Radio Detection of Astrophysical Neutrinos

Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 1 Radio Detection of Astrophysical Neutrinos Jörg R. Hörandel! Radboud Universiteit, Nijmegen - Vrije Universiteit, Brussel Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 1 Radio Detection of Astrophysical Neutrinos Jörg R. Hörandel! Radboud Universiteit, Nijmegen - Vrije Universiteit, Brussel visible astronomy Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 1 Radio Detection of Astrophysical Neutrinos “New messengers” Visual astronomy Energy spectrum => birth of astrophysics! (Chemical composition, velocities, temperatures, etc.) 1st revolution (19th c.) 2nd revolution (20th c.) 3rd revolution (21st c.) ! "!" #n (radio, infrared, UV, X, gamma, TeV…) Spectroscopy Multi-wavelength astronomy Multi-messenger astronomy Non-thermal astronomy, High-energy astrophysics A brief history of astronomy… 1st revolution (19th c.): spectroscopy birth of astrophysics! chem. composition, temperatures, velocities Jörg R. Hörandel! Radboud Universiteit, Nijmegen - Vrije Universiteit, Brussel visible astronomy Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 1 Radio Detection of Astrophysical Neutrinos “New messengers” Visual astronomy Energy spectrum => birth of astrophysics! (Chemical composition, velocities, temperatures, etc.) 1st revolution (19th c.) 2nd revolution (20th c.) 3rd revolution (21st c.) ! "!" #n (radio, infrared, UV, X, gamma, TeV…) Spectroscopy Multi-wavelength astronomy Multi-messenger astronomy Non-thermal astronomy, High-energy astrophysics A brief history of astronomy… 1st revolution (19th c.): spectroscopy birth of astrophysics! chem. composition, temperatures, velocities Jörg R. Hörandel! Radboud Universiteit, Nijmegen - Vrije Universiteit, Brussel “New messengers” Visual astronomy Energy spectrum => birth of astrophysics! (Chemical composition, velocities, temperatures, etc.) 1st revolution (19th c.) 2nd revolution (20th c.) 3rd revolution (21st c.) ! "!" #n (radio, infrared, UV, X, gamma, TeV…) Spectroscopy Multi-wavelength astronomy Multi-messenger astronomy Non-thermal astronomy, High-energy astrophysics A brief history of astronomy… 2nd revolution (20th c.): ! multi-wavelength astronomy non-thermal astronomy, high-energy astrophysics visible astronomy Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 1 Radio Detection of Astrophysical Neutrinos “New messengers” Visual astronomy Energy spectrum => birth of astrophysics! (Chemical composition, velocities, temperatures, etc.) 1st revolution (19th c.) 2nd revolution (20th c.) 3rd revolution (21st c.) ! "!" #n (radio, infrared, UV, X, gamma, TeV…) Spectroscopy Multi-wavelength astronomy Multi-messenger astronomy Non-thermal astronomy, High-energy astrophysics A brief history of astronomy… 1st revolution (19th c.): spectroscopy birth of astrophysics! chem. composition, temperatures, velocities Jörg R. Hörandel! Radboud Universiteit, Nijmegen - Vrije Universiteit, Brussel “New messengers” Visual astronomy Energy spectrum => birth of astrophysics! (Chemical composition, velocities, temperatures, etc.) 1st revolution (19th c.) 2nd revolution (20th c.) 3rd revolution (21st c.) ! "!" #n (radio, infrared, UV, X, gamma, TeV…) Spectroscopy Multi-wavelength astronomy Multi-messenger astronomy Non-thermal astronomy, High-energy astrophysics A brief history of astronomy… 2nd revolution (20th c.): ! multi-wavelength astronomy non-thermal astronomy, high-energy astrophysics “New messengers” Visual astronomy Energy spectrum => birth of astrophysics! (Chemical composition, velocities, temperatures, etc.) 1st revolution (19th c.) 2nd revolution (20th c.) 3rd revolution (21st c.) ! "!" #n (radio, infrared, UV, X, gamma, TeV…) Spectroscopy Multi-wavelength astronomy Multi-messenger astronomy Non-thermal astronomy, High-energy astrophysics A brief history of astronomy… 3rd revolution (21th c.): ! multi-messenger astronomy new messengers, using all 4 fundamental forces visible astronomy Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 2 P.S. Shawhan, PoS(ICHEP2018)695 Radio Detection of Astrophysical Neutrinos Jörg R. Hörandel! Radboud Universiteit, Nijmegen - Vrije Universiteit, Brussel Multi-Messenger Astroparticle Physics Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 3 Radio Emission in Air Showers e+ e+ e+ e+eeeee-e+ edrift e+ drift air shower atmospheric nucleus cosmic ray coherent radio pulse deflection of particles in geomagnetic field B ee+ atmosphere of Earth is transparent in tens of MHz band Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 4 First radio detection of air showers 1965 Blackett’s Field ~1967, Porter MSc Jelley et al Nature 1965, R. A. Porter MSc Thesis 1967, Jelley et al Nature 1965 R. A. Porter MSc Thesis 1967 Blackett’s Field ~1967 Porter MSc JRH, EPJ Web of Conferences 216 (2019) 01003 Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 5 2005: understanding the signal 2014: understanding the emission processes 2016: radio technique mature: properties of cosmic rays 2018: beyond capabilities of standard installations 2024: Auger ! Radio Detector Auger Engineering Radio Array AERA Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 final bias-free selection of 594 showers (see Sec. IV). We find that the median resolution in Xmax shows a clear relation with shower energy, reaching a resolution of better than 15 g cm−2in the highest energy bin. We parametrize the resolution as δXmax ¼a·!!!!!!!!!!!!!!!! 1018 eV E r⊕b·1018 eV E⊕c; ð8Þ inspired by the energy resolution of electromagnetic calorimeters [47], but also functioning as a generic expansion in terms of energy. Here, a¼14.0$6.8g cm−2, b¼12.7$2.5g cm−2, and c¼11.2$4.7g cm−2are fitted free parameters, and ⊕indicates the quadratic sum. The constant term cprovides an indication of the resolution that might potentially be obtained for AERA with this method at the highest energies (given this parametrization). The change in resolution of Xmax is dominated by the uncertainty on the measured radio signals and hence becomes less accurate at lower energy. Comparing our resolution to the resolution achieved by FD, we achieve similar values at the highest energies where the FD reaches 15 g cm−2[40]. Furthermore, our method remains competitive down to lower energies where, for example, at E¼1017.8eV, the FD achieves the same resolution of 25 g cm−2. The most recent results by the LOFAR radio array, where a similar simulation-fitting method is used to determine Xmax, report an average resolution of 19 g cm−2between 1016.8and 1018.3eV [48]. Despite the much denser antenna spacing of LOFAR, AERA achieves similar resolutions considering the respective energy regimes to which the two experiments are sensitive. We note that up till now it has been common to quote a single resolution value for Xmax reconstruction methods for radio experiments, mainly because of a limited number of measured showers being available. Here, we show the resolution in Xmax depends strongly on the shower energy, driven primarily by the strength of radio signals measured in the antennas. Hence, the resolution is a function of detector sensitivity and shower energy and thus heavily depends on the shower selection criteria. As such, any direct comparison of methods is less straightforward if obtained at sufficiently dissimilar detectors. VI. Xmax MOMENTS AND THE DISTRIBUTION OF Xmax From the Xmax distribution, for each of our six energy bins, we now also calculate the first two moments of the distribution, the mean hXmaxiand the width σðXmaxÞ.To obtain the latter, we first subtract in quadrature the width caused by the method uncertainty, such that only the width caused by shower-to-shower fluctuations σðXmaxÞremains. The method uncertainty cannot simply be characterized by a single value since the uncertainties on Xmax for our air showers do not necessarily follow a Gaussian distribution. A bootstrap resampling procedure is applied for this reason, and with this, we then also calculate the uncertainty on σðXmaxÞ. This procedure is further described in Appendix C. The resulting mean and width of the true Xmax distribution are shown in Fig. 15, where we also compare this to the results from the FD (gray) and theoretical predictions of three different hadronic interaction models for a mass composition of just protons (red) or just iron nuclei (blue) [31,49–51]. The systematic uncertainties determined in the previous section are shown with capped markers. Table IV in Appendix Dlists the values for the two moments of the distributions for the six energy bins, together with their statistical and systematic uncertainties. With these AERA results, we show good agreement between the Auger radio and fluorescence measurements of hXmaxi, both pointing toward a (mixed)-light composition of cosmic rays at around E¼1017.5eV. Note that the two measurements share the systematic uncertainty on the energy scale, which is constructed from calibration of the SD energy to the FD energy scale [44]. Taking this contribution out reduces both systematic uncertainty bands by about 3g cm−2. Second, the determination of the systematic uncertainties due to the reconstruction method for AERA Xmax data depends on the assumed composition. FIG. 14. Resolution of the Xmax reconstruction method, δXmax , as a function of energy in units of column density. Shown per energy bin are the median values of the uncertainties on Xmax (circles with uncertainties σbfrom bootstrap resampling) for all showers in the bias-free sample (Table I) and a parametrized fit [Eq. (8)] of the resolution in Xmax (solid line with 1σ-confidence bands). Also shown are the resolutions achieved by the Auger fluorescence telescopes [40]. The black hatched region at low energy indicates the cut on energy applied earlier. The extent of each energy bin, including the number of showers per bin, is inset at the bottom of the figure. A. ABDUL HALIM et al. PHYS. REV. D 109, 022002 (2024) 022002-14 final bias-free selection of 594 showers (see Sec. IV). We find that the median resolution in Xmax shows a clear relation with shower energy, reaching a resolution of better than 15 g cm−2in the highest energy bin. We parametrize the resolution as δXmax ¼a·!!!!!!!!!!!!!!!! 1018 eV E r⊕b·1018 eV E⊕c; ð8Þ inspired by the energy resolution of electromagnetic calorimeters [47], but also functioning as a generic expansion in terms of energy. Here, a¼14.0$6.8g cm−2, b¼12.7$2.5g cm−2, and c¼11.2$4.7g cm−2are fitted free parameters, and ⊕indicates the quadratic sum. The constant term cprovides an indication of the resolution that might potentially be obtained for AERA with this method at the highest energies (given this parametrization). The change in resolution of Xmax is dominated by the uncertainty on the measured radio signals and hence becomes less accurate at lower energy. Comparing our resolution to the resolution achieved by FD, we achieve similar values at the highest energies where the FD reaches 15 g cm−2[40]. Furthermore, our method remains competitive down to lower energies where, for example, at E¼1017.8eV, the FD achieves the same resolution of 25 g cm−2. The most recent results by the LOFAR radio array, where a similar simulation-fitting method is used to determine Xmax, report an average resolution of 19 g cm−2between 1016.8and 1018.3eV [48]. Despite the much denser antenna spacing of LOFAR, AERA achieves similar resolutions considering the respective energy regimes to which the two experiments are sensitive. We note that up till now it has been common to quote a single resolution value for Xmax reconstruction methods for radio experiments, mainly because of a limited number of measured showers being available. Here, we show the resolution in Xmax depends strongly on the shower energy, driven primarily by the strength of radio signals measured in the antennas. Hence, the resolution is a function of detector sensitivity and shower energy and thus heavily depends on the shower selection criteria. As such, any direct comparison of methods is less straightforward if obtained at sufficiently dissimilar detectors. VI. Xmax MOMENTS AND THE DISTRIBUTION OF Xmax From the Xmax distribution, for each of our six energy bins, we now also calculate the first two moments of the distribution, the mean hXmaxiand the width σðXmaxÞ.To obtain the latter, we first subtract in quadrature the width caused by the method uncertainty, such that only the width caused by shower-to-shower fluctuations σðXmaxÞremains. The method uncertainty cannot simply be characterized by a single value since the uncertainties on Xmax for our air showers do not necessarily follow a Gaussian distribution. A bootstrap resampling procedure is applied for this reason, and with this, we then also calculate the uncertainty on σðXmaxÞ. This procedure is further described in Appendix C. The resulting mean and width of the true Xmax distribution are shown in Fig. 15, where we also compare this to the results from the FD (gray) and theoretical predictions of three different hadronic interaction models for a mass composition of just protons (red) or just iron nuclei (blue) [31,49–51]. The systematic uncertainties determined in the previous section are shown with capped markers. Table IV in Appendix Dlists the values for the two moments of the distributions for the six energy bins, together with their statistical and systematic uncertainties. With these AERA results, we show good agreement between the Auger radio and fluorescence measurements of hXmaxi, both pointing toward a (mixed)-light composition of cosmic rays at around E¼1017.5eV. Note that the two measurements share the systematic uncertainty on the energy scale, which is constructed from calibration of the SD energy to the FD energy scale [44]. Taking this contribution out reduces both systematic uncertainty bands by about 3g cm−2. Second, the determination of the systematic uncertainties due to the reconstruction method for AERA Xmax data depends on the assumed composition. FIG. 14. Resolution of the Xmax reconstruction method, δXmax , as a function of energy in units of column density. Shown per energy bin are the median values of the uncertainties on Xmax (circles with uncertainties σbfrom bootstrap resampling) for all showers in the bias-free sample (Table I) and a parametrized fit [Eq. (8)] of the resolution in Xmax (solid line with 1σ-confidence bands). Also shown are the resolutions achieved by the Auger fluorescence telescopes [40]. The black hatched region at low energy indicates the cut on energy applied earlier. The extent of each energy bin, including the number of showers per bin, is inset at the bottom of the figure. A. ABDUL HALIM et al. PHYS. REV. D 109, 022002 (2024) 022002-14 resolution 10 Measuring mass/particle type A. Abdul Halim et al., Phys. Rev. D, 109 (2024) 022002 A. Abdul Halim et al., Phys. Rev. Lett., 132 (2024) 021001 ⊕indicates the quadratic sum. The cparameter provides a prediction of the potential resolution that our method might be able to reach for AERA data. For radio experiments with a denser antenna spacing or experiments with lower ambient noise conditions one might reasonably expect this resolution to be even better. For example, LOFAR reported an average resolution of 19 g cm−2using a similar method [31] and simulation studies for the upcoming Square Kilometer Array suggest an average resolution of 6–8g cm−2could be reached [32]. In all likelihood their respective resolutions will improve with energy similar to the trend shown for AERA, making the radio technique very competitive for precision Xmax measurements. Comparison to other experiments.—In Fig. 4we show our hXmaxiresults together with various results from previous works. Measurements by other experiments that use the radio technique to measure Xmax are highlighted in color. In the past, Tunka-Rex [33], Yakutsk-Radio [34], and LOFAR [31] (and its prototype LOPES [35]) have shown Xmax measurements, but it has been challenging to make significant statements on the compatibility of the radio technique with fluorescence and air-Cherenkov light measurements. This is because these experiments either didn’t have a second technique to directly compare to or due to a combination of large statistical uncertainties and limited investigation of detector-specific systematic uncertainties. It is difficult to make statements on the compatibility of AERA and Tunka-Rex or Yakutsk-Radio without a full picture of those systematic uncertainties, but there do not seem to be significant discrepancies within their statistical uncertainties (note that the highest energy bin of TunkaRex only contains ten showers, hence its deviation with AERA is arguably not significant). However, the LOFAR measurements include a detailed estimation of systematic uncertainties, have much smaller statistical uncertainties than Tunka-Rex or Yakutsk-Radio, and share many similarities with AERA in the method to reconstruct Xmax, so we can compare these results to the FD and AERA results. We note that the difference between the Auger FD and LOFAR measurements, as can be seen in Fig. 4, previously left open the possibility of a systematic shift in Xmax due to an inherent difference between radio and fluorescence techniques. However, the AERA Xmax results now show no significant bias with respect to the fluorescence results, not when comparing their full datasets nor on an event-toevent level. Additionally, a study of the compatibility of the full shape of the Xmax distribution as measured by AERA and the Auger FD, available in [26], also finds no significant discrepancies within uncertainties. This strongly suggests that the differences between Auger and LOFAR must be either physical (e.g., due to differences in the magnetic field or atmospheric conditions, their altitudes, or their southern versus northern exposure) or systematic (e.g., due to the event selection or reconstruction), but not inherent to either the radio or fluorescence techniques. At higher energies, a seemingly similar difference in hXmaxi(both in magnitude and direction) can be observed between the fluorescence results of Auger (gray squares) and TA (gray plus markers). However, a detailed comparison by an Auger-TA working group has found that, given the known selection bias in the TA data, this difference is compatible within uncertainties [36]. This comparison only covers energies above 1018.2eV, so does not overlap with the LOFAR data. An AERA-LOFAR working group has started looking into their apparent differences, investigating, for example, differences in event selection, Xmax reconstruction method, and energy scale. Regardless, a deeper comparison of AERA and LOFAR data opens a new way to try to understand and reduce systematic uncertainties on air-shower and cosmic-ray parameters. Furthermore, the combination of fluorescence and radio measurements, linked by hybrid detectors such as at Auger, might resolve or constrain differences even more. Conclusions.—In this work, we have used seven years of AERA data to investigate the depth of maximum of extensive air showers at energies where the cosmic-ray origin is expected to transition from Galactic to extragalactic sources. We show our Xmax results to be in agreement with the results of the fluorescence telescopes at the Pierre FIG. 4. Mean of the Xmax distribution as measured by AERA in this work (black). The results are compared to predictions from air-shower simulations for multiple hadronic interaction models (lines) for proton (red) and iron (blue) mass compositions [15–18] and compared to measurements by LOFAR [31], TunkaRex [33], Yakutsk-Radio [34], and Auger FD [18]. Note that the Yakutsk-Radio results do not account for aperture effects on the same level as the other experiments. Colors have been used to highlight the measurements with the radio technique. The statistical uncertainties on the measurements are shown as vertical bars and for radio the systematic uncertainties, if available, are shown with caps. PHYSICAL REVIEW LETTERS 132, 021001 (2024) 021001-7 thus be used to probe the cosmic-ray mass composition. A heavier primary nucleus (which acts roughly as a superposition of multiple lower-energy nucleons) will, on average, interact higher up in the atmosphere and hence produce a wider footprint on the ground than a lighter primary particle. This is illustrated in Fig. 2. The particle type itself is not a direct observable, but the atmospheric depth where the shower is maximally developed, the depth of the shower maximum Xmax, which depends on the particle type, can be related to the shower footprint shape, making Xmax a probe for the primary particle type. Several methods have been used over the past years to reconstruct the particle type from radio signals, most of those relying on determining either the slope, width, or full shape of the lateral distribution function (LDF) of the radio footprint to determine Xmax [11–15]. In addition, also other methods using for example the slope of the frequency spectrum [16,17] and shape of the shower wavefront have been attempted [18]. Out of all these methods, the highest resolution in Xmax has been thus far achieved by using the LDF of the radio footprint by fitting of simulated air showers to measured air showers [19,20]. In this work, the simulation-fitting method has been further developed for the Auger Engineering Radio Array (AERA), by accounting for the effects of the sparse (compared to other radio experiments) and irregular array of radio stations. Also, a thorough investigation of systematic uncertainties has been made. We present the details of the Xmax reconstruction method and quantify the resolution as a function of cosmic-ray energy. Next, we apply the method to the set of air showers measured by AERA to determine the distributions of Xmax and interpret this in terms of the cosmic-ray mass composition. We then compare the composition to the results of the fluorescence detector (FD) at the Pierre Auger Observatory. Furthermore, we use a subset of air showers, simultaneously measured and independently reconstructed with both AERA and FD to directly evaluate our method and place bounds on the total systematic uncertainty between the two Xmax detection techniques. This paper will start with a description of the AERA airshower reconstruction and the selection of showers in Sec. II. Then, the Xmax reconstruction method will be described in Sec. III. In Sec. IV, we make an inventory of systematic uncertainties on the reconstruction of the Xmax distribution of the selected air showers. The resolution with which Xmax is reconstructed is then shown in Sec. V. Finally, the resulting Xmax distribution as measured by AERA will be presented in Sec. VI. In an accompanying publication [21], these results are discussed in the context of the larger field of other measurement techniques and experiments. II. AERA DATA RECONSTRUCTION The Pierre Auger Observatory [22] is located near the town of Malargüe in Argentina. It aims at detecting UHECRs up to the highest energies. The observatory covers an area of 3000 km2, making it the largest of its FIG. 1. Example of a footprint of the radio emission on the ground for a simulated air shower with an energy of 8.2×1017 eV, a zenith angle of 50.2°, and a depth of the shower maximum of 749 g cm−2. The strength of the emission is evaluated at simulated antenna positions (markers) and interpolated in between for visibility (background). The footprint has been projected into the shower plane, i.e., tilted into the plane perpendicular to the shower axis  vand rotated to project the magnetic field  Balong the x axis. FIG. 2. Schematic view of three air showers that started at different heights in the atmosphere and their radio emission footprints on the ground. It illustrates that the depth of the shower maximum affects the radio emission footprint in both width and general shape. The asymmetry is a consequence of how the geomagnetic and charge excess radio emission mechanisms interfere during the shower development. This figure has been previously published in [11]. RADIO MEASUREMENTS OF THE DEPTH OF AIR-SHOWER …PHYS. REV. D 109, 022002 (2024) 022002-5 Auger Engineering Radio Array AERA Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 PoS(ICRC2017)944 Earth Skimming Tau-neutrino in Antarcicatle 3 difficulty in other radio experiments. Since we can only detect the tau out of the three flavors, it one can point out that there is a disadvantage that the flux is reduced to one-third. However, when the flux of the tau neutrinos delivered from this experiment is studied with a total flux obtained from other experiments, a precise measurement of the neutrino flavor ratios will be a sharp probe to understand on many excellent physics topics such as the neutrino source, neutrino oscillation, neutrino decay, and the mass hierarchy [10]. An accurate pointing resolution of the neutrino arrival direction is also a major advantage in this experiment. The Askaryan signal in the cold ice forms a wide Cherenkov cone of about 56 degrees which introduces an inevitable ambiguity in determining the direction of the detected neutrinos. On the other hand, since the radio signal from the tau shower is shaped like a pencil beam (the Cherenkov angle in air is less than 1.5 degrees), the pointing resolution in our experiment is intrinsically excellent, which is crucially helpful for point source search. In fact, the detection of the mountain skimming tau neutrino has been proposed for a long time, but mainly in experiments based on the optical signal detection [11]; For instance, the lowest limit obtained by PAO in the energy band of the GZK neutrinos was also based on this method [12]. On the other hand, we use the radio technique which can realize a cost effective experiment instead of using high-cost optical equipment. The radio method has another great advantage of being able to achieve almost 100% of the lowest duty cycle, which is the most vulnerable to the optical based experiments. Our ground based approach can provide a long exposure which is limited in balloon based experiments due to limited flight time. Furthermore, by placing the antennas close to the showers in front of the mountain, we can reduce the 1/r 2 propagation loss, and therefore can lower the energy threshold than the balloon experiment, and utilize it in the GZK energy band. Figure 1. Concept of the tau neutrino detection via radio. 2. Tau-neutrino detection in Antarctica ANITA’s observation of EAS has brought valuable lessons. First, while other radio-based EAS experiments use low frequencies below 100 MHz, ANITA has proved the strong emissions even at high frequencies above 200 MHz where the Galactic noise is relatively low. This has influenced our experimental design which uses a relatively small antennas designed for the high frequencies. It provides a practical way to create a large scale antenna array by simplifying the antenna support structure and making installation easier. ANITA also proved that Antarctica is the excellent place for the experiment, where the ambient anthropogenic noise is the quietest and has the strongest geo-magnetic field required for the geo-synchrotron radiation. Noting that 11 Radio neutrino detection Two main principles neutrino interacts —> shower •Detection in thin medium (air) J. Nam, POS (ICRC2017) 944 Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 PoS(ICRC2017)944 Earth Skimming Tau-neutrino in Antarcicatle 3 difficulty in other radio experiments. Since we can only detect the tau out of the three flavors, it one can point out that there is a disadvantage that the flux is reduced to one-third. However, when the flux of the tau neutrinos delivered from this experiment is studied with a total flux obtained from other experiments, a precise measurement of the neutrino flavor ratios will be a sharp probe to understand on many excellent physics topics such as the neutrino source, neutrino oscillation, neutrino decay, and the mass hierarchy [10]. An accurate pointing resolution of the neutrino arrival direction is also a major advantage in this experiment. The Askaryan signal in the cold ice forms a wide Cherenkov cone of about 56 degrees which introduces an inevitable ambiguity in determining the direction of the detected neutrinos. On the other hand, since the radio signal from the tau shower is shaped like a pencil beam (the Cherenkov angle in air is less than 1.5 degrees), the pointing resolution in our experiment is intrinsically excellent, which is crucially helpful for point source search. In fact, the detection of the mountain skimming tau neutrino has been proposed for a long time, but mainly in experiments based on the optical signal detection [11]; For instance, the lowest limit obtained by PAO in the energy band of the GZK neutrinos was also based on this method [12]. On the other hand, we use the radio technique which can realize a cost effective experiment instead of using high-cost optical equipment. The radio method has another great advantage of being able to achieve almost 100% of the lowest duty cycle, which is the most vulnerable to the optical based experiments. Our ground based approach can provide a long exposure which is limited in balloon based experiments due to limited flight time. Furthermore, by placing the antennas close to the showers in front of the mountain, we can reduce the 1/r 2 propagation loss, and therefore can lower the energy threshold than the balloon experiment, and utilize it in the GZK energy band. Figure 1. Concept of the tau neutrino detection via radio. 2. Tau-neutrino detection in Antarctica ANITA’s observation of EAS has brought valuable lessons. First, while other radio-based EAS experiments use low frequencies below 100 MHz, ANITA has proved the strong emissions even at high frequencies above 200 MHz where the Galactic noise is relatively low. This has influenced our experimental design which uses a relatively small antennas designed for the high frequencies. It provides a practical way to create a large scale antenna array by simplifying the antenna support structure and making installation easier. ANITA also proved that Antarctica is the excellent place for the experiment, where the ambient anthropogenic noise is the quietest and has the strongest geo-magnetic field required for the geo-synchrotron radiation. Noting that 11 Radio neutrino detection Two main principles neutrino interacts —> shower •Detection in thin medium (air) J. Nam, POS (ICRC2017) 944 •Detection in dense medium (ice) firn ice ~ Xint <latexit sha1_base64="4+x9NZsI47i6AQRd0Ta8tbFYMFA=">AAAB/XicbVDLSsNAFJ34rPUVHzs3g0VwVZIq6LLoxmUF+4AmhMl00g6dmYSZSaGG4K+4caGIW//DnX/jpM1CWw8MHM65lzn3hAmjSjvOt7Wyura+sVnZqm7v7O7t2weHHRWnEpM2jlkseyFShFFB2ppqRnqJJIiHjHTD8W3hdydEKhqLBz1NiM/RUNCIYqSNFNjH3oTgrJcHHkd6JHlGhc4Du+bUnRngMnFLUgMlWoH95Q1inHIiNGZIqb7rJNrPkNQUM5JXvVSRBOExGpK+oQJxovxslj6HZ0YZwCiW5gkNZ+rvjQxxpaY8NJNFRrXoFeJ/Xj/V0bVvDkpSTQSefxSlDOoYFlXAAZUEazY1BGFJTVaIR0girE1hVVOCu3jyMuk06u5FvXF/WWvelHVUwAk4BefABVegCe5AC7QBBo/gGbyCN+vJerHerY/56IpV7hyBP7A+fwCCn5Xo</latexit> ~v⌫ <latexit sha1_base64="gf7R24uxlFwC2mw9oShPRWeaF/k=">AAAB8nicbVBNS8NAEJ3Ur1q/qh69LBbBU0mqoMeiF48V7AckoWy2m3bpZjfsbgol9Gd48aCIV3+NN/+N2zYHbX0w8Hhvhpl5UcqZNq777ZQ2Nre2d8q7lb39g8Oj6vFJR8tMEdomkkvVi7CmnAnaNsxw2ksVxUnEaTca38/97oQqzaR4MtOUhgkeChYzgo2V/GBCST6Z9QOR9as1t+4ugNaJV5AaFGj1q1/BQJIsocIQjrX2PTc1YY6VYYTTWSXINE0xGeMh9S0VOKE6zBcnz9CFVQYolsqWMGih/p7IcaL1NIlsZ4LNSK96c/E/z89MfBvmTKSZoYIsF8UZR0ai+f9owBQlhk8twUQxeysiI6wwMTalig3BW315nXQade+q3ni8rjXvijjKcAbncAke3EATHqAFbSAg4Rle4c0xzovz7nwsW0tOMXMKf+B8/gDF+5GS</latexit> bedrock 200m 2.7km 10m 1-2 km 1-2 km Figure 1: Sketch of Askaryan radiation in ice. The dimensions are shown for a typical case at the South Pole. obviously preferred compared for example to air which is a thousand times thinner and would require a thousand times larger volume to yield the same number of neutrino interactions. When a high-energy neutrino interacts in the ice it creates a cascade of millions of secondary particles which is often referred to as particle shower. These showers consist almost exclusively of photons, electrons and positrons. Comparing the neutrino energy of interest of around 1018 eV with the rest mass energy of an electron or positron of 511 keV reveals that about a trillion of electrons-positron pairs could be created if all kinetic energy of the neutrino is transferred into the creation of new particles, i.e., the particle shower. 1.1.2 Radio emission from in-ice showers When a particle shower develops in a dielectric medium, such as ice, it develops a charge asymmetry. As the cascade develops, the number of electrons increasingly exceeds the number of positrons. This time-varying negative charge excess generates coherent radiation in the radio frequency range. The radiation is typically referred to as charge-excess radiation or Askaryan radiation as it was postulated by Gurgen Askaryan in 1962 [4]. The Askaryan radiation exhibits a characteristic feature of being emitted on a cone around the shower direction as illustrated in Fig. 1. The shower propagates with the vacuum speedof-light c0whereas the radio signal propagates at a slower speed-of-light of c0/n where the index-of-refraction of deep ice is n=1.78. This leads to constructive interference of all emission 3 firn ice ~ Xint <latexit sha1_base64="4+x9NZsI47i6AQRd0Ta8tbFYMFA=">AAAB/XicbVDLSsNAFJ34rPUVHzs3g0VwVZIq6LLoxmUF+4AmhMl00g6dmYSZSaGG4K+4caGIW//DnX/jpM1CWw8MHM65lzn3hAmjSjvOt7Wyura+sVnZqm7v7O7t2weHHRWnEpM2jlkseyFShFFB2ppqRnqJJIiHjHTD8W3hdydEKhqLBz1NiM/RUNCIYqSNFNjH3oTgrJcHHkd6JHlGhc4Du+bUnRngMnFLUgMlWoH95Q1inHIiNGZIqb7rJNrPkNQUM5JXvVSRBOExGpK+oQJxovxslj6HZ0YZwCiW5gkNZ+rvjQxxpaY8NJNFRrXoFeJ/Xj/V0bVvDkpSTQSefxSlDOoYFlXAAZUEazY1BGFJTVaIR0girE1hVVOCu3jyMuk06u5FvXF/WWvelHVUwAk4BefABVegCe5AC7QBBo/gGbyCN+vJerHerY/56IpV7hyBP7A+fwCCn5Xo</latexit> ~v⌫ <latexit sha1_base64="gf7R24uxlFwC2mw9oShPRWeaF/k=">AAAB8nicbVBNS8NAEJ3Ur1q/qh69LBbBU0mqoMeiF48V7AckoWy2m3bpZjfsbgol9Gd48aCIV3+NN/+N2zYHbX0w8Hhvhpl5UcqZNq777ZQ2Nre2d8q7lb39g8Oj6vFJR8tMEdomkkvVi7CmnAnaNsxw2ksVxUnEaTca38/97oQqzaR4MtOUhgkeChYzgo2V/GBCST6Z9QOR9as1t+4ugNaJV5AaFGj1q1/BQJIsocIQjrX2PTc1YY6VYYTTWSXINE0xGeMh9S0VOKE6zBcnz9CFVQYolsqWMGih/p7IcaL1NIlsZ4LNSK96c/E/z89MfBvmTKSZoYIsF8UZR0ai+f9owBQlhk8twUQxeysiI6wwMTalig3BW315nXQade+q3ni8rjXvijjKcAbncAke3EATHqAFbSAg4Rle4c0xzovz7nwsW0tOMXMKf+B8/gDF+5GS</latexit> bedrock 200m 2.7km 10m 1-2 km 1-2 km Figure 1: Sketch of Askaryan radiation in ice. The dimensions are shown for a typical case at the South Pole. obviously preferred compared for example to air which is a thousand times thinner and would require a thousand times larger volume to yield the same number of neutrino interactions. When a high-energy neutrino interacts in the ice it creates a cascade of millions of secondary particles which is often referred to as particle shower. These showers consist almost exclusively of photons, electrons and positrons. Comparing the neutrino energy of interest of around 1018 eV with the rest mass energy of an electron or positron of 511 keV reveals that about a trillion of electrons-positron pairs could be created if all kinetic energy of the neutrino is transferred into the creation of new particles, i.e., the particle shower. 1.1.2 Radio emission from in-ice showers When a particle shower develops in a dielectric medium, such as ice, it develops a charge asymmetry. As the cascade develops, the number of electrons increasingly exceeds the number of positrons. This time-varying negative charge excess generates coherent radiation in the radio frequency range. The radiation is typically referred to as charge-excess radiation or Askaryan radiation as it was postulated by Gurgen Askaryan in 1962 [4]. The Askaryan radiation exhibits a characteristic feature of being emitted on a cone around the shower direction as illustrated in Fig. 1. The shower propagates with the vacuum speedof-light c0whereas the radio signal propagates at a slower speed-of-light of c0/n where the index-of-refraction of deep ice is n=1.78. This leads to constructive interference of all emission 3 S. Barwick, C. Glaser; World Scientific Series in Astrophysics, The Encyclopedia of Cosmology, pp. 237-302 (2023) Askaryan effect n=1,78 —> Č ~56° Figure 11: Attenuation length vs. depth for three di↵erent sites of a radio neutrino detector at South Pole, Moore’s Bay on the Ross Ice Shelf and Greenland at Summit station as implemented in NuRadioMC [17]. The solid lines show the attenuation length at 200 MHz and the dashed lines show the attenuation length at 400 MHz. thus, the ice quality close to the antennas matters most, especially at low neutrino energies where the neutrino interaction happens closer to the antennas. The site with the shortest attenuation lengths is Moore’s Bay on the Ross Ice Shelf but because of the reflective properties of the bottom of the ice shelf and the resulting large sky coverage, the site provides important physics capabilities for a future detector. 5.2 Index-of-refraction profile All sites have in common that the density changes in the upper O(100 m). Due to the increasing gravitational pressure, the firn gets compressed into clear ice with increasing depth. To first order, the density at a given depth is directly dependent on the overburden of ice at that depth. Then, the density profile is continuous and follows an exponential [66]. An empirical depth–density relation has been given [67] as ⇢(z)=⇢ice (⇢ice ⇢surface)exp(Cz), where ⇢ice is the asymptotic density of polar ice (917 kg/m3) and ⇢surface is the density at the surface. The index-of-refraction nscales linearly with the density and is often parameterized as n(z)= 1+0.86 cm3/g⇥⇢(z) and thus follow the same shape as the density profile [66]. A compilation of measurements from di↵erent sites across Antarctica is shown in Fig. 12 that can be described well with an exponential fit to the data points which is also shown in the figure. As a consequence of the changing index-of-refraction, radio waves propagating through the firn will bend downwards due to continuous Fresnel refraction which will result in curved signal trajectories. A signal can reach the receiver often via a direct path, and another path where the trajectory is bent downwards and reaches the receiver from above (refracted trajectory). We show a few typical trajectories in Fig. 13. The firn-air interface also reflects radio waves back into the firn, and for most expected geometries of neutrino radio signals, the condition of totalinternal-reflection (TIR) is fulfilled which results in a lossless reflection of the signal leading to additional reflected signal trajectories. In general, between any two positions exist either zero or two signal trajectories, where the two signal trajectories are a combination of direct, refracted, and reflected trajectories. 18 attenuation length in ice Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 12 Observations of the Askaryan Effect in Ice P. W. Gorham, 1 S. W. Barwick, 2 J. J. Beatty, 3 D. Z. Besson, 4 W. R. Binns, 5 C. Chen, 6 P. Chen, 6 J. M. Clem, 7 A. Connolly, 8 P. F. Dowkontt, 5 M. A. DuVernois, 9 R. C. Field, 6 D. Goldstein, 2 A. Goodhue, 8 C. Hast, 6 C. L. Hebert, 1 S. Hoover, 8 M. H. Israel, 5 J. Kowalski, 1 J. G. Learned, 1 K. M. Liewer, 10 J. T. Link, 1,11 E. Lusczek, 9 S. Matsuno, 1 B. Mercurio, 3 C. Miki, 1 P. Mioc ˇinovic ´, 1 J. Nam, 2 C. J. Naudet, 10 J. Ng, 6 R. Nichol, 3 K. Palladino, 3 K. Reil, 6 A. Romero-Wolf, 1 M. Rosen, 1 L. Ruckman, 1 D. Saltzberg, 8 D. Seckel, 7 G. S. Varner, 1 D. Walz, 6 and F. Wu 2 (ANITA Collaboration) 1 Department of Physics and Astronomy, University of Hawaii, Manoa, Hawaii 96822, USA 2 University of California, Irvine California 92697, USA 3 Department of Physics, Ohio State University, Columbus, Ohio 43210, USA 4 Department of Physics and Astronomy, University of Kansas, Lawrence, Kansas 66045, USA 5 Department of Physics, Washington University in St. Louis, Missouri 63130, USA 6 Stanford Linear Accelerator Center, Menlo Park, California, 94025, USA 7 University of Delaware, Newark, Delaware 19716, USA 8 Department of Physics and Astronomy, University of California, Los Angeles, California 90095, USA 9 School of Physics and Astronomy, University of Minnesota, Minneapolis, Minnesota 55455, USA 10 Jet Propulsion Laboratory, Pasadena, California 91109, USA 11 Currently at NASA Goddard Space Flight Center, Greenbelt, Maryland, 20771, USA (Received 13 December 2006; revised manuscript received 15 June 2007; published 25 October 2007) We report on observations of coherent, impulsive radio Cherenkov radiation from electromagnetic showers in solid ice. This is the first observation of the Askaryan effect in ice. As part of the complete validation process for the ANITA experiment, we performed an experiment at the Stanford Linear Accelerator Center in June 2006 using a 7.5 metric ton ice target. We measure for the first time the largescale angular dependence of the radiation pattern, a major factor in determining the solid-angle acceptance of ultrahigh-energy neutrino detectors. DOI: 10.1103/PhysRevLett.99.171101 PACS numbers: 95.55.Vj, 29.40.Ka, 98.70.Sa Very large-scale detectors, such as the Antarctic Muon and Neutrino Detector Array and its successor IceCube, have demonstrated the excellent utility of Cherenkov radiation in detection of neutrino interactions at >TeV energies [1,2] with ice as a target medium. However, at neutrino energies above 100 PeV, the cubic-kilometer scale of such detectors is inadequate to detect more than a handful of events from the predicted cosmogenic neutrino fluxes [3] which represent the most compelling models at these energies. The relevant detector volume for convincing detection and characterization of these neutrinos is in the range of hundreds to thousands of km3 of ice, and the economic constraints of scaling up the optical Cherenkov technique almost certainly preclude extending it much beyond the size of the current IceCube detector, which will be completed early in the next decade. Given the need for an alternative technique with a more tractable economy of scale to reach into the EeV (!1000 PeV) energy regime, a new method has emerged within the last decade. This method, the radio Cherenkov technique, relies on properties of electromagnetic cascades in a dielectric medium. It was first hypothesized by Askaryan [4] and confirmed in 2001 at SLAC [5]. High energy processes such as Compton, Bhabha, and Møller scattering, and positron annihilation rapidly lead to a "20% negative charge asymmetry in the electronphoton part of a cascade. In dense media the shower charge bunch is largely contained within a several cm radius. At wavelengths of #10 cm,muchlargerthan the characteristic shower bunch size, the relativistic shower bunch appears as a single charge moving through the dielectric over a distance of several meters or more. As an example, a typical shower with mean Bjorken inelasticity hyi!0:2, initiated by a E!! 100 PeV neutrino creates a total number of charged particles at shower maximum of order ne$$ne%! hyiE!=1 GeV "2&107.Thenetchargeisthusne$% ne%"4&106e. Since the radiated power for Cherenkov emission grows quadratically with the charge of the emitter, the coherent power in the cm-to-m wavelength regime is "1013 times greater than the single-charge emission, far exceeding any other secondary emission in optical or longer-wave bands, and dominating other coherent radio emission processes in solids, such as transition radiation [6] and synchrotron radiation in the Earth’s magnetic field [7]. PRL 99, 171101 (2007) PHYSICAL REVIEW LETTERS week ending 26 OCTOBER 2007 0031-9007=07=99(17)=171101(5) 171101-1 2007 The American Physical Society photons via a bremsstrahlung radiator. Such methods were used in earlier Askaryan discovery experiments to avoid any initial excess charge in the shower development. In our case, the typical shower had a total composite energy of 3!1019 eV, with a total of "2!1010e#e$pairs at shower maximum. Simulations of the charge excess development indicate a net charge asymmetry of about 20%. Thus the initial electrons contribute at most "15% of the total negative charge excess in the shower, and we have corrected for this bias in the results we show here. In addition, radio-absorbing foam was in place on the front face of the ice, and very effectively suppressed rf noise from the upstream metal beam vacuum windows and air gaps. Our previous experiments [5,16] measured accelerator rf backgrounds at SLAC, and they are a negligible contribution to our measurements. A schematic and perspective view of the experiment layout is shown in Fig. 1. The ice was contained in a 10 cm thick insulating foamlined box, and a 10 cm foam lid was used during operation, along with a freezer unit, to maintain temperatures of between $5to $20 %C. Such temperatures are adequate to avoid significant rf absorption over the several m pathlengths of the radiation through the ice [19]. The ANITA payload (Fig. 2, consisting of an array of 32 dual-polarization quad-ridged horns along with eight omnidirectional conical antennas, was used to receive the emission at a location about 15 m away from the center of the target, as shown in Fig. 1.Theantennafrequency range is from 200–1200 MHz, a frequency range over which the rf transmissivity of ice is at its highest [19]. ANITA horn antennas are arranged so that adjacent antennas in both the lower and upper payload sections respond well even to a signal directed along their nearest neighbors’ bore sights. This allows multiple antennas (typically 4 to 6 horns and 3 to 4 of the conical antennas) to sample the arriving wave front. The signals are digitized by custom compact-PCI-based 8-channel digitizer modules [20], 9 of which are used to record all 72 antenna signals simultaneously at 2:6 G samples=sec . Figure 3(left) presents results for the absolute field strength of the radio impulse in the several different antennas employed here, with uncertainties of about &40% in field strength (&3 dB), dominated by systematic rather than statistical errors, arising from a combination of the 1–2 dB uncertainty in the gain calibration of the antennas, and by comparable uncertainties in removing secondary reflections from the measured impulse power. The field strengths are compared to a parametrization based on shower #electrodynamics simulations for ice [9], and the agreement is well within our experimental errors. Figure 3(right) shows results of the scaling of the pulse power with shower energy. The dependence is completely consistent with quadratic scaling, indicating that the radiation is coherent over the 200–1200 MHz frequency window. Figure 4shows the measured and predicted angular dependence of the radiation, for data up to 800 MHz; beyond this the SNR of the data was inadequate for such sub-band analysis. The Cherenkov cone refracts into the forward direction out of the ice, and is clearly delineated by the data. Here we show statistical #systematic errors within a measurement run; the overall normalization (with separate systematic error) is taken from Fig. 3. We scale these data within the overall systematic errors to match the peak of the field strength. For T486, L"1:2mfor Eq. (1) above. The measured angular dependence follows closely the expectations for Cherenkov radiation, including the narrowing of the Cherenkov cone with higher frequencies. This behavior arises from the kL term in the exponential of Eq. (1), and is important since the width of the Cherenkov cone determines the detection solid angle, or acceptance of ahighenergyneutrinodetector.Thismeasurementconfirms the predicted behavior for the first time. FIG. 1 (color). Top: Side view schematic of the target and receiver arrangement in ESA. Bottom: Perspective view of the setup, showing the key elements. PRL 99, 171101 (2007) PHYSICAL REVIEW LETTERS week ending 26 OCTOBER 2007 171101-3 In summary, Askaryan’s hypothesis has now been confirmed in detail by laboratory experiments for virtually all of the dielectrics (ice, salt, sand—the latter approximating the Lunar regolith) that Askaryan envisioned as the best media in which to exploit the coherent radio Cherenkov emission from high energy particle showers. Askaryan’s intent was to illuminate a methodology by which low fluxes of ultrahigh-energy particles could be made observable through exploitation of huge volumes of natural materials. With the recent sharpening of predictions for the fluxes of ultrahigh-energy neutrinos, and the growth in the number of experiments that make use of it, we expect that Askaryan’s hope will be soon fulfilled. FIG. 4 (color). Top: Angular dependence of the radiation for both the in-ice and refracted case, for a frequency range from 200–800 MHz, compared to data. The data errors are combined statistical and systematic, but with an arbitrary overall normalization-see Fig. 3for the normalization factor. The inice and refracted curves are the theoretical expectation for a shower in ice at a beam current of 109e!per bunch and 28.5 GeV electrons, and the refraction includes only geometric optics. Bottom: Same as top for three different subfrequency bands. FIG. 3 (color). Left: Field strength vs frequency of radio Cherenkov radiation in the T486 experiment, for several different antennas used, including a theoretical curve [9]. Right: Pulse power vs total shower energy (number of particles " mean energy=particle), curve is for completely coherent radio Cherenkov emission. FIG. 2 (color). Left: The ANITA payload (center) above and downstream of the ice target (here covered). Right top, target with cover removed, in ambient light. Right bottom: ice target illuminated from interior scattered optical Cherenkov radiation. PRL 99, 171101 (2007) PHYSICAL REVIEW LETTERS week ending 26 OCTOBER 2007 171101-4 In summary, Askaryan’s hypothesis has now been confirmed in detail by laboratory experiments for virtually all of the dielectrics (ice, salt, sand—the latter approximating the Lunar regolith) that Askaryan envisioned as the best media in which to exploit the coherent radio Cherenkov emission from high energy particle showers. Askaryan’s intent was to illuminate a methodology by which low fluxes of ultrahigh-energy particles could be made observable through exploitation of huge volumes of natural materials. With the recent sharpening of predictions for the fluxes of ultrahigh-energy neutrinos, and the growth in the number of experiments that make use of it, we expect that Askaryan’s hope will be soon fulfilled. FIG. 4 (color). Top: Angular dependence of the radiation for both the in-ice and refracted case, for a frequency range from 200–800 MHz, compared to data. The data errors are combined statistical and systematic, but with an arbitrary overall normalization-see Fig. 3for the normalization factor. The inice and refracted curves are the theoretical expectation for a shower in ice at a beam current of 109e!per bunch and 28.5 GeV electrons, and the refraction includes only geometric optics. Bottom: Same as top for three different subfrequency bands. FIG. 3 (color). Left: Field strength vs frequency of radio Cherenkov radiation in the T486 experiment, for several different antennas used, including a theoretical curve [9]. Right: Pulse power vs total shower energy (number of particles " mean energy=particle), curve is for completely coherent radio Cherenkov emission. FIG. 2 (color). Left: The ANITA payload (center) above and downstream of the ice target (here covered). Right top, target with cover removed, in ambient light. Right bottom: ice target illuminated from interior scattered optical Cherenkov radiation. PRL 99, 171101 (2007) PHYSICAL REVIEW LETTERS week ending 26 OCTOBER 2007 171101-4 Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 13 The Radio Neutrino Observatory in Greenland (RNO-G) A single RNO-G Station Antennas types: Log-periodic dipole antenna (LPDA), Horizontally-polarized (Hpol), and Vertically-polarized (Vpol) 5 A single RNO-G Station Antennas types: Log-periodic dipole antenna (LPDA), Horizontally-polarized (Hpol), and Vertically-polarized (Vpol) 5 RNO-G: a new experimental effort • Deployed near Summit Station, Greenland • Hardware is fully funded to reach 35(+) stations: already the largest in-ice radio neutrino detector by area! • Currently in building phase: holes for 7 more stations are being drilled this year, with DAQs installed next year • Science team is also working on calibration, simulation, instrument performancesee following talks from RNO-G team members! Stations are all solar powered and send data over LTE 4 ARENA conference, Chicago (2024) Modeling our antennas Credit: Mohammad F. H. Seikh 9 LPDAs The RNO-G Collaboration 2 Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 14 The Radio Neutrino Observatory in Greenland (RNO-G) detection of solar flares Calibrated Antenna Locations (ω→10–15 cm) Figure: Enter Caption https://arxiv.org/pdf/2404.14995 Solar Flares 15 Figure: Enter Caption https://arxiv.org/pdf/2404.14995 Solar Flares 11 ARENA conference, Chicago (2024) Close-up image of CME along B-lines https://arxiv.org/pdf/2404.14995 Solar Flares 4 Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 Marco Muzio (Penn State) Askaryan Radiation •Neutrino interaction in dense medium initiates particle cascade" •Particle cascade emits pulse of coherent radio emission along Cherenkov cone — Askaryan radiation" •Radio has ~1 km attenuation length in ice" •Radio antenna embedded in ice = efficient monitor of enormous volume 3 vertex Askaryan Radiation ν forward view E-field polarization direction side view Marco Muzio (Penn State) ARA Detector Overview •5 independent stations on hexagonal grid at South Pole" •Each station has 4 strings embedded in ice" •Each string has 4 radio antennas (2 VPols & 2 HPols) at ~200 m depth" •Trigger condition:" •3 like-polarization antennas with integrated power 5x ambient noise within 170 ns coincidence" •~6 Hz trigger (+1 Hz software trigger) 4 2 km IceCube 31 2 South Pole Station South Pole 5 4 WT3 Skiway Deployed 2012 Deployed 2013 Deployed 2018 Cable Calibration antennas Calibration antennas antenna clusters Central station electronics Power and communications to ICL Downhole instrumentation FO transmitter Top Hpol Bottom Vpol Bottom Hpol Top Vpol Depth: 180 m surface antennas not shown 15 Askaryan Radio Array ARENA conference, Chicago (2024) Marco Muzio (Penn State) ARA Detector Overview •5 independent stations on hexagonal grid at South Pole" •Each station has 4 strings embedded in ice" •Each string has 4 radio antennas (2 VPols & 2 HPols) at ~200 m depth" •Trigger condition:" •3 like-polarization antennas with integrated power 5x ambient noise within 170 ns coincidence" •~6 Hz trigger (+1 Hz software trigger) 4 2 km IceCube 31 2 South Pole Station South Pole 5 4 WT3 Skiway Deployed 2012 Deployed 2013 Deployed 2018 Cable Calibration antennas Calibration antennas antenna clusters Central station electronics Power and communications to ICL Downhole instrumentation FO transmitter Top Hpol Bottom Vpol Bottom Hpol Top Vpol Depth: 180 m surface antennas not shown Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 16 Askaryan Radio Array ARENA conference, Chicago (2024) Marco Muzio (Penn State) Summary •ARA has accumulated ~24 station-years of livetime through 2021" •Conducting first-ever array-wide neutrino search in deep stations •Highly-coordinated, multi-institution analysis" •Improved analysis methods & detector characterization" •Proof of concept for next-generation large in-ice radio arrays" •e.g. RNO-G (35 stations) # & IceCube-Gen2 Radio (361 stations) " •Will yield either: •First UHE neutrino candidates •Strongest flux limit up to 100 EeV from any in-ice radio experiment 16 working on analysis with 5 stations Marco Muzio (Penn State) ν 5 <1° resolution on vertex reconstruction J. Torres (2021) - Neutrino Astrophysics with the Askaryan Radio Array •Cross-correlating signal in each antenna allows for interaction vertex reconstruction" •Vertex reconstruction allows for background CR and anthropogenic signals to be discarded Vertex Reconstruction Δt *Calibration pulser for illustration See talk from Alan Salcedo Gomez this session Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 Arrays of IceCube-Gen2 Radio Array targets the Ultrahigh-Energy Regime • 169 hybrid stations, 1.75 km spacing on square • 192 shallow stations, interspersed 1.24 km spacing IceCube-Gen2 Collaboration TDR 2 Station Designs A Hybrid Approach Shallow Hybrid •Robust discovery-level instrument that combines shallow and deep antennas to mitigate against systematics with two approaches •Hybrid: 24 channels (17 deep cylindrical antennas up to 150 m maximum depth, 7 LPDAs surface) •Shallow: 7 LPDAs, one 10-m deep dipole 4 17 The Radio Array of IceCube-Gen2 ARENA conference, Chicago (2024) 500 km2 Science Case •Astrophysical neutrinos: •Resolve the high-energy neutrino sky from TeV to EeV •Multimessenger observations •Energy, spectrum, and fiavor •Cosmogenic neutrinos: origin of cosmic ray accelerators •Targeting discovery at fiux level where 10% of the UHECR are protons, with flve years of data •Fundamental physics at UHE energies •Expect 3 deg. angular resolution and 65% energy resolution (68% containment) Discovery-level array at Ultrahigh Energies IceCube-Gen2 Collaboration TDR 3 V. B . Valera, M. Bustamante and C. Glaser, JHEP 06 (2022) 105 I. Esteban, S. Prohira, J. Beacom, Phys. Rev. D 106, 023021 D. Fiorillo, V. B. Valera, M. Bustamante JCAP 03 (2023) 026 I. Plaisier, S. Bouma, A. Nelles, EPJ-C 83, 443 (2023) S. Bouma et al., PoS(ICRC2023)1045 A. Coleman et al. arXiv:2402.02432 IceCube-Gen2 TDR 2023/2024 M. Muzio, M. Unger, S. Wissel PRD 107, 103030 (2023) V. B . Valera, M. Bustamante and C. Glaser, Phys. Rev. D 107, 043019 (2023) V. B . Valera, M. Bustamante and C. Glaser, JHEP 06 (2022) 105 I. Esteban, S. Prohira, J. Beacom, Phys. Rev. D 106, 023021 A. Coleman et al. arXiv:2402.02432 A. Garcia-Soto PRD 107, 033009 (2023) Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 18 BEACON: Beamforming Elevated Array for Cosmic Neutrinos Tau Exit ⟨𝑨𝜴⟩ Tau Decay Radio Signal Elevation >2km Pointing Array Trigger Array Station L~𝓞(100 m) BEACON: Beamforming Elevated Array for COsmic Neutrinos •Concept: 𝓞(1000) independent radio interferometers on mountaintops, designed to detect the radio emission of upgoing air showers created by earth-skimming 𝜈𝜏 •Goal: measure the flux of 𝜈𝜏 at E > 100 PeV •Advantages: + radio = low cost, high duty cycle + high elevation = large detector volume + phased array trigger = greater sensitivity and directional rejection of background Extensive Air Shower Concept paper: S. Wissel et al. JCAP11(2020)065 5 ARENA conference, Chicago (2024) Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 25 Upgraded Surface Detector of Auger Observatory radio antenna 30-80 MHz two orthogonal polarizations 250 MHz sampling plastic scintillator 120 MHz sampling read-out electronics e/ µ µ e/ water-Cherenkov detector 120 MHz sampling atmosphere of Earth is transparent in 30-80 MHz band Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 26 A measured air shower Hörandel Part B2 COSMICISFONTIBUS [85]) in the field with a drone, carrying a reference antenna in a defined pattern above a RD station. RD WCD Azimuth (deg) 156.99±0.01 157±0.1 Zenith (deg) 84.7±0.01 84.7±0.1 Energy (EeV) 36.23 !± !3.34 38.55 !± !2.92 Core X (km) -19.8 -17.40±0.88 Core Y (km) -8.73 -9.78±0.45 Figure 5: An air shower measured with the RD. The diffuse Galactic radio emission is well measured [86], this is used as a standard reference signal to calibrate the RD in-situ [87]. The Galactic radio emission is recorded periodically on each station locally and is used to correct for potential timedependent changes in the parameters (such as e.g. potential gain changes of the electronics as a function of temperature). It should be noted that the absolute calibration of the full electronics chain (performed in the laboratory) agrees within 5% uncertainty with the parameters obtained from the Galactic emission. This demonstrates our excellent understanding of the complete signal chain. The atmospheric electric fields are continuously monitored at 5 positions in the SD array. This allows to generate a veto against strong atmospheric electric field during thunderstorms [88], which would distort the energy measurements of a shower. This yields absolutely calibrated time traces for each antenna and polarisation direction. A study with AERA demonstrates the long-term stability of the radio detection technique [89]. Only marginal deviations have been found over a period of 10 years. PoS(ICRC2025)294 First Data of the 3000 km2Radio Detector at the Pierre Auger Observatory Bjarni Pont Figure 7: (Left): Comparison of the cosmic ray energy as measured by the WCD and the electromagnetic energy as measured by the RD. Highlighted in color are the zenith angles of the air showers. Statistical uncertainties of reconstructed parameters are shown as gray bars. (Right): An example of a near-horizontal air shower with an energy of approximately 32 EeV arriving from west. The gray markers are SD positions and the green markers show where the RD was deployed at the time of the event. Stations that measured a significant radio signal are shown with star markers. The color indicates the arrival time of the pulse. The underlying red markers show the signals measured by the WCD, where the size is proportional to the signal. 3.2 Extremely extended radio footprints In Fig. 7(Right), we illustrate the detection of a near-horizontal radio footprint measured in early 2024. The air shower had a zenith angle of 85→and an electromagnetic energy of about 32 EeV, compatible with the estimate of the WCD. This event demonstrates the ability of the RD to measure in this near-horizontal regime, illustrating the potential to measure ultra-high-energy neutral particles. Of particular interest are showers coming from the west where the Andes mountain range provides the potential for earth-skimming neutrinos to interact and to produce an air shower (the event shown does not have a su!cient zenith angle to be considered a candidate). 4. Outlook The Radio Detector is expected to operate for at least a decade, providing a substantial increase in the number of cosmic rays with an estimation of the mass [14]. Together with the mass measurements by the WCD-SSD at low zenith angles, it will be able to cover most of the southern sky. This will allow for studies of the mass composition of ultra-high-energy cosmic rays. In addition, the measurements of the amount of muons in the shower as a function of energy by this WCD-RD hybrid approach will contribute to addressing the muon puzzle [16]. In parallel, the Radio Detector will extend the radio-based energy scale to energies beyond 1018.5eV, building on the results and method developed for the AERA radio detector at the Pierre Auger Observatory [13]. This extension will provide an independent way to access the cosmic-ray energy up to the highest observed energies. Furthermore, the use of both amplitude and phase information in the radio signal enables interferometric reconstruction techniques, allowing for a three-dimensional and time7 Figure 6: Measured e/m energy (RD) in EAS as a function of total CR energy (WCD) [90]. Highlighted in colour are the zenith angles of the EAS. The installation of the RD in the whole array has been completed end of 2024. At the time of writing, commissioning of the RD systems is almost completed. First RD data are being analysed in parallel. The first data look very promising [90]. For illustration, in Fig.5 a measured air shower is shown with an energy of almost 40 EeV, coming in just 5above the horizon. The energy fluence as a function of distance to the shower axis has been fitted with a specially adapted function for HAS [91]. This function describes the data very well and yields the total e/m energy delivered to the ground in form of radio waves (in the 30 80 MHz band). The shower has been reconstructed independently, using the information from the RDs and the WCDs. Both, the direction (azimuth and zenith angles) as well as the energy agree perfectly within their uncertainties. The location of the shower axis, intersecting with the ground is slightly offset between the particle content of the shower (WCD) and the radiation (RD), this is caused by refractive effects during the propagation of the radio waves in the atmosphere [92]. For illustration and to facilitate cross checks an independent reconstruction has been applied for Fig.5. Goal of the proposed work is to establish a hybrid reconstruction, combining the information from the RDs and WCDs (sub project #1). Another data-driven approach to verify the performance of the RD system is to investigate the measured e/m energy (from RD) as a function of the total shower energy (from WCD), as illustrated in Fig.6. This is a very important plot since it illustrates the end-to-end verification of the complete signal chain. The absolute energy scale obtained by the RD is fully compatible with the well established energy scale of the WCD. It should be noted that the data shown are taken during the installation, i.e. only with a fraction of the full 3000 km2array. Therefore, only loose quality cuts have been applied (to keep some showers for the analysis), which results in a larger scattering of points around the main diagonal. In the future, when we can apply stricter quality cuts, we expect a much narrower distribution along the main diagonal. It should also be noted that the e/m energy is 7 Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 27 Deflection of cosmic rays in magnetic fields DRAFT VERSION MAY 28, 2024 Typeset using L A T EXtwocolumn style in AASTeX631 The Coherent Magnetic Field of the Milky Way MICHAEL UNGER 1, 2 AND GLENNYS R. FARRAR 3 1Institute for Astroparticle Physics (IAP), Karlsruhe Institute of Technology (KIT), Karlsruhe, Germany 2Institutt for fysikk, Norwegian University of Science and Technology (NTNU), Trondheim, Norway 3Center for Cosmology and Particle Physics, Department of Physics, New York University, New York, NY 10003, USA ABSTRACT We present a suite of models of the coherent magnetic field of the Galaxy (GMF) based on new divergencefree parametric functions describing the global structure of the field. The model parameters are fit to the latest full-sky Faraday rotation measures of extragalactic sources (RMs) and polarized synchrotron intensity (PI) maps from WMAP and Planck. We employ multiple models for the density of thermal and cosmic-ray electrons in the Galaxy, needed to predict the skymaps of RMs and PI for a given GMF model. The robustness of the inferred properties of the GMF is gauged by studying many combinations of parametric field models and electron density models. We determine the pitch angle of the local magnetic field ((11 ±1)), explore the evidence for a granddesign spiral coherent magnetic field (inconclusive), determine the strength of the toroidal and poloidal magnetic halo fields below and above the disk (magnitudes the same for both hemispheres within ⇡10%), set constraints on the half-height of the cosmic-ray diffusion volume (2.9 kpc), investigate the compatibility of RMand PI-derived magnetic field strengths (compatible under certain assumptions) and check if the toroidal halo field could be created by the shear of the poloidal halo field due to the differential rotation of the Galaxy (possibly). A set of eight models is identified to help quantify the present uncertainties in the coherent GMF spanning different functional forms, data products and auxiliary input. We present the corresponding skymaps of rates for axion-photon conversion in the Galaxy, and deflections of ultra-high energy cosmic rays. Keywords: Galactic magnetic field, Galactic physics, Milky Way, Cosmic rays 1. INTRODUCTION Spiral galaxies are known to be permeated by large-scale magnetic fields, with energy densities comparable to the turbulent and thermal energy densities of the interstellar medium; see e.g. Beck (2016) for a recent review. A good knowledge of the global structure of these fields is important for understanding their origin, infering their effect on galactic dynamics, estimating the properties of diffuse motion of low-energy Galactic cosmic rays, and studying the impact of magnetic deflections on the arrival directions of extragalactic ultrahigh-energy cosmic rays. The GMF is also important for new physics studies, for instance axion-photon conversion in the GMF or the interpretation of possible signatures of astrophysical dark matter annihilation. The determination of the large-scale structure of the magnetic field of our Galaxy is particularly challenging since one [email protected] [email protected] must infer it from the vantage point of Earth, located inside the field. Previous attempts to model the Galactic magnetic field (GMF) are summarized by Jaffe (2019). In this paper, we focus on the coherent magnetic field of the Galaxy, leaving the study of its turbulent component to the near future. Following Jansson & Farrar (2012a) (hereafter JF12), we derive the GMF by fitting suitably general parametric models of its structure to the two astrophysical data sets which are the most constraining of the coherent magnetic fields: the rotation measures (RMs) of extragalactic polarized radio sources and the polarized intensity (PI) of the synchrotron emission of cosmic-ray electrons in the Galaxy. The relation of these two astrophysical observables to the magnetic field is detailed in Sec. 2, followed by a description of the RM and PI data in Sec. 3. The interpretation of this data relies on the knowledge of the three-dimensional density of thermal electrons and cosmic-ray electrons in the Galaxy. We discuss these auxiliary models in Sec. 4. The parametric models of the GMF investigated in this paper are introduced in Sec. 5and the model optimization is described in Sec. 6. arXiv:2311.12120v3 [astro-ph.GA] 26 May 2024 DRAFT VERSION MAY 28, 2024 Typeset using L A T EXtwocolumn style in AASTeX631 The Coherent Magnetic Field of the Milky Way MICHAEL UNGER 1, 2 AND GLENNYS R. FARRAR 3 1Institute for Astroparticle Physics (IAP), Karlsruhe Institute of Technology (KIT), Karlsruhe, Germany 2Institutt for fysikk, Norwegian University of Science and Technology (NTNU), Trondheim, Norway 3Center for Cosmology and Particle Physics, Department of Physics, New York University, New York, NY 10003, USA ABSTRACT We present a suite of models of the coherent magnetic field of the Galaxy (GMF) based on new divergencefree parametric functions describing the global structure of the field. The model parameters are fit to the latest full-sky Faraday rotation measures of extragalactic sources (RMs) and polarized synchrotron intensity (PI) maps from WMAP and Planck. We employ multiple models for the density of thermal and cosmic-ray electrons in the Galaxy, needed to predict the skymaps of RMs and PI for a given GMF model. The robustness of the inferred properties of the GMF is gauged by studying many combinations of parametric field models and electron density models. We determine the pitch angle of the local magnetic field ((11 ±1)), explore the evidence for a granddesign spiral coherent magnetic field (inconclusive), determine the strength of the toroidal and poloidal magnetic halo fields below and above the disk (magnitudes the same for both hemispheres within ⇡10%), set constraints on the half-height of the cosmic-ray diffusion volume (2.9 kpc), investigate the compatibility of RMand PI-derived magnetic field strengths (compatible under certain assumptions) and check if the toroidal halo field could be created by the shear of the poloidal halo field due to the differential rotation of the Galaxy (possibly). A set of eight models is identified to help quantify the present uncertainties in the coherent GMF spanning different functional forms, data products and auxiliary input. We present the corresponding skymaps of rates for axion-photon conversion in the Galaxy, and deflections of ultra-high energy cosmic rays. Keywords: Galactic magnetic field, Galactic physics, Milky Way, Cosmic rays 1. INTRODUCTION Spiral galaxies are known to be permeated by large-scale magnetic fields, with energy densities comparable to the turbulent and thermal energy densities of the interstellar medium; see e.g. Beck (2016) for a recent review. A good knowledge of the global structure of these fields is important for understanding their origin, infering their effect on galactic dynamics, estimating the properties of diffuse motion of low-energy Galactic cosmic rays, and studying the impact of magnetic deflections on the arrival directions of extragalactic ultrahigh-energy cosmic rays. The GMF is also important for new physics studies, for instance axion-photon conversion in the GMF or the interpretation of possible signatures of astrophysical dark matter annihilation. The determination of the large-scale structure of the magnetic field of our Galaxy is particularly challenging since one [email protected] [email protected] must infer it from the vantage point of Earth, located inside the field. Previous attempts to model the Galactic magnetic field (GMF) are summarized by Jaffe (2019). In this paper, we focus on the coherent magnetic field of the Galaxy, leaving the study of its turbulent component to the near future. Following Jansson & Farrar (2012a) (hereafter JF12), we derive the GMF by fitting suitably general parametric models of its structure to the two astrophysical data sets which are the most constraining of the coherent magnetic fields: the rotation measures (RMs) of extragalactic polarized radio sources and the polarized intensity (PI) of the synchrotron emission of cosmic-ray electrons in the Galaxy. The relation of these two astrophysical observables to the magnetic field is detailed in Sec. 2, followed by a description of the RM and PI data in Sec. 3. The interpretation of this data relies on the knowledge of the three-dimensional density of thermal electrons and cosmic-ray electrons in the Galaxy. We discuss these auxiliary models in Sec. 4. The parametric models of the GMF investigated in this paper are introduced in Sec. 5and the model optimization is described in Sec. 6. arXiv:2311.12120v3 [astro-ph.GA] 26 May 2024 THE COHERENT MAGNETIC FIELD OF THE MILKY WAY 29 o +180 o -180 = 20 EVR base expX neCL spur nebCor twistX cre10 synCG JF12 Figure 19. Angular deflections of ultrahigh-energy cosmic rays in the eight model variations derived in this paper and JF12. The cosmic-ray rigidity is 20 EV (2⇥1019 V). Filled circles denote a grid of arrival directions and the open symbols are the back-tracked directions at the edge of the Galaxy. Figure 20. Left: Rigidity threshold such that the angular deflection in the given direction is 20in all models. Right: Rigidity threshold such that the model predictions of the angular deflection differ by 20. (1 EV =1018 V) JF12 model are generally within the range of deflections predicted for the GMF models derived in this work. This is not the case for the deflections calculated with the GMF model of Pshirkov et al. (2011), due to the absence of a poloidal component in that model (c.f., Sec. 7.6). Current studies of the anisotropies of ultrahigh-energy cosmic rays indicate the presence of “hot spots” of cosmic-ray clusters at intermediate angular scales of 20(Abbasi et al. 2014;Abreu et al. 2022). For the identification of extragalactic sources related to these overdensities, a precision in backtracking through the GMF at least as good as their angular size, qmax, is needed. Figure 20 aims to illustrate this requirement. In the left panel, we show the minimum rigidity such that the deflection for a CR arriving in the given direction is less than qmax =20in all 8 models. Requiring that the deflections in half of the sky are less than qmax =20, according to all of these models, requires the rigidity to be greater than or equal to Rnocorr 50 =20 EV. The minimum rigidity requirement improves considerably if the arrival directions are corrected for their expected deflection in the GMF. The limit on the precision with which we infer the source position arises from the difference between the models, and not the overall magnitude of the deflection. The differences of predicted deflections within the model ensemble are smaller than the deflections themselves. Therefore, as shown in the right panel of Fig. 20, the required minimum rigidity is lower when the deflections are corrected for. With corrections, the rigidity quantile at which half of the sky can be observed at qmax =20or better, decreases to Rcorr 50 =11 EV giving a much greater observational reach. Note that this discussion is indicative only, since the minimal rigidity requirement may change when random fields are included in the analysis. 8.2. Axions Another important application of the model ensemble presented in this paper is the prediction of the conversion of need to know rigidity (mass) of incoming cosmic rays R=E Z⇡E A/2 Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 28 Neutrino flux limits Hörandel Part B2 COSMICISFONTIBUS Figure 2: Left: Upper limits on the integral photon flux from the PAO alongside limits from other experiments. Shaded bands show cosmogenic flux predictions from UHECR interactions with Galactic matter (grey), background radiation fields (violet, green, orange), and hot gas in the Galactic halo (blue). Dashed lines denote super-heavy Dark Matter predictions [26]. Right: upper limits on the diffuse flux of neutrinos for two Auger analyses together with the recent KM3NeT neutrino detection. IceCube limits are scaled for a E2 nflux assumption. The predicted fluxes from cosmogenic and astrophysical neutrino models are illustrated as well [27]. from the SBG catalogue of around 20% at 40 EeV1with a magnetic field blurring of around 20for a rigidity of 10 EV1provides a fair simultaneous description of all three observables, see Fig.1. The SBG model is favoured with a significance of 4.5scompared to a reference model with only homogeneously distributed background sources. By investigating a scenario with Centaurus A as a single source in combination with the homogeneous background, it is confirmed that this region of the sky provides the dominant contribution to the observed anisotropy signal. Models containing a catalogue of jetted active galactic nuclei whose flux scales with the g-ray emission are, however, disfavoured as they cannot adequately describe the measured arrival directions. Figure 3: The spectral flux of gamma rays, neutrinos, and UHECRs [25]. The PAO data are also used to contribute to multimessenger astroparticle physics through strong flux limits on ultra-highenergy neutrino and photon fluxes [26–31], see Fig.2. Those can be used to provide an independent, indirect measurement of the mass composition of UHECRs, they exclude that UHECRs are only protons at the highest energies. Combining the multimessenger information from gs, ns, and CRs is a powerful tool to understand the origin of UHECRs [25,32]. This is illustrated in Fig.3: The spectral flux (F)of neutrinos inferred from the IceCube eight-year up-going track analysis (red fit) and the six-year high-energy starting event analysis (magenta fit) compared to the flux of unresolved extragalactic g-ray sources (blue data) and UHECRs (green data). The neutrino spectra are indicated by the best-fit power-law (solid line) and 1suncertainty range (shaded range). The multi-messenger connections are highlighted: A: The joined production of charged pions (p±)and neutral pions (p0)in CR interactions leads to the emission of neutrinos (dashed blue) and g-rays (solid blue), respectively. B: CR emission models (solid green) of UHECRs imply a maximal flux (calorimetric limit) of neutrinos from the same sources (green dashed). C: The same CR model predicts the emission of cosmogenic neutrinos from the collision with cosmic microwave background photons (GZK mechanism). 11 EeV=1018 eV. Particle rigidity: energy/charge E/Z. 1 EV=1018 V. 2 S. Sehgal et al. Proceedings of Science ICRC2025 (2025) 1170 see Marcus Niechciol, this conference Jörg R. Hörandel - XXI Workshop on Neutrino Telescopes - Padova 2025 28 Neutrino flux limits CR detected from 0.5° above horizon with E~52 EeV! crosses complete array ~65 km footprint Hörandel Part B2 COSMICISFONTIBUS Milestones and Deliverables: •M21 Algorithm to select CR nuclei of a certain charge and energy, i.e. rigidity E/Z⇡2·E/A(with Z⇡A/2 for most nuclei) implemented in the PAO software framework. •P21 Publication on "rigidity-resolved sky maps of UHECRs and implications on models of the origin of UHECRs". In case point sources or enhancements are found on the sky, this would be a groundbreaking way forward and would enable the •P22 Publication "charged-particle astronomy: a new window to the high-energy Universe". Sub project #3: Hybrid air shower reconstruction for neutral particles (photons and neutrinos) (PhD2+PI). Neutral particles, such as photons and neutrinos provide valuable insight to the extreme processes in the high-energy Universe in addition to nuclear CRs. Detecting EeV photons and neutrinos would be a break through for the field. Objective of the proposed work is to improve the photon and neutrino detection at the highest energies >1019 eV, using the unique facility of the PAO RD, with its strong sensitivity to the e/m shower component. Already now (see Fig.2) the upper limits from the PAO start to constrain models for the production of photons and neutrinos at the highest energies. Increasing the sensitivity by an order of magnitude would either allow to further constrain the production models, or in an ideal case, detect the first photons and neutrinos ever in the EeV energy range. Bjarni Pont for the Pierre Auger Collaboration — ICRC2025 — July 2025 10 Earth-skimming potential •65km footprint (larger than Lake Geneva)" •Nearly from the horizon: 89 degrees." •Around 2x1019 eV NAC Radio signal: Energy fluency [eV/m2] 65km Geneva S Figure 11: Measured HAS arriving 0.5 above the horizon. The footprint has a length of ⇠65 km, crossing the complete array. The energy reconstructed form the RD is 52 EeV. Air showers initiated by high-energy photons in the atmosphere differ significantly from those from nuclei [102]. For a photoninduced shower, the first interactions and generations are purely electromagnetic, since the radiation length is more than two orders of magnitude smaller than the mean free path for photonuclear interactions. Yet, the development of the shower is delayed by the typically small multiplicity of electromagnetic interactions. Thus the maximum development of the shower is reached at a slant atmospheric depth Xmax larger for photon primaries than for nuclei, with a difference of ⇡200 g/cm2between photons and protons at 1019 eV and even larger between photons and heavy nuclei. Since the mean free path for photonuclear interactions is much larger than the radiation length, the transfer of energy to the hadron and muon channels is reduced hence only a small fraction of the electromagnetic component in a photon-induced shower is injected into the hadronic cascade. Showers induced by photons are thus characterised by a lower content of muons: on average, simulations show that photon showers have nearly one order of magnitude less muons than proton showers of the same energy. g-induced air showers need to be treated differently in the reconstruction software. An extra algorithm needs to be developed to convert the measured quantities (with a strong emphasise on the e/m component) to the energy of the incoming photon. For hadron-induced HAS most particles are absorbed while traversing a relatively large amount of atmosphere, only high-energy muons are expected at detector level. In contrast, neutrinos will interact close to the array either in the nearby Andes mountains or the Earth crust (for upwardgoing ns). This results in air showers with a large(r) e/m component and a HAS with a large e/m-to-muon ratio is an interesting ncandidate. A HAS coming in 5above the horizon from the direction of the Andes is shown in Fig.5. If this shower would have a slightly lower elevation (the Andes are seen from the PAO at elevations <2) this would be a very interesting ncandidate. Another interesting, recently recorded shower, coming in only 0.5above the horizon with an energy of 52 EeV is depicted in Fig.11. If such a shower would come from the direction of the Andes (left in the figure) it would be a serious neutrino candidate. The different e/m-to-muon ratio in n-induced showers compared to nuclei-induced ones requires a different treatment of n-showers in the reconstruction software. We aim to implement a reconstruction algorithm which combines the WCD and RD information to select neutrino candidates and to establish an energy scale for neutrinos in the PAO reconstruction software. 11 Hörandel Part B2 COSMICISFONTIBUS Figure 2: Left: Upper limits on the integral photon flux from the PAO alongside limits from other experiments. Shaded bands show cosmogenic flux predictions from UHECR interactions with Galactic matter (grey), background radiation fields (violet, green, orange), and hot gas in the Galactic halo (blue). Dashed lines denote super-heavy Dark Matter predictions [26]. Right: upper limits on the diffuse flux of neutrinos for two Auger analyses together with the recent KM3NeT neutrino detection. IceCube limits are scaled for a E2 nflux assumption. The predicted fluxes from cosmogenic and astrophysical neutrino models are illustrated as well [27]. from the SBG catalogue of around 20% at 40 EeV1with a magnetic field blurring of around 20for a rigidity of 10 EV1provides a fair simultaneous description of all three observables, see Fig.1. The SBG model is favoured with a significance of 4.5scompared to a reference model with only homogeneously distributed background sources. By investigating a scenario with Centaurus A as a single source in combination with the homogeneous background, it is confirmed that this region of the sky provides the dominant contribution to the observed anisotropy signal. Models containing a catalogue of jetted active galactic nuclei whose flux scales with the g-ray emission are, however, disfavoured as they cannot adequately describe the measured arrival directions. Figure 3: The spectral flux of gamma rays, neutrinos, and UHECRs [25]. The PAO data are also used to contribute to multimessenger astroparticle physics through strong flux limits on ultra-highenergy neutrino and photon fluxes [26–31], see Fig.2. Those can be used to provide an independent, indirect measurement of the mass composition of UHECRs, they exclude that UHECRs are only protons at the highest energies. Combining the multimessenger information from gs, ns, and CRs is a powerful tool to understand the origin of UHECRs [25,32]. This is illustrated in Fig.3: The spectral flux (F)of neutrinos inferred from the IceCube eight-year up-going track analysis (red fit) and the six-year high-energy starting event analysis (magenta fit) compared to the flux of unresolved extragalactic g-ray sources (blue data) and UHECRs (green data). The neutrino spectra are indicated by the best-fit power-law (solid line) and 1suncertainty range (shaded range). The multi-messenger connections are highlighted: A: The joined production of charged pions (p±)and neutral pions (p0)in CR interactions leads to the emission of neutrinos (dashed blue) and g-rays (solid blue), respectively. B: CR emission models (solid green) of UHECRs imply a maximal flux (calorimetric limit) of neutrinos from the same sources (green dashed). C: The same CR model predicts the emission of cosmogenic neutrinos from the collision with cosmic microwave background photons (GZK mechanism). 11 EeV=1018 eV. Particle rigidity: energy/charge E/Z. 1 EV=1018 V. 2 S. Sehgal et al. Proceedings of Science ICRC2025 (2025) 1170 see Marcus Niechciol, this conference