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Given our state of knowledge on the sources of ultrahigh energy cosmic rays. the detection of such a halo would have a lasting impact on this field of research. | Given our state of knowledge on the sources of ultrahigh energy cosmic rays, the detection of such a halo would have a lasting impact on this field of research. |
The detectability of such signals thus deserves close serutiny. | The detectability of such signals thus deserves close scrutiny. |
The production of gamma ray signatures of ultrahigh energy cosmic rays have been studied numerically by ? and ?.. | The production of gamma ray signatures of ultrahigh energy cosmic rays have been studied numerically by \cite{Ferrigno04} and \cite{ASM06}. |
However. the former authors have focused their study on the Compton cascades following the production of ultrahigh energy photons and pairs close to the source and they have neglected the deflection imparted by the surrounding magnetic fields which dilutes the gamma-ray signal (2).. while ?. have addressed both the synchrotron emission of secondary pairs and the Compton cascading down to TeV energies. albeit for the particular case of a source located in à magnetized cluster of galaxies. | However, the former authors have focused their study on the Compton cascades following the production of ultrahigh energy photons and pairs close to the source and they have neglected the deflection imparted by the surrounding magnetic fields which dilutes the gamma-ray signal \citep{GA05}, while \cite{ASM06} have addressed both the synchrotron emission of secondary pairs and the Compton cascading down to TeV energies, albeit for the particular case of a source located in a magnetized cluster of galaxies. |
The high magnetic field that prevails in such environments increases the residence time of primary and secondary charged particles and thus also increases the gamma ray flux. | The high magnetic field that prevails in such environments increases the residence time of primary and secondary charged particles and thus also increases the gamma ray flux. |
The present paper aims at examining the prospects for the detectability of gamma ray halos around ultrahigh energy cosmic ray sources. relaxing most of the assumptions made in the above previous studies. | The present paper aims at examining the prospects for the detectability of gamma ray halos around ultrahigh energy cosmic ray sources, relaxing most of the assumptions made in the above previous studies. |
In particular. we discuss the more general case of a source located in the field. outside clusters of galaxies. | In particular, we discuss the more general case of a source located in the field, outside clusters of galaxies. |
We focus our discussion on the synchrotron signal emitted by secondary pairs. which offers a possibility of unambiguous detection: nevertheless. the deflection and the dilution of the Compton cascading gamma ray signal at TeV energies is also discussed. | We focus our discussion on the synchrotron signal emitted by secondary pairs, which offers a possibility of unambiguous detection; nevertheless, the deflection and the dilution of the Compton cascading gamma ray signal at TeV energies is also discussed. |
Going further than ?. and ?.. we take into account the inhomogeneous distribution of the magnetic fields in the source environment. | Going further than \cite{Aharonian02} and \cite{GA05}, we take into account the inhomogeneous distribution of the magnetic fields in the source environment. |
We also relax the assumption of a pure proton composition of ultra high energy cosmic rays. underlying to the above studies. | We also relax the assumption of a pure proton composition of ultra high energy cosmic rays, underlying to the above studies. |
The chemical composition of ultrahigh energy cosmic rays indeed remains an open question. | The chemical composition of ultrahigh energy cosmic rays indeed remains an open question. |
While experiments such as the Fly’s Eye and HiRes have suggested a transition from heavy to light above ~10/5? eV (???).. the most recent measurements made with the Pierre Auger Observatory rather point towards a heavy composition above 101 eV (?22).. | While experiments such as the Fly's Eye and HiRes have suggested a transition from heavy to light above $\sim 10^{18.5}$ eV \citep{Fly, Hires, Hires10}, the most recent measurements made with the Pierre Auger Observatory rather point towards a heavy composition above $10^{19}\,$ eV \citep{U07,Abraham09_comp,Auger10}. |
As the energy losses and magnetic deflection of high energy nuclei differ from those of a proton of a same energy. one should naturally expect different gamma ray signatures. | As the energy losses and magnetic deflection of high energy nuclei differ from those of a proton of a same energy, one should naturally expect different gamma ray signatures. |
The lay-out of the present paper is as follows. | The lay-out of the present paper is as follows. |
In Section 2.. we first test the dependence of the gamma ray flux produced by ultrahigh energy cosmic rays on the type. intensity and structure of magnetized environments. | In Section \ref{section:B}, we first test the dependence of the gamma ray flux produced by ultrahigh energy cosmic rays on the type, intensity and structure of magnetized environments. |
We also discuss the effects of various chemical compositions and injection spectra. | We also discuss the effects of various chemical compositions and injection spectra. |
We conclude on the robustness of the gamma ray signature. according to these parameters and find that the normalization and thus the detectability of this flux ultimately depends on the energy injected in the primary cosmic rays. | We conclude on the robustness of the gamma ray signature according to these parameters and find that the normalization and thus the detectability of this flux ultimately depends on the energy injected in the primary cosmic rays. |
In Section ??.. we discuss the detectability of ultrahigh energy cosmic ray signatures in. gamma rays. | In Section \ref{section:detectability}, we discuss the detectability of ultrahigh energy cosmic ray signatures in gamma rays. |
Applying the results of our calculations. we show that the average type of sources contributing to the ultrahigh energy cosmic ray spectrum produces a gamma ray flux more than two orders of magnitudes lower than the sensitivity of the current and upcoming instruments. | Applying the results of our calculations, we show that the average type of sources contributing to the ultrahigh energy cosmic ray spectrum produces a gamma ray flux more than two orders of magnitudes lower than the sensitivity of the current and upcoming instruments. |
We then explore the case of rare powerful sources with cosmic ray luminosity over energy 10 eV of Lig>10776 eres”! | We then explore the case of rare powerful sources with cosmic ray luminosity over energy $10^{19}$ eV of $L_{19}>10^{44-46}$ $\,$ $^{-1}$. |
We assume throughout this paper that sources emit isotropically and discuss how the conclusions are modified for beamed emission in Section ??.. | We assume throughout this paper that sources emit isotropically and discuss how the conclusions are modified for beamed emission in Section \ref{section:conclusion}. |
The gamma ray signatures of those sources could be detectable provided that they are located far enough not to overshoot the observed cosmic ray spectrum. | The gamma ray signatures of those sources could be detectable provided that they are located far enough not to overshoot the observed cosmic ray spectrum. |
Finally. we also briefly discuss the detection of nearby sources. considering the radiogalaxy Centaurus A as a protoypical example. | Finally, we also briefly discuss the detection of nearby sources, considering the radiogalaxy Centaurus A as a protoypical example. |
We draw our conclusions in Section ??.. | We draw our conclusions in Section \ref{section:conclusion}. |
As mentioned above. we focus our discussion on the synchrotron signal that can be produced close to the source by very high energy electrons and positrons. that result themselves from the interactions of primary cosmic. rays. | As mentioned above, we focus our discussion on the synchrotron signal that can be produced close to the source by very high energy electrons and positrons, that result themselves from the interactions of primary cosmic rays. |
The signal associated to inverse Compton cascades of these electrons/positrons on radiation backgrounds will be discussed in Section ??.. | The signal associated to inverse Compton cascades of these electrons/positrons on radiation backgrounds will be discussed in Section \ref{sec:compton-cascade}. |
The secondary electrons and positrons are ereated through one of the three following channels: 1) by the decay of a charged pion produced during a photo-hadronie interaction (Ayamrt+... then >μμv, and pro—eyevs. With A a cosmic ray nucleus). i1) by photo pair production during an interaction with a background photon (Ay—ο+...). or ii) by the disintegration of a neutral ptor into ultrahigh energy photons which then interact with CMB and radio backgrounds to produce electron and positron pairs (Ay—2°+... then zx"2y. and yYe—€€. With ype a cosmic background photon). | The secondary electrons and positrons are created through one of the three following channels: i) by the decay of a charged pion produced during a photo-hadronic interaction $A\,\gamma\rightarrow
\pi^+ +...$, then $\pi^+\rightarrow \mu^+ \, \nu_\mu$, and $\mu^+\rightarrow e^+ \, \nu_e \, \bar{\nu_\mu}$, with $A$ a cosmic ray nucleus), ii) by photo pair production during an interaction with a background photon $A\,\gamma\rightarrow e^+\,e^- +...$ ), or iii) by the disintegration of a neutral pion into ultrahigh energy photons which then interact with CMB and radio backgrounds to produce electron and positron pairs $A\,\gamma\rightarrow \pi^0 +...$, then $\pi^0\rightarrow 2\gamma$, and $\gamma\,\gamma_{\rm bg}\rightarrow
e^+ \, e^-$, with $\gamma_{\rm bg}$ a cosmic background photon). |
In all cases. the resulting electrons and. positrons typically carry up to à few percents of the initial cosmic ray energy. | In all cases, the resulting electrons and positrons typically carry up to a few percents of the initial cosmic ray energy. |
While propagating in the intergalactic medium. these ultrahigh energy electrons and positrons up-scatter CMB or radio. photons through inverse Compton processes and/or they lose energy through synchrotron radiation. | While propagating in the intergalactic medium, these ultrahigh energy electrons and positrons up-scatter CMB or radio photons through inverse Compton processes and/or they lose energy through synchrotron radiation. |
Following ?.. the effective inverse Compton cooling on the CMB and radio backgrounds can be written as x5Mpe(£,/10'*eV)", with ajc=Lif the electron energy E,€10 eV and aie=0.25 if 10eV.€E,<10°’ eV. Above 10'S eV. the sealing of the cooling length actually depends on the assumptions made for the radio background. which ts unfortunately not very well known. | Following \cite{GA05}, the effective inverse Compton cooling on the CMB and radio backgrounds can be written as $x_{e\gamma}\,\approx\,5\,{\rm Mpc}\,(E_e/10^{18}\,{\rm
eV})^{\alpha_{\rm IC}}$ , with $\alpha_{\rm IC}=1$ if the electron energy $E_e\lesssim 10^{18}\,$ eV and $\alpha_{\rm IC}\simeq 0.25$ if $10^{18}\,{\rm eV}\,\lesssim E_e\lesssim 10^{20}\,$ eV. Above $10^{18}$ eV, the scaling of the cooling length actually depends on the assumptions made for the radio background, which is unfortunately not very well known. |
The possible differences that such uncertainties could introduce should however not | The possible differences that such uncertainties could introduce should however not |
Transiting- extrasolar planets are of. great scientific.↼↔ value. | Transiting extrasolar planets are of great scientific value. |
While the radial velocity method continues to. be very successful in finding planets and characterising their orbits. only transits can currently reveal the properties of the planets themselves. | While the radial velocity method continues to be very successful in finding planets and characterising their orbits, only transits can currently reveal the properties of the planets themselves. |
In addition to the basic planetary parameters that can be determined. such as planet mass. size. and average density. the atmospheres of transiting planets can be probed through either secondary eclipse observations (e.g. Charbonneau et al. | In addition to the basic planetary parameters that can be determined, such as planet mass, size, and average density, the atmospheres of transiting planets can be probed through either secondary eclipse observations (e.g. Charbonneau et al. |
2005; Deming et al. | 2005; Deming et al. |
2005: Knutson et al. | 2005; Knutson et al. |
2007) or atmospheric transmission. spectroscopy. the subject of this paper | 2007) or atmospheric transmission spectroscopy, the subject of this paper. |
In transmission spectroscopy. the depth of a planet transit is measured as function of wavelength. | In transmission spectroscopy, the depth of a planet transit is measured as function of wavelength. |
It is expected that at certain wavelengths. a transit will be slightly deeper due to absorption in. the planets atmosphere. | It is expected that at certain wavelengths, a transit will be slightly deeper due to absorption in the planet's atmosphere. |
In the optical transmission spectrum of hot Jupiters. the strongest of these absorption features was predicted. to come from ↴↳the sodium. D lines. at 5889-oo and5896on, and (Girown2001: 5 eager&S asselov2000). | In the optical transmission spectrum of hot Jupiters, the strongest of these absorption features was predicted to come from the sodium D lines at 5889 and 5896 $ $ (Brown 2001; Seager Sasselov 2000). |
Indeed. Charbonneauetal | Indeed, Charbonneau et al. ( |
dqand tenia 33b iio Fthetransi ina wideband. usingtheSTLS spectrographontheHubbleS Gaselesco | 2002) detected D absorption in the transmission spectrum of the transiting exoplanet HD209458b, at a level of $\pm$ in a $ $ wide band, using the STIS spectrograph on the Hubble Space Telescope (HST). |
dletectepe | This constitutes the first detection of an atmosphere around an extrasolar planet. |
l ο)HT thoughttobecausedbyaneva poratingexos phereVidal— jaretal. | Subsequently, strong absorption features have been detected in HD209458b from hydrogen, oxygen, and carbon at a level of, thought to be caused by an evaporating exosphere (Vidal-Madjar et al. |
2003. 2004) | 2003; 2004). |
Recently. hothydrogenhasbeenderectedby B | Recently, hot hydrogen has been detected by Ballester et al. ( |
ofjUg | 2007), also using HST data. |
eiera | A claim by Tinetti et al. ( |
tioisdita aedauetiud àdd/aisaikeiniby | 2007) of the detection of water vapour from a comparison of transit depths at several wavelengths in the infrared, as measured with Spitzer, has been disputed by Ehrenreich et al. ( |
null | 2007). |
T inertietal ( | This, while Swain et al. ( |
2007jo f | 2008) identify both water and methane from NICMOS/HST data. |
Despite several attempts from ground-based observatories. no confirmation has yet been obtained of the NaDp planetary absorption feature in the transmission spectrum of HD209458b. | Despite several attempts from ground-based observatories, no confirmation has yet been obtained of the D planetary absorption feature in the transmission spectrum of HD209458b. |
In Pondeeneral Peround-basede transmissionunoe spectroscopy has not been a great success. | In general, ground-based transmission spectroscopy has not been a great success. |
Typically. upper limits to à NaDD absorption signal of 0.1-1% have been reached (Moutou et al. | Typically, upper limits to a D absorption signal of $-$ have been reached (Moutou et al. |
2001: Snellen 2004: Narita et al. | 2001; Snellen 2004; Narita et al. |
2005). implying that systematic effects dominate the error budgets. | 2005), implying that systematic effects dominate the error budgets. |
A modern Echelle spectrograph on a 8-10 m. class telescope can provide spectra from the brightest transiting exoplanet systems with signal-to-noise ratios. SNR> 100. within a few minutes of exposure time. | A modern Echelle spectrograph on a $-$ 10 m. class telescope can provide spectra from the brightest transiting exoplanet systems with signal-to-noise ratios, $>$ 100, within a few minutes of exposure time. |
Integrating over a few Angstrom and over the duration of a transit. this would mean that photon noise statistics should allow detections down to a few times 1071. | Integrating over a few Angstrom and over the duration of a transit, this would mean that photon noise statistics should allow detections down to a few times $^{-4}$. |
Although the HST detection of the NaDD absorption feature is only just at this level. it is expected that its width is only a fraction of the (2002). passbandusedbyCharbonneauetal. | Although the HST detection of the D absorption feature is only just at this level, it is expected that its width is only a fraction of the $ $ passband used by Charbonneau et al. ( |
analysisoftheS TISdatas ingetal. | 2002), as recently shown by of the STIS data (Sing et al. |
2008). | 2008). |
ThismeansthattheNaabsorptioi band should be at the ~1077 level. | This means that the Na absorption within a $ $ band should be at the $\sim10^{-3}$ level. |
In this paper we re-analyse a data-set from the High Dispersion Spectrograph on the Subaru telescope. that covers one transit of HD209458b. | In this paper we re-analyse a data-set from the High Dispersion Spectrograph on the Subaru telescope, that covers one transit of HD209458b. |
The aim is to identify and correct for possible systematic effects. and to improve on the results previously presented by Narita et al. ( | The aim is to identify and correct for possible systematic effects, and to improve on the results previously presented by Narita et al. ( |
2005: NAROS). | 2005; NAR05). |
Their analysis resulted in spectra with an SNR of afew hundred in the stellar continuum. | Their analysis resulted in spectra with an SNR of a few hundred in the stellar continuum. |
However. near strong absorption lines. such as the NaDD doublet. clear coherent structures were visible (Figure 2 in NAROS). well in excess of the photon shot noise. | However, near strong absorption lines, such as the D doublet, clear coherent structures were visible (Figure 2 in NAR05), well in excess of the photon shot noise. |
While NAROS argued that the positions of the spectral lines are stable to within0.01À.. spectral shifts at only a fraction of this level (e.g. due to slit-centering variations) could cause these effects. | While NAR05 argued that the positions of the spectral lines are stable to within, spectral shifts at only a fraction of this level (e.g. due to slit-centering variations) could cause these effects. |
This encouraged us to analyse this data again. | This encouraged us to analyse this data again. |
In Section 2. the observations. data reduction and analysis are described. | In Section 2, the observations, data reduction and analysis are described. |
The result on the DD absorption are presented discussed. in ↼∖↼Section 3.4 together with. à comparison. with. 02) iesJsID with a recent detection of NaDD paceT(the, BIISIHSTASdi1c VFhisciaegOPE SIe stdetectionofana | The result on the D absorption are presented and discussed in Section 3, together with a comparison with the STIS/HST results, and with a recent detection of D absorption in exoplanet HD189733b (Redfield et al. |
tmospheredi = HD209458- was observed on the night. of- October 24. 2007. "Ing the High Dispersion Spectrograph (HDS: Noguchi et b | 2008) HD209458 was observed on the night of October 24, 2002, using the High Dispersion Spectrograph (HDS; Noguchi et al. |
2002) on the Subaru telescope. | 2002) on the Subaru telescope. |
We obtained the data Usingno SMOKATVA archiveueshbis system. (Baba"m et al. | We obtained the data using the SMOKA archive system (Baba et al. |
; 2002). | 2002). |
nno The observationsthe have been described in Winn et al. ( | The observations have been described in Winn et al. ( |
2004) and Narita et al.(2005: NAROS). | 2004) and Narita et al.(2005; NAR05). |
Thirty-two spectra were taken inmode (without the todine cell). of which the last thirty were made with an exposure time of 500 seconds. | Thirty-two spectra were taken in (without the iodine cell), of which the last thirty were made with an exposure time of 500 seconds. |
The entrance slit was 4” long and 0.8” wide. oriented with à constant position angle. resulting in a spectral resolution of R~45 0000 with 0.9 km s! per pixel. | The entrance slit was $''$ long and $''$ wide, oriented with a constant position angle, resulting in a spectral resolution of $\sim$ 000 with 0.9 km $^{-1}$ per pixel. |
We concentrated on the data from the red CCD. which contains 21 orders of 4100 pixels covering 200Ai4. | We concentrated on the data from the red CCD, which contains 21 orders of 4100 pixels covering $<\lambda<$. |
of the thirty spectra fall outside the transit (nine <O800A..before Twelveingress and three after egress). and eighteen during the transit. | Twelve of the thirty spectra fall outside the transit (nine before ingress and three after egress), and eighteen during the transit. |
The uncertainty in the transit timing 15 neglegible. | The uncertainty in the transit timing is neglegible. |
For the initial data reduction we followed the procedure asrecentlyshownbyre-of NAROS. | For the initial data reduction we followed the procedure of NAR05. |
First the frames were processed using the IRAF software package. including the extraction of the | First the frames were processed using the IRAF software package, including the extraction of the |
rrates on the stellar structure at Z=0. | rates on the stellar structure at $Z=0$. |
A model comparable with our 0.8M; model is Model 1 of Picardietal.(2004, P04),, of composition Y=0.23 and Z=O. | A model comparable with our $0.8 \msun$ model is Model 1 of \citet[][P04]{pic04}, , of composition $Y = 0.23$ and $Z = 0$. |
In their computations, release 4.98 of FRANEC was used and time-dependent convective mixing was calculated; neutrino energy-loss rates by plasma-neutrino emission were modified, with consequences being reported as minimal. | In their computations, release 4.98 of FRANEC was used and time-dependent convective mixing was calculated; neutrino energy-loss rates by plasma-neutrino emission were modified, with consequences being reported as minimal. |
Reaction rates and conductive opacities are, respectively, common with our models. | Reaction rates and conductive opacities are, respectively, common with our models. |
The neutrino energy-loss rates are common with our model for photo- and pair-neutrino processes, but they use energy-loss rate of Dicusetal.(1976) for bremsstrahlung and of Beaudetetal.(1967) for recombination processes. | The neutrino energy-loss rates are common with our model for photo- and pair-neutrino processes, but they use energy-loss rate of \citet{dic76} for bremsstrahlung and of \citet{bea67} for recombination processes. |
For plasma neutrino energy losses, they adopt an energy-loss rate (Espositoetal.2003) which differs only slightly from 196 in the temperature and density ranges relevant to the ignition of the helium flash. | For plasma neutrino energy losses, they adopt an energy-loss rate \citep{esp03} which differs only slightly from I96 in the temperature and density ranges relevant to the ignition of the helium flash. |
Although there are some differences in adopted input physics, the evolution of the helium core flash of the P04 model is similar to that of our 08cf model. | Although there are some differences in adopted input physics, the evolution of the helium core flash of the P04 model is similar to that of our 08cf model. |
The helium core mass at the onset of the core helium flash is the same in both models, namely, Μι=0.52Mo. | The helium core mass at the onset of the core helium flash is the same in both models, namely, $M_{1} = 0.52 \msun$. |
When CNO-cycle reactions are the dominant contributors to the hydrogen-burning luminosity, so that the hydrogen profile is very steep, the quantity My. defined by P04 as the mass of the helium core when the maximum hydrogen-burning luminosity is reached is nearly the same as the quantity Μι we have defined as the mass of the helium core when the core helium flash begins. | When CNO-cycle reactions are the dominant contributors to the hydrogen-burning luminosity, so that the hydrogen profile is very steep, the quantity $M_{\rm{He}}$ defined by P04 as the mass of the helium core when the maximum hydrogen-burning luminosity is reached is nearly the same as the quantity $M_{1}$ we have defined as the mass of the helium core when the core helium flash begins. |
The maximum helium-burning luminosities differ by less than a factor of two, being Lye=1.2x1019L5 in the P04 model and Lye=7.6x10?L5 in our model 08cf. | The maximum helium-burning luminosities differ by less than a factor of two, being $L_{\rm He} = \pow{1.2}{10} L_{\sun}$ in the P04 model and $L_{\rm He} = \pow{7.6}{9} L_{\sun}$ in our model 08cf. |
'The mass at the outer edge of the convective shell at the onset of hydrogen mixing is the same in both cases, namely, 0.506Mo. | The mass at the outer edge of the convective shell at the onset of hydrogen mixing is the same in both cases, namely, $0.506 \msun$. |
The values of Atmix (see Table 2)) and Xc, the mass fraction of carbon in the helium convective shell, are also comparable: Atmix=2.1x10? sec in the P04 model versus 1.87x10? sec in model 08cf, and Xo=4.15x107? in the P04 model versus Xc=4.28x107? in model 08cf. | The values of $\Delta t_{\rm{mix}}$ (see Table \ref{tab:he-flash}) ) and $X_{\rm C}$, the mass fraction of carbon in the helium convective shell, are also comparable: $\Delta t_{\rm{mix}} = \pow{2.1}{5}$ sec in the P04 model versus $\pow{1.87}{5}$ sec in model 08cf, and $X_{\rm C} = \pow{4.15}{-2}$ in the P04 model versus $X_{\rm C} = \pow{4.28}{-2}$ in model 08cf. |
There is a large difference in model characteristics when the helium convective shell first appears; in the P04 model, the convective shell driven by helium burning appears at a mass shell Mpcgs=0.348Mo when Lye=0.658Lo, while, in our model 08cf, Mpcs=0.3825Mo when Lye=1.57x10?Lo. | There is a large difference in model characteristics when the helium convective shell first appears; in the P04 model, the convective shell driven by helium burning appears at a mass shell $M_{\rm BCS} = 0.348 \msun$ when $L_{\rm{He}} = 0.658 L_{\sun}$, while, in our model 08cf, $M_{\rm BCS} = 0.3825 \msun$ when $L_{\rm He} = \pow{1.57}{2} L_{\sun}$. |
The factor of 200 difference in the helium-burning luminosity when shell convection begins is probably simply a typographical error in P04, an interpretation reinforced by the fact that the time for the helium-burning luminosity to reach its maximum value is almost the same in the P04 model (723 yr) and in ours (715 yr). | The factor of 200 difference in the helium-burning luminosity when shell convection begins is probably simply a typographical error in P04, an interpretation reinforced by the fact that the time for the helium-burning luminosity to reach its maximum value is almost the same in the P04 model (723 yr) and in ours (715 yr). |
It takes more than ~3x104 yr for the helium-burning rate to increase by a factor of 200 in this range. | It takes more than $\sim 3\times 10^4$ yr for the helium-burning rate to increase by a factor of 200 in this range. |
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