Relativistic iron emission lines in neutron star low-mass X-ray binaries as probes of neutr

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308nm准分子激光治疗白癜风皮肤病进展马跃;杨春俊;邓赞红;林颖;江海河;方晓东【摘要】临床研究表明波长308 nm准分子激光治疗白癜风皮肤病效果明显.治疗效果与病人的性别、年龄、病程没有明显关联,但与皮肤类型、皮损部位、治疗频度和疗程等有明显关联.在一定剂量范围内随治疗剂量和治疗时间增加治疗效果越来越明显,呈线性变化.而皮损部位的影响,疗效从明显到不明显的顺序是:面部,颈部和头皮,生殖器,四肢,躯干,手脚或肢端关节.皮肤类型对治疗效果影响明显.另外,在一定的观察时间内发现308 nm准分子激光疗效明显高于NB UVB.%This study is a retrospective review of the clinical efficacy of 308 nm excimer laser in vitiligo treatment. Clinical studies have reported that 308 nm excimer laser shows effective treatment of vitiligo. The treatment efficacy has showed no obvious correlation with the patients' sex, age and course of disease, but obviously correlated with skin type, skin lesions, treatment frequency and duration. The treatment efficacy increased linearly with the dose and duration of treatment in a certain dose range. The treatment response showed anatomical preferences, in a decline order of face, neck, scalp, genitals, limbs, trunk, extremities joints. Another factor on the efficacy was skin type. In addition, 308 nm excimer laser treatment appears to be more effective than NB UVB phototherapy in a certain observation period.【期刊名称】《激光生物学报》【年(卷),期】2012(021)004【总页数】6页(P294-298,339)【关键词】308 nm准分子激光;激光疗法;白癜风皮肤病【作者】马跃;杨春俊;邓赞红;林颖;江海河;方晓东【作者单位】中国科学院安徽光学精密机械研究所,安徽省光子器件与材料重点实验室,安徽合肥230031中南大学湘雅二医院普外科,湖南长沙410011;安徽医科大学第一附属医院皮肤性病科,安徽合肥230022;中国科学院安徽光学精密机械研究所,安徽省光子器件与材料重点实验室,安徽合肥230031中南大学湘雅二医院普外科,湖南长沙410011;中国科学院安徽光学精密机械研究所,安徽省光子器件与材料重点实验室,安徽合肥230031中南大学湘雅二医院普外科,湖南长沙410011;中国科学院安徽光学精密机械研究所,安徽省光子器件与材料重点实验室,安徽合肥230031中南大学湘雅二医院普外科,湖南长沙410011;中国科学院安徽光学精密机械研究所,安徽省光子器件与材料重点实验室,安徽合肥230031;中国科学技术大学,安徽合肥230029【正文语种】中文【中图分类】R751.05白癜风是一种常见的后天性表皮色素脱失皮肤病，发病人口分布世界各地，平均发病率0.1% ～4%［1］。

【在研科研项目】1.科技部国际合作项目：新型金刚石结构高能粒子探测器的合作研究（2015DFG02100）， 2014年-2017年。

2.国家自然科学基金项目，氩等离子体弧处理金刚石自支撑膜强度提高的机制（51272024），2013年-2016年。

3.国家自然科学基金项目，氢致金刚石表面载流子输运沟道形成与稳定机制研究（51402013），2015年-2017年。

【代表性学术论文】1. S. Cao ，C. Li ，et al. Long-lived and well-resolved Mn2+ ion emissions in CuInS-ZnS quantum dots[J]. Scientific Reports, 2014, 4: 7510.2. S. Liu, J.L. Liu, C.M. Li, et al. The mechanical enhancement of chemical vapor deposited diamond film by plasma low pressure/ high-temperature treatment [J]. Carbon, 2013, 65: 365-370.3. C.M. Li, et al. Effect of arc characteristics on the properties of large size diamond wafers prepared by DC arc plasma jet CVD [J]. Diamond and Related Materials, 2013, 39: 47-52.4. J.L. Liu, C.M. Li, et al. RF characteristic of MESFET on H-terminated DC arc jet CVD diamond film [J]. Applied Surface Science, 2013, 284: 798-803.5. L. Wang, C. M. Li, et al. Low-aspect ratio graphite hollow nano- structures[J]. CrystEngComm, 2013, 15：8907-8911.李成明，1962年9月出生，新材料技术研究院教授，博士生导师。

Observation of Gravitational Waves from a Binary Black Hole MergerB.P.Abbott et al.*(LIGO Scientific Collaboration and Virgo Collaboration)(Received21January2016;published11February2016)On September14,2015at09:50:45UTC the two detectors of the Laser Interferometer Gravitational-Wave Observatory simultaneously observed a transient gravitational-wave signal.The signal sweeps upwards in frequency from35to250Hz with a peak gravitational-wave strain of1.0×10−21.It matches the waveform predicted by general relativity for the inspiral and merger of a pair of black holes and the ringdown of the resulting single black hole.The signal was observed with a matched-filter signal-to-noise ratio of24and a false alarm rate estimated to be less than1event per203000years,equivalent to a significance greaterthan5.1σ.The source lies at a luminosity distance of410þ160−180Mpc corresponding to a redshift z¼0.09þ0.03−0.04.In the source frame,the initial black hole masses are36þ5−4M⊙and29þ4−4M⊙,and the final black hole mass is62þ4−4M⊙,with3.0þ0.5−0.5M⊙c2radiated in gravitational waves.All uncertainties define90%credible intervals.These observations demonstrate the existence of binary stellar-mass black hole systems.This is the first direct detection of gravitational waves and the first observation of a binary black hole merger.DOI:10.1103/PhysRevLett.116.061102I.INTRODUCTIONIn1916,the year after the final formulation of the field equations of general relativity,Albert Einstein predicted the existence of gravitational waves.He found that the linearized weak-field equations had wave solutions: transverse waves of spatial strain that travel at the speed of light,generated by time variations of the mass quadrupole moment of the source[1,2].Einstein understood that gravitational-wave amplitudes would be remarkably small;moreover,until the Chapel Hill conference in 1957there was significant debate about the physical reality of gravitational waves[3].Also in1916,Schwarzschild published a solution for the field equations[4]that was later understood to describe a black hole[5,6],and in1963Kerr generalized the solution to rotating black holes[7].Starting in the1970s theoretical work led to the understanding of black hole quasinormal modes[8–10],and in the1990s higher-order post-Newtonian calculations[11]preceded extensive analytical studies of relativistic two-body dynamics[12,13].These advances,together with numerical relativity breakthroughs in the past decade[14–16],have enabled modeling of binary black hole mergers and accurate predictions of their gravitational waveforms.While numerous black hole candidates have now been identified through electromag-netic observations[17–19],black hole mergers have not previously been observed.The discovery of the binary pulsar system PSR B1913þ16 by Hulse and Taylor[20]and subsequent observations of its energy loss by Taylor and Weisberg[21]demonstrated the existence of gravitational waves.This discovery, along with emerging astrophysical understanding[22], led to the recognition that direct observations of the amplitude and phase of gravitational waves would enable studies of additional relativistic systems and provide new tests of general relativity,especially in the dynamic strong-field regime.Experiments to detect gravitational waves began with Weber and his resonant mass detectors in the1960s[23], followed by an international network of cryogenic reso-nant detectors[24].Interferometric detectors were first suggested in the early1960s[25]and the1970s[26].A study of the noise and performance of such detectors[27], and further concepts to improve them[28],led to proposals for long-baseline broadband laser interferome-ters with the potential for significantly increased sensi-tivity[29–32].By the early2000s,a set of initial detectors was completed,including TAMA300in Japan,GEO600 in Germany,the Laser Interferometer Gravitational-Wave Observatory(LIGO)in the United States,and Virgo in binations of these detectors made joint obser-vations from2002through2011,setting upper limits on a variety of gravitational-wave sources while evolving into a global network.In2015,Advanced LIGO became the first of a significantly more sensitive network of advanced detectors to begin observations[33–36].A century after the fundamental predictions of Einstein and Schwarzschild,we report the first direct detection of gravitational waves and the first direct observation of a binary black hole system merging to form a single black hole.Our observations provide unique access to the*Full author list given at the end of the article.Published by the American Physical Society under the terms of the Creative Commons Attribution3.0License.Further distri-bution of this work must maintain attribution to the author(s)and the published article’s title,journal citation,and DOI.properties of space-time in the strong-field,high-velocity regime and confirm predictions of general relativity for the nonlinear dynamics of highly disturbed black holes.II.OBSERVATIONOn September14,2015at09:50:45UTC,the LIGO Hanford,W A,and Livingston,LA,observatories detected the coincident signal GW150914shown in Fig.1.The initial detection was made by low-latency searches for generic gravitational-wave transients[41]and was reported within three minutes of data acquisition[43].Subsequently, matched-filter analyses that use relativistic models of com-pact binary waveforms[44]recovered GW150914as the most significant event from each detector for the observa-tions reported here.Occurring within the10-msintersite FIG.1.The gravitational-wave event GW150914observed by the LIGO Hanford(H1,left column panels)and Livingston(L1,rightcolumn panels)detectors.Times are shown relative to September14,2015at09:50:45UTC.For visualization,all time series are filtered with a35–350Hz bandpass filter to suppress large fluctuations outside the detectors’most sensitive frequency band,and band-reject filters to remove the strong instrumental spectral lines seen in the Fig.3spectra.Top row,left:H1strain.Top row,right:L1strain.GW150914arrived first at L1and6.9þ0.5−0.4ms later at H1;for a visual comparison,the H1data are also shown,shifted in time by this amount and inverted(to account for the detectors’relative orientations).Second row:Gravitational-wave strain projected onto each detector in the35–350Hz band.Solid lines show a numerical relativity waveform for a system with parameters consistent with those recovered from GW150914[37,38]confirmed to99.9%by an independent calculation based on[15].Shaded areas show90%credible regions for two independent waveform reconstructions.One(dark gray)models the signal using binary black hole template waveforms [39].The other(light gray)does not use an astrophysical model,but instead calculates the strain signal as a linear combination of sine-Gaussian wavelets[40,41].These reconstructions have a94%overlap,as shown in[39].Third row:Residuals after subtracting the filtered numerical relativity waveform from the filtered detector time series.Bottom row:A time-frequency representation[42]of the strain data,showing the signal frequency increasing over time.propagation time,the events have a combined signal-to-noise ratio(SNR)of24[45].Only the LIGO detectors were observing at the time of GW150914.The Virgo detector was being upgraded, and GEO600,though not sufficiently sensitive to detect this event,was operating but not in observational mode.With only two detectors the source position is primarily determined by the relative arrival time and localized to an area of approximately600deg2(90% credible region)[39,46].The basic features of GW150914point to it being produced by the coalescence of two black holes—i.e., their orbital inspiral and merger,and subsequent final black hole ringdown.Over0.2s,the signal increases in frequency and amplitude in about8cycles from35to150Hz,where the amplitude reaches a maximum.The most plausible explanation for this evolution is the inspiral of two orbiting masses,m1and m2,due to gravitational-wave emission.At the lower frequencies,such evolution is characterized by the chirp mass[11]M¼ðm1m2Þ3=5121=5¼c3G596π−8=3f−11=3_f3=5;where f and_f are the observed frequency and its time derivative and G and c are the gravitational constant and speed of light.Estimating f and_f from the data in Fig.1, we obtain a chirp mass of M≃30M⊙,implying that the total mass M¼m1þm2is≳70M⊙in the detector frame. This bounds the sum of the Schwarzschild radii of thebinary components to2GM=c2≳210km.To reach an orbital frequency of75Hz(half the gravitational-wave frequency)the objects must have been very close and very compact;equal Newtonian point masses orbiting at this frequency would be only≃350km apart.A pair of neutron stars,while compact,would not have the required mass,while a black hole neutron star binary with the deduced chirp mass would have a very large total mass, and would thus merge at much lower frequency.This leaves black holes as the only known objects compact enough to reach an orbital frequency of75Hz without contact.Furthermore,the decay of the waveform after it peaks is consistent with the damped oscillations of a black hole relaxing to a final stationary Kerr configuration. Below,we present a general-relativistic analysis of GW150914;Fig.2shows the calculated waveform using the resulting source parameters.III.DETECTORSGravitational-wave astronomy exploits multiple,widely separated detectors to distinguish gravitational waves from local instrumental and environmental noise,to provide source sky localization,and to measure wave polarizations. The LIGO sites each operate a single Advanced LIGO detector[33],a modified Michelson interferometer(see Fig.3)that measures gravitational-wave strain as a differ-ence in length of its orthogonal arms.Each arm is formed by two mirrors,acting as test masses,separated by L x¼L y¼L¼4km.A passing gravitational wave effec-tively alters the arm lengths such that the measured difference isΔLðtÞ¼δL x−δL y¼hðtÞL,where h is the gravitational-wave strain amplitude projected onto the detector.This differential length variation alters the phase difference between the two light fields returning to the beam splitter,transmitting an optical signal proportional to the gravitational-wave strain to the output photodetector. To achieve sufficient sensitivity to measure gravitational waves,the detectors include several enhancements to the basic Michelson interferometer.First,each arm contains a resonant optical cavity,formed by its two test mass mirrors, that multiplies the effect of a gravitational wave on the light phase by a factor of300[48].Second,a partially trans-missive power-recycling mirror at the input provides addi-tional resonant buildup of the laser light in the interferometer as a whole[49,50]:20W of laser input is increased to700W incident on the beam splitter,which is further increased to 100kW circulating in each arm cavity.Third,a partially transmissive signal-recycling mirror at the outputoptimizes FIG. 2.Top:Estimated gravitational-wave strain amplitude from GW150914projected onto H1.This shows the full bandwidth of the waveforms,without the filtering used for Fig.1. The inset images show numerical relativity models of the black hole horizons as the black holes coalesce.Bottom:The Keplerian effective black hole separation in units of Schwarzschild radii (R S¼2GM=c2)and the effective relative velocity given by the post-Newtonian parameter v=c¼ðGMπf=c3Þ1=3,where f is the gravitational-wave frequency calculated with numerical relativity and M is the total mass(value from Table I).the gravitational-wave signal extraction by broadening the bandwidth of the arm cavities [51,52].The interferometer is illuminated with a 1064-nm wavelength Nd:Y AG laser,stabilized in amplitude,frequency,and beam geometry [53,54].The gravitational-wave signal is extracted at the output port using a homodyne readout [55].These interferometry techniques are designed to maxi-mize the conversion of strain to optical signal,thereby minimizing the impact of photon shot noise (the principal noise at high frequencies).High strain sensitivity also requires that the test masses have low displacement noise,which is achieved by isolating them from seismic noise (low frequencies)and designing them to have low thermal noise (intermediate frequencies).Each test mass is suspended as the final stage of a quadruple-pendulum system [56],supported by an active seismic isolation platform [57].These systems collectively provide more than 10orders of magnitude of isolation from ground motion for frequen-cies above 10Hz.Thermal noise is minimized by using low-mechanical-loss materials in the test masses and their suspensions:the test masses are 40-kg fused silica substrates with low-loss dielectric optical coatings [58,59],and are suspended with fused silica fibers from the stage above [60].To minimize additional noise sources,all components other than the laser source are mounted on vibration isolation stages in ultrahigh vacuum.To reduce optical phase fluctuations caused by Rayleigh scattering,the pressure in the 1.2-m diameter tubes containing the arm-cavity beams is maintained below 1μPa.Servo controls are used to hold the arm cavities on resonance [61]and maintain proper alignment of the optical components [62].The detector output is calibrated in strain by measuring its response to test mass motion induced by photon pressure from a modulated calibration laser beam [63].The calibration is established to an uncertainty (1σ)of less than 10%in amplitude and 10degrees in phase,and is continuously monitored with calibration laser excitations at selected frequencies.Two alternative methods are used to validate the absolute calibration,one referenced to the main laser wavelength and the other to a radio-frequencyoscillator(a)FIG.3.Simplified diagram of an Advanced LIGO detector (not to scale).A gravitational wave propagating orthogonally to the detector plane and linearly polarized parallel to the 4-km optical cavities will have the effect of lengthening one 4-km arm and shortening the other during one half-cycle of the wave;these length changes are reversed during the other half-cycle.The output photodetector records these differential cavity length variations.While a detector ’s directional response is maximal for this case,it is still significant for most other angles of incidence or polarizations (gravitational waves propagate freely through the Earth).Inset (a):Location and orientation of the LIGO detectors at Hanford,WA (H1)and Livingston,LA (L1).Inset (b):The instrument noise for each detector near the time of the signal detection;this is an amplitude spectral density,expressed in terms of equivalent gravitational-wave strain amplitude.The sensitivity is limited by photon shot noise at frequencies above 150Hz,and by a superposition of other noise sources at lower frequencies [47].Narrow-band features include calibration lines (33–38,330,and 1080Hz),vibrational modes of suspension fibers (500Hz and harmonics),and 60Hz electric power grid harmonics.[64].Additionally,the detector response to gravitational waves is tested by injecting simulated waveforms with the calibration laser.To monitor environmental disturbances and their influ-ence on the detectors,each observatory site is equipped with an array of sensors:seismometers,accelerometers, microphones,magnetometers,radio receivers,weather sensors,ac-power line monitors,and a cosmic-ray detector [65].Another∼105channels record the interferometer’s operating point and the state of the control systems.Data collection is synchronized to Global Positioning System (GPS)time to better than10μs[66].Timing accuracy is verified with an atomic clock and a secondary GPS receiver at each observatory site.In their most sensitive band,100–300Hz,the current LIGO detectors are3to5times more sensitive to strain than initial LIGO[67];at lower frequencies,the improvement is even greater,with more than ten times better sensitivity below60Hz.Because the detectors respond proportionally to gravitational-wave amplitude,at low redshift the volume of space to which they are sensitive increases as the cube of strain sensitivity.For binary black holes with masses similar to GW150914,the space-time volume surveyed by the observations reported here surpasses previous obser-vations by an order of magnitude[68].IV.DETECTOR VALIDATIONBoth detectors were in steady state operation for several hours around GW150914.All performance measures,in particular their average sensitivity and transient noise behavior,were typical of the full analysis period[69,70]. Exhaustive investigations of instrumental and environ-mental disturbances were performed,giving no evidence to suggest that GW150914could be an instrumental artifact [69].The detectors’susceptibility to environmental disturb-ances was quantified by measuring their response to spe-cially generated magnetic,radio-frequency,acoustic,and vibration excitations.These tests indicated that any external disturbance large enough to have caused the observed signal would have been clearly recorded by the array of environ-mental sensors.None of the environmental sensors recorded any disturbances that evolved in time and frequency like GW150914,and all environmental fluctuations during the second that contained GW150914were too small to account for more than6%of its strain amplitude.Special care was taken to search for long-range correlated disturbances that might produce nearly simultaneous signals at the two sites. No significant disturbances were found.The detector strain data exhibit non-Gaussian noise transients that arise from a variety of instrumental mecha-nisms.Many have distinct signatures,visible in auxiliary data channels that are not sensitive to gravitational waves; such instrumental transients are removed from our analyses [69].Any instrumental transients that remain in the data are accounted for in the estimated detector backgrounds described below.There is no evidence for instrumental transients that are temporally correlated between the two detectors.V.SEARCHESWe present the analysis of16days of coincident observations between the two LIGO detectors from September12to October20,2015.This is a subset of the data from Advanced LIGO’s first observational period that ended on January12,2016.GW150914is confidently detected by two different types of searches.One aims to recover signals from the coalescence of compact objects,using optimal matched filtering with waveforms predicted by general relativity. The other search targets a broad range of generic transient signals,with minimal assumptions about waveforms.These searches use independent methods,and their response to detector noise consists of different,uncorrelated,events. However,strong signals from binary black hole mergers are expected to be detected by both searches.Each search identifies candidate events that are detected at both observatories consistent with the intersite propa-gation time.Events are assigned a detection-statistic value that ranks their likelihood of being a gravitational-wave signal.The significance of a candidate event is determined by the search background—the rate at which detector noise produces events with a detection-statistic value equal to or higher than the candidate event.Estimating this back-ground is challenging for two reasons:the detector noise is nonstationary and non-Gaussian,so its properties must be empirically determined;and it is not possible to shield the detector from gravitational waves to directly measure a signal-free background.The specific procedure used to estimate the background is slightly different for the two searches,but both use a time-shift technique:the time stamps of one detector’s data are artificially shifted by an offset that is large compared to the intersite propagation time,and a new set of events is produced based on this time-shifted data set.For instrumental noise that is uncor-related between detectors this is an effective way to estimate the background.In this process a gravitational-wave signal in one detector may coincide with time-shifted noise transients in the other detector,thereby contributing to the background estimate.This leads to an overestimate of the noise background and therefore to a more conservative assessment of the significance of candidate events.The characteristics of non-Gaussian noise vary between different time-frequency regions.This means that the search backgrounds are not uniform across the space of signals being searched.To maximize sensitivity and provide a better estimate of event significance,the searches sort both their background estimates and their event candidates into differ-ent classes according to their time-frequency morphology. The significance of a candidate event is measured against the background of its class.To account for having searchedmultiple classes,this significance is decreased by a trials factor equal to the number of classes [71].A.Generic transient searchDesigned to operate without a specific waveform model,this search identifies coincident excess power in time-frequency representations of the detector strain data [43,72],for signal frequencies up to 1kHz and durations up to a few seconds.The search reconstructs signal waveforms consistent with a common gravitational-wave signal in both detectors using a multidetector maximum likelihood method.Each event is ranked according to the detection statistic ηc ¼ffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffiffi2E c =ð1þE n =E c Þp ,where E c is the dimensionless coherent signal energy obtained by cross-correlating the two reconstructed waveforms,and E n is the dimensionless residual noise energy after the reconstructed signal is subtracted from the data.The statistic ηc thus quantifies the SNR of the event and the consistency of the data between the two detectors.Based on their time-frequency morphology,the events are divided into three mutually exclusive search classes,as described in [41]:events with time-frequency morphology of known populations of noise transients (class C1),events with frequency that increases with time (class C3),and all remaining events (class C2).Detected with ηc ¼20.0,GW150914is the strongest event of the entire search.Consistent with its coalescence signal signature,it is found in the search class C3of events with increasing time-frequency evolution.Measured on a background equivalent to over 67400years of data and including a trials factor of 3to account for the search classes,its false alarm rate is lower than 1in 22500years.This corresponds to a probability <2×10−6of observing one or more noise events as strong as GW150914during the analysis time,equivalent to 4.6σ.The left panel of Fig.4shows the C3class results and background.The selection criteria that define the search class C3reduce the background by introducing a constraint on the signal morphology.In order to illustrate the significance of GW150914against a background of events with arbitrary shapes,we also show the results of a search that uses the same set of events as the one described above but without this constraint.Specifically,we use only two search classes:the C1class and the union of C2and C3classes (C 2þC 3).In this two-class search the GW150914event is found in the C 2þC 3class.The left panel of Fig.4shows the C 2þC 3class results and background.In the background of this class there are four events with ηc ≥32.1,yielding a false alarm rate for GW150914of 1in 8400years.This corresponds to a false alarm probability of 5×10−6equivalent to 4.4σ.FIG.4.Search results from the generic transient search (left)and the binary coalescence search (right).These histograms show the number of candidate events (orange markers)and the mean number of background events (black lines)in the search class where GW150914was found as a function of the search detection statistic and with a bin width of 0.2.The scales on the top give the significance of an event in Gaussian standard deviations based on the corresponding noise background.The significance of GW150914is greater than 5.1σand 4.6σfor the binary coalescence and the generic transient searches,respectively.Left:Along with the primary search (C3)we also show the results (blue markers)and background (green curve)for an alternative search that treats events independently of their frequency evolution (C 2þC 3).The classes C2and C3are defined in the text.Right:The tail in the black-line background of the binary coalescence search is due to random coincidences of GW150914in one detector with noise in the other detector.(This type of event is practically absent in the generic transient search background because they do not pass the time-frequency consistency requirements used in that search.)The purple curve is the background excluding those coincidences,which is used to assess the significance of the second strongest event.For robustness and validation,we also use other generic transient search algorithms[41].A different search[73]and a parameter estimation follow-up[74]detected GW150914 with consistent significance and signal parameters.B.Binary coalescence searchThis search targets gravitational-wave emission from binary systems with individual masses from1to99M⊙, total mass less than100M⊙,and dimensionless spins up to 0.99[44].To model systems with total mass larger than 4M⊙,we use the effective-one-body formalism[75],whichcombines results from the post-Newtonian approach [11,76]with results from black hole perturbation theory and numerical relativity.The waveform model[77,78] assumes that the spins of the merging objects are alignedwith the orbital angular momentum,but the resultingtemplates can,nonetheless,effectively recover systemswith misaligned spins in the parameter region ofGW150914[44].Approximately250000template wave-forms are used to cover this parameter space.The search calculates the matched-filter signal-to-noiseratioρðtÞfor each template in each detector and identifiesmaxima ofρðtÞwith respect to the time of arrival of the signal[79–81].For each maximum we calculate a chi-squared statisticχ2r to test whether the data in several differentfrequency bands are consistent with the matching template [82].Values ofχ2r near unity indicate that the signal is consistent with a coalescence.Ifχ2r is greater than unity,ρðtÞis reweighted asˆρ¼ρ=f½1þðχ2rÞ3 =2g1=6[83,84].The final step enforces coincidence between detectors by selectingevent pairs that occur within a15-ms window and come fromthe same template.The15-ms window is determined by the10-ms intersite propagation time plus5ms for uncertainty inarrival time of weak signals.We rank coincident events basedon the quadrature sumˆρc of theˆρfrom both detectors[45]. To produce background data for this search the SNR maxima of one detector are time shifted and a new set of coincident events is computed.Repeating this procedure ∼107times produces a noise background analysis time equivalent to608000years.To account for the search background noise varying acrossthe target signal space,candidate and background events aredivided into three search classes based on template length.The right panel of Fig.4shows the background for thesearch class of GW150914.The GW150914detection-statistic value ofˆρc¼23.6is larger than any background event,so only an upper bound can be placed on its false alarm rate.Across the three search classes this bound is1in 203000years.This translates to a false alarm probability <2×10−7,corresponding to5.1σ.A second,independent matched-filter analysis that uses adifferent method for estimating the significance of itsevents[85,86],also detected GW150914with identicalsignal parameters and consistent significance.When an event is confidently identified as a real gravitational-wave signal,as for GW150914,the back-ground used to determine the significance of other events is reestimated without the contribution of this event.This is the background distribution shown as a purple line in the right panel of Fig.4.Based on this,the second most significant event has a false alarm rate of1per2.3years and corresponding Poissonian false alarm probability of0.02. Waveform analysis of this event indicates that if it is astrophysical in origin it is also a binary black hole merger[44].VI.SOURCE DISCUSSIONThe matched-filter search is optimized for detecting signals,but it provides only approximate estimates of the source parameters.To refine them we use general relativity-based models[77,78,87,88],some of which include spin precession,and for each model perform a coherent Bayesian analysis to derive posterior distributions of the source parameters[89].The initial and final masses, final spin,distance,and redshift of the source are shown in Table I.The spin of the primary black hole is constrained to be<0.7(90%credible interval)indicating it is not maximally spinning,while the spin of the secondary is only weakly constrained.These source parameters are discussed in detail in[39].The parameter uncertainties include statistical errors and systematic errors from averaging the results of different waveform models.Using the fits to numerical simulations of binary black hole mergers in[92,93],we provide estimates of the mass and spin of the final black hole,the total energy radiated in gravitational waves,and the peak gravitational-wave luminosity[39].The estimated total energy radiated in gravitational waves is3.0þ0.5−0.5M⊙c2.The system reached apeak gravitational-wave luminosity of3.6þ0.5−0.4×1056erg=s,equivalent to200þ30−20M⊙c2=s.Several analyses have been performed to determine whether or not GW150914is consistent with a binary TABLE I.Source parameters for GW150914.We report median values with90%credible intervals that include statistical errors,and systematic errors from averaging the results of different waveform models.Masses are given in the source frame;to convert to the detector frame multiply by(1þz) [90].The source redshift assumes standard cosmology[91]. Primary black hole mass36þ5−4M⊙Secondary black hole mass29þ4−4M⊙Final black hole mass62þ4−4M⊙Final black hole spin0.67þ0.05−0.07 Luminosity distance410þ160−180MpcSource redshift z0.09þ0.03−0.04。

中子星内部中微子辐射的相对论修正戚湛强;丁文波;张承民;侯佳为【摘要】Through the relativistic mean field theory and relevant weak interactional cooling theories,the relativistic cooling properties in conventional and hyperonic neutron star matter are studied.Also a comparison between relativistic and non-relativistic results with the consideration of the gravity correction is performed.Results show that the correction of relativistic effect in neutrino emission makes the neutrino emissivity,neutrino luminosity,and cooling rate be lower,in comparison with the non-relativistic case.Due to the correction of relativistic effect in neutrino emission,the cooling rate of neutron star has the largest decline,for the conventional neutron star matter with the consideration of gravity correction.The rate of decline reaches 56％ for the conventional neutron star with a mass of two solar masses,however,the decline in hyperonic matter is the smallest,about 38％.%通过相对论平均场理论和相关的弱相互作用冷却理论,在有、无超子两种情况下的中子星物质中研究含相对论效应的中子星冷却性质,并且与非相对论情况进行对比分析.考虑到极微小尺度上的引力会偏离牛顿引力,引入引力修正效应.结果表明中微子辐射的相对论效应降低了中微子发射率、发光度以及星体的冷却速度.在考虑引力修正的无超子的中子星物质中,相对论效应所引起的星体冷却速度降幅最大,对于两倍太阳质量的传统物质中子星可达56％,而超子物质中降幅最小,约为38％.【期刊名称】《天文学报》【年(卷),期】2017(058)001【总页数】10页(P98-107)【关键词】恒星:中子;恒星:相对论效应;恒星:引力修正;恒星:超子【作者】戚湛强;丁文波;张承民;侯佳为【作者单位】渤海大学数理学院锦州121013;中国科学院国家天文台北京100012;渤海大学数理学院锦州121013;中国科学院国家天文台北京100012;渤海大学数理学院锦州121013【正文语种】中文【中图分类】P145中子星的质量约为1—2倍太阳质量[1−2],半径为10—20km,这样高度致密的物质环境为研究致密物质的性质、中微子辐射、致密天体演化等提供了一个理想的天然载体.近几十年来,因为探测中子星温度的技术得以发展,所以中子星结构、冷却性质等方面的研究已经成为天体物理和核物理领域的研究热点.中子星从诞生之初1011K的极高温度下降到较低的108K主要是通过辐射中微子实现的,而冷却效率最高的中微子辐射机制当属核子的直接Urca过程(以下简称dUrca过程),其次是超子、反K介子、夸克等的直接Urca过程[3−4].在中子星内部的中微子辐射过程中,中微子发射率是一个较为关键的物理量,它给出单位时间单位体积内中微子辐射所带走的能量.目前学者们使用较多的中微子发射率公式主要有两个：Lattimer等[5]于1991年推导的不包含相对论效应的公式,Leinson等[6]于2001年提出的考虑相对论效应的公式.显然,后者所得的结果更为可信,因为通常情况下,dUrca过程发生的临界密度在2ρ0附近(ρ0为饱和核密度),在这样高的密度下,参与dUrca过程的粒子的运动速度很高,需要考虑相对论效应.然而以前关于中微子辐射的讨论多数都采用不含相对论效应的公式,若重新计算则工作量较大.因此,定量地给出不同中子星物质组成下不同质量中子星的冷却性质在考虑相对论效应下的修正值将有助于前人科研成果的推广和精确应用,具有一定的现实意义.近期的观测数据显示中子星的质量比人们预想的略大.比如,毫秒脉冲星PSRJ1614-2230的质量被精确测定为(1.97±0.04)M⊙(M⊙表示太阳质量),半径约为11—15km[7],而2013年观测到的脉冲星PSR J0348+0432质量更大[8],为(2.01±0.04)M⊙.如此大的质量与人们理论预测的中子星正则质量1.4 M⊙存在一定差距,学者们猜测这主要是由两方面原因造成的：一是中子星物质的物态方程(Equation of State,EoS)应该较硬,而从前的理论结果偏软;二是理论计算中可能忽略了某些需要考虑的因素,比如引力修正[9−11]、强磁场效应[12]、电场效应[13]等,而其中的引力修正正逐渐被学者们认可.引力虽然是人类认识最早的力,但是至今仍未被完全研究清楚,仍有一些现象是现在的引力理论所不能解释的[14−16].而近年的实验和理论也表明在中子星这样的高密环境下,研究极微小尺度下的物理问题时,引力可能偏离牛顿的平方反比定律[17−18].陈列文等人的研究表明,引力修正效应可硬化物态方程,并对中子星的结构产生显著影响[9−11].因此,本文也考虑了引力修正对中微子辐射所带来的影响.本文将采用相对论平均场理论和相关弱相互作用理论对GM1参数下的无超子中子星物质(主要含中子、质子、电子、少量µ子,即传统中子星物质)、含超子的中子星物质进行计算,并考虑引力修正效应,给出考虑相对论效应后中子星内部的中微子辐射性质、星体的冷却曲线等与未考虑相对论效应结果的具体差异.本文第2部分介绍所用到的理论模型和计算方法;第3部分给出相关的模型参数;第4部分对计算结果进行细致的讨论;第5部分总结本文的结论.强相互作用下,关于中子星物质方面的计算,我们采用相对论平均场理论(RMFT),拉氏量的表达式为：各字母的含义及相关细节见文献[19-20].通过求解由场方程、化学平衡、总重子数守恒、电中性条件等构成的非线性方程组,可以得到各介子场的场量、质子、中子、电子等微观粒子的费米动量、数密度、化学势、体系的EoS等信息.把EoS作为输入量,求解Tolman-Oppenheimer-Volko ff(TOV)方程可以得到一系列静态的中子星序列,给出中子星的质量-半径关系,对于确定质量的中子星,可以得到其核心密度以及密度、压强等随径向的变化规律.关于引力修正,我们考虑的是非牛顿引力所引起的修正效应.早在1971年,Fujii[21]就指出在极微小的尺度下,非牛顿引力势可写为：其中G和λ分别为引力常数和特征长度,α代表无量纲的强度参数,m代表重子质量.按照矢量玻色子交换模型,与核子弱耦合的矢量U玻色子是目前公认最适合传递相互作用的候选粒子.在相对论平均场理论中,引力修正(U玻色子)的引入使得能量密度和压强的公式中各需要加入一项[22−24]：其中g为耦合常数,mU和nB分别代表U玻色子质量和总重子数密度.因此,考虑引力修正后中子星的EoS发生改变,总能量密度ε和总压强P分别变为：ε0和P0是未考虑引力修正时的EoS.从(4)式可以看出,引力修正使得体系物态方程变硬,进而增大中子星的最大质量.关于电子的核子Urca过程(dUrcae过程)的反应式为：其中n、p、e−和分别表示中子、质子、电子和反电子中微子.与反应式(5)式对应的非相对论的中微子发射率公式[5]为：公式中的GF为弱相互作用的耦合常数,θC为Cabibbo角,T表示星体内部温度,Θ为阶跃函数.CV=1,CA=1.26,它们分别是矢量耦合常数和轴矢量耦合常数.µe为电子化学式,而分别为质子、电子和中子的费米动量.(6)式中的分别代表中子和质子的朗道有效质量,在相对论平均场理论中,它们的表达式为：其中i代表核子,mi为核子裸质量,gσi为核子与σ场的耦合常数,而σ代表σ场的场量.实际上(6)式中的Θ给出的是电子dUrca过程的发生条件,即而考虑相对论效应的中微子发射率(关于电子)的公式[6]为：其中分别为中子和质子的狄拉克有效质量,子的Urca过程对应的中微子发射率与电子类似,只要把电子相关的物理量变为µ子的即可,所以实际上若µ子的Urca过程发生,则总的中微子发射率增加1倍.中子星的冷却曲线通过计算冷却方程得到.(10)式中的t代表时间,即中子星的年龄,Cυ是体系的热容.Lυ代表星体的中微子发光度,通过把中微子发射率按星体体积积分计算得到,而Lγ是星体的光子发光度,我们采用经验公式来进行近似计算.上式中Te,7=Te/(107K),而Te代表星体的表面有效温度,它与星体内部温度T的关系为：相对论平均场理论(Relativistic Mean-Field Theory,RMFT)框架下的核子-介子耦合常数是通过饱和核物质的5大性质来确定的,RMFT发展至今已经拥有多套成熟的核物质参数组,本文采用EoS偏硬的一组参数—GM1,相关细节请参见文献[25].本文考虑了全部的重子八重态,为了全面地考察超子所带来的影响,只考虑Λ、Σ和Ξ超子在对称核物质中均表现为吸引势的情况,且设它们与σ场的耦合常数与核子-σ耦合常数相同,即χσH=gσH/gσN=1(H代表超子,N代表核子).而超子与ω场耦合不敏感,因此设χωH=gωH/gωN=1,超子与ρ介子的耦合常数可通过SU(6)对称模型获得χρΛ=0, χρΣ=2,χρΞ=1(χωH的定义与σ和ω场的超子耦合常数类似).引力修正方面引入U玻色子来传递相互作用,通常认为(3)式中的(简称U玻色子有效耦合常数)应小于200 GeV−2,而对于中子星,其值应在50 GeV−2附近或略大于此值.本文为了研究引力修正所带来的一般影响,U玻色子有效耦合常数选取30 GeV−2.值得指出的是,在拟合参数时,引力修正在饱和核密度处也将产生影响,但是此影响与同位旋极度不对称的中子星物质相比较为微弱,因此我们暂时将其忽略. 采用上述理论和参数进行计算,通过求解TOV方程可得传统中子星物质(简称np物质)和超子物质(简称npH物质)中子星的质量-半径关系,见图1.从图中可以看出考虑引力修正后,中子星的质量和半径均增大,而超子的引入则降低了星体的质量,参见表1.观察图中的阴影部分,可以发现中子星PSR J0348+0432可以很好地存在于本文计算的4种情况之下,即组成PSR J0348+0432的物质不仅可以是传统中子星物质,超子也可能存在,并且引力修正效应也适用于PSR J0348+0432.但是,对于中子星PSR J1614-2230,含引力修正效应的传统中子星物质就不适合了,这是因为修正的引力使得中子星的半径变大,而对于传统中子星物质,半径增大的幅度最大,致使理论计算的星体半径大于PSR J1614-2230的观测半径.中子星PSR J1614-2230中完全可能存在超子(无论是否考虑引力修正),而这与文献[7]的结论不同.根据(8)式进行计算,可得到不同的中子星物质下dUrca过程发生的密度范围,如表1所示.引力修正对费米子的费米动量不产生影响,因此电子和µ子的dUrca过程(简称dUrcaµ过程)发生的临界密度不受U玻色子影响.从表1可以看出超子的出现使dUrca过程发生的临界密度降低,密度范围增大,以µ子的dUrca过程最为明显. 通过中微子发射率公式(6)式和(9)式可以算得不同质量中子星内部中微子辐射的性质,其中(6)式对应的结果不含相对论效应,(9)式考虑了相对论修正.图2给出了两颗质量分别为1.4 M⊙(a)和2.0 M⊙(b)中子星内部中微子发射率随径向距离的变化关系,图中标记“R-”对应考虑相对论效应的结果(采用(9)式进行计算).总地来看,相对论效应降低了中微子发射率,但是对于不同的星体质量、不同的物质,以及是否考虑引力修正,降低的幅度明显不同.图2(a)中,对于1.4 M⊙的中子星,不含引力修正的传统中子星物质(np)和超子物质(npH)能够发生dUrca过程,但是考虑引力修正之后dUrca过程在1.4 M⊙星体中则未开启.这是因为引力修正降低了中子星的核心密度,使得1.4 M⊙中子星的核心密度低于dUrca过程发生的临界密度,见表1和表2.而对于图2(a)中的4条曲线,我们很容易发现在径向距离约小于4.3km(中高密度下)时,npH物质的中微子发射率曲线明显高于np物质,这主要是因为np物质构成的中子星内部只发生了电子的dUrca过程,而含超子的npH物质中子星中dUrcae 和dUrcaµ过程都发生了(参见表2),µ子dUrca过程的加入使得总中微子发射率加倍.经过计算,我们发现对于np物质,相对论效应使中微子发射率降低了约34.7%,而npH物质降低了约32.4%.对于大质量中子星,由于其核心密度较大,即使是对于考虑引力修正的中子星,其内部电子和µ子的dUrca过程也都发生了,见图2(b).从图2(b)可以看出,在较高密度的核心处,np物质和npU物质(含引力修正的传统中子星物质)对应的中微子发射率最高,而含超子的npH物质的中微子发射率则最低,这说明超子实际上降低了中微子发射率,尤其是在高密度处,超子的种类和数量较多,中微子发射率降低的幅度也较大.引力修正效应没有降低中微子发射率,反而略微起增加作用.关于相对论与非相对论中微子发射率的差异,我们进行了计算,发现对于2.0M⊙中子星,在其密度较高的核心处,np物质中相对论效应使中微子发射率降低了33.2%,考虑引力修正效应后(npU)降低了36.1%,超子物质(npH)降低了15.5%,考虑引力修正后(npHU)降低了27.9%.由此可见引力修正使得中微子辐射的相对论效应加强,而超子则明显弱化了相对论效应所带来的影响,并且密度越高,星体质量越大,超子的这种弱化影响越显著,而引力修正所带来的影响对密度不敏感.结合图2的(a)和(b),可以发现对于有dUrca过程发生的中子星,除npH物质外的各种情况下中微子发射率相对论与非相对论的差值相差不大,相对论所引起的发射率降幅均在30%左右,而npH物质的降幅最小,约为16%.表2给出1.4 M⊙和2.0 M⊙中子星的核心密度、半径、dUrcae和dUrcaµ过程发生的临界径向距离,以及星体的中微子发光度.可以看出超子促进dUrca过程发生,并且拓宽星体内部dUrca过程发生的范围.所以虽然超子降低了中微子发射率,但是由于dUrca过程发生的范围较广,星体整体中微子发光度却并不是最低的(2.0 M⊙时, npHU情况最低).而引力修正的作用与超子正好相反,虽然引力修正提高了中微子发射率,但是却减小了星体内部dUrca过程发生的范围,因此考虑引力修正效应后,中微子发光度降低.对于2.0 M⊙中子星,np、npH物质中相对论效应引起中微子发光度降低率依次为35.4%和27.6%,而考虑引力修正后,对应的中微子发光度降低率依次为35.6%和32.9%,说明引力修正放大了中微子辐射的相对论效应,而超子所起的作用正好相反,这一结论主要是由中微子发射率决定的(见图2).需要说明的是,表2中“npU”和“npHU”的1.4M⊙中子星实际上是相同的,因为考虑引力修正后1.4 M⊙中子星的核心密度小于超子出现的临界密度,即“npHU”实际上并无超子存在.图3给出2.0 M⊙中子星的冷却曲线.可以看出“npU”曲线在最下方,说明考虑引力修正的传统中子星物质构成的中子星冷却速度最快,其次是“npHU”、“np”等,而冷却最慢的是考虑相对论效应的超子星(R-npH),其次是不考虑相对论效应的超子星(npH)、考虑相对论效应的传统中子星(R-np)等.这说明引力修正加速中子星冷却,而相对论效应则减缓星体的冷却速度.为了仔细探究相对论效应的影响,我们以4×108K的温度为考察点来进行说明.np物质从起始温度降到4×108K需要的时间约为359 yr,而加入相对论效应后约为555 yr;考虑引力修正后(npU),分别需要264 yr和412 yr;超子物质(npH)分别需要814 yr和1124 yr;考虑引力修正后(npHU),分别需要307 yr和457 yr.也就是说,对于2.0 M⊙中子星,若构成物质为传统中子星物质,则考虑中微子辐射的相对论效应后,星体冷却速度下降了54.9%,而考虑引力修正后,下降了56.1%;若星体由超子物质构成,则相对论的效应导致冷却速度下降了38.1%,考虑引力修正后变为48.9%.实际上,上述结果是由中微子辐射特性和星体的热容共同作用产生的.需要说明的是,由于小质量中子星的冷却曲线区分度不大,相对论效应不十分显著,因此文章对1.4 M⊙中子星的冷却曲线未做详细讨论.通过相对论平均场理论和相关弱作用理论,我们在考虑引力修正的情况下,对传统中子星物质和超子物质下的中子星结构、含相对论修正的中微子发射率、发光度、冷却曲线等,进行了计算,并且把结果与非相对论的情况进行对比分析,所得主要结论如下：(1)引力修正增大中子星的最大质量和其对应半径,却减小星体核心密度;(2)相对论的中微子发光度总是低于非相对论的情况,两者的差异在超子物质中最小,相对论效应使中微子发射率下降了约16%,而考虑引力修正的传统中子星物质中,两者的差异最大,下降了约36%(此结论主要通过2.0 M⊙中子星的相关图像得出),中微子发光度也有类似结论;(3)引力修正放大了中微子辐射的相对论效应,而超子却明显弱化了相对论与非相对论结果的差异;(4)引力修正加速中子星冷却;(5)中微子辐射的相对论效应减慢星体的冷却速度,对于2.0 M⊙的中子星,超子物质中冷却速度的降幅约为38%,而在含引力修正的传统中子星物质中这种降幅可达56%.【相关文献】[1]Zhang C M,Wang J,Zhao Y H,et al.A&A,2011,527：A83[2]Cheng Z,Zhang C M,Zhao Y H,et al.ChA&A,2014,38：294[3]Gangopadhyay T,Ray S,Li X D,et al.MNRAS,2013,431：3216[4]Xu R X.arXiv：astro-ph/0211348,2002[5]Lattimer J M,Pethick C J,Prakash M,et al.PhRvL,1991,66：2701[6]Leinson L B,Perez A.PhLB,2001,518：15[7]Demorest P B,Pennucci T,Ransom S M.Nature,2010,467：1081[8]Antoniadis J,Freire P C C,Wex N.Science,2013,340：448[9]Wen D H,Li B A,Chen L W.PhRvL,2009,103：211102[10]Zheng H,Chen L W.PhRvD,2012,85：043013[11]Zhang D R,Yin P L,Wang W,et 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射电天文学英文词典Radio Astronomy Terminology Dictionary.1. Aperture Synthesis: A technique in radio astronomy that combines signals from multiple antennas to create a single, larger virtual antenna, improving angular resolution.2. Beamforming: The process of combining signals from multiple antennas to create a directional beam of radio waves, enhancing the sensitivity of a radio telescope towards a specific direction.3. Bolometric Detector: A device that measures thetotal power of radiation received at a given frequency, regardless of its polarization.4. Flux Density: The measure of the power received froma source per unit area, usually expressed in units of watts per square meter.5. Interferometer: A device that combines signals from two or more antennas to measure the phase differences between the received signals, enabling the determination of the angular position of radio sources.6. Radio Telescope: A telescope designed to receive and analyze radio waves from celestial objects.7. Spectral Line: A narrow emission or absorption feature in the radio spectrum of a celestial object, resulting from transitions between quantum states of atoms or molecules.8. Synchrotron Radiation: Radiation emitted by charged particles moving in circular paths at nearly the speed of light, a common phenomenon in astrophysical objects with strong magnetic fields.9. Telescope Beam: The region of space through which a radio telescope is sensitive to radiation, typically described by its shape and size.10. Very Large Array (VLA): A radio telescopeconsisting of multiple antennas arranged in a large array, used for high-resolution observations of celestial objects.11. Zenith Angle: The angle between the vertical direction and a line from the observer to a celestial object, measured from the zenith point overhead.12. Angular Resolution: The ability of a radiotelescope to distinguish between two closely spaced sources, measured as the smallest angular separation at which two sources can be resolved.13. Astrometry: The branch of astronomy dealing withthe measurement of positions and movements of celestial objects.14. Bandwidth: The range of frequencies over which a radio receiver is sensitive to incoming radiation.15. Continuum Radiation: The smooth, featurelesscomponent of the radio spectrum of a celestial object, resulting from the superposition of many spectral lines.16. Cosmic Microwave Background (CMB): Radiation left over from the Big Bang, detected as a faint microwave signal filling all directions in the sky.17. Dynamic Range: The ratio between the strongest and weakest signals that a radio telescope can detect simultaneously.18. Faraday Rotation: The rotation of the plane of polarization of electromagnetic waves as they pass through a magnetized plasma, caused by the interaction between the magnetic field and the plasma's charged particles.19. Frequency Synthesiser: A device that generates precise radio frequencies for use in radio telescopes, ensuring accurate measurements of celestial objects.20. Noise Temperature: A measure of the random fluctuations in a radio receiver's output signal, expressedas an equivalent temperature.21. Point Source: A celestial object that appears as a single point of radiation in the sky, with no angular structure resolvable by a radio telescope.22. Radio Galaxy: A galaxy that emits significant amounts of radio radiation, often due to synchrotron radiation from relativistic electrons in its magnetic field.23. Redshift: The increase in wavelength of electromagnetic radiation emitted by a celestial object due to its motion away from the observer, resulting in a shift towards longer wavelengths in the spectrum.24. Single-Dish Telescope: A radio telescope consisting of a single large antenna, used for observations ofcelestial objects.25. Spectral Index: A measure of the steepness or flatness of a celestial object's radio spectrum, indicating the relative amounts of radiation at different frequencies.26. Synthesis Imaging: The process of combining signals from multiple antennas in a radio telescope array to create an image of a celestial object, improving both angular resolution and sensitivity.27. Telescope Time: The amount of time allocated to a specific research project for observations with a radio telescope.28. Telescope Sensitivity: The ability of a radio telescope to detect weak radio signals from celestial objects, measured as the minimum flux density detectable.29. Tidal Force: The force acting on a body due to the differential gravitational attraction of two other bodies, such as the Moon and Sun on Earth, causing tides.30. Total Power Map: A map of the sky created by measuring the total power of radio radiation received from different directions, used to identify and study celestial objects.This is a brief overview of some key terms related to radio astronomy. The field is vast and diverse, and this list is not exhaustive. However, it provides a solid foundation for understanding the language and concepts of radio astronomy.。

a r X i v :a s t r o -p h /0412513v 1 20 D e c 2004What is the origin of the soft excess in AGN?Małgorzata Sobolewska ∗and Chris Done †∗Copernicus Astronomical Center,Warsaw,Poland †University of Durham,Durham,UK Abstract.We investigate the nature of the soft excess below 1keV observed in AGN.We use the XMM-Newton data of the low redshift,optically bright quasar,PG 1211+143,and we compare it with the Narrow Line Seyfert 1galaxy,1H 0707-495,which has one of the strongest soft excesses seen.We test various ideas for the origin of the soft X-ray excess,including a separate spectral com-ponent (for example a low temperature Comptonized component),a reflection-dominated model,or a complex absorption model.All three can give good fits to the data,and χ2fitting criteria are not sufficient to discriminate among them.Instead,we favor the complex absorption model on the grounds that it is the most physically plausible.INTRODUCTION Several models are considered in the literature to explain the origin of the soft excess below 1keV in AGN,which cannot be explained by either of the two emission com-ponents usually considered in accreting black holes,namely the accretion disk and the high energy power law Comptonized emission.Spectra with a strong soft excess also often show a strong deficit at ∼7keV which again has no obvious identification.Partial covering models allow for a spatially non-uniform cold absorber attenuating only a fraction of intrinsic radiation.They can explain the sharp features around 7keV .However,they require an enormous overabundance of iron (5-30times the Solar value [1];5times the Solar value [2]),and they cannot simultaneously explain the soft X-ray emission.Thus,we do not consider partial covering models in our analysis as they only provide a method to explain the sharp feature around 7keV .The soft excess may originate in a distinct unknown physical process which manifests itself as a separate spectral component (Model 1).It can be modeled e.g.with low temperature Comptonization models [3]).The apparent temperature of this component is rather similar (∼0.1−0.3keV)across a diverse selection of AGN (e.g.[4]).An obvious way to produce a fixed soft excess energy is to relate it to atomic processes rather than continuum emission.In reflection models,the dramatic strength of the soft excess requires that the source is reflection dominated (Model 2).This might happen if the disc fragments at high accretion rates,hiding the hard X-ray source among many clumps.Such models can successfully fit the spectra,but need a large range in ionization states of the reflecting material,and supersolar abundances ([5]).Alternatively,the strong jump in opacity for partially ionized material could result in a soft excess from absorption (Model 3,[6]).The underlying continuum may be modeled by only one component.The soft excess and hardening of the spectrum at high energies would be only artifacts of absorption acting at 1-2keV ,while the H-like iron K alpharesonance absorption from the same material could produce the deficit at∼7keV.RESULTSThe bestfit models(and their components)together with the data to model ratios are shown in Figure1.The models are additionally modified by galactic absorption, cold absorption at the redshift of the source,and one(in1H0707-495)or two(in PG1211+143)warm absorbers modeling the narrow absorption features.Thus,the only thing that differentiates the three models is the description of the origin of the soft excess.FIGURE1.Modeling the soft excess in PG1211+143and1H0707-495.Red:the bestfit spectra.Left: Model1–a low temperature Comptonization(blue),power law(magenta)and ionized reflection(green). In1H0707-495there is no need for the intrinsic power law.Middle:Model2–three reflectors(green, magenta and blue)with different column densities and ionization states,located at different radii.Right: Model3–a power law(blue)and reflection(green)subject to relativistically smeared absorption.TABLE1.Results of spectralfits to the PG1211+143and1H0707-495dataPG1211+143989(939)972(937)998(939)1H0707-495311(288)292(283)318(287)χ2of1.10and1.08,respectively).The bestfit is obtained with the model of complex reflection(with iron abundance set to2times the Solar abundance)with reducedχ2of 1.03.The values ofχ2and number of degrees of freedom infits are presented in Table1.DISCUSSION AND CONCLUSIONSBased only on theχ2criterion it is not possible to uniquely distinguish between the three kinds of models as each of them provides an adequate description.Moreover, many model uncertainties can contribute toχ2,e.g.in both the reflection and absorption scenarios we expect an(unknown!)range of ionization states to be present.Hence direct comparison ofχ2can be misleading when the models are known to be incomplete. Instead,we should be guided as well by physical plausibility.The model with a distinct component involves an emission from an unknown physical process,with unknown processfixing its typical energy,which does not seem to have an analogue in the spectra of galactic black holes(GBHs)[7],and cannot simultaneously explain the7keV feature.By contrast,both reflection and absorption are much more plausible,as they give a physical reasonfixed energy for the soft excess and can reproduce the structure around iron K.However,the reflection models require quite strong supersolar abundances,and there is no evidence for reflection dominated spectra in the GBH systems.The absorption model seems to reproduce the strong soft excesses well with moderate column densities(∼1022−23cm−2)of solar abundance material.Such material would not be seen in the GBHs because it would be completely ionized by the much higher accretion disc temperatures seen in the stellar mass black hole systems.Thus the com-plex absorption model provides an interesting alternative because(1)it does not require a separate soft component in the spectra,(2)the hard radiation slope resulting from the fits is similar to that of GBHs in the high/soft state,(3)it describes the spectra of AGNs in terms of atomic processes,explaining the universal shape of the soft excess.ACKNOWLEDGMENTSThis work was partially supported by the Polish Committee for Scientific Research through grants1P03D01626and2P03D00322.REFERENCES1.Tanaka,Y.,Boller,T.,Gallo,L.,Keil,R.,&Ueda,Y.2004,PASJ,56,L92.Gallo,L.C.,Tanaka,Y.,Boller,T.,et al.,W.N.2004,MNRAS,353,10643.Magdziarz,P.,Blaes,O.M.,Zdziarski,A.A.,Johnson,W.N.,&Smith,D.A.1998,MNRAS,301,1794.Czerny,B.,Nikołajuk,M.,Ró˙z a´n ska,A.,et al.2003,A&A,412,3175.Fabian,A.C.,Ballantyne,D.R.,Merloni,A.,et al.2002,MNRAS,331,L356.Gierli´n ski,M.&Done,C.2004,MNRAS,349,L77.Gierli´n ski,M.&Done,C.2004,MNRAS,347,885。