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2000). to be low the line. | 2000), to be below the line. |
This is understandable: Faster pulsus teud o have larger acceleration potentials aud hence. produce nore energetic photous than the slower pulsars. aud for üeher energv photons. even higher maguetic fields are required for photon splitting to dominate. | This is understandable: Faster pulsars tend to have larger acceleration potentials and hence, produce more energetic photons than the slower pulsars, and for higher energy photons, even higher magnetic fields are required for photon splitting to dominate. |
For example. OY a resonant inverse Coniptou-conutrolled vacuna eap (vlüch is more possible iu the high D regine). the typical »hoton cherev nav be expressed as ε(1οςVO)225.5&104PUTB's 97, and the photon splitting dominant line is (Zhang Tarcding 2001) D,ο --Ὁ or P»2.414«10H(p1s)0879819 I adopting B2pη...10:96VPP. | For example, for a resonant inverse Compton-controlled vacuum gap (which is more possible in the high B regime), the typical photon energy may be expressed as $\epsilon_c({\rm ICS-VG}) \simeq 5.5\times 10^4 P^{1/14} {B'}\xi^{-3/7}$ , and the photon splitting dominant line is (Zhang Harding 2001) B_p (P/ 1 ), or $\dot P\geq 2.44\times 10^{-12} (P/ 1 {\rm s}) ^{-69/49}
\xi^{8/49}$ by adopting $B_p=6.4\times 10^{19} {\rm G} \sqrt{P\dot
P}$. |
This is also the deathBue for the anti-parallel rotators in which vacuuuctyvpe ΠΙΟΣ gaps are expected. | This is also the deathline for the anti-parallel rotators in which vacuum-type inner gaps are expected. |
For parallel rotators. a simular Lue for the space-charec-linitedtlow type inner gap max be obtained after detailed umuerical calculations (Warding et al. | For parallel rotators, a similar line for the space-charge-limited-flow type inner gap may be obtained after detailed numerical calculations (Harding et al. |
2001. in preparation). | 2001, in preparation). |
This is a Lue to define whether delaved pair formation is necessary. rather than a deathline. which is then defined according to the binding condition at the surface (Zhang Tarding 2000). | This is a line to define whether delayed pair formation is necessary, rather than a deathline, which is then defined according to the binding condition at the surface (Zhang Harding 2000). |
We lave shown that if ANPs/SCRs ie indecd iiaesnetars. even if photon splittius could completely suppress one-photon pair production. 5-rayvs in the magnetar nmaeguetosphere may still eenerate clectron-positron pairs via two-photom pair production. mainly because the ANP/SCR euvironmnenuts are very hot. | We have shown that if AXPs/SGRs are indeed magnetars, even if photon splitting could completely suppress one-photon pair production, $\gamma$ -rays in the magnetar magnetosphere may still generate electron-positron pairs via two-photon pair production, mainly because the AXP/SGR environments are very hot. |
Nou-detection of radio ciission from ANPs/SCRs max be because the low energvpairs generate radio enission withtoo low a frequency to be observable iu the bands above several hundred MITIz. | Non-detection of radio emission from AXPs/SGRs may be because the low energypairs generate radio emission withtoo low a frequency to be observable in the bands above several hundred MHz. |
Searching for low frequency Cluission frou these objects is of ereat interest. and if detected. the low frequency cussion will rule out the | Searching for low frequency emission from these objects is of great interest, and if detected, the low frequency emission will rule out the |
(Satoetal.2007:Johusonct (Butleretal.2006b:2005Endl2006). | \citep{sat07, joh07}, \citep{but06, end06}, \citep{pau06, set07}. |
. 100 (e.JugeJohusIiull2007) (c.e..Bouvieretal.2007:ITuertac | $>$ \citep[e.g.,][]{cmj07}
\citep{saa97}
\citep[e.g.,][]{bou07, hue08}. |
t2008).. Paulsonetal.(2006) »1 Ανν c Setiawanotal.(2007) ju, Setiawanetal.ond(2008). 10 Mu Unelamoetal. | \citet{pau06} $-$ $>$ $-$ $_{Jup}$ $\sigma$ \citet{set07} $_{Jup}$ $-$ \citet{set08} $\sim$ $_{Jup}$ \citet{huel08} \citep{tul95}. |
(2008). (Tullet:al.1995). 1.2" R~GO.000. ~LSO0 2 " ~ i2 i IHucrta(20 | $''$ $\sim$ $\sim$ $\sim$ $''$ $\sim$ $^{-1}$\citep{nid02, but96, cum99} $^{-1}$ \citet{hue07} \citet{hue08}. |
07) | $-$ |
references therein. | references therein. |
The code has been used to investigate the possible binary progenitors of gamma ray bursts, and also to the rates and properties of binary gravitational wave sources (??).. ? | The code has been used to investigate the possible binary progenitors of gamma ray bursts, and also to the rates and properties of binary gravitational wave sources \citep{2003ApJ...589L..37B,2004A&A...415..407B}. . |
presented a calculation of the expected mass spectrum of merging compact object binaries including the double neutron stars. | \citet{2004MNRAS.352.1372B} presented a calculation of the expected mass spectrum of merging compact object binaries including the double neutron stars. |
They hinted at a possibility of detecting a large number of non equal mass NS-NS binaries, and showed that they may not show up in the radio sample due to selection effects (?). | They hinted at a possibility of detecting a large number of non equal mass NS-NS binaries, and showed that they may not show up in the radio sample due to selection effects \citep{2002ApJ...572..407B} . |
In this paper we present the merger of two models: binary population synthesis and pulsar evolution models and compare them with the available data. | In this paper we present the merger of two models: binary population synthesis and pulsar evolution models and compare them with the available data. |
In section 2 we present the model of the binary evolution and the pulsar evolution we use. | In section 2 we present the model of the binary evolution and the pulsar evolution we use. |
Section 3 contains the results and comparison with the observations. | Section 3 contains the results and comparison with the observations. |
Section 4 is devoted to comparison of the radio and gravitational wave selection effects, while Section 5 contains the conclusions. | Section 4 is devoted to comparison of the radio and gravitational wave selection effects, while Section 5 contains the conclusions. |
There are currently 10 known pulsars in nine NS-NS binaries. | There are currently 10 known pulsars in nine NS-NS binaries. |
In one case, both stars in the binary are visible as radio pulsars (J0737-3039). | In one case, both stars in the binary are visible as radio pulsars (J0737-3039). |
At Fig. | At Fig. |
1 we present the spin period -spin period derivative diagram (P— P) for this binaries. | \ref{pulsobs} we present the spin period -spin period derivative diagram $P-\dot{P}$ ) for this binaries. |
Most of the objects are concentrated in the millisecond pulsar region. | Most of the objects are concentrated in the millisecond pulsar region. |
However there are two noticeable outliers: J07373-3039B, the companion of the J07373-3039A in the binary pulsars system; and J1906-0746, a young pulsar with a likely neutron star companion. | However there are two noticeable outliers: J07373-3039B, the companion of the J07373-3039A in the binary pulsars system; and J1906-0746, a young pulsar with a likely neutron star companion. |
Detailed properties of those systems can be found in Tab. | Detailed properties of those systems can be found in Tab. |
1 (?,,? and ?)). | \ref{pulsobswlas}
\citet{2004Sci...304..547S}, \citet{2005ApJ...618L.119F} and \citet{2008AIPC..983..485K}) ). |
The binaries in Tab. | The binaries in Tab. |
1 are ordered according to the merger time (time remaining to the coalescence). | \ref{pulsobswlas} are ordered according to the merger time (time remaining to the coalescence). |
Three binaries at the bottom of the table have merger times much longer than 10Gyr. | Three binaries at the bottom of the table have merger times much longer than $10\;{\rm Gyr}$. |
We list the spin period and its derivative, the masses of the neutron star, as well as the present orbital parameters. | We list the spin period and its derivative, the masses of the neutron star, as well as the present orbital parameters. |
The distribution of the orbital periods spreads over two orders of magnitude, with no clear evidence of clustering in these range, and the orbits of all systems are significantly eccentric. | The distribution of the orbital periods spreads over two orders of magnitude, with no clear evidence of clustering in these range, and the orbits of all systems are significantly eccentric. |
The masses of the seven objects with the merger times below 10 Gyrs are determined very well and are in the range between 1.25 and 1.44 Mo. | The masses of the seven objects with the merger times below 10 Gyrs are determined very well and are in the range between 1.25 and 1.44 $_{\odot}$. |
The masses of the neutron star in the remaining three binaries are not so well constrained and may even lie outside this range. | The masses of the neutron star in the remaining three binaries are not so well constrained and may even lie outside this range. |
Due to the selection effects observed sample of NS-NS binaries might differ from the intrinsic population. | Due to the selection effects observed sample of NS-NS binaries might differ from the intrinsic population. |
There is also a possibility that the population observed in radio will have different properties that the one observed in gravitational waves. | There is also a possibility that the population observed in radio will have different properties that the one observed in gravitational waves. |
The most obvious selection effect in radio is drop of observed flux proportionally to distance squared from the observer to the pulsar. | The most obvious selection effect in radio is drop of observed flux proportionally to distance squared from the observer to the pulsar. |
Our radio-telescopes have limited sensitivity so we will not see the weakest and/or farthest pulsars. | Our radio-telescopes have limited sensitivity so we will not see the weakest and/or farthest pulsars. |
The binary pulsar population observed in radio is restricted so far to our Galaxy. | The binary pulsar population observed in radio is restricted so far to our Galaxy. |
In this paper we assume that all pulsars with observed flux below ~1mJy at 400MHz are not detectable in radio. | In this paper we assume that all pulsars with observed flux below $\sim
1\;{\rm mJy}$ at $400\;{\rm MHz}$ are not detectable in radio. |
We also try to incorporate in our calculations a second important selection effect directly connected with the fact that the pulsars considered are in binaries. | We also try to incorporate in our calculations a second important selection effect directly connected with the fact that the pulsars considered are in binaries. |
According to ?,, while searching pulsars with stack search and phase modulation methods it is difficult to detect radio emission from neutron stars which are in binaries with orbital period within range [0.3h;4h]. In general it is due to the fact that signal-to-noise ratio drops drastically due to variation of pulsar period in during the observation. | According to \citet{Faulkner}, while searching pulsars with stack search and phase modulation methods it is difficult to detect radio emission from neutron stars which are in binaries with orbital period within range $[0.3\;{\rm h};4\;{\rm h}].$ In general it is due to the fact that signal-to-noise ratio drops drastically due to variation of pulsar period in during the observation. |
There are of course different methods of detecting radio pulsars which are sensitive for pulsars of such properties. | There are of course different methods of detecting radio pulsars which are sensitive for pulsars of such properties. |
Furthermore this orbital period range was calculated for a pulsar with the spin period P—9ms (?),, and it may vary as a function of P. However, our aim is to see how this selection effect influences the observed properties of the population. | Furthermore this orbital period range was calculated for a pulsar with the spin period $P=9\;{\rm ms}$ \citep{Faulkner}, and it may vary as a function of $P.$ However, our aim is to see how this selection effect influences the observed properties of the population. |
Binaries detectable in gravitational waves will be those with relatively small orbital separation. | Binaries detectable in gravitational waves will be those with relatively small orbital separation. |
Namely, we assume that binaries with coalescence time longer than 10!? years are not visible in gravitational waves. | Namely, we assume that binaries with coalescence time longer than $10^{10}$ years are not visible in gravitational waves. |
In both ? and ? the authors suggest that there will be differences between the two observed populations. | In both \citet{2004MNRAS.352.1372B} and \citet{2005MmSAI..76..513G} the authors suggest that there will be differences between the two observed populations. |
In particular they reached the result that the mass ratio distributions will not be the same. | In particular they reached the result that the mass ratio distributions will not be the same. |
The population modelling in those papers was quite crude. | The population modelling in those papers was quite crude. |
The binaries with lifetimes above the Hubble time were neglected, and it has been assumed that observability of pulsars is proportional to its lifetime. | The binaries with lifetimes above the Hubble time were neglected, and it has been assumed that observability of pulsars is proportional to its lifetime. |
Such difference between the observed population in the radio in gravitational waves would be quite important as we need to know how the gravitational waves signal looks like to detect it. | Such difference between the observed population in the radio in gravitational waves would be quite important as we need to know how the gravitational waves signal looks like to detect it. |
In order to know this signal we need the initial parameters for numerical relativity calculations. | In order to know this signal we need the initial parameters for numerical relativity calculations. |
Until recently most calculations of coalescing neutron star binaries and black hole binaries assumed that mass ratio q is close to unity (e.g. ?,, ?)). | Until recently most calculations of coalescing neutron star binaries and black hole binaries assumed that mass ratio $q$ is close to unity (e.g. \citet{2007AdSpR..39..271G}, , \citet{2006gr.qc....10122B}) ). |
In this paper we present carefully model the binary pulsar population withthe updated stellar evolution code (?) combined with the detailed pulsar evolution model. | In this paper we present carefully model the binary pulsar population withthe updated stellar evolution code \citep{2008ApJS..174..223B} combined with the detailed pulsar evolution model. |
1961. Araudo ct al. | 1964, Araudo et al. |
2007). | 2007). |
We considered that the cucrey density of relativistic partickss πακος Έπος contributions (addeuds): where el. p. and 62 stand for relativistic primary electrons. protous. aux SCCOLLCary electron-positron pairs (i.c. pairs conine from charexc pion decays). respectively. aud is the ummber desity. | We considered that the energy density of relativistic particles makes three contributions (addends): where $e1$ , $p$ , and $e2$ stand for relativistic primary electrons, protons, and secondary electron-positron pairs (i.e. pairs coming from charged pion decays), respectively, and $n$ is the number density. |
T1ο relation between primary clectrous and protons d::5 4au... With a>0. | The relation between primary electrons and protons is $u_{p}=a u_{e1}$, with $a\geq0$. |
Three Cases were considered: just electrous). ¢=1 (equal enerev deusitv iu οκtL Species). and «100 (protou-dominated case. as Oserved in the galactic cosnic ravs). | Three cases were considered: $a=0$ (just electrons), $a=1$ (equal energy density in both species), and $a=100$ (proton-dominated case, as observed in the galactic cosmic rays). |
The order of maguitide of he equipartition magnetic Ποια led to B~5«N10 oq (three times this value for the case a Loo). | The order of magnitude of the equipartition magnetic field led to $B\sim 5 \times 10^{-5}$ G (three times this value for the case $a=100$ ). |
T1ο παπα value for t10 CLCOYSV of the particCS Was cternuned through the balance of enerev ean and losses. | The maximum value for the energy of the particles was determined through the balance of energy gain and losses. |
Different loss mechauisus were considered: svuchrotrou radiution. muverse Conrxton (IC) scattering of IR. stelar. cosmüc microwave backerounud photons. reativistic Dreimisstralhluug. aud particle escape frou the radiation region cne to convection by the stellar wind. | Different loss mechanisms were considered: synchrotron radiation, inverse Compton (IC) scattering of IR, stellar, cosmic microwave background photons, relativistic Bremsstrahlung, and particle escape from the radiation region due to convection by the stellar wind. |
Iu the case of protoris. the only relevant losses are proton-protou (pp) meastic collisions and couvective escape. | In the case of protons, the only relevant losses are proton-proton $pp$ ) inelastic collisions and convective escape. |
Diffusion is negligible In coniparison to convection in this situation. | Diffusion is negligible in comparison to convection in this situation. |
Both primary clectrous and ootons reach energies up to ~Lob! oV. which is imposed by uonradiative losses. except ora=LOO. where svucirotron losses dominate for electrons. | Both primary electrons and protons reach energies up to $\sim 10^{13}$ eV, which is imposed by nonradiative losses, except for $a=100$, where synchrotron losses dominate for electrons. |
Iu Fig. | In Fig. |
L we show the losses for electrons iu the case a=1l. | \ref{elosses} we show the losses for electrons in the case $a=1$. |
Values of magneic field and maxinnun energies obtained for electrous aud Yotons are given in Tab el. | Values of magnetic field and maximum energies obtained for electrons and protons are given in Table 1. |
The presence of highly relativistic partices dn a dense imediuu with high photon density can result iu the cficicut eeicration of eamuna-ravs. | The presence of highly relativistic particles in a dense medium with high photon density can result in the efficient generation of gamma-rays. |
Although protous cau be effectively accelerated up to the highest energies only in the shocked wind. where he deusitv is low. they can diffuse or be couvected 1pstreal up to tio region with the swept material aud densities of n~LOO ὃν | Although protons can be effectively accelerated up to the highest energies only in the shocked wind, where the density is low, they can diffuse or be convected upstream up to the region with the swept material and densities of $n\sim 100$ $^{-3}$. |
The COLYYCSDOLKling ο ¢wissivity cau be calculaxd using the delta-functional apxoxination (c.g, Aliaroniau Atovan 2000. απο ο al. | The corresponding gamma-ray emissivity can be calculated using the delta-functional approximation (e.g. Aharonian Atoyan 2000, Kelner et al. |
2006). | 2006). |
For the case &=I. the total huuinosity from pp iteractions is simular to what is Obtained from relativisic Bremesstrabhme of electrons. since the cross sections are siuular. | For the case $a=1$, the total luminosity from $pp$ interactions is similar to what is obtained from relativistic Bremsstrahlung of electrons, since the cross sections are similar. |
Iu Figure 5 we row the spectral energv distribution obtained for the case a=1. with all conutribitions included. (svuchrotrou pAelf-Conipton losses are ucellegible). | In Figure \ref{SED-1} we show the spectral energy distribution obtained for the case $a=1$, with all contributions included (synchrotron self-Compton losses are negligible). |
It can be seen that oei this case the inverse Conptonu up-scattering of IR photons is the major contribution at high euergies. with a peak around 100 GeV. The detectability of the source bvi istriuenuts like the Fermi x-ray observatory LAT will depend ou the actual particle density aud the contribution relaed to the secondary electrous at large e. | It can be seen that in this case the inverse Compton up-scattering of IR photons is the major contribution at high energies, with a peak around 100 GeV. The detectability of the source by instruments like the Fermi $\gamma$ -ray observatory LAT will depend on the actual particle density and the contribution related to the secondary electrons at large $a$. |
Detailed calculations for a set of main parameters will be eiveu elsewhere. | Detailed calculations for a set of main parameters will be given elsewhere. |
The pp coutribution extends well iuto the TeV reeluc. but it is weaker aud will be difficult to detect with erOlud-based. Cherenkov telescope arrays like VERITAS or MAGIC IL | The $pp$ contribution extends well into the TeV regime, but it is weaker and will be difficult to detect with ground-based Cherenkov telescope arrays like VERITAS or MAGIC II. |
Iu contrast. if the relativistic particle contain is protou-«oninated (o6=100). gamma-rays fron neutral pion decay* dominate the high-enerew spectrum. | In contrast, if the relativistic particle contain is proton-dominated $a=100$ ), gamma-rays from neutral pion decay dominate the high-energy spectrum. |
The planned CTA North observatory might detect t1C source. easily vielΠιο nuformation on the cutoff at high enereies. | The planned CTA North observatory might detect the source, easily yielding imformation on the cutoff at high energies. |
Olservations of the spectral slope in this regime ciun be use to idenifv the proton content throieh t DIuuinositv level axd the proton spectral mdex. since is preserved iu the corresponding photon imdex. | Observations of the spectral slope in this regime can be used to identify the proton content through the luminosity level and the proton spectral index, since it is preserved in the corresponding photon index. |
Radi polarization data wi] provide additional iufoiiialol ¢ the maenetic field. | Radio polarization data will provide additional information on the magnetic field. |
οbservatious of 1373651 with X-rav observaorjes like aud call very useful for deterniuσαre the cutoff of the svuciota 2οςπια. which Is directly relatedtothe Wand ierev of the electrons. | Observations of $^{\circ}\,3654$ with X-ray observatories like and can be very useful for determining the cutoff of the synchrotron spectrum, which is directly relatedtothe maximum energy of the electrons. |
Tus. in turn. would vield valualle information on the actia value of t1e magnetic field axd +je correctness of the equixwtitiou hivpothesis. | This, in turn, would yield valuable information on the actual value of the magnetic field and the correctness of the equipartition hypothesis. |
High-mass stars are usually defined as those exceeding ~8M... based on the fact that stars above this mass limit do not have a pre-main sequence phase (Palla Stahler 1993)). | High-mass stars are usually defined as those exceeding $\sim$ 8, based on the fact that stars above this mass limit do not have a pre-main sequence phase (Palla Stahler \cite{past}) ). |
This means that accretion is ongoing until the star ignites hydroge burning and reaches the zero-age nain sequence (ZAMS). | This means that accretion is ongoing until the star ignites hydrogen burning and reaches the zero-age main sequence (ZAMS). |
At this point the strong radiation pressure may halt and eve reverse infall and thus stop further growth of the stellar mass. | At this point the strong radiation pressure may halt and even reverse infall and thus stop further growth of the stellar mass. |
This led to the so-called "radiation pressure problem". | This led to the so-called “radiation pressure problem”. |
Recent studies have demonstrated that this limitation holds only 1 spherical symmetry. | Recent studies have demonstrated that this limitation holds only in spherical symmetry. |
As first envisionedby Nakano (1987)) and recently demonstrated by Krumholz et al. (2009)) | As first envisionedby Nakano \cite{nakano}) ) and recently demonstrated by Krumholz et al. \cite{krum}) ) |
anc Kuiper et al. (2010)). | and Kuiper et al. \cite{kuip}) ), |
accretion through a circumstellar disk can explain the formation of stars up to the upper limit of the initial. mass function. by allowing part of the photons to escape along the disk axis and boosting the ram pressure of the accreting gas through the small disk solid angle. | accretion through a circumstellar disk can explain the formation of stars up to the upper limit of the initial mass function, by allowing part of the photons to escape along the disk axis and boosting the ram pressure of the accreting gas through the small disk solid angle. |
It also appears that the powerful ionizing fluxes from these OB-type stars are not sufficient to destroy the disk. which eventually turns into an ionized. rotating accretion flow close to the star (Sollins et al. 2005:: | It also appears that the powerful ionizing fluxes from these OB-type stars are not sufficient to destroy the disk, which eventually turns into an ionized, rotating accretion flow close to the star (Sollins et al. \cite{sollins}; |
Keto 20071). | Keto \cite{keto07}) ). |
For these reasons it seems established that circumstellar accretion disks play a crucial role in the formation ofa// stars and not only solar-type stars. | For these reasons it seems established that circumstellar accretion disks play a crucial role in the formation of stars and not only solar-type stars. |
This theoretical result contrasts with the limited observational evidence of disks in high-mass (proto)stars. | This theoretical result contrasts with the limited observational evidence of disks in high-mass (proto)stars. |
Only in recent years the number of disk candidates associated with luminous young stellar objects (YSOs) has significantly increased. mostly owing to the improvement of (subymillimeter interferometers in terms of angular resolution and sensitivity. | Only in recent years the number of disk candidates associated with luminous young stellar objects (YSOs) has significantly increased, mostly owing to the improvement of (sub)millimeter interferometers in terms of angular resolution and sensitivity. |
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