modes. The solid lines represent primary CMB anisotropies, the dotted lines represent magneticallyinduced compensated scalar modes (except for the BB panel, where it represents the lensing contribution), the dashed lines represent vector modes, whereas dotdashed lines represent magneticallyinduced passive tensor modes. We consider PMFs with B_{1 Mpc} = 4.5 nG and n_{B} = −2.9. 

Open with DEXTER 
Fig. 4 Dependence of the magneticallyinduced CMB power spectrum due to passive tensor modes on the spectral index for a GUTscale PMF (left) and comparison between the two extremes for the time ratio τ_{ν}/τ_{B} (right). The black lines show the primary CMB anisotropies; for the other colours we refer to the legend. Solid lines represent PMFs generated at the GUT scale, τ_{ν}/τ_{B} = 10^{17}, whereas dashed lines represent PMFs generated at late times, τ_{ν}/τ_{B} = 10^{6}. 

Open with DEXTER 
2.3. Impact of helical PMFs on CMB anisotropies
In addition to the ubiquitous presence of magnetic fields in the Universe, astrophysical observations show that some galaxies might have a helical magnetic field structure (Widrow 2002; Vallée 2004). Following the hypothesis that the amplification of PMFs may have played a role in the generation of largescale magnetic fields, the observed magnetic helicity may be related to helicity of the PMFs. A helical intergalactic magnetic field could also be related to a possible CP violation recently hypothesized by Tashiro et al. (2014) in an indirect study of cosmological largescale magnetic fields in voids using secondary γray data. Helicity of the PMFs influences MHD processes in the early plasma, as well as cosmological perturbation dynamics, allowing different processes of energy transport, for example the inverse cascade mechanism (Biskamp 2003). These processes of energy transport play a role in the early time evolution of the PMFs and may have an impact on our understanding of their generation mechanisms, especially if PMFs with a nonzero helicity are generated. Moreover, the presence of helicity would test possible modifications of Maxwell’s theory by constraining parameters describing the gauge invariance (i.e., mass of the photon) and Lorentz invariance (i.e., existence of a preferred frame of reference) as discussed by Carroll et al. (1990). Thus it would carry information about particle physics at very high temperatures (above 1 TeV).
Fig. 5 CMB TT (top left), TE (top right), EE (bottom left), and BB (bottom right) power spectra due to helical PMFs compared to the ones that are due to nonhelical PMFs. Solid lines are nonhelical predictions, while dashed lines are helical predictions. Blue are the scalar modes, green the vector, and red the compensated tensor modes. We consider PMFs with B_{1 Mpc} = 4.5 nG and n_{B} = −1. 

Open with DEXTER 
A possible way to detect magnetic helicity directly from CMB data is to study the polarized CMB (cross) power spectra. A nonzero helicity in the PMFs changes the amplitudes of the parityeven power spectra and induces parityodd crosscorrelations between the temperature and Bpolarization anisotropies and E and Bpolarization anisotropies (Pogosian et al. 2002; Caprini et al. 2004; Kahniashvili & Ratra 2005; Ballardini et al. 2015). Such parityodd crosscorrelators are also generated by Faraday rotation, but only owing to homogeneous PMFs, not owing to the stochastic model for PMFs considered in this paper (Kosowsky et al. 2005). Parityodd signals may therefore give a more direct possibility for studying helical PMFs.
Our present study for a stochastic background of helical PMFs is an extension of the model described in the previous section. We consider the impact of such PMFs on CMB anisotropies in temperature and polarization. Using the exact expressions for the energymomentum tensor components including the helical contribution discussed by Ballardini et al. (2015), we derive the predictions for the impact of a helical PMF on the CMB power spectra in temperature and polarization.
The most general ansatz for the twopoint correlation function, built on Eq. (1), but taking into account an antisymmetric part (Pogosian et al. 2002), is (12)with P_{H}(k) = A_{H}k^{nH}. From a geometrical point of view P_{B}(k) denotes the symmetric part and P_{H}(k) the antisymmetric part of the correlator. The totally antisymmetric tensor ϵ_{ijl} is related to the parity violation under a transformation k → −k. The realizability condition gives P_{H}(k) <P_{B}(k).
In this case the model is described by four parameters, two amplitudes and two spectral indices, where A_{B} is the amplitude of the power spectrum of the fields, the same as defined in Eq. (2), and A_{H} is the amplitude of the power spectrum of the helical part of the PMFs. These amplitudes can be expressed in terms of meansquare values of the magnetic field and of the helical component, respectively. As done for the amplitude of the magnetic field in Eq. (2), we can express the amplitude of the helical component on a comoving scale of λ as (Ballardini et al. 2015) (13)The helical spectral index needs to satisfy n_{H}> −4 for convergence.
The helical term in Eq. (12)generates new sources that contribute to the energy momentum tensor. All components of the energy momentum tensor are quadratic in the fields, so the Fourierspace twopoint correlation function generates symmetric and antisymmetric sources: symmetric ones that are due to products of P_{B} with P_{B} (the components found in the nonhelical case) and products of P_{H} with P_{H}; antisymmetric ones, which generate oddparity angular power spectra, from products of P_{B} with P_{H}. The Fourier components of the energy momentum tensor are where the nonhelical parts are given by Eqs. (8)–(11). The system of equations for the background and the perturbations, as well as the initial conditions, are unmodified.
In Fig. 5 we show the predictions for the magneticallyinduced compensated modes, considering the additional contributions to the sources owing to helicity. We consider the maximally helical case, A_{H} = A_{B}, with equal spectral indices, n_{B} = n_{H}, using the solutions to the helical energy momentum tensor components derived by Ballardini et al. (2015). The predictions show no difference in the shape and in the slope of the angular power spectra, with a small shift in the amplitude, which is always smaller for the helical case, at least for scales relevant for the CMB. Magnetic helicity induces parityodd crosscorrelations between the E and Bpolarization anisotropies, as well as between temperature and Bpolarization anisotropies. The parityodd crosscorrelations are sourced by the mixed terms in the correlation function of the energy momentum tensor, proportional to . These terms, after decomposition, contribute to the vector and tensor sources as In the limit of small momenta, i.e., k ≪ k_{D}, the spectra never show a whitenoise behaviour, A^{(X)}(k) = const, contrary to what happens for the symmetric and nonhelical parts. For n_{B} + n_{H}> −2 the behaviour is proportional to and for n_{B} + n_{H}< −2 it does not depend on the damping scale and is proportional to k^{nB + nH + 2}.
2.4. Constraints from the CMB temperature and polarization power spectra
Here we present the constraints from Planck on helical and nonhelical PMFs. In the literature there are several previous studies that already derived constraints on PMFs using different combinations of observed CMB power spectra (Caprini 2010; Yamazaki et al. 2010; Paoletti & Finelli 2011; Shaw & Lewis 2012; Paoletti & Finelli 2013; Ade et al. 2015). We use an extended version of the CosmoMC code (Lewis & Bridle 2011), modified to include the magnetic contributions to the CMB power spectra, as described in the previous subsections, and to include the parameters characterizing the PMFs in the Markov chain Monte Carlo analysis.
We assume a flat Universe and a CMB temperature T_{0} = 2.7255 K, and we use the BBN consistency condition (Ichikawa & Takahashi 2006; Hamann et al. 2008). We restrict our analysis to three massless neutrinos. A nonvanishing neutrino mass would not modify the results since it would only enhance the power on large scales in the presence of PMFs for the compensated modes, where the PMF contribution is less relevant (Shaw & Lewis 2010). The pivot scale of the primordial scalar is set to k_{∗} = 0.05 Mpc^{1}. We consider the lensing effect for the primary CMB power spectrum and follow the method implemented in the Planck likelihood to marginalize over astrophysical residuals and secondary anisotropy contamination of the smallangularscale data (Planck Collaboration XI 2016). This contamination is particularly relevant for the PMF scenario, since PMFs impact mainly small angular scales. If this contamination is not properly considered it may lead to biased constraints on PMFs (Paoletti & Finelli 2013). We sample the posterior using the MetropolisHastings algorithm (Hastings 1970), generating between four and sixteen parallel chains and imposing a conservative GelmanRubin convergence criterion (Gelman & Rubin 1992) of R−1 < 0.01^{8}. We vary the baryon density ω_{b} = Ω_{b}h^{2}, the cold dark matter density ω_{c} = Ω_{c}h^{2} , the reionization optical depth τ_{reion}, the ratio of the sound horizon to the angular diameter distance at decoupling θ, the scalar amplitude ln(A_{s}10^{10}), and the scalar slope n_{s}. In the Markov Chain Monte Carlo (MCMC) analysis, we include the magnetic parameters B_{1 Mpc} and n_{B} for the compensated modes, and add the parameter τ_{rat} = τ_{ν}/τ_{B} whenever we also consider the passive tensor mode. We use flat priors for the magnetic parameters in the ranges [0,10] for B_{1 Mpc}/ nG, [− 2.9,3] for n_{B} (n_{B}> −3 to avoid infrared divergence in the PMF energy momentum tensor correlations). We sample τ_{rat} logarithmically, with a flat prior on log _{10}τ_{rat} in the range [4,17].
2.4.1. Likelihood
We derive the constraints on PMFs using the Planck likelihood, which is described in detail in Planck Collaboration XI (2016). Here we give a brief summary of the main points. The Planck likelihood is based on the Planck 2015 data and considers both temperature and polarization. As in 2013, we use a hybrid approach with the combination of two likelihoods, one dedicated to low ℓ and the other to high ℓ.
The Planck lowℓ likelihood is a fully pixelbased likelihood with temperature and polarization treated jointly and at the same resolution, N_{side} = 16. The ℓrange is 2 <ℓ< 29 in TT, TE, EE, and BB. The likelihood is based on the foregroundcleaned LFI maps at 70 GHz and the temperature map derived by the component separation method Commander using 94% of the sky at frequencies from 30 to 353 GHz (Planck Collaboration IX 2016). The polarization map covers 54% of the sky and is derived from the 70 GHz Q and U maps cleaned with the 30 GHz map as a synchrotron template and the 353 GHz map as a dust template (see Planck Collaboration XI 2016). This likelihood is denoted as “lowP” throughout the paper. Contrary to the 2013 analysis, where a combination of Planck temperature and WMAP9 polarization data was used, the 2015 lowℓ likelihood is based entirely on Planck data for both temperature and polarization.
The Planck highℓ likelihood is based on a Gaussian approximation (Planck Collaboration XV 2014 and Planck Collaboration XI 2016 for polarization) and covers the ℓrange 30 <ℓ< 2500. It uses the halfmission crosspower spectra of the 100 GHz, 143 GHz, and 217 GHz channels, measured in the cleanest region of the sky far from the Galactic plane and bright point sources. The sky fractions considered are 66% of the sky for 100 GHz, 57% for 143 GHz, and 47% for 217 GHz in temperature, whereas in polarization they are 70%, 50%, and 41%, respectively. The likelihood takes foregrounds and secondary anisotropies into account. In particular, for the temperature spectra it considers the contributions of dust, clustered Cosmic Infrared Background (CIB), thermal and kinetic Sunyaev Zeldovich effect (tSZ and kSZ), the crosscorrelation between tSZ and CIB, and a Poissonian term for unresolved point sources for the temperature spectra. In polarization, only the dust contribution is considered. Each model is parameterized as a template contribution to the C_{ℓ} with a free amplitude. The dominant contribution for ℓ< 500 is dust, whereas high ℓmodes are dominated by point sources and in particular the CIB for the 217 GHz autocorrelation. For the details of the foreground modelling see Planck Collaboration XI (2016). This highℓ likelihood is denoted as “Planck TT”, for temperature only, or “Planck TT, TE, EE”, for temperature plus polarization, throughout the paper.
2.4.2. Constraints with compensated scalar and vector contributions
We perform an analysis with the Planck 2015 baseline likelihood. In Table 1 we report the derived constraints. The constraint on the PMF amplitude is B_{1 Mpc}< 4.4 nG at the 95% confidence limit (CL) for the case that includes temperature and polarization data both at low and high multipoles. The same constraint results when including the polarization only at low ℓ. As in previous analyses, PMFs with positive spectral indices are constrained to lower amplitudes than PMFs with negative spectral indices. In Fig. 6 we present the results of this analysis compared with the Planck 2013 constraints (Planck Collaboration XVI 2014). The upper limits from current Planck data are slightly higher than those obtained in 2013 (Planck Collaboration XVI 2014), where the constraint from Planck data alone was B_{1 Mpc}< 4.1 nG. The weaker constraint can be explained by several changes of the 2015 data and likelihood with respect to 2013. The change in the calibration (see Planck Collaboration I 2016), the different likelihood implementation, and the different models for the foreground residuals are all factors that contribute to the changed upper limit. In fact, all these factors, including also the slightly higher spectral index n_{s} with respect to 2013, are mimicking a slightly larger signal in the temperature anisotropies, which is compatible with larger values of the PMF amplitude.
Mean parameter values and bounds of the central 68% CL from Planck TT, TE, EE (left column) and Planck TT (right column).
Fig. 6 Comparison of the constraints on the smoothed PMF amplitude (top) and the spectral index (bottom) from the 2015 temperature and temperature plus polarization data with the 2013 results for magneticallyinduced compensated initial conditions only. 

Open with DEXTER 
We now include the polarization data in the analysis. Although the impact of PMFs on TE and EE polarization is less important than on temperature anisotropies, we show the results for the case of Planck data, which also includes highℓTE and EE polarization, in Fig. 6. Although the shape of the posterior changes slightly, there is no net improvement on the 95% CL upper bound on B_{1 Mpc} with the addition of the highℓTE and EE polarization.
2.4.3. Constraints with passive tensor contributions
As described in the previous subsection, in addition to regular magneticallyinduced compensated modes, the presence of PMFs prior to neutrino decoupling generates passive modes. In Fig. 3 we show that the passive tensor modes may give the dominant magnetic contribution to the CMB power spectra for a nearly scaleinvariant PMF power spectrum. Their inclusion in the analysis may therefore be relevant for the constraints on PMFs. We include the passive tensor contribution in the MCMC code with the addition of the parameter τ_{rat}, with the settings described above. We perform MCMC analyses with the Planck 2015 likelihood, combining the lowℓ temperature and polarization data either with highℓ temperature data or with highℓ temperature and polarization data, i.e., TT+lowP and TT, TE, EE+lowP.
Figure 7 and Table 2 present the results, compared to the results of the case that includes only compensated contributions. The 95% CL constraint on the PMF amplitude is B_{1 Mpc}< 4.5 nG, which implies that the addition of the passive tensor contribution does not improve the constraint on the amplitude of the PMFs. This result is expected on the basis of the shape of the angular power spectra that are due to passive tensor modes and their strong dependence on the PMF spectral index. This mode is basically a primary tensor mode with an amplitude that depends on the PMFs. Its spectrum flattens on large angular scales and then decays on intermediate ones for red PMF spectra, whereas it acquires a steeper shape for blue PMF spectra, but with a much lower amplitude than for compensated vector modes. Therefore, passive tensor modes only contribute significantly for nearly scaleinvariant indices. In Fig. 8 we present the twodimensional plot for the PMF amplitude and the spectral index. It shows the strong degeneracy between the two parameters, meaning that the same magneticallyinduced power spectrum can be realized with different pairs of amplitude and spectral index. The effect of passive tensor modes on the CMB temperature angular power spectrum is dominant over the primary CMB anisotropies only for a very limited range of spectral indices, near the scaleinvariant case. Therefore, the degeneracy between amplitude and spectral index reduces the influence of the passive tensor mode contribution on the constraints on the amplitude and we do not see any improvement in adding the passive tensor mode. While the constraint on the amplitude is almost unchanged, the spectral index in the MCMC analysis is sensitive to the contribution of the passive mode. In the lower panel of Fig. 7 we show the different shapes of the posterior distributions for the PMF spectral index for different data combinations. The inclusion of the passive tensor mode influences the lowspectral index part of the posterior, while the compensated modes influence the highspectral index part. In addition to the analyses that consider the combination of passive and compensated modes, we perform an analysis including only the contribution of the passive tensor mode. We obtain B_{1 Mpc}< 6.5 nG at 95% CL using PlanckTT+lowP. This result shows that the contribution of the passive term alone, when considering the spectral index and the generation epoch as free parameters, does not have the constraining power of the combination of passive and compensated modes.
Upper bounds of the central 95% CL for the PMF amplitude.
Fig. 7 Constraints on the smoothed PMF amplitude (top) and spectral index (bottom) from Planck temperature data with and without the passive tensor contribution. Constraints including both compensated and passive modes are indicated with C+P in the legend, constraints using only compensated modes are marked with C. 

Open with DEXTER 
Fig. 8 PMF amplitude versus the spectral index for the baseline Planck 2015 case. C+P denotes the case where both compensated and passive modes are considered, whereas C indicates the case with only compensated modes. The two contours represent the 68% and 95% confidence levels. 

Open with DEXTER 
Fig. 9 Twodimensional posterior distributions of the PMF amplitude versus the parameter that describes the Poissonian term of unresolved point sources for the three frequencies considered in the likelihood. The two contours represent the 68% and 95% confidence levels. 

Open with DEXTER 
2.4.4. Impact of astrophysical residuals
We adopt the Planck likelihood treatment of astrophysical contaminants as described in the Planck likelihood paper. Considering the complexity of this model and the number of nuisance parameters involved, we investigate whether foreground residuals have an impact on the constraints on PMFs. Specifically, we investigate possible degeneracies with the foreground parameters by considering the twodimensional distributions of the magnetic parameters and foreground parameters. The relevant cases are shown in Fig. 9. There, we present the twodimensional distributions of the PMF amplitude and the Poissonian amplitudes for the three frequencies considered in the Planck highℓ likelihood: 100 GHz, 143 GHz, 217 GHz, and the 143 × 217 GHz crossspectrum. We note that especially for the 143 GHz and the 143 × 217 GHz analyses, a weak degeneracy between the two parameters is seen. This result may indicate an impact of the astrophysical residual modelling on the PMF constraints. To investigate this issue, we perform an analysis by fixing the four parameters for the Poissonian amplitudes to their bestfit values from the standard Lambda Cold Dark Matter model. This analysis yields a limit of B_{1 Mpc}< 3.0 nG, which is smaller than the constraint obtained in the case where the parameters associated to astrophysical residuals are free to vary. This result has no statistical significance, but demonstrates that there is an impact of the astrophysical residuals on the PMF constraints when the data considered require a complex model for the residuals. The shape of the dominant PMF contributions to the angular power spectrum on small angular scales is responsible for this degeneracy. In fact, the steep slope of the vector mode may be degenerate with astrophysical residual contributions, as shown by Paoletti & Finelli (2013). For comparison we show an analogous plot for the 2013 analysis in Fig. A.1, which considers the degeneracy for the foreground parameters of the Planck 2013 likelihood. We note how, in contrast to the 2015 analysis, there is only a small degeneracy with the Poissonian amplitude at 143 GHz. This result shows the importance of the foreground residual modelling for the PMF constraints.
In contrast to the Poissonian terms, we do not observe any degeneracy with the other foreground components, as shown in Fig. A.2, including the clustering component of the foreground residuals, which is the other dominant contribution on small angular scales at the frequencies considered in this analysis. The fact that we do not observe a degeneracy in this case is due to the difference in the spectral shape of the PMF contribution and the clustering term. Although both are relevant on small angular scales the slightly different shapes break the degeneracy.
2.4.5. Constraints for specific PMF models
Planck 2015 results confirm what has been observed in previous analyses, namely that CMB data allow negative PMF spectral indices with larger field amplitudes than positive indices. The spectral index of the PMFs is the main discriminating factor among possible generation mechanisms.
Some cases are of particular interest owing to their connection with specific classes of generation mechanisms. In particular, PMFs generated during phase transitions or via secondorder perturbative effects, vector perturbations, etc., are characterized by positive spectral indices, equal or greater than 2. To investigate the maximal amplitude allowed by Planck data for fields of this type, we perform two dedicated analyses, the first with fixed index n_{B} = 2, and the second only restricted to positive spectral indices for PMFs. We include both compensated and passive modes in the analysis, giving B_{1 Mpc}< 0.011 nG at 95% CL (B_{1 Mpc}< 0.012 nG at 95% CL, when considering PlanckTT, TE, EE+lowP) for the n_{B} = 2 case and B_{1 Mpc}< 0.55 nG at 95% CL for n_{B}> 0.
We consider a third case of interest, the almost scaleinvariant fields with n_{B} = −2.9. This specific case is connected to PMF generation from inflation and is studied to test the strength of the passive tensor modes in constraining the amplitude of the PMFs. Moreover, we want to compare the results obtained from the Planck power spectra with those coming from the nonGaussianity analysis, presented in the next section, which is performed for this spectral index as well. We obtain B_{1 Mpc}< 2.0 nG at the 95% CL. Note that this nearly scaleinvariant case is dominated by the tensor passive mode. In fact, when we only consider the tensor passive contributions, excluding the compensated ones, we obtain the same result as in the passivecompensated combined case, B_{1 Mpc}< 2.0 nG at the 95% CL.
Together with the passive tensor mode there is also a scalar passive mode, as shown by Shaw & Lewis (2010). When we include this scalar passive contribution in our analysis we obtain again B_{1 Mpc}< 2.0 nG at the 95% CL. We can therefore conclude that the passive scalar contribution is subdominant with respect to the tensor one. As will become clear in the next section, this result shows that the constraining power of the angular power spectrum is comparable to the one of the nonGaussianity.
In addition to these specific types of PMFs we have performed some analyses with fixed spectral index, choosing a grid of values covering the full range we sample. With respect to the case where the spectral index is a free variable, the cases with fixed spectral index are expected to give stronger constraints on the PMF amplitude, with a trend in agreement with the general case, because one of the two parameters describing the PMF is fixed. In Table 3 we present the results of these analyses. We note how, as expected, the trend of the results with fixed spectral index is in agreement with the one of the twodimensional plot of Fig. 8 obtained with the generic sampling of the index. We also note how the constraint is weakest for n_{B} = −1.5 as expected from the impact on the angular power spectrum.
95% CL upper bounds of the PMF amplitude for fixed spectral index with compensated plus passive tensor modes.
2.4.6. Constraints from the BICEP2/KeckPlanck joint analysis
We perform an analysis using the recent BICEP2/KeckPlanck crosscorrelation (indicated as BKP; BICEP2/Keck Array and Planck Collaborations 2015) in addition to the Planck 2015 data. The BKP likelihood is obtained from the BB and EE bandpowers for all crossspectra between the BICEP2/Keck maps and the Planck maps at all frequencies (BICEP2/Keck Array and Planck Collaborations 2015). We consider both compensated and passive contributions and study two cases, one in which we leave the spectral index free to vary and another in which we fix it to n_{B} = −2.9. In the latter case, there is a contribution to the Bmode polarization on large angular scales from the passive tensor mode. Figure 10 shows the comparison of the results of these two analyses with the results obtained from Planck data alone. The constraints are B_{1 Mpc}< 4.7 nG for the case with free spectral index and B_{1 Mpc}< 2.2 nG at the 95% CL for the case with n_{B} = −2.9. These are slightly higher upper bounds for the amplitude of PMFs, but they are fully compatible with the results derived from Planck data alone. We note that the posterior distribution for the nearly scaleinvariant case changes with the addition of the BKP data but it does not show any significant deviation from the posterior based only on Planck data.
Fig. 10 Probability distributions for the PMF amplitude including the BICEP2/KeckPlanck crosscorrelation, compared with the one based only on Planck data. Top: the case in which the spectral index is free to vary, bottom: the case with n_{B} = −2.9. 

Open with DEXTER 
2.4.7. Constraints with maximally helical contributions
Fig. 11 PMF amplitude constraint for the helical case (solid black) compared with the nonhelical case (dashed red). The dotted blue line shows the constraint on the amplitude of the helical component as an alternative interpretation of the constraints on the amplitude of PMFs with a helical component. 

Open with DEXTER 
We perform an MCMC analysis including the maximally helical contribution. We restrict our analysis to the case of temperature and polarization with only even crosscorrelations. The odd crosscorrelations TB and EB are present only in the lowP likelihood, therefore only for very low multipoles where the signal from helical PMFs is negligible. Thus we do not include odd crosscorrelators in our analysis.
We perform an analysis using the PlanckTT+lowP likelihood. The constraint on the PMF amplitude in the maximally helical case is B_{1 Mpc}< 5.6 nG at the 95% CL. A comparison with the corresponding results for the nonhelical case is shown in Fig. 11. The analysis with the Planck TT, TE, EE+lowP likelihood gives B_{1 Mpc}< 5.8 nG at the 95% CL. As in the nonhelical case, the inclusion of highℓ polarization does not improve the constraints. Figure 5 shows that magnetic fields with a maximally helical component produce smaller CMB fluctuations in temperature and polarization than nonhelical fields of the same strength. As a result of this, the amplitude of maximally helical magnetic fields is less constrained than the one of nonhelical fields for this Planck 2015 data release. When considering helical PMFs, we have two components that contribute to the magneticallyinduced perturbations, as shown in Eq. (12), a symmetric and an antisymmetric part, represented by P_{B} and P_{H}, respectively. These power spectra can be associated with two amplitudes of the field, B_{1 Mpc} associated with the symmetric part and ℬ_{1 Mpc} associated with the antisymmetric part (see Eq. (13)). In the maximally helical case the two amplitudes are not independent from each other, they are related through the conditions A_{H} = A_{B} and n_{B} = n_{H}. Therefore we constrain a single amplitude, which can be expressed either through B_{1 Mpc} or ℬ_{1 Mpc}. The constraint B_{1 Mpc}< 5.6 nG can thus be converted into the constraint ℬ_{1 Mpc}< 4.6 nG at the 95% CL. Figure 11 shows the posterior distribution for the amplitude expressed as ℬ_{1 Mpc} in blue.
2.4.8. Constraints from the impact of PMFs on the CMB anisotropies via their impact on the thermal history of the Universe
Primordial magnetic fields are damped on scales smaller than the photon diffusion and freestreaming scale. This leads to heating of ordinary matter (electrons and baryons), which affects both the thermal and ionization history of the Universe (Subramanian & Barrow 1998a; Jedamzik et al. 2000; Sethi & Subramanian 2005; Schleicher et al. 2008; Kunze & Komatsu 2014; Chluba et al. 2015), leading to a Comptony distortion of the CMB and changes in the CMB power spectra through modifications of the Thomson visibility function around decoupling.
Two heating mechanisms have been discussed in the literature, one due to decaying magnetic turbulence at very small scales and the other due to ambipolar diffusion (e.g., Sethi & Subramanian 2005). In this paper, we follow the approach described by Chluba et al. (2015) to incorporate these heating mechanisms^{9}.
We perform an analysis considering the combination of the heating terms with the gravitational contribution of PMFs. Considering ambipolar diffusion, decaying magnetic turbulence, and gravitational effects we obtain an upper limit of B_{1 Mpc}< 0.90 nG at 95% CL for nearly scaleinvariant PMFs with n_{B} = −2.9. We obtain the same result, namely B_{1 Mpc}< 0.90 nG at 95% CL, when dropping the gravitational effect and considering only the impact of PMFs on the primary CMB anisotropies through their heating effect. These results show that the dominant contribution is given by the heating terms. We have also performed analyses with the two terms of ambipolar diffusion and decaying magnetic turbulence considered separately. The results show that the two terms are roughly at the same level in constraining PMFs, with a slightly stronger contribution from the decaying magnetic turbulence term (see also Chluba et al. 2015).
Together with the Planck TT+ lowP likelihood combination, we have performed an analysis including highℓ polarization. In particular, we have considered the case of Planck TT, TE, EE+ lowP. The result is: B_{1 Mpc}< 0.86 nG at the 95% CL. Owing to the nature of the effect of PMFs on the thermal history of the Universe and its impact on the CMB angular power spectra, the polarization data on small angular scales tighten the constraints of this analysis.
3. Magneticallyinduced nonGaussianities
The CMB anisotropies induced by PMFs are nonGaussian. This is because magnetic forcing (as described by the magnetic energy momentum tensor) is quadratic in the magnetic fields and therefore the resulting fluctuations are nonGaussian even for Gaussian fields^{10} (Brown & Crittenden 2005). There are already published theoretical studies of the passivemode bispectra (Trivedi et al. 2010; Shiraishi et al. 2011, 2012; Shiraishi 2013), as well as studies of the compensatedmode bispectra (Seshadri & Subramanian 2009; Caprini et al. 2009; Cai et al. 2010; Shiraishi et al. 2010; Kahniashvili & Lavrelashvili 2010) and of trispectra (Trivedi et al. 2012, 2014). This illustrates that it is possible to use CMB nonGaussianities to constrain the PMF amplitude for different generation mechanisms. Several nonGaussianity constraints have previously been used for this purpose (Caprini et al. 2009; Seshadri & Subramanian 2009; Trivedi et al. 2010, 2012; Shiraishi et al. 2012). The nonGaussianity constraints on PMFs are complementary to those derived from the angular power spectra. In this section we present three different methods for constraining PMFs using nonGaussianity measurements, all involving the first of the higherorder stastical moments, the bispectrum. The methods can be applied to either the passive or the compensated modes.
3.1. Magneticallyinduced passivetensor bispectrum
The goal of this subsection is to derive an observational limit on the PMF strength from the passive bispectrum. The dominant contribution to the passive bispectrum is the largescale tensor mode, while the scalar mode contributes subdominantly to the small scales. According to Shiraishi et al. (2012) and Shiraishi (2013), the signaltonoise ratio (integrated over ℓ) is expected to be almost saturated beyond ℓ = 500 in estimates based on the temperature bispectrum. Including higher multipoles would therefore not bring significant improvements. Thus, we take into account the tensormode contribution for ℓ ≤ 500 in the following. Here we concentrate on the almost scaleinvariant case, n_{B} = −2.9. In this case, the passivetensor bispectrum is amplified in the squeezedlimit configuration with ℓ_{1} ≪ ℓ_{2} ≈ ℓ_{3}, as a consequence of the localtype structure of the nonGaussian gravitational waves induced by the PMF, given by (20)Here τ_{ν} and τ_{B} are defined in the same way as in the previous section, while ρ_{γ,0} is the present photon energy density. The projection tensor is given by the products of the spin± transversetraceless tensors as , normalized as ℳ_{ijij} = 4. This projection induces a tangled angular dependence on k_{i} in the primordial gravitational wave bispectrum and the resultant CMB bispectrum is given by a nonfactorizable combination of ℓmodes (Shiraishi et al. 2011, 2012; Shiraishi 2012). The resultant CMB temperature and Emode bispectra are almost uncorrelated with the usual scalarmode bispectra because of their different CMB transfer functions. To derive constraints, we introduce an amplitude parameter proportional to the amplitude of the magneticallyinduced bispectrum, (21)where B_{1 Mpc} and τ_{B} are treated as free parameters. The normalization factors in the last equation are chosen to be comparable to current upper bounds on B_{1 Mpc} for a PMF created at the GUT epoch, i.e., τ_{ν}/τ_{B} = 10^{17}. It can be seen that the magneticallyinduced bispectrum, which is proportional to (B_{1 Mpc})^{6}, has a logarithmic dependence on τ_{B}. An analysis of the constraints from WMAP data is presented by Shiraishi & Sekiguchi (2014), yielding (68% CL).
In order to constrain the nonfactorizable magneticallyinduced bispectrum, we use an optimal estimator derived within the socalled separable modal methodology (see Fergusson et al. 2010, 2012; Shiraishi et al. 2014, 2015 for autobispectra; and Fergusson 2014, and Liguori et al., in prep. for crossbispectra), where the theoretical bispectrum templates are decomposed in finite subsets of the separable eigenbasis. Thus, the bispectrum estimator remains factorizable like in the usual KSW approach (Komatsu et al. 2005). Our tensor bispectrum template can be reconstructed well in this modal decomposition with about 400 eigenvectors composed of polynomials and a few special functions modelling the CMB temperature transfer functions.
From the foregroundcleaned SMICA temperature map, we obtain observational constraints on the amplitude of the passivetensor bispectrum. The observational data and the (Gaussian) simulation maps used in the computation of the linear term and the error bars are inpainted in the same manner as for the Planck tensor nonGaussianity analysis (Planck Collaboration XXIV 2014; Planck Collaboration XVII 2016), after including experimental aspects (beam, mask, and anisotropic noise). Our final result is (T only) at 68% CL, which gives no evidence for a signal at the 2σ level. This Planck temperature constraint is in good agreement with the WMAP one (Shiraishi & Sekiguchi 2014). Analogous results have been derived for the combination of T and E −modes and the Eonly case, but since these results are still preliminary, we use the Tonly mode for the present analysis.
Assuming that the bispectrum is generated by PMFs, its amplitude is given by Eq. (21). The amplitude of the bispectrum depends on the amplitude of the fields to the sixth power and on the logarithm of the ratio τ_{ν}/τ_{B}, which is greater than unity. Therefore we have . The result obtained in this analysis therefore leads to an upper bound on the strength of GUT generated PMFs (ln(τ_{ν}/τ_{B}) = 10^{17}) of B_{1 Mpc}< 2.8 nG (95% CL).
3.2. Magneticallyinduced anisotropic passive scalar bispectrum
Anisotropic stress from magnetic fields leads to curvature perturbations on superhorizon scales according to (Shaw & Lewis 2010) Shiraishi (2012) showed that the threepoint correlation of the curvature perturbation ζ_{k} sourced by magnetic fields is (22)where and k_{∗} denotes the pivot wavenumber. We can investigate this primordial nonGaussianity by estimating the expansion coefficient c_{L} (Shiraishi et al. 2013), where c_{0} is related to the localform f_{NL} as and P_{L} is the Lth order Legendre polynomial.
If the magnetic field is generated at the GUT scale with a nearly scaleinvariant spectrum, the Legendre coefficients are related to the field amplitude via where B_{1 Mpc} and ℬ_{1 Mpc} are the amplitudes of the nonhelical and helical magnetic field components (smoothed on a scale of 1 Mpc), respectively. Estimating c_{2} allows us to constrain B_{1 Mpc}, but estimating c_{1} does not lead to a useful constraint owing to its dependence on the helical component of the PMF, which is not considered in this analysis (cf. Eq. (24)). By the central limit theorem, the estimated value of c_{2} follows a Gaussian distribution. Therefore, the loglikelihood is given by (26)where with ĉ_{2}(d) being the estimated value from the data d, and σ^{2} corresponding to the variance of its estimation, which includes cosmic variance and noise variance. In Eq. (26), we have dropped an irrelevant constant term. We estimate c_{2} from SMICA, NILC, SEVEM, and Commander foregroundcleaned maps. The variance is estimated from the realistic Planck simulations for each foregroundcleaning method (Planck Collaboration IX 2016; Planck Collaboration X 2016). To determine B_{1 Mpc} and its confidence region, we use the CosmoMC package (Lewis & Bridle 2002) as a generic sampler and obtain the posterior probability of B_{1 Mpc}, given the likelihood. This analysis yields upper bounds on the amplitude of the nonhelical magnetic field component, B_{1 Mpc}, which are presented in Table 4. Using the Planck data constraint on the localform f_{NL} (Planck Collaboration XVII 2016) and , we impose an additional constraint of B_{1 Mpc}< 5.5 nG at 95% CL, which is weaker than the c_{2} constraints.
Planck constraints on the amplitude of the nonhelical magnetic field component, B_{1 Mpc} [nG], from the SMICA, NILC, SEVEM, and Commander foregroundcleaned maps at 95% CL.
The constraint from the passivescalar bispectrum with SMICA maps is B_{1 Mpc}< 4.5 nG (95% CL). Thus, the addition of the polarized bispectrum leads to an improved constraint with respect to the previous limit, B_{1 Mpc}< 5.2 nG, from the Planck 2013 analysis (Shiraishi et al. 2013; Planck Collaboration XXIV 2014).
3.3. Magneticallyinduced compensatedscalar bispectrum
Now we derive the magneticallyinduced scalar bispectrum on large and intermediate angular scales using a semianalytical method. We compute an effective f_{NL} based on the comparison between the bispectrum and the power spectrum and derive the constraints on the amplitude of the PMF using Planck measurements.
We derive the magneticallyinduced scalar bispectrum on large angular scales for compensated initial conditions, basing our analysis on the treatment presented by Caprini et al. (2009). For simplicity we redefine the parameter describing the amplitude of the PMF in this section. Instead of using the smoothed amplitude we directly use the root mean square value of the field. This quantity is finite thanks to the sharp cutoff inserted in the PMF power spectrum to model the smallscale suppression of the field. The mean square of the field is then defined as (27)The magneticallyinduced bispectrum on large angular scales depends on the temperature anisotropy on large angular scales and therefore on the SachsWolfe signal induced by the PMFs (Caprini et al. 2009; Bonvin & Caprini 2010). We use the expression derived by Paoletti et al. (2009) and Caprini et al. (2009), (28)where is the temperature anisotropy, Ω_{B}(k) = ⟨ B^{2}(k) ⟩ /ρ_{rel}, and τ_{0} and τ_{dec} are the conformal time at present and at decoupling, respectively. For simplicity we have used an approximated expression for the initial conditions instead of the exact one, which also involves the Lorentz force Caprini et al. (in prep.). We therefore introduce a correction factor, α = 0.5, which numerically includes all the contributions that we do not consider in the expression. The above relation holds for the compensated mode initial conditions. The CMB bispectrum on large scales can then be written as (29)Substituting Eq. (28), we note that the bispectrum depends on the 3point correlation function of magnetic energy density, The characteristic feature of the magneticallyinduced bispectrum, generated by compensated modes, is that, contrary to what often happens for inflationary nonGaussianities, it is not possible to identify an a priori dominant geometric configuration. It is therefore necessary to analyse the bispectrum independently of the geometric configuration. Caprini et al. (2009) derive an approximate expression for the threepoint correlation function of magnetic energy density, which is independent of the geometric configuration and is tested against the analytic results for the flattened case. Using this expression, we derive the magneticallyinduced bispectrum and specify a geometric configuration only after integrating the magnetic energy density bispectrum in kspace. We show the results for the case corresponding to a local f_{NL}. The magneticallyinduced bispectrum for the nearly scaleinvariant case, n_{B} = −2.9, is given by (30)For n_{B}> −1, we find (31)For the specific case n_{B} = −2 it is (32)where we assume as the amplitude of the primordial gravitational potential power spectrum (). In all numerical estimates we have taken ℓ_{D} = k_{D}τ_{0} = 3000, ℓ_{max} = 750, and ℓ_{min} = 10. The Planck limit on local f_{NL} is given in Planck Collaboration XVII (2016). We use the result of the SMICA KSW T+E ISWlensingsubtracted analysis, namely f_{NL}< 5.8 at 68% CL. This limit on f_{NL} translates into constraints on the PMF amplitude of nG for n_{B} = −2.9, nG for n_{B} = −2, and nG for n_{B} = 2. The constraints on the smoothed amplitude of the field are B_{1 Mpc}< 3 nG for n_{B} = −2.9, B_{1 Mpc}< 0.07 nG for n_{B} = −2, and B_{1 Mpc}< 0.04 nG for n_{B} = 2. These results show how the constraints are competitive with, in addition to being complementary to, the ones given by the CMB angular power spectrum.
4. Faraday rotation
4.1. Constraints on PMFs from the Faraday rotation power spectrum
The presence of a PMF at the last scattering surface induces a rotation of the polarization plane of the CMB photons (Kosowsky & Loeb 1996). This effect is known as Faraday rotation (hereafter FR). The Faraday depth Φ is proportional to the integral along the line of sight of the magnetic field component along this direction, B_{∥}, and the thermal electron density, n_{e}, i.e., (33)The constant is K = 0.81 rad m^{2} pc^{1} cm^{3} μG^{1}= 2.6 × 10^{26} rad nG^{1}. The unit of the Faraday depth Φ is rad m^{2}. In this equation, is the unit vector in the lineofsight direction.
Here we analyse the impact of FR on the polarized CMB power spectra. We assume that the magnetic field is generated at some predecoupling epoch. We do not consider the generation mechanism itself, but only investigate the observable effects caused at recombination. Several works have derived the modification of the Boltzmann equation for the Stokes parameters in the presence of a homogeneous PMF (see, e.g., Scóccola et al. 2004) and for a stochastic distribution (see, e.g., Kosowsky et al. 2005). Here we explore the case of a stochastic distribution.
As discussed above, a PMF may induce scalar, vector, and tensor perturbations. At recombination, FR mixes the signatures of different perturbations. Some previous attempts to constrain FR signatures in the WMAP power spectrum are presented by Kahniashvili et al. (2009) and Pogosian et al. (2011). Both assume that Emodes are converted into Bmodes via FR. They obtain magnetic field strength limits of B_{1 Mpc} ≲ 100 nG and suggest an almost scaleinvariant spectrum, i.e., n_{B} ≈ −2.9 for a powerlaw distribution.
As in previous works, we assume a PMF distribution described by a power law as in Eq. (2). Note that any helical part of the field does not contribute to the Faraday rotation (Campanelli et al. 2004). The generation of magneticallyinduced Bmodes through FR of Emodes is described by (Kosowsky et al. 2005) (34)and the rotation of primordial Bmodes into magneticallyinduced Emodes by (35)where N_{ℓ} = (2(ℓ−2) ! / (ℓ + 2) ! )^{1 / 2} is a normalization factor, with L ≡ ℓ(ℓ + 1), L_{1} ≡ ℓ_{1}(ℓ_{1} + 1), L_{2} ≡ ℓ_{2}(ℓ_{2} + 1), and is a ClebschGordan coefficient. The power spectrum of the rotation angle is related to the one of the Faraday depth through (36)where ν_{0} is the observed frequency, and (37)Here, x_{D} = τk_{D}, where τ is the conformal time and k_{D} is given by Eq. (3).
In Eqs. (34)and (35), and are the primordial power spectra, whereas and are the ones including the effect of Faraday rotation, i.e., the observed ones. We use the observed Emode spectrum at 70 GHz as a proxy for the primordial one, , in Eq. (34)to calculate predicted Bmode spectra to be compared with the observed one.
The 70 GHz observations give B_{1 Mpc}< (1040,1380) nG (68%, 95% CL). The reduced χ^{2} is 1.35. The magnetic field spectral index for the stochastic distribution remains unconstrained. In Fig. 12, we show the probability contours derived from the 70 GHz data for the magnetic field strength B_{1 Mpc} and the spectral index of the PMF power spectrum, n_{B}. The upper bounds obtained from the FR analysis are very high compared to the other methods.
Fig. 12 Probability contours of PMF strength vs. spectral index of the PMF power spectrum as constrained by the 70 GHz observations. 

Open with DEXTER 
They are also slightly weaker than those previously obtained from FR analyses with WMAP5 (Kahniashvili et al. 2009), WMAP7 (Pogosian et al. 2011), and WMAP9 data.
4.2. Robustness of the results in the presence of foregrounds
The magnetic fields of the Milky Way contribute to the net Faraday rotation. Although the precise geometry of these magnetic fields remains uncertain (see, e.g., RuizGranados et al. 2010; Jansson & Farrar 2012), the Galactic Faraday depth could be a foreground for the primordial Faraday depth, at least on large scales.
To quantify the impact of the Galactic FR on the detection of primordial magnetic fields, we use Galactic observations of polarized synchrotron emission at 1.4 GHz and 23 GHz and the synthesized allsky Faraday rotation map derived from extragalactic radio source emission provided by Oppermann et al. (2015). In addition, we use simulations of the Galactic Faraday rotation obtained by using an axisymmetric Galactic magnetic field model for the halo field described by RuizGranados et al. (2010). Maps of Stokes Q and U are provided by Wolleben et al. (2006) at 1.4 GHz and by Bennett et al. (2013) at 23 GHz. Both frequencies are dominated by polarized synchrotron emission from within the Milky Way and are used to obtain the Galactic Faraday depth. For computing the power spectrum of the FR coming from simulations and observations at 1.4 and 23 GHz, we use the polarization processing mask provided by WMAP9^{11}.
In Fig. 13, we show the power spectra of Galactic Φ derived from polarized measurements at 1.4 GHz and 23 GHz, simulations, and for the allsky Faraday rotation map provided by Oppermann et al. (2015)^{12}. The fluctuations in the Galactic Faraday sky are not isotropic. Therefore, their statistics are not completely described by a power spectrum. Oppermann et al. (2015) model the Galactic Faraday depth as the product of an isotropic Gaussian random field and a latitudedependent function. In Fig. 13, we show two angular power spectra derived from the results of Oppermann et al. (2015). For the first one, we generate Gaussian realizations from their angular power spectrum, multiply them with their latitude profile, and pass them through the anafast routine of HEALPix. Averaging the result over 1000 realizations gives the blue dashed line in Fig. 13. For comparison, we also show the angular power spectrum of Oppermann et al. (2015) multiplied with the square of their profile function at a latitude of  b  = 45°, which gives the strength of the foreground Faraday rotation at a typical latitude used for CMB analysis. We also plot the power spectrum of the primordial Faraday rotation for PMFs of strength 1, 10, and 100 nG, respectively, and a spectral index of n_{B} = −2.9.
Fig. 13 Power spectra of the primordial Faraday depth Φ for magnetic field strengths of 100 (solid line), 10 (dotted line), and 1 nG (dashed line), the Galactic Faraday depth from the allsky Faraday map of Oppermann et al. (2015) (blue dashed line), the Galactic Faraday depth from Oppermann et al. (2015) at 45° latitude (green dashed line), the Galactic Faraday depth derived from Galactic emission at 1.4 and 23 GHz (red circles), and the Galactic Faraday depth for a Galactic magnetic field model (black crosses). 

Open with DEXTER 
Our main conclusion, as indicated by Fig. 13, is that observations of the Galactic Faraday depth show that primordial magnetic field strengths lower than 10 nG would require a detailed knowledge of the Galactic FR, at least at multipoles lower than ℓ ≈ 50, where the foreground rotation would be dominant. The exact detection limit for PMFs achievable through a study of Faraday rotation depends on Galactic latitude and on the extent to which the correction for foreground rotation is possible.
The upper limit found from the FR analysis is B_{1 Mpc}< 1380 nG and no restriction can be obtained for the spectral index n_{B}. In any case, our FR constraint on PMFs is well above the expected contamination level by Galactic FR, and therefore this contamination is currently not an issue. This is consistent with the prediction by De et al. (2013) that the Galactic Faraday depth would not be measurable with Planck.
5. Constraints on PMFs from Alfvén waves
Here we investigate the signature statistical anisotropy induced by Alfvén waves, which delivers yet another constraint on PMFs. In Planck Collaboration XXIV (2014) we constrained the Alfvén waves in the early Universe, where some arbitrary origin (including stochastic PMFs) was assumed for primordial vector perturbations. Given no evidence for Alfvén waves from that analysis, we now consider stochastic PMFs as the source of primordial vector perturbations and constrain an average background magnetic field and the energy density of stochastic PMFs^{13}. PMFs may produce Alfvén waves in the early Universe, which leave observable imprints on the CMB via the Doppler and integrated SachsWolfe effects. Durrer et al. (1998) show that Alfvén waves in the early Universe generate a fractional CMB anisotropy (38)where k denotes a Fourier mode vector, a sky direction, the unit vector in the direction of the homogeneous background magnetic field , and T_{0} is taken again as 2.7255 K (Fixsen 2009). Here Ω(k,τ_{last}) and Ω_{0}(k) denote the gauge invariant linear combination of vector perturbations at last scattering and at an initial time, respectively. In this analysis, we assume a nonhelical stochastic PMF, B, to be the sole source of initial vector fluctuations, (Durrer et al. 1998; Kahniashvili et al. 2008). The Alfvén wave velocity, ν_{A}, is given by (Durrer et al. 1998) (39)where ρ_{r} and p_{r} are the comoving density and pressure of the photons.
Kahniashvili et al. (2008) show that Alfvén waves in the early Universe produce correlations between harmonic modes separated by Δℓ = 0, ± 2, and Δm = 0, ± 1, ± 2. We give the explicit form of the correlations in Appendix B. Investigating these imprints, we impose a constraint on the Alfvén waves in the early Universe. In the weak Alfvén wave limit, the CMB data loglikelihood ℒ can be expanded as where . The first term on the righthand side is simply equal to the likelihood of the standard cosmological model and the first and second derivatives of the likelihood are obtained by (40)where ⟨ ... ⟩ denotes the ensemble average of signal and noise. The Hessian ℋ is given by (41)where a is the vector consisting of the spherical harmonic coefficients, a_{ℓm}, of the CMB anisotropy data, and C is their covariance matrix.
In our analysis, we consider the foregroundcleaned SMICA map, where we apply the common mask (Planck Collaboration X 2016). We assume the fiducial Planck cosmological model and use realistic Planck simulations to estimate the ensemble average values for signal and noise, as required in Eq. (40). The quantity C^{1}a, required in Eq. (41), is determined via the messenger field method (Elsner & Wandelt 2013). Some of the parameters (n_{B},θ_{B},φ_{B}) influence the signature correlation nonlinearly. Owing to these nonlinear parameters, we use the CosmoMC package (Lewis & Bridle 2002) as a generic sampler for the loglikelihood in Eq. (40)and obtain the posterior probability for the Alfvén wave parameters . As discussed previously, we assume the initial vector fluctuations to be entirely sourced by a nonhelical stochastic PMF, B. In Table 5, we show upper bounds on this combination of parameters at 68% and 95% CL, after marginalizing over the spectral index n_{B} and the direction θ_{B},φ_{B}.
Other theoretical models with correlations across multipoles with Δℓ = ± 1, ± 2 are investigated in Planck Collaboration XVI (2016) and Planck Collaboration XX (2016). The Planck data show no evidence in favour of these models.
Planck constraints on the Alfvén wave amplitude .
6. Conclusions
6.1. Methodology
In this paper, we have presented the constraints on a stochastic background of primordial magnetic fields using Planck data. PMFs may have left different types of imprints on the CMB. This is why the CMB can be regarded as one of the best laboratories for investigating and constraining PMFs. The richness of Planck data, which provide several different probes based on the statistics of CMB anisotropies, allows us to constrain PMFs using different methods, deriving constraints that are complementary to each other. Aiming for a broad perspective on PMFs, we have taken advantage of these different possibilities offered by Planck. In particular, we have exploited the impact of PMFs on the CMB anisotropy angular power spectra in temperature and polarization, both through the magneticallyinduced modes and their effect on the ionization history of the Universe. In addition we have considered the Faraday rotation induced by them on the CMB polarization. Beyond the twopoint statistics probed by the angular power spectrum, we have investigated higherorder statistical moments of the CMB imprinted by PMFs. In particular, we have analysed the CMB bispectrum for the presence of magneticallyinduced passive and compensated modes. Finally, we have considered the correlations between harmonic modes induced by PMFs.
6.2. Constraints for nonhelical fields from the angular power spectra
The constraints based on the CMB angular power spectra have been derived using the Planck likelihood. The general analysis, which only considers the contributions from compensated modes, provides the constraint B_{1 Mpc}< 4.4 nG at 95% CL, with positive spectral indices constrained to lower amplitudes of the fields, which we have shown to be robust under the inclusion of highℓ polarization data. In fact, the impact of PMFs on the EE and TE polarization spectra is negligible compared to the dominant contribution given by the magneticallyinduced vector modes on the TT spectra on small angular scales. The inclusion of the passive tensor contribution does not improve the constraints on the amplitude but affects the posterior distribution for the PMF spectral index. The results confirm that CMB data constrain PMFs with positive spectral indices to have smaller amplitudes. The inclusion of the contribution of passive modes, thanks to their sensitivity to the slope of the PMF spectrum, strongly disfavours nearly scaleinvariant magnetic spectra for amplitudes greater than very few nanoGauss (see Fig. 8). This is due to the large contribution that magneticallyinduced passive tensor modes with nearly scaleinvariant spectra give to the CMB anisotropies on large angular scales.
We have also performed an analysis using both the Planck 2015 data and the BICEP2/KeckPlanck crosscorrelation. The results are fully compatible with the analysis based on Planck data alone, with only a slightly higher upper limit on the PMF amplitude.
Our likelihood analysis for PMF properties is sensitive to CMB foreground residuals in the data, since these contribute to small angular scales in the CMB spectra as well. In our analysis we have used the foreground residual treatment provided by the Planck likelihood. Since some of the foreground models assume an angular power spectrum with a shape similar to the one given by magneticallyinduced perturbations, we have investigated the issue of possible degeneracies between PMFs and foreground residuals. In particular, we have noticed that there is a degeneracy with the Poissonian terms modelling unresolved point sources at 100, 143, 143 × 217 and 217 GHz. Fortunately, other foreground contributions do not show any degeneracy thanks to their different spectral shapes. We have further tested how severely foreground residuals might affect the obtained constraints on PMF amplitudes. Since PMFs that are consistent with the data do not significantly affect other cosmological parameters, we have fixed the amplitudes of foreground residuals to the values obtained under the Planck ΛCDM cosmology. The limit in this case, B_{1 Mpc}< 3.0 nG at 95% CL, is slightly tighter than the limit obtained without fixing the foreground residual amplitudes, B_{1 Mpc}< 4.4 nG at 95% CL. This test has no statistical significance in the constraints on PMFs but illustrates the impact foreground residuals can have on constraints on the PMF amplitude.
Our new constraints are compatible with previous constraints from other experiments and from the previous Planck release (Planck Collaboration XVI 2014). The slightly higher upper limits with respect to the 2013 Planck release are due to changes in the 2015 data. In particular, the changed calibration and slightly different slope of the power spectrum of cosmological perturbations allow for stronger PMFs, with possible contributions from the different foreground residual treatment.
6.3. Constraints on maximally helical PMFs
We also constrain maximally helical PMFs. We restrict our analysis to the maximally helical case because of the absence of TB and EB information in the Planck 2015 highℓ likelihood. Maximal helicity decreases the amplitude of the magnetically generated CMB fluctuations and, as a consequence, we obtain B_{1 Mpc}< 5.6 nG at 95% CL in this case.
6.4. Selected scenarios
We have further investigated two specific PMF models of interest: causally generated fields with a spectral index of n_{B} = 2 and fields with an almost scaleinvariant power spectrum with n_{B} = −2.9. The constraints for these extreme cases are nG and nG at 95% CL, respectively.
The impact of PMFs on the ionization history of the Universe directly affects the CMB temperature and polarization power spectra. In particular, we have also considered the main dissipative effects operating during and after recombination, namely ambipolar diffusion and energy cascading in MHD turbulence, in the prediction for the CMB spectra in temperature and polarization. These modify the primary CMB power spectra in addition to the gravitational contributions of the magnetic modes. For the nearly scaleinvariant case we have obtained the constraint nG at 95% CL. This limit is tighter than when neglecting the effect on the ionization history. However, uncertainties related to the modelling of the heating mechanism (see discussion by Chluba et al. 2015) suggests that further investigation of this promising avenue is needed.
6.5. NonGaussianitybased constraints
For the nonGaussianity analyses we have focused on the passive modes with a nearly scaleinvariant power spectrum, n_{B} = −2.9, and the compensated scalar modes. These are the dominant contributions on large angular scales, where the nonGaussianity analyses are performed.
In our first CMB nonGaussianity analysis, we have considered passive tensor modes for PMFs with nearly scaleinvariant spectra. These contribute predominantly to the CMB fluctuations on large angular scales. For this case, we have calculated the resulting CMB bispectrum and compared it with the observational limit. We have used a bimodal decomposition to estimate the amplitude of the Planck bispectrum in the squeezed configuration, in which the observational limit on the amplitude of the bispectrum can be translated into a constraint on the amplitude of PMFs. Using the temperature bispectrum we obtain nG for fields that were generated at the GUT phase transition.
For our second nonGaussianity analysis we have used a different approach to the magneticallyinduced bispectrum. We have considered the passive contributions by tensor and scalar modes for nearly scaleinvariant fields, but instead of using the bispectrum amplitude we have used the local type of nonGaussianity induced by PMFs, considering only the nonhelical part. The local bispectrum from multipoles with even (where ℓ_{n} stands for the multipoles ℓ_{1},ℓ_{2},ℓ_{3} normally used to express the bispectrum, see for example Eq. (29)) has been decomposed with coefficients determined by the amplitude of the PMF. A likelihood estimation allows us to constrain the PMF amplitude with CMB foregroundcleaned maps. We have applied this method to all four component separation methods used by Planck. The SMICA map, which is expected to contain the least foreground residuals (Planck Collaboration IX 2016), gives the constraint nG, improving previous constraints derived from the scalar bispectrum for WMAP data.
Our third nonGaussianity analysis focuses on the compensated scalar modes. In this case, we have used an analytic estimate of the bispectrum on large angular scales. We have used an improved estimate of the source term Caprini, et al. (in prep.) with respect to previous results (Caprini et al. 2009). This analytic estimate can be compared with the observed local f_{NL} from Planck, giving nG.
The results from the different nonGaussianity analyses (although coming from different methods) are all consistent and at the level of those derived with the likelihood analysis using only the CMB angular power spectra.
6.6. Constraints from Faraday rotation
We have further considered the effects of Faraday rotation on the primary CMB polarization anisotropies. In this context, we have used the EE and BBpolarization power spectra. We have derived the constraints on the PMF amplitude using a χ^{2} analysis based on the LFI 70 GHz lowℓ (ℓ< 30) polarization power spectra. The resulting constraint is B_{1 Mpc}< 1380 nG. The upper limits from Faraday rotation are larger than those derived from magneticallyinduced perturbations, thermal effects, and nonGaussianity. On one hand, the Faraday rotation signal rapidly vanishes with increasing frequency (see Eq. (36)) and thus is strong only for lower frequencies. On the other hand, since the BB spectrum is the result of the rotation of the EE spectrum, it has a stronger contribution on smaller angular scales. Our analysis includes only the low multipoles of the 70 GHz data, where the signal is lower owing to its spectral shape. The combination of data sets and type of signal results in less stringent constraints. However, even with the restricted subset of Planck data available, the constraints are only slightly weaker than derived in previous analyses performed with WMAP (Kahniashvili et al. 2009; Pogosian et al. 2011).
To estimate the impact of Galactic Faraday rotation on the results, we have analysed synthetic Galactic Faraday maps as well as radio synchrotron data at 1.4 GHz and 23 GHz. We have accounted for the fact that the signal is not isotropic on the sky but depends on the latitude of the observations. We have derived an estimate of the power spectrum for the Faraday depth, shown in Fig. 13, and compared it with the predictions for different values of the PMF amplitude. Our results show that the threshold for which the Galactic contamination may become relevant is around 10 nG. This amplitude is much below our current constraints, which can therefore be considered clean from Galactic contamination.
6.7. Constraints on Alfvén waves
To complete the round of different types of analyses involving different probes, we have investigated the correlation induced between different modes in harmonic space by Alfvén waves produced by the presence of PMFs (Kahniashvili et al. 2008). This correlation has not been used to constrain the PMF amplitude directly, but the Alfvén wave parameter, which is a combination of the stochastic background amplitude, the Alfvén velocity, and the assumed mean background field. We have again used the SMICA foregroundcleaned map to derive the upper limit . From this constraint we deduce that the data do not show any evidence of Alfvén waves, a conclusion that was also reached in previous analyses (Planck Collaboration XXIV 2014) carried out with more generic assumptions on the origin of the Alfvén waves. The absence of Alfvén waves is also compatible with the results from other models for the harmonic space correlations that are not related to the PMFs (Planck Collaboration XVI 2016; Planck Collaboration XX 2016).
6.8. Concluding summary
The results presented show that CMB anisotropies are one of the best probes for investigating the nature of PMFs. The Planck 2015 data offer the possibility to use the PMFs signatures either in the angular power spectra in temperature and polarization or in higherorder statistics, where both measurements can be tackled by different methodologies. All the independent constraints we obtain are consistent with each other.
The Planck 2015 data constrain the PMF amplitude at the nanoGauss level. Different signatures are sensitive to different contributions and may be optimal for specific types of PMF. In particular, the analysis that uses the gravitational impact of PMFs on the CMB angular power spectra is dominated by the compensated vector contribution on small angular scales and therefore is able to constrain PMFs without any assumptions on their generation mechanism. On the contrary, two of the three analyses of nonGaussianities are dominated by passive tensor modes, which can provide significant constraints only for nearly scaleinvariant PMFs.
The future of both classes of methods, the angular power spectra and the nonGaussianities, is bright, but three avenues are particularly promising. The helicity of PMFs will be constrained by TB and EB crosscorrelations, which will be included in the next Planck release. The study of the PMF’s impact on the ionization history is expected to further improve with future Planck polarization data. NonGaussianities are a distinctive signature of PMFs and further studies may provide more and more refined predictions of the magneticallyinduced passive and compensated CMB bispectra and trispectra, which will improve the future Planck analyses.
Planck (http://www.esa.int/Planck) is a project of the European Space Agency (ESA) with instruments provided by two scientific consortia funded by ESA member states and led by Principal Investigators from France and Italy, telescope reflectors provided through a collaboration between ESA and a scientific consortium led and funded by Denmark, and additional contributions from NASA (USA).
Note that, unlike in previous analyses (Finelli et al. 2008; Paoletti et al. 2009; Paoletti & Finelli 2011), we explicitly include the dependence of the damping scale on the baryon density to account for the updated cosmology with respect to the previous treatments.
For a quasiscaleinvariant PMF power spectrum (n_{B} ≈ −3), the estimated value for k_{D} varies slightly, by about 30–40%, for different approaches (Subramanian & Barrow 1998a; Jedamzik et al. 2000; Kunze & Komatsu 2015). However, this does not affect the main results of this paper significantly.
Note that we use the notation of Ma & Bertschinger (1995) for the scalar anisotropic stress σ_{B}.
This convergence criterion is based on the analysis of the variance within a chain and between chains. The posterior marginal variance is a weighted average of the different variances. If all the chains have reached convergence, this value will be very close to the variance within each single chain. The parameter R is defined as the square root of the ratio of the posterior marginal variance and the variance within a chain. With this definition, the closer R is to unity, the closer to convergence the chains are.
As explained by Chluba et al. (2015), the effect of PMFs on the ionization history was previously overestimated using an approach similar to the one of RECFAST. The reason is that in RECFAST the photoionization rates are evaluated assuming temperature T ≡ T_{e}, although from a physical point of view the radiation temperature should be used. This reduces the effect on the ionization history by up to one order of magnitude and a consistent treatment is implemented both in Recfast++ and CosmoRec (Chluba & Thomas 2011).
This peculiarity is in common with topological defects (see Figueroa et al. 2010).
Acknowledgments
The Planck Collaboration acknowledges the support of: ESA; CNES and CNRS/INSUIN2P3INP (France); ASI, CNR, and INAF (Italy); NASA and DoE (USA); STFC and UKSA (UK); CSIC, MINECO, JA, and RES (Spain); Tekes, AoF, and CSC (Finland); DLR and MPG (Germany); CSA (Canada); DTU Space (Denmark); SER/SSO (Switzerland); RCN (Norway); SFI (Ireland); FCT/MCTES (Portugal); ERC and PRACE (EU). A description of the Planck Collaboration and a list of its members, indicating which technical or scientific activities they have been involved in, can be found at http://www.cosmos.esa.int/web/planck/planckcollaboration. This research used resources of the National Energy Research Scientific Computing Center, a DOE Office of Science User Facility supported by the Office of Science of the US Department of Energy under Contract No. DEAC0205CH11231. Some of the results in this paper have been derived using the HEALPix package.
References
 Adamek, J., Durrer, R., Fenu, E., & Vonlanthen, M. 2011, J. Cosmol. Astropart. Phys., 6, 17 [NASA ADS] [CrossRef] (In the text)
 Ade, P. A. R., Arnold, K., Atlas, M., et al. 2015, Phys. Rev. D, 92, 123509 [NASA ADS] [CrossRef] (In the text)
 Ballardini, M., Finelli, F., & Paoletti, D. 2015, J. Cosmol. Astropart. Phys., 10, 031 [NASA ADS] [CrossRef] (In the text)
 Barrow, J. D., Ferreira, P. G., & Silk, J. 1997, Phys. Rev. Lett., 78, 3610 [NASA ADS] [CrossRef] (In the text)
 Battaner, E., Florido, E., & JimenezVicente, J. 1997, A&A, 326, 13 [NASA ADS] (In the text)
 Beck, R. 2000, in IAU Joint Discussion, 14 (In the text)
 Bennett, C. L., Larson, D., Weiland, J. L., et al. 2013, ApJS, 208, 20 [NASA ADS] [CrossRef] (In the text)
 Bernet, M. L., Miniati, F., Lilly, S. J., Kronberg, P. P., & DessaugesZavadsky, M. 2008, Nature, 454, 302 [NASA ADS] [CrossRef] [PubMed] (In the text)
 BICEP2/Keck Array and Planck Collaborations 2015, Phys. Rev. Lett., 114, 101301 [NASA ADS] [CrossRef] [PubMed] (In the text)
 Biskamp, D. 2003, Magnetohydrodynamic Turbulence (Cambridge, UK: Cambridge University Press) (In the text)
 Bonvin, C. 2010, ArXiv eprints [arXiv:1005.3332] (In the text)
 Bonvin, C., & Caprini, C. 2010, J. Cosmol. Astropart. Phys., 5, 22 [NASA ADS] [CrossRef] (In the text)
 Bonvin, C., Caprini, C., & Durrer, R. 2013, Phys. Rev. D, 88, 083515 [NASA ADS] [CrossRef] (In the text)
 Broderick, A. E., Chang, P., & Pfrommer, C. 2012, ApJ, 752, 22 [NASA ADS] [CrossRef] (In the text)
 Brown, I., & Crittenden, R. 2005, Phys. Rev. D, 72, 063002 [NASA ADS] [CrossRef] (In the text)
 Burigana, C., & Zizzo, A. 2006, Astron. Nachr., 327, 424 [NASA ADS] [CrossRef] (In the text)
 Burigana, C., Danese, L., & de Zotti, G. 1991, A&A, 246, 49 [NASA ADS] (In the text)
 Cai, R.G., Hu, B., & Zhang, H.B. 2010, J. Cosmol. Astropart. Phys., 1008, 025 [NASA ADS] [CrossRef] (In the text)
 Campanelli, L., Dolgov, A. D., Giannotti, M., & Villante, F. L. 2004, ApJ, 616, 1 [NASA ADS] [CrossRef] (In the text)
 Caprini, C. 2006, Astron. Nachr., 327, 422 [NASA ADS] [CrossRef] (In the text)
 Caprini, C. 2010, PoS (Texas 2010) 222 (In the text)
 Caprini, C., & Durrer, R. 2002, Phys. Rev. D, 65, 023517 [NASA ADS] [CrossRef] (In the text)
 Caprini, C., Durrer, R., & Kahniashvili, T. 2004, Phys. Rev. D, 69, 063006 [NASA ADS] [CrossRef] (In the text)
 Caprini, C., Finelli, F., Paoletti, D., & Riotto, A. 2009, J. Cosmol. Astropart. Phys., 6, 21 [NASA ADS] [CrossRef] (In the text)
 Carroll, S. M., Field, G. B., & Jackiw, R. 1990, Phys. Rev. D, 41, 1231 [NASA ADS] [CrossRef] (In the text)
 Chluba, J., & Sunyaev, R. A. 2012, MNRAS, 419, 1294 [NASA ADS] [CrossRef] (In the text)
 Chluba, J., & Thomas, R. M. 2011, MNRAS, 412, 748 [NASA ADS] (In the text)
 Chluba, J., Paoletti, D., Finelli, F., & RubiñoMartín, J. A. 2015, MNRAS, 451, 2244 [NASA ADS] [CrossRef] (In the text)
 De, S., Pogosian, L., & Vachaspati, T. 2013, Phys. Rev. D, 88, 063527 [NASA ADS] [CrossRef] (In the text)
 Dermer, C. D., Cavadini, M., Razzaque, S., et al. 2011, ApJ, 733, L21 [NASA ADS] [CrossRef] (In the text)
 Durrer, R., & Caprini, C. 2003, J. Cosmol. Astropart. Phys., 11, 10 [NASA ADS] [CrossRef] (In the text)
 Durrer, R., & Neronov, A. 2013, A&ARv, 21, 62 [NASA ADS] [CrossRef] (In the text)
 Durrer, R., Kahniashvili, T., & Yates, A. 1998, Phys. Rev. D, 58, 123004 [NASA ADS] [CrossRef] (In the text)
 Durrer, R., Ferreira, P. G., & Kahniashvili, T. 2000, Phys. Rev. D, 61, 043001 [NASA ADS] [CrossRef] (In the text)
 Elsner, F., & Wandelt, B. D. 2013, A&A, 549, A111 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Fedeli, C., & Moscardini, L. 2012, J. Cosmol. Astropart. Phys., 11, 55 [NASA ADS] [CrossRef] (In the text)
 Feretti, L., Giovannini, G., Govoni, F., & Murgia, M. 2012, A&ARv, 20, 54 [NASA ADS] [CrossRef] (In the text)
 Fergusson, J. R. 2014, Phys. Rev. D, 90, 043533 [NASA ADS] [CrossRef] (In the text)
 Fergusson, J., Liguori, M., & Shellard, E. 2010, Phys. Rev. D, 82, 023502 [NASA ADS] [CrossRef] (In the text)
 Fergusson, J., Liguori, M., & Shellard, E. 2012, J. Cosmol. Astropart. Phys., 1212, 032 [NASA ADS] [CrossRef] (In the text)
 Ferrari, C., Govoni, F., Schindler, S., Bykov, A. M., & Rephaeli, Y. 2008, Space Sci. Rev., 134, 93 [NASA ADS] [CrossRef] (In the text)
 Figueroa, D. G., Caldwell, R. R., & Kamionkowski, M. 2010, Phys. Rev. D, 81, 123504 [NASA ADS] [CrossRef] (In the text)
 Finelli, F., Paci, F., & Paoletti, D. 2008, Phys. Rev. D, 78, 023510 [NASA ADS] [CrossRef] (In the text)
 Fixsen, D. J. 2009, ApJ, 707, 916 [NASA ADS] [CrossRef] (In the text)
 Gelman, A., & Rubin, D. 1992, Statistical Science, 7, 457, http://www.stat.columbia.edu/~gelman/research/published/itsim.pdf [NASA ADS] [CrossRef] (In the text)
 Giovannini, M. 2004a, Phys. Rev. D, 70, 123507 [NASA ADS] [CrossRef] (In the text)
 Giovannini, M. 2004b, Inter. J. Mod. Phys. D, 13, 391 [NASA ADS] [CrossRef] [MathSciNet] (In the text)
 Giovannini, M. 2008, in String Theory and Fundamental Interactions, eds. M. Gasperini, & J. Maharana (Berlin Springer Verlag), Lect. Notes Phys., 737, 863 (In the text)
 Giovannini, M., & Kunze, K. E. 2008a, Phys. Rev. D, 77, 123001 [NASA ADS] [CrossRef] (In the text)
 Giovannini, M., & Kunze, K. E. 2008b, Phys. Rev. D, 77, 063003 [NASA ADS] [CrossRef] (In the text)
 Giovannini, M., & Kunze, K. E. 2008c, Phys. Rev. D, 77, 061301 [NASA ADS] [CrossRef] (In the text)
 Govoni, F., & Feretti, L. 2004, Int. J. Mod. Phys. D, 13, 1549 [NASA ADS] [CrossRef] (In the text)
 Grasso, D., & Riotto, A. 1998, Phys. Lett. B, 418, 258 [NASA ADS] [CrossRef] (In the text)
 Grasso, D., & Rubinstein, H. R. 1995, Astropart. Phys., 3, 95 [NASA ADS] [CrossRef] (In the text)
 Hamann, J., Lesgourgues, J., & Mangano, G. 2008, J. Cosmol. Astropart. Phys., 3, 004 [NASA ADS] [CrossRef] (In the text)
 Hastings, W. 1970, Biometrika, 57, 97 [CrossRef] [MathSciNet] (In the text)
 Hu, W., & Silk, J. 1993, Phys. Rev. D, 48, 485 [NASA ADS] [CrossRef] (In the text)
 Hu, W., & White, M. 1997, ApJ, 479, 568 [NASA ADS] [CrossRef] (In the text)
 Ichikawa, K., & Takahashi, T. 2006, Phys. Rev. D, 73, 063528 [NASA ADS] [CrossRef] (In the text)
 Ichiki, K., Takahashi, K., Ohno, H., Hanayama, H., & Sugiyama, N. 2006, Science, 311, 827 [NASA ADS] [CrossRef] (In the text)
 Jain, R. K., & Sloth, M. S. 2012, Phys. Rev. D, 86, 123528 [NASA ADS] [CrossRef] (In the text)
 Jansson, R., & Farrar, G. R. 2012, ApJ, 761, L11 [NASA ADS] [CrossRef] (In the text)
 Jedamzik, K., Katalinić, V., & Olinto, A. V. 1998, Phys. Rev. D, 57, 3264 [NASA ADS] [CrossRef] (In the text)
 Jedamzik, K., Katalinić, V., & Olinto, A. V. 2000, Phys. Rev. Lett., 85, 700 [NASA ADS] [CrossRef] [PubMed] (In the text)
 Kahniashvili, T. 2005, New A Rev., 49, 79 [NASA ADS] [CrossRef] (In the text)
 Kahniashvili, T., & Lavrelashvili, G. 2010, ArXiv eprints [arXiv:1010.4543] (In the text)
 Kahniashvili, T., & Ratra, B. 2005, Phys. Rev. D, 71, 103006 [NASA ADS] [CrossRef] (In the text)
 Kahniashvili, T., & Ratra, B. 2007, Phys. Rev. D, 75, 023002 [NASA ADS] [CrossRef] (In the text)
 Kahniashvili, T., Kosowsky, A., Mack, A., & Durrer, R. 2001, in Cosmology and Particle Physics, eds. R. Durrer, J. GarciaBellido, & M. Shaposhnikov, AIP Conf. Ser., 555, 451 (In the text)
 Kahniashvili, T., Lavrelashvili, G., & Ratra, B. 2008, Phys. Rev. D, 78, 063012 [NASA ADS] [CrossRef] (In the text)
 Kahniashvili, T., Maravin, Y., & Kosowsky, A. 2009, Phys. Rev. D, 80, 023009 [NASA ADS] [CrossRef] [MathSciNet] (In the text)
 Kahniashvili, T., Tevzadze, A. G., Sethi, S. K., Pandey, K., & Ratra, B. 2010, Phys. Rev. D, 82, 083005 [NASA ADS] [CrossRef] (In the text)
 Kahniashvili, T., Maravin, Y., Lavrelashvili, G., & Kosowsky, A. 2014, Phys. Rev. D, 90, 083004 [NASA ADS] [CrossRef] (In the text)
 Kim, E.J., Olinto, A. V., & Rosner, R. 1996, ApJ, 468, 28 [NASA ADS] [CrossRef] (In the text)
 Kim, J., & Naselsky, P. 2009, J. Cosmol. Astropart. Phys., 7, 41 [NASA ADS] [CrossRef] (In the text)
 Komatsu, E., Spergel, D. N., & Wandelt, B. D. 2005, ApJ, 634, 14 [NASA ADS] [CrossRef] (In the text)
 Kosowsky, A., & Loeb, A. 1996, ApJ, 469, 1 [NASA ADS] [CrossRef] (In the text)
 Kosowsky, A., Kahniashvili, T., Lavrelashvili, G., & Ratra, B. 2005, Phys. Rev. D, 71, 043006 [NASA ADS] [CrossRef] (In the text)
 Kuchar, P., & Enßlin, T. A. 2011, A&A, 529, A13 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Kunze, K. E. 2011, Phys. Rev. D, 83, 023006 [NASA ADS] [CrossRef] (In the text)
 Kunze, K. E. 2013, Plasma Physics and Controlled Fusion, 55, 124026 [NASA ADS] [CrossRef] (In the text)
 Kunze, K. E., & Komatsu, E. 2014, J. Cosmol. Astropart. Phys., 1, 9 [NASA ADS] [CrossRef] (In the text)
 Kunze, K. E., & Komatsu, E. 2015, J. Cosmol. Astropart. Phys., 6, 027 [NASA ADS] [CrossRef] (In the text)
 Lewis, A. 2004, Phys. Rev. D, 70, 043011 [NASA ADS] [CrossRef] (In the text)
 Lewis, A., & Bridle, S. 2002, Phys. Rev. D, 66, 103511 [NASA ADS] [CrossRef] (In the text)
 Lewis, A., & Bridle, S. 2011, Astrophysics Source Code Library [record ascl:1106.025] (In the text)
 Lewis, A., & Challinor, A. 2011, Astrophysics Source Code Library [record ascl:1102.026] (In the text)
 Ma, C.P., & Bertschinger, E. 1995,ApJ, 455, 7 (In the text)
 Mack, A., Kahniashvili, T., & Kosowsky, A. 2002, Phys. Rev. D, 65, 123004 [NASA ADS] [CrossRef] (In the text)
 Moss, A., Scott, D., Zibin, J. P., & Battye, R. 2011, Phys. Rev. D, 84, 023014 [NASA ADS] [CrossRef] (In the text)
 Neronov, A., & Vovk, I. 2010, Science, 328, 73 [NASA ADS] [CrossRef] [PubMed] (In the text)
 Neronov, A., Semikoz, D., & Banafsheh, M. 2013a, ArXiv eprints [arXiv:1305.1450] (In the text)
 Neronov, A., Taylor, A. M., Tchernin, C., & Vovk, I., 2013b, A&A, 554, A31 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Oppermann, N., Junklewitz, H., Greiner, M., et al. 2015, A&A, 575, A118 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Paoletti, D., & Finelli, F. 2011, Phys. Rev. D, 83, 123533 [NASA ADS] [CrossRef] (In the text)
 Paoletti, D., & Finelli, F. 2013, Phys. Lett. B, 726, 45 [NASA ADS] [CrossRef] (In the text)
 Paoletti, D., Finelli, F., & Paci, F. 2009, MNRAS, 396, 523 [NASA ADS] [CrossRef] (In the text)
 Planck Collaboration XV. 2014, A&A, 571, A15 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XVI. 2014, A&A, 571, A16 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XXIV. 2014, A&A, 571, A24 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration I. 2016, A&A, 594, A1 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration II. 2016, A&A, 594, A2 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration III. 2016, A&A, 594, A3 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration IV. 2016, A&A, 594, A4 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration V. 2016, A&A, 594, A5 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration VI. 2016, A&A, 594, A6 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration VII. 2016, A&A, 594, A7 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration VIII. 2016, A&A, 594, A8 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration IX. 2016, A&A, 594, A9 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration X. 2016, A&A, 594, A10 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XI. 2016, A&A, 594, A11 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XII. 2016, A&A, 594, A12 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XIII. 2016, A&A, 594, A13 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XIV. 2016, A&A, 594, A14 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XV. 2016, A&A, 594, A15 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XVI. 2016, A&A, 594, A16 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XVII. 2016, A&A, 594, A17 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XVIII. 2016, A&A, 594, A18 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XIX. 2016, A&A, 594, A19 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XX. 2016, A&A, 594, A20 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XXI. 2016, A&A, 594, A21 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XXII. 2016, A&A, 594, A22 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XXIII. 2016, A&A, 594, A23 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XXIV. 2016, A&A, 594, A24 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XXV. 2016, A&A, 594, A25 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XXVI. 2016, A&A, 594, A26 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XXVII. 2016, A&A, 594, A27 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Planck Collaboration XXVIII. 2016, A&A, 594, A28 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Pogosian, L., Vachaspati, T., & Winitzki, S. 2002, Phys. Rev. D, 65, 083502 [NASA ADS] [CrossRef] (In the text)
 Pogosian, L., Yadav, A. P. S., Ng, Y.F., & Vachaspati, T. 2011, Phys. Rev. D, 84, 043530 [NASA ADS] [CrossRef] (In the text)
 Ratra, B. 1992, ApJ, 391, L1 [NASA ADS] [CrossRef] (In the text)
 RuizGranados, B., RubiñoMartín, J. A., & Battaner, E. 2010, A&A, 522, A73 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Ryu, D., Schleicher, D. R. G., Treumann, R. A., Tsagas, C. G., & Widrow, L. M. 2012, Space Sci. Rev., 166, 1 [NASA ADS] [CrossRef] (In the text)
 Schleicher, D. R. G., Banerjee, R., & Klessen, R. S. 2008, Phys. Rev. D, 78, 083005 [NASA ADS] [CrossRef] (In the text)
 Scóccola, C., Harari, D., & Mollerach, S. 2004, Phys. Rev. D, 70, 063003 [NASA ADS] [CrossRef] (In the text)
 Seshadri, T., & Subramanian, K. 2009, Phys. Rev. Lett., 103, 081303 [NASA ADS] [CrossRef] (In the text)
 Sethi, S. K., & Subramanian, K. 2005, MNRAS, 356, 778 [NASA ADS] [CrossRef] (In the text)
 Shaw, J. R., & Lewis, A. 2010, Phys. Rev. D, 81, 043517 [NASA ADS] [CrossRef] (In the text)
 Shaw, J. R., & Lewis, A. 2012, Phys. Rev. D, 86, 043510 [NASA ADS] [CrossRef] (In the text)
 Shiraishi, M. 2012, J. Cosmol. Astropart. Phys., 6, 15 [NASA ADS] [CrossRef] (In the text)
 Shiraishi, M. 2013, J. Cosmol. Astropart. Phys., 1311, 006 [NASA ADS] [CrossRef] (In the text)
 Shiraishi, M., & Sekiguchi, T. 2014, Phys. Rev. D, 90, 103002 [NASA ADS] [CrossRef] (In the text)
 Shiraishi, M., Nitta, D., Yokoyama, S., Ichiki, K., & Takahashi, K. 2010, Phys. Rev. D, 82, 121302 [NASA ADS] [CrossRef] (In the text)
 Shiraishi, M., Nitta, D., Yokoyama, S., Ichiki, K., & Takahashi, K. 2011, Phys. Rev. D, 83, 123003 [NASA ADS] [CrossRef] (In the text)
 Shiraishi, M., Nitta, D., Yokoyama, S., & Ichiki, K. 2012, J. Cosmol. Astropart. Phys., 1203, 041 [NASA ADS] [CrossRef] (In the text)
 Shiraishi, M., Komatsu, E., Peloso, M., & Barnaby, N. 2013, J. Cosmol. Astropart. Phys., 5, 2 [NASA ADS] [CrossRef] (In the text)
 Shiraishi, M., Liguori, M., & Fergusson, J. R. 2014, J. Cosmol. Astropart. Phys., 1405, 008 [NASA ADS] [CrossRef] (In the text)
 Shiraishi, M., Liguori, M., & Fergusson, J. R. 2015, J. Cosmol. Astropart. Phys., 1, 7 [NASA ADS] [CrossRef] (In the text)
 Subramanian, K., & Barrow, J. D. 1998a, Phys. Rev. D, 58, 083502 [NASA ADS] [CrossRef] (In the text)
 Subramanian, K., & Barrow, J. D. 1998b, Phys. Rev. Lett., 81, 3575 [NASA ADS] [CrossRef] (In the text)
 Subramanian, K., & Barrow, J. D. 2002, MNRAS, 335, L57 [NASA ADS] [CrossRef] (In the text)
 Subramanian, K., Seshadri, T. R., & Barrow, J. D. 2003, MNRAS, 344, L31 [NASA ADS] [CrossRef] (In the text)
 Tashiro, H., Chen, W., Ferrer, F., & Vachaspati, T. 2014, MNRAS, 445, L41 [NASA ADS] [CrossRef] (In the text)
 Tavecchio, F., Ghisellini, G., Foschini, L., et al. 2010, MNRAS, 406, L70 [NASA ADS] (In the text)
 Tavecchio, F., Ghisellini, G., Bonnoli, G., & Foschini, L. 2011, MNRAS, 414, 3566 [NASA ADS] [CrossRef] (In the text)
 Taylor, A. M., Vovk, I., & Neronov, A. 2011, A&A, 529, A144 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Trivedi, P., Subramanian, K., & Seshadri, T. 2010, Phys. Rev. D, 82, 123006 [NASA ADS] [CrossRef] (In the text)
 Trivedi, P., Seshadri, T., & Subramanian, K. 2012, Phys. Rev. Lett., 108, 231301 [NASA ADS] [CrossRef] (In the text)
 Trivedi, P., Subramanian, K., & Seshadri, T. 2014, Phys. Rev. D, 89, 043523 [NASA ADS] [CrossRef] (In the text)
 Turner, M. S., & Widrow, L. M. 1988, Phys. Rev. D, 37, 2743 [NASA ADS] [CrossRef] (In the text)
 Vachaspati, T. 1991, Phys. Lett. B, 265, 258 [NASA ADS] [CrossRef] (In the text)
 Vallée, J. P. 2004, New Astron. Rev., 48, 763 [NASA ADS] [CrossRef] (In the text)
 Vovk, I., Taylor, A. M., Semikoz, D., & Neronov, A. 2012, ApJ, 747, L14 [NASA ADS] [CrossRef] (In the text)
 Wasserman, I. 1978, ApJ, 224, 337 [NASA ADS] [CrossRef] (In the text)
 Widrow, L. M. 2002, Rev. Mod. Phys., 74, 775 [NASA ADS] [CrossRef] (In the text)
 Widrow, L. M., Ryu, D., Schleicher, D. R. G., et al. 2012, Space Sci. Rev., 166, 37 [NASA ADS] [CrossRef] (In the text)
 Wolfe, A. M., Jorgenson, R. A., Robishaw, T., Heiles, C., & Prochaska, J. X. 2008, Nature, 455, 638 [NASA ADS] [CrossRef] (In the text)
 Wolleben, M., Landecker, T. L., Reich, W., & Wielebinski, R. 2006, A&A, 448, 411 [NASA ADS] [CrossRef] [EDP Sciences] (In the text)
 Yamazaki, D. G., Ichiki, K., Kajino, T., & Mathews, G. J. 2007, in Particles, Strings, and CosmologyPASCOS, eds. A. Rajantie, C. Contaldi, P. Dauncey, & H. Stoica, AIP Conf. Ser., 957, 449 (In the text)
 Yamazaki, D. G., Ichiki, K., Kajino, T., & Mathews, G. J. 2008, Phys. Rev. D, 77, 043005 [NASA ADS] [CrossRef] (In the text)
 Yamazaki, D. G., Ichiki, K., Kajino, T., & Mathews, G. J. 2010, Phys. Rev. D, 81, 023008 [NASA ADS] [CrossRef] (In the text)
 Zizzo, A., & Burigana, C. 2005, New Astron., 11, 1 [NASA ADS] [CrossRef] (In the text)
Appendix A: Impact of foregrounds on PMF constraints from the angular power spectra
In Fig. A.1 we present the twodimensional proability distributions of PMF amplitude and foreground parameters for the Planck 2013 likelihood, which show only a mild degeneracy with the Poissonian amplitude for the 143 GHz channel.
Fig. A.1 Probability contours for the PMF amplitude and the foreground parameters for the Planck 2013 likelihood. 

Open with DEXTER 
In Fig. A.2, we plot the twodimensional probability distributions of the PMF amplitude and the foreground parameters (for their description see Planck Collaboration XI 2016), except for the Poissonian terms, which have been discussed in Sect. 2.4.4. These plots do not show any degeneracy.
Fig. A.2 Probability contours for the PMF amplitude and the foreground parameters. 

Open with DEXTER 
Appendix B: Statistical anisotropy induced by Alfvén waves
It has been shown that the presence of Alfvén waves in the early Universe leads to specific correlations of the CMB in harmonic space (Kahniashvili et al. 2008). The signature correlations induced by Alfvén waves are as follows:
Here C_{ℓ} is the power spectrum in the absence of Alfvén waves, θ_{B} and φ_{B} are the spherical angles of the direction of the background magnetic field , and is given by (B.1)where k_{D} is the comoving wave number of the dissipation scale owing to photon viscosity and given approximately by 10 / (cτ_{last}) (Durrer et al. 1998). The dissipative damping effect becomes significant at multipoles ℓ ≳ 500 (Durrer et al. 1998). We note that the damping scale considered in this context is different from the one considered in Sect. 2. The damping effects considered in each case are related to two different aspects and are specific to the two topics treated; in the study of the impact of PMFs on CMB anisotropies, the damping scale considered is due to the dissipation of the PMFs themselves as investigated by Subramanian & Barrow (1998a), in this section instead, we consider a damping scale derived from the damping of the vector perturbations generated by the PMFs, not the PMFs themselves. The latter damping scale is derived by Durrer et al. (1998). In Eq. (B.1), denotes the total power of the nonhelical PMF smoothed at the spatial scale of λ. There are general cases for dipole and quadrupole coupling, where cosmological parameters are assumed to vary with position (Moss et al. 2011). If we allow a positiondependent parameter to be a vector and treat the PMF as a positiondependent parameter, the correlation induced by Alfvén waves may be incorporated into the framework of this approach.
All Tables
Mean parameter values and bounds of the central 68% CL from Planck TT, TE, EE (left column) and Planck TT (right column).
95% CL upper bounds of the PMF amplitude for fixed spectral index with compensated plus passive tensor modes.
Planck constraints on the amplitude of the nonhelical magnetic field component, B_{1 Mpc} [nG], from the SMICA, NILC, SEVEM, and Commander foregroundcleaned maps at 95% CL.
All Figures
Fig. 1 Magneticallyinduced CMB TT (top left), TE (top right), EE (bottom left), and BB (bottom right) power spectra. The solid lines represent primary CMB anisotropies, the dotted lines represent magneticallyinduced compensated scalar modes (except for the BB panel, where it represents the lensing contributions and the solid line represents primary tensor modes with a tensortoscalar ratio of r = 0.1), the dashed lines represent vector modes, whereas the dotdashed lines represent magneticallyinduced compensated tensor modes. We consider PMFs with B_{1 Mpc} = 4.5 nG and n_{B} = −1. 

Open with DEXTER  
In the text 
Fig. 2 Dependence of the magneticallyinduced CMB power spectrum on the spectral index. For all plotted cases, the amplitude is B_{1 Mpc} = 4.5 nG. The black lines show primary CMB anisotropies; for the other colours we refer to the legend. Left: scalar contributions, right: vector contributions. 

Open with DEXTER  
In the text 
Fig. 3 Magneticallyinduced CMB TT (top left), TE (top right), EE (bottom left), and BB (bottom right) power spectra due to passive tensor modes, compared with the ones that are due to compensated modes. The solid lines represent primary CMB anisotropies, the dotted lines represent magneticallyinduced compensated scalar modes (except for the BB panel, where it represents the lensing contribution), the dashed lines represent vector modes, whereas dotdashed lines represent magneticallyinduced passive tensor modes. We consider PMFs with B_{1 Mpc} = 4.5 nG and n_{B} = −2.9. 

Open with DEXTER  
In the text 
Fig. 4 Dependence of the magneticallyinduced CMB power spectrum due to passive tensor modes on the spectral index for a GUTscale PMF (left) and comparison between the two extremes for the time ratio τ_{ν}/τ_{B} (right). The black lines show the primary CMB anisotropies; for the other colours we refer to the legend. Solid lines represent PMFs generated at the GUT scale, τ_{ν}/τ_{B} = 10^{17}, whereas dashed lines represent PMFs generated at late times, τ_{ν}/τ_{B} = 10^{6}. 

Open with DEXTER  
In the text 
Fig. 5 CMB TT (top left), TE (top right), EE (bottom left), and BB (bottom right) power spectra due to helical PMFs compared to the ones that are due to nonhelical PMFs. Solid lines are nonhelical predictions, while dashed lines are helical predictions. Blue are the scalar modes, green the vector, and red the compensated tensor modes. We consider PMFs with B_{1 Mpc} = 4.5 nG and n_{B} = −1. 

Open with DEXTER  
In the text 
Fig. 6 Comparison of the constraints on the smoothed PMF amplitude (top) and the spectral index (bottom) from the 2015 temperature and temperature plus polarization data with the 2013 results for magneticallyinduced compensated initial conditions only. 

Open with DEXTER  
In the text 
Fig. 7 Constraints on the smoothed PMF amplitude (top) and spectral index (bottom) from Planck temperature data with and without the passive tensor contribution. Constraints including both compensated and passive modes are indicated with C+P in the legend, constraints using only compensated modes are marked with C. 

Open with DEXTER  
In the text 
Fig. 8 PMF amplitude versus the spectral index for the baseline Planck 2015 case. C+P denotes the case where both compensated and passive modes are considered, whereas C indicates the case with only compensated modes. The two contours represent the 68% and 95% confidence levels. 

Open with DEXTER  
In the text 
Fig. 9 Twodimensional posterior distributions of the PMF amplitude versus the parameter that describes the Poissonian term of unresolved point sources for the three frequencies considered in the likelihood. The two contours represent the 68% and 95% confidence levels. 

Open with DEXTER  
In the text 
Fig. 10 Probability distributions for the PMF amplitude including the BICEP2/KeckPlanck crosscorrelation, compared with the one based only on Planck data. Top: the case in which the spectral index is free to vary, bottom: the case with n_{B} = −2.9. 

Open with DEXTER  
In the text 
Fig. 11 PMF amplitude constraint for the helical case (solid black) compared with the nonhelical case (dashed red). The dotted blue line shows the constraint on the amplitude of the helical component as an alternative interpretation of the constraints on the amplitude of PMFs with a helical component. 

Open with DEXTER  
In the text 
Fig. 12 Probability contours of PMF strength vs. spectral index of the PMF power spectrum as constrained by the 70 GHz observations. 

Open with DEXTER  
In the text 
Fig. 13 Power spectra of the primordial Faraday depth Φ for magnetic field strengths of 100 (solid line), 10 (dotted line), and 1 nG (dashed line), the Galactic Faraday depth from the allsky Faraday map of Oppermann et al. (2015) (blue dashed line), the Galactic Faraday depth from Oppermann et al. (2015) at 45° latitude (green dashed line), the Galactic Faraday depth derived from Galactic emission at 1.4 and 23 GHz (red circles), and the Galactic Faraday depth for a Galactic magnetic field model (black crosses). 

Open with DEXTER  
In the text 
Fig. A.1 Probability contours for the PMF amplitude and the foreground parameters for the Planck 2013 likelihood. 

Open with DEXTER  
In the text 
Fig. A.2 Probability contours for the PMF amplitude and the foreground parameters. 

Open with DEXTER  
In the text 