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The latter has ev¢lical (imine residuals. as is shown in Figure 3 of Hobbsetal.(2010).. that can imply ev¢lical changes in the rotation [requency. | The latter has cyclical timing residuals, as is shown in Figure 3 of \citet{hob10}, that can imply cyclical changes in the rotation frequency. |
It is thought Chat pulsar elitches reflect a variable coupling between the solid crust of a neutron star and (he superthiud interior rotating more rapidly than the solid crust. | It is thought that pulsar glitches reflect a variable coupling between the solid crust of a neutron star and the superfluid interior rotating more rapidly than the solid crust. |
In terms of vortex pinning moclels. (he origin of elitehes can be explained bx the catastrophic unpinning of superfluid vortices (Anderson&Hoh1975:Alparetal.1984.1959.1993: 1985). | In terms of vortex pinning models, the origin of glitches can be explained by the catastrophic unpinning of superfluid vortices \citep{and75,alp84,alp89,alp93,pin85}. |
. These models provide a satisfactory interpretation of large elitches in pulsars. | These models provide a satisfactory interpretation of large glitches in pulsars. |
Now the interpretation of a phenomenon of slow glitches is uncertain. | Now the interpretation of a phenomenon of slow glitches is uncertain. |
According to llobbsetal.(2010).. the slow glitch. phenomenon is a quasi-periodic component of timing noise. unrelated (ο normal elitches. | According to \citet{hob10}, the slow glitch phenomenon is a quasi-periodic component of timing noise, unrelated to normal glitches. |
At the same time. the slow elitches revealed in PSR D1642—03 possess the properties which meet (he requirements of the elitch models: 1) all the slow elitehes observed. have significant exponential decay after elitch which is characterized bv the same parameter (~0.9. 2) (he size ol elitches ancl (he time interval to the following elitch obev a strong Lnear relation. | At the same time, the slow glitches revealed in PSR $-$ 03 possess the properties which meet the requirements of the glitch models: 1) all the slow glitches observed have significant exponential decay after glitch which is characterized by the same parameter $Q\sim0.9$, 2) the size of glitches and the time interval to the following glitch obey a strong linear relation. |
The third property. (he presence of a modulation process which forces the glitch: amplitides and Che inter-elitch intervals to change wilh a discrete step. was not considered vel by any theory of pulsar elitches. | The third property, the presence of a modulation process which forces the glitch amplitudes and the inter-glitch intervals to change with a discrete step, was not considered yet by any theory of pulsar glitches. |
In the case of slow elitches. il is necessary to account for the cause of a continuous generation of slow glitches. | In the case of slow glitches, it is necessary to account for the cause of a continuous generation of slow glitches. |
For PSR BOOL9+06. a sawtooth-like modulation of the rotation frequency will a period οἱ 600 davs could be interpreted by the fee precession of an isolated pulsar if this modulation was accompanied by correlated observable changes in (he pulse profile (Shaham1977:Nelsonetal.1990:Cordes.1993:Stairsοἱ 2000). | For PSR B0919+06, a sawtooth-like modulation of the rotation frequency with a period of 600 days could be interpreted by the free precession of an isolated pulsar if this modulation was accompanied by correlated observable changes in the pulse profile \citep{sha77,nel90,cor93,sta00}. |
. We detected no pulse profile changes curing our observations al 112 MlIz. | We detected no pulse profile changes during our observations at 112 MHz. |
A precession model requires a strictly periodie modulation of the pulsar's rotation Ireequency. | A precession model requires a strictly periodic modulation of the pulsar's rotation frequency. |
. The presence of the phase shift. equal to 400 days. in the modulation curve B between cveles 9 and 10. as is shown in Figure 2((6). contradicts this requirement. | The presence of the phase shift, equal to 400 days, in the modulation curve B between cycles 9 and 10, as is shown in Figure \ref{resid1}( (c), contradicts this requirement. |
These (wo arguments testify against the interpretation of this sawtooth-like modulation in (he terms of a free precession model. | These two arguments testify against the interpretation of this sawtooth-like modulation in the terms of a free precession model. |
Besides. (his pulsar has experienced a lavee gliteh. | Besides, this pulsar has experienced a large glitch. |
As discussed by Link(2007).. a slowly precessing neutron star cannot produce a glitch. | As discussed by \citet{lin07}, a slowly precessing neutron star cannot produce a glitch. |
The useful information for explanation of anorigin of slow glitches can be obtained | The useful information for explanation of anorigin of slow glitches can be obtained |
index of fluctuations n.. the scalar fluctuation amplitude AZ. and the optical reionization depth 7 (J.DunklevandWright2009:D.Larsonaud2011). | index of fluctuations $n_s$, the scalar fluctuation amplitude $\Delta_\cR^2$, and the optical reionization depth $\tau$ \citep{Dunkley2009,Larson2011}. |
. The model assumes a nearly Gaussian spectrum of initial fluctuations with statistical isotropy over (he sky (PartridgeandWilkinson1967).. | The model assumes a nearly Gaussian spectrum of initial fluctuations with statistical isotropy over the sky \citep{PW1967CMBR}. |
Any deviation from isotropy would be a challenge for a new plwsies(Elstathiou2003:Teemarkοἱal.2003). | Any deviation from isotropy would be a challenge for a new \citep{Efstathiou2003,Tegmark2003}. |
. Although upto the recent WMADP data the ACDM model remains a good fit of the observed microwave sky (E.IXomatsuandWright2011).. the Κον feature displaved by WMADP (C.L.BennettandWright2011) and previous experiments (Bennettetal.al.1992:Strukovet 1993).. is that the Universe is isotropic in the mean. but anisotropic in correlations. | Although upto the recent WMAP data the $\Lambda$ CDM model remains a good fit of the observed microwave sky \citep{Komatsu2011}, the key feature displayed by WMAP \citep{BennetEtAl2011} and previous experiments \citep{Bennet1992,Smoot1992,strukov1993anisotropy}, is that the Universe is isotropic in the mean, but anisotropic in correlations. |
This means there are no preferable directions. but there are preferable angles of correlations. | This means there are no preferable directions, but there are preferable angles of correlations. |
The observation of such anisotropy is contr-intuitive Iron classical physics point of view. but seems quite natural in quantum mechanics. like that of (he Einstein-Podolsky- (EPR) correlations (Einsteinetal.1935). | The observation of such anisotropy is contr-intuitive from classical physics point of view, but seems quite natural in quantum mechanics, like that of the Einstein-Podolsky-Rosen (EPR) correlations \citep{EPR1935}. |
. The anisotropy. in angle correlations has been receiving constant attention since discovered (Smootοἱal.1992).. | The anisotropy in angle correlations has been receiving constant attention since discovered \citep{Smoot1992}. |
WMAP mission itself was designed to use the observed correlations of fluctuations to put narrower constraints on cosmological parameters after those obtained by (he previous COsmic Dackeround Explore (CODE) mission (Dennettetal.1996.2003). | WMAP mission itself was designed to use the observed correlations of fluctuations to put narrower constraints on cosmological parameters after those obtained by the previous COsmic Background Explore (COBE) mission \citep{Bennet1992,Bennet2003}. |
. Two main tvpes of data are available to constrain the cosmological parameters: the observed correlations in the distribution of galaxies (PeeblesandGroth1975): and the observed. correlations of the relic radiation (C.L.BennettanclWright2011). | Two main types of data are available to constrain the cosmological parameters: the observed correlations in the distribution of galaxies \citep{PG1975}; and the observed correlations of the relic radiation \citep{BennetEtAl2011}. |
. The main instrument for the analvsis of both data tvpe remains the decomposition with respect to the representations of the SO(3) group of rotations in IE*. | The main instrument for the analysis of both data type remains the decomposition with respect to the representations of the $SO(3)$ group of rotations in $\R^3$. |
This means the n-point correlation function du"!...u7) transforms under the spacial rotations according to the law where A is the matrix of SO(3) rotation in appropriate representation. | This means the $n$ -point correlation function $\mean{{u}^{\beta_1}\ldots{u}^{\beta_n} }$ transforms under the spacial rotations according to the law where $\Lambda$ is the matrix of $SO(3)$ rotation in appropriate representation. |
Thus a scalar remains invariuit under rotations: the n-point correlation function of a vector field thus transforms according to the direct product ào...A.elc. | Thus a scalar remains invariant under rotations; the $n$ -point correlation function of a vector field thus transforms according to the direct product $\underbrace{\Lambda \otimes \ldots \otimes \Lambda}_{n\ times}$,etc. |
In the CODE and the WMAP data the full skv map was decomposed into a series of spherical harmonies Y5,(). which form irrecucible representations of SO(3) group in R*: where n is (he unit direction vector. | In the COBE and the WMAP data the full sky map was decomposed into a series of spherical harmonics $Y_{lm}(\vn)$ , which form irreducible representations of $SO(3)$ group in $\R^3$ : where $\vn$ is the unit direction vector. |
If the CAIBR. anisotropy is driven by a Gaussian | If the CMBR anisotropy is driven by a Gaussian |
mPhere has been determined. cllort⋅ over the past several vears (o understand. the history. of. luminous. matter in. the Universe. | There has been determined effort over the past several years to understand the history of luminous matter in the Universe. |
u n.Ultimately. one wishes: to have a consistent. understanding. which. would tic. together the detailed. physical processes at work in stars ancl ISM in the Milky Wav and local galaxies with the integrated properties of⋅ more distant. svstemis. | Ultimately, one wishes to have a consistent understanding which would tie together the detailed physical processes at work in stars and ISM in the Milky Way and local galaxies with the integrated properties of more distant systems. |
. .“Phe spectral properties. and energy budget of. the distant.. galaxies 2.in turn are crucial. in. understanding. the universal. history. of. star. formation.. the very faintestD. source counts. and the extragalactic. background radiation. | The spectral properties and energy budget of the distant galaxies in turn are crucial in understanding the universal history of star formation, the very faintest source counts, and the extragalactic background radiation. |
In particular. the infrared and. sub-nim regimes have | In particular, the infrared and sub-mm regimes have |
a peak in the counts at the positions taugcutial to the line of sight. | a peak in the counts at the positions tangential to the line of sight. |
If the ring is relatively thin. then away frou. the tangential points the counts should die away quickly. | If the ring is relatively thin, then away from the tangential points the counts should die away quickly. |
If the vine is verv broad (e.g. ax proposed by Iseut et al. | If the ring is very broad (e.g. as proposed by Kent et al. |
1991). then the peak becomes a lot broader. | 1991), then the peak becomes a lot broader. |
Figure G shows some possible forms that the ring counts could produce for au axially svuuuetric svstem. | Figure \ref{Fig:ring} shows some possible forms that the ring counts could produce for an axially symmetric system. |
If the rine were elliptical then the longitudes of the peaks would no ouecr be sviuuetrie: however. the shapes of the peaks would remain basically unaltered. | If the ring were elliptical then the longitudes of the peaks would no longer be symmetric; however, the shapes of the peaks would remain basically unaltered. |
Clearly. it is possible hat a contrived ring (particularly a patchy rine). couple with a lLiehly improbable cistyibution of extinction. couk reproduce the form of the in-plane counts. | Clearly, it is possible that a contrived ring (particularly a patchy ring), coupled with a highly improbable distribution of extinction, could reproduce the form of the in-plane counts. |
However. as wil ve shown. this distribution of extinction does not exist ane rence the rine v dtself camnot explain the in-plane counts. | However, as will be shown, this distribution of extinction does not exist and hence the ring by itself cannot explain the in-plane counts. |
The other acruative to explain the asvuumetry in the star counts is re existence of a bar. | The other alternative to explain the asymmetry in the star counts is the existence of a bar. |
The work in II9 ancl 00 limits the possible orieutatious of he bar. | The work in H94 and H00 limits the possible orientations of the bar. |
The predicted bar has the near end at /=27° at a clistance of 5.7 ο aud the far eux at |=12" at a distance of about 11 spe. | The predicted bar has the near end at $l=27^\circ$ at a distance of 5.7 kpc and the far end at $l=-12^\circ$ at a distance of about 11 kpc. |
We shall limit the discussion to whether the data presented here are consistent with the previous results rather than try o re-determimne all the parameters. as the arguments preseutec would be almost identical. | We shall limit the discussion to whether the data presented here are consistent with the previous results rather than try to re-determine all the parameters, as the arguments presented would be almost identical. |
Whilst simplistically oue would expect the near eud of he bar to give more counts than the far end. this is not recessarily the case in the plane. | Whilst simplistically one would expect the near end of the bar to give more counts than the far end, this is not necessarily the case in the plane. |
Blitz Sperecl (19915) showed that the far cud of the bar cau eive a higher surface xiehtuess than the rear end in the plane: Unavaue ct al. ( | Blitz Spergel (1991b) showed that the far end of the bar can give a higher surface brightness than the near end in the plane; Unavane et al. ( |
1998) and Unavane Cülimiore (1998) predict a similar result for star counts. | 1998) and Unavane Gilmore (1998) predict a similar result for star counts. |
There are. however. a series of effects hat have to be taken iuto account: Figure 7? shows the measured and model counts distribution for three latitude ranges. | There are, however, a series of effects that have to be taken into account: Figure \ref{Fig:model} shows the measured and model counts distribution for three latitude ranges. |
The mocel coutaius the truncated disc. but with the bulee and extinction based on Wainscoat et al. ( | The model contains the truncated disc, but with the bulge and extinction based on Wainscoat et al. ( |
1992). | 1992). |
The simplest possible bar distribution iu aerecineut with IIO0 was then added. | The simplest possible bar distribution in agreement with H00 was then added. |
some additional component in the UV making up of the flux, but this suggestion is entirely dependent on the absolute flux calibrations of the SDSS spectrum and the GALEX fluxes. | some additional component in the UV making up of the flux, but this suggestion is entirely dependent on the absolute flux calibrations of the SDSS spectrum and the GALEX fluxes. |
In their study of the SDSS CV population as a whole, Gánsickeal.(2009) showed that a characterstic feature of this sample is an accumulation of intrinsically faint CVs in the period range mmin. | In their study of the SDSS CV population as a whole, \citet{Gansicke+09mn} showed that a characterstic feature of this sample is an accumulation of intrinsically faint CVs in the period range min. |
The vast majority of these systems have optical spectra dominated by the emission from their white dwarfs, with no noticeable contribution from the secondary stars, implying extremely late secondary spectral types. | The vast majority of these systems have optical spectra dominated by the emission from their white dwarfs, with no noticeable contribution from the secondary stars, implying extremely late secondary spectral types. |
These findings are consistent with theoretical predictions that the companion stars are whittled down as the CVs evolve towards | These findings are consistent with theoretical predictions that the companion stars are whittled down as the CVs evolve towards |
galaxies are common in cluster environments. but emission—line AGN could plausibly have a range of host galaxy types. | galaxies are common in cluster environments, but emission–line AGN could plausibly have a range of host galaxy types. |
However. as Figure 5. shows. there are no significant differences between the emission and absorption-line AGN in terms of their radio and optical luminosities. suggesting that there are no fundamental differences in host galaxies. | However, as Figure \ref{absrad} shows, there are no significant differences between the emission and absorption–line AGN in terms of their radio and optical luminosities, suggesting that there are no fundamental differences in host galaxies. |
Table provides mean values for various properties of these AGN: there is only a 0.24 magnitude Ss 2a) difference in mean host galaxy absolute magnitude between the Aa and Ae classes. whilst the difference in mean local galaxy surface density is at the <5c level. | Table \ref{tabprop} provides mean values for various properties of these AGN: there is only a 0.24 magnitude $\lta
2\sigma$ ) difference in mean host galaxy absolute magnitude between the Aa and Ae classes, whilst the difference in mean local galaxy surface density is at the $\gta 5\sigma$ level. |
Possible explanations for the differences between the emission and absorption line AGN. and the consequences of these results. are discussed in Section 5.3.. | Possible explanations for the differences between the emission and absorption line AGN, and the consequences of these results, are discussed in Section \ref{absemis}. |
Comparing the properties of the AGN with those of the elliptical galaxy sample selected from the SDSS data (Table 13). the mean local surface densities are similar. | Comparing the properties of the AGN with those of the elliptical galaxy sample selected from the SDSS data (Table \ref{tabprop}) ), the mean local surface densities are similar. |
Figure 6 shows the fraction of all elliptical galaxies (statistically constructed from the morphological mix as a function of environment derived from the SDSS subsample) that host absorption—line AGN as a function of galaxy surface density. | Figure \ref{agnellip} shows the fraction of all elliptical galaxies (statistically constructed from the morphological mix as a function of environment derived from the SDSS subsample) that host absorption–line AGN as a function of galaxy surface density. |
Because of the prevalence of early—type galaxies in clusters this distribution is flattened from the equivalent fraction in all galaxies. but it is still evident that absorption—line AGN have a weak preference for richer environments. even compared to elliptical galaxies in general. | Because of the prevalence of early–type galaxies in clusters this distribution is flattened from the equivalent fraction in all galaxies, but it is still evident that absorption--line AGN have a weak preference for richer environments, even compared to elliptical galaxies in general. |
The fractions of galaxies that host radio-loud AGN activity or radio—selected star formation activity are shown in Figure 7 as a 'unction of the size of group or cluster in which that galaxy lies (as determined by the friends—of—friends analysis: see Section 3.2)). | The fractions of galaxies that host radio–loud AGN activity or radio–selected star formation activity are shown in Figure \ref{sfgpsize}
as a function of the size of group or cluster in which that galaxy lies (as determined by the friends–of–friends analysis; see Section \ref{gpclus}) ). |
Star forming galaxies prefer to avoid richer environments. although his dependence is clearly weaker than that on local environment ound above. | Star forming galaxies prefer to avoid richer environments, although this dependence is clearly weaker than that on local environment found above. |
Star forming galaxies are also relatively rare in isolated galaxies: this latter result is most likely related to the radio selection. which picks out only galaxies with relatively high star ormation rates (2SAL. /yr). | Star forming galaxies are also relatively rare in isolated galaxies: this latter result is most likely related to the radio selection, which picks out only galaxies with relatively high star formation rates $\gta 5
M_{\odot}$ /yr). |
AGN are most common in moderate groups and poor clusters. | AGN are most common in moderate groups and poor clusters. |
The fraction of galaxies with AGN activity is lower in relatively isolated environments (1—3 galaxies) and also in the richest clusters €50 galaxies): a KS test gives a probability of that the AGN are not simply randomly drawn from all galaxies. | The fraction of galaxies with AGN activity is lower in relatively isolated environments (1–3 galaxies) and also in the richest clusters $>50$ galaxies); a KS test gives a probability of that the AGN are not simply randomly drawn from all galaxies. |
By sub-classifying the AGN into emission-line AGN and absorption-line AGN (cf. | By sub-classifying the AGN into emission–line AGN and absorption–line AGN (cf. |
Figure 8)) it is apparent that the fall- in the AGN fraction in rich environments happens earlier for emission-line AGN than absorption-line AGN. | Figure \ref{gpsize}) ) it is apparent that the fall-off in the AGN fraction in rich environments happens earlier for emission–line AGN than absorption–line AGN. |
The former are essentially absent from all clusters of more than 20 galaxies. whilst the latter still strongly populate the moderate clusters | The former are essentially absent from all clusters of more than 20 galaxies, whilst the latter still strongly populate the moderate clusters |
the 6 star. | the 6 star. |
This causes a clip in at 8M. | This causes a dip in at 8. |
.. Superimposed on the same figure are IMES from Salpeter (1955) fine)) and WKWroupa. Tout Gilmore (1993) fine) normalized over the mass range Q.08-100 M... | Superimposed on the same figure are IMFs from Salpeter (1955) ) and Kroupa, Tout Gilmore (1993) ) normalized over the mass range 0.08-100 $_{\odot}$. |
The Salpeter LME has the form of a single slope power-law (this model uses an index of 1.31) and places a higher proportion of the mass of a stellar generation into both the lower and upper extremes of the mass distribution when compared with the Ixroupa et al. ( | The Salpeter IMF has the form of a single slope power-law (this model uses an index of 1.31) and places a higher proportion of the mass of a stellar generation into both the lower and upper extremes of the mass distribution when compared with the Kroupa et al. ( |
1993) IME. | 1993) IMF. |
Lt is apparent that the role of stars between 0.3 22 mM. > 6 is emphasised. by adopting a Kroupa ct al. | It is apparent that the role of stars between 0.3 $\ge$ $m$ $_{\odot}$ $\ge$ 6 is emphasised by adopting a Kroupa et al. |
EME: whereas the frequeney of high mass stars is increased with the Salpeter LAL. | IMF, whereas the frequency of high mass stars is increased with the Salpeter IMF. |
One can anticipate that a Ixroupa ct al. ( | One can anticipate that a Kroupa et al. ( |
1993) LM woulcl give rise to more ACB stars than a Salpeter (1955) function. leading to hiejer values of at low metallicity. | 1993) IMF would give rise to more AGB stars than a Salpeter (1955) function, leading to higher values of at low metallicity. |
At metallicities approaching solar however. the Salpeter LAL should. generate the highest ratios. since it favours the sith of massive stars when compared wit1 the Ixroupa et al. | At metallicities approaching solar however, the Salpeter IMF should generate the highest ratios, since it favours the birth of massive stars when compared with the Kroupa et al. |
law. | law. |
The influence of the LALF can be seen by comxwing the results presented in Figure 6. which were derived using the Salpeter function. with those fron: Figure 2. | The influence of the IMF can be seen by comparing the results presented in Figure 6, which were derived using the Salpeter function, with those from Figure 2. |
A ratio of ~0.3 at Fe/H = 0 15 obtained with a Salpeter LAL. | A ratio of $\sim$ 0.3 at [Fe/H] = 0 is obtained with a Salpeter IMF. |
Εις is higher than in 1¢ Ixroupa et al. | This is $\sim$ higher than in the Kroupa et al. |
case ancl over twice the solar value. | case and over twice the solar value. |
Owing to the production of fewer AGB stars. at Fe/H) = 2 in the Sal»eter case is about half the value derivec with the Ixroupa IME. | Owing to the production of fewer AGB stars, at [Fe/H] = $-$ 2 in the Salpeter case is about half the value derived with the Kroupa IMF. |
Phe Ixroupa et al. | The Kroupa et al. |
model is in better agreement with the Gay Lambert (2000) data set. | model is in better agreement with the Gay Lambert (2000) data set. |
Although the [arge values of measured. by. Yonο¢ (2003) in hisher metallicity stars could. be attained. by increasing the role of massive stars. this comes a the expense of satisiving observations a low Fefl]. | Although the large values of measured by Yong (2003) in higher metallicity stars could be attained by increasing the role of massive stars, this comes at the expense of satisfying observations at low [Fe/H]. |
Multi-xcomponent IMPs with steeper slopes at high nxiss are favoured over the Salpeter single yower law on both observational and theoretical grounds (e.g. Ixrouva οἱ al. | Multi-component IMFs with steeper slopes at high mass are favoured over the Salpeter single power law on both observational and theoretical grounds (e.g. Kroupa et al. |
1993: Scalo 1986). | 1993; Scalo 1986). |
Lt has been suggested that the mass distribution of a stellar generation is ΠΕοσους by factors including the thermal energyIp and. chemical composition of the star-forming OOgas (Larson 1998). | It has been suggested that the mass distribution of a stellar generation is influenced by factors including the thermal energy and chemical composition of the star-forming gas (Larson 1998). |
Accorclinely. the LAL might be expected to evolve over time. | Accordingly, the IMF might be expected to evolve over time. |
There are theoretical arguments hat the IME of primordial gas would be biased toward higher mass stars. while higher metallicity environments would form relatively more stars (Ixroupa 2001). | There are theoretical arguments that the IMF of primordial gas would be biased toward higher mass stars, while higher metallicity environments would form relatively more low-mass stars (Kroupa 2001). |
Ho the. LAL followed this trend then the fit. between our model predictions and | If the IMF followed this trend then the fit between our model predictions and |
by less than an order of magnitude. | by less than an order of magnitude. |
Governatoetal.(2002) report the formation of realistic disk galaxies in ACDM and AWDM cosmologies using simulations that include standard prescriptions for cooling. a UV background. and star formation. | \cite{governato02a}
report the formation of realistic disk galaxies in $\Lambda$ CDM and $\Lambda$ WDM cosmologies using simulations that include standard prescriptions for cooling, a UV background, and star formation. |
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